{
    "0606/astro-ph0606664_arXiv.txt": {
        "abstract": "We present numerical simulations investigating the interaction of AGN jets with galaxy clusters, for the first time taking into account the dynamic nature of the cluster gas and detailed cluster physics.  The simulations successfully reproduce the observed morphologies of radio sources in clusters.  We find that cluster inhomogeneities and large scale flows have significant impact on the morphology of the radio source and cannot be ignored a-priori when investigating radio source dynamics.  Morphological comparison suggests that the gas in the centres of clusters like Virgo and Abell 4059 show significant shear and/or rotation.  We find that shear and rotation in the intra-cluster medium move large amounts of cold material back into the path of the jet, ensuring that subsequent jet outbursts encounter a sufficient column density of gas to couple with the inner cluster gas, thus alleviating the problem of evacuated channels discussed in the recent literature.  The same effects redistribute the excess energy $\\Delta E$ deposited the jet, making the distribution of $\\Delta E$ at late times consistent with being isotropic. ",
        "introduction": "\\label{sec:intro} Recent observations show a multitude of physical effects that occur when active galactic nuclei (AGN) interact with the ambient intracluster medium (ICM). While these effects are widely believed to be crucial for the formation of structure in the universe, they are still poorly understood.  The central galaxy in almost every strong cooling core contains an active nucleus and a jet--driven radio galaxy. The radio power of these cooling cores is somewhat correlated with the X-ray luminosity, although the range of the radio power is much greater than the range of the X-ray core power. This is supported by a recently discovered correlation between the Bondi accretion rates and the jet power in nearby, X-ray luminous elliptical galaxies \\citep{allen:06}. These results show that the AGN at the centres of large elliptical galaxies feed back enough energy to quench cooling and star formation, thus providing a possible explanation for the observed cut-off at the bright end of the galaxy luminosity function \\citep{benson:03, croton:05, bower:06}.  The study of AGN-ICM interactions is of great interest in a much broader context.  Given the Magorrian relation \\citep{magorrian:98}, clusters hosting a giant cD galaxy should have a massive black hole in their centre.  The growth of these black holes is an integral part of feedback--regulated cooling of the ICM.  Radio-loud AGN inflate kpc-sized bubbles of relativistic, underdense plasma (a.k.a. radio lobes) that displace the hot intra-cluster medium and thus appear as depressions in the X-ray surface brightness. High-resolution X-ray observations of cooling flow clusters with {\\em Chandra} have revealed a multitude of X-ray holes often coincident with patches of radio emission, a compiled list of which is presented by \\citet{birzan:04}.  Numerical simulations of hot, underdense bubbles in clusters of galaxies have been performed by a number of authors \\citep[e.g.][]{churazov:01, bruggen:02, bruggen:02a, reynolds:01, ruszkowski:04, dallavecchia:04,omma:04,omma:04b}. Common to all of these simulations is that they use a spherically symmetric, analytical profile for the ICM.  \\cite{quilis:01} simulated AGN feedback in a cosmologically evolved cluster but did not model the jets.  Most recently, \\citep{vernaleo:05} performed three-dimensional simulations of a jet in a hydrostatic, spherically symmetric cluster model. In their simulation, the jet power was modulated by the mass accretion rate across the inner boundary. In all of their models, jet heating failed to prevent the catastrophic cooling of gas at the centre because the jet preferentially heated gas lying along the jet axis but failed to heat matter in the equatorial plane. However, in reality the dynamics of the jet may be strongly affected by the ambient medium. Bulk velocities may advect and distort the jet and the radio lobe.  Density inhomogeneities can affect both the jet propagation and the deposition of entropy in the ICM by the jet. For the purpose of studying the extent of cluster heating by AGN, the crucial question is {\\em exactly where and how much} energy is deposited by the jet in the ICM. The primary aim of this work is therefore to study the dynamics of the jet and the extent of ICM heating subject to the motions and inhomogeneities of the ICM.  In this letter, we present first results from hydrodynamical simulations of AGN heating in a realistic galaxy cluster.  The primary objectives of this work is to study the interaction of the jet with a dynamic, inhomogeneous ICM.  The letter is organised as follows: \\S\\ref{sec:code} describes the technical setup of the simulations, \\S\\ref{sec:discussion} presents the results and discusses the implications for cluster physics, and \\S\\ref{sec:summary} summarised and concludes. ",
        "conclusions": "\\subsection{Cluster and Radio Source Morphology} While the jet is active, the jet-inflated cocoon and the surrounding swept up shell generally follow the morphological evolution discussed previously in the literature: The jets inflate two oblong cocoons, which push thermal gas aside.  The early expansion is supersonic, shocking the surrounding gas directly, while at later times the expansion slows down and eventually becomes sub-sonic. Fig.~\\ref{fig:maps1} shows a time series of density and entropy cuts as well as simulated X-ray maps (using the Chandra ACIS-I response and assuming an APEC thermal model) and low-frequency radio maps (assuming equipartition and neglecting radiative cooling).  The aspect ratio of the lobes is roughly 3 and decreases with time, indicating the lateral spreading and shear of the radio plasma.  \\begin{figure} \\resizebox{\\columnwidth}{!}{\\includegraphics{fig1.eps}}\\ \\caption{Top-to-bottom: Time series of snapshots at 0Myrs, 5Myrs, 10Myrs, 20Myrs, 40Myrs, 80Myrs, and 160Myrs after jet onset. Left-to-right: (a) density cut through cluster centre, (b) entropy cut through cluster centre. The images are 450 kpc in width and 335 kpc in height.\\label{fig:maps1}} \\end{figure}  \\begin{figure} \\resizebox{\\columnwidth}{!}{\\includegraphics{fig2.eps}} \\caption{Time series of snapshots (0Myrs, 5Myrs, 10Myrs, 20Myrs, 40Myrs, 80Myrs, and 160Myrs after jet). Left-to-right: (a) Top panel: velocity map through cluster center; below: low-frequency radio synchrotron map, (b) X-ray maps (red: 0.3-2\\,keV, green: 2-5\\,keV, blue:\\,5-10 keV).\\label{fig:maps2}} \\end{figure}  During the active source evolution, while the jet is on, the source closely resembles the best studied example of a source of comparable power: Cyg A.  The fact that aspect ratio and X-ray and radio morphology of the simulated source so closely resemble Cyg A indicates that we properly modeled the effect of the dentist drill effect (which can, in fact, be directly observed in Cyg A, where the jet axis is clearly mis-aligned with the radio and X-ray hot spots).  We verified this by running a test case with half the opening angle of the jitter cone ($10^{\\circ}$ instead of $20^{\\circ}$) which produced too narrow a cocoon (with an aspect ratio of around 7).  The X-ray maps show a clear increase in the surface brightness in the outgoing shock/compression wave and a wispy, turbulent wake behind the wave (Fig.~\\ref{fig:maps2}).  This is best demonstrated by un-sharp masking of the images.  An unsharp masked movie of the X-ray images can be viewed at {\\tt http://space.mit.edu/$\\sim$heinzs/jetpower/}, showing a network of roughly spherical outgoing sound waves that are excited behind the major initial shock wave, consistent with the complex network of sound waves seen in Perseus \\citep{fabian:03}. They also show bud-like bubbles appear at the edges of the main shell when the jet axis moves due to ``denstist-drill'', reminiscent of the budding bubbles in M87 \\citep{forman:05}.  Two distinctions to previously published results are important to point out regarding the morphological evolution of the source:  First, it is clear from the series of snapshots shown in Fig.~\\ref{fig:maps1} that the non-sphericity of the atmosphere induces significant asymmetries in the two sides of the cocoon.  Given that the difference is so strong, it is clear that the impact on the morphology as a whole must be significant.  This justifies our initial motivation for this study: In order to investigate the evolution of radio sources and their impact on galaxy cluster atmospheres, one cannot neglect the dynamic nature of the atmospheres prior to jet injection.  Second, the dynamic nature of the cluster atmosphere significantly alters the late stage evolution, after the jet switches off (between the fourth and the fifth frame at $3.3\\times 10^{6}\\,{\\rm yrs}$) and the source evolution becomes sub-sonic.  During the sub-sonic stage, the impact of the pressure and density of the surrounding gas become more and more important.  In particular, the rotation that is present in the cluster (see Fig.~\\ref{fig:maps2}) has clearly sheared the plasma away from the jet axis.  We will quantify these results below. It is already clear, however, that this effect can solve the dilemma of launching multiple jet episodes into the same cluster: Any evacuated channel created by previous outbursts will have been buffeted around sufficiently to move enough cluster gas back into the path of the new jet to provide a sufficient cross section for interaction and to couple the jet to the inner cluster gas.  The spiral-like morphology of the radio source at late stages is very reminiscent of the large scale ($>30\\,{\\rm kpc}$) morphology of M87 and Abell 4059 \\citep{heinz:02}, where the large scale radio bubbles are clearly mis-aligned with the inner jet.  This leads us to suggest that there might be significant rotation or shear present in the gas of these clusters.  Before proceeding to a more quantitative analysis, it is worth pointing out that the episode of jet activity {\\em does not} disrupt the cluster or blow the central region apart, despite the fact that we are simulating a powerful source (certainly at the upper end of the power range expected from central cluster radio sources).  Contrary to what might be expected naively, the expansion shock weakens and does not super-heat the central atmosphere beyond convective stability. This does not mean, however, that the jet does not induce significant turbulent motion in the central cluster.  \\subsection{Quantitative results} As found in earlier investigations of jet-cluster interactions, the impact of the jet on the cluster leads to an increase in the central entropy and a decrease in density, i.e, a net heat input.  This is qualitatively apparent form the entropy panels in Fig.~\\ref{fig:maps1}.  Figures \\ref{fig:massprofiles} and \\ref{fig:lagrangian} quantify this result.  The bottom panel of Fig.~\\ref{fig:massprofiles} shows the cumulative radial mass, i.e., the mass $M(<r)$ contained within radius $r$ from the cluster centre. The initial time step lies above all other curves out to a radius of about 60 kpc, indicating that the cluster has been inflated (i.e., mass moved out).  The bottom panel of Fig.~\\ref{fig:lagrangian} shows the increase in the mean entropy $\\Delta s=\\Delta \\log{(p/\\rho^{5/3})}$ of the cluster gas at different time steps over the control simulation, plotted against the Lagrangian coordinate $M(<r)$ from Fig.~\\ref{fig:massprofiles}.  This plot clearly shows that the entropy of the central $10^{12}\\,{\\rm M_{\\odot}}$ has been increased by about 0.1, corresponding to an increase in $p/\\rho^{5/3}$ by about 25\\%.  The top panel of Fig.~\\ref{fig:lagrangian} shows the ratio of the cooling times with and without jet.  At late times ($t \\gtrsim 80\\,{\\rm Myrs}$) the cooling time for the inner few$\\times 10^{12}\\,{\\rm M_{\\odot}}$ has been increased by about 50\\%, more at earlier times.  This indicates that, on average, the jet can halt cooling for a time significantly longer than the time the jet is active.  Future simulations including radiative cooling will be necessary to investigate whether this increase is sufficient to offset cooling entirely and whether this result is consistent with {\\em XMM}-Newton and {\\em Chandra} constraints on cluster cooling rates.  \\begin{figure} \\resizebox{0.95\\columnwidth}{!}{\\includegraphics{fig3.eps}} \\caption{{\\em bottom panel:} cumulative mass $M_{\\rm R}$ within radius $r$ as a function of $r$ at different times (grey scales); note the general inflation of the cluster atmosphere by the energy injection into the centre; {\\em top panel:} ratio of target mass $M_{20}$ (mass contained within 20$^{\\circ}$ of jet axis) over $M_{20}$ for the control simulation, indicating the relative gas depletion around the jet axis.  The residual $\\sim 20\\%$ reduction after 160 Myrs is consistent with the general reduction in density seen in the $M_{\\rm R}$.\\label{fig:massprofiles}} \\end{figure}  \\begin{figure} \\resizebox{0.95\\columnwidth}{!}{\\includegraphics{fig4.eps}} \\caption{{\\em bottom panel:} average local increase in thermal cluster entropy $s=log(p/\\rho^{5/3})$ relative to control simulation without jet, plotted as a function of cumulative mass $M(<r)$ (see Fig.~\\ref{fig:massprofiles}).  {\\em top panel:} ratio of average local cooling time of the thermal cluster gas with and without jet, showing the temporary suspension of cooling.\\label{fig:lagrangian}} \\end{figure}  \\begin{figure} \\resizebox{0.95\\columnwidth}{!}{\\includegraphics{fig5.eps}} \\caption{Angular distribution (relative to mean jet axis) of excess energy $\\Delta E$ injected by jet into thermal gas phase as percentage relative to total injected energy.  For comparison, the dotted line shows the isotropic case.\\label{fig:angle}} \\end{figure}  \\subsection{The isotropisation of energy and density} One of the criticisms of models of cluster heating by jets is the non-spherical nature of the energy input.  The problem is twofold: Firstly, if the energy is predominantly injected into the axial direction, the equatorial cluster gas does not benefit from the energy injection and can cool unimpeded, while the axial gas is super-heated, making AGN heating very inefficient.  Secondly, the channel of jet exhaust and super-heated gas along the jet axis would allow subsequent episodes of jet activity to punch easily through the inner cluster and deposit their energy well outside the central region, thus allowing the inner regions to cool unimpeded.  We already suggested that the shear and rotation in the central cluster can rectify the second problem.  The top panel of Fig.~\\ref{fig:massprofiles} shows the radial profile of the target mass $M_{20}$ contained in a $20^{\\circ}$ angle from the jet axis, relative to that of the control simulation.  Clearly, the amount of mass is significantly reduced during the early stages of the simulation.  However, at late times, $M_{20}$ increases back up to 80\\% of the unperturbed value, consistent with the general reduction of density in the central cluster shown in the bottom panel of Fig.~\\ref{fig:massprofiles}. Such a small reduction is dynamically insignificant, i.e., the amount of target mass is more than sufficient for a subsequent jet episode to couple to the cluster at {\\em any} radius.  This proves quantitatively that cluster dynamics can solve the problem of evacuated channels raised by \\cite{vernaleo:05}.  If cluster mass can be moved around to re-fill the jet channel and to buffet the fossil radio plasma around, it is plausible to assume that the energy input of the jet might be isotropised more rapidly as well. The entropy panel of Fig.~\\ref{fig:maps1} shows this qualitatively: there is no clear high entropy channel along the axis of the jet.  In fact, the draft of the rising bubbles draws low entropy material out of the central cluster, as previously discussed by \\citep{churazov:03}.  This is borne out by a quantitative analysis: Fig.~\\ref{fig:angle} shows the excess energy in the thermal cluster gas (once again excluding non-thermal jet exhaust) as a function of polar angle $\\theta$ away from the jet axis (measured around the cluster centre).  At early times, the excess energy is clearly peaked around the jet axis (within about 30$^{\\circ}$),but at late times (160 Myrs, solid black line), the distribution is consistent with being isotropic (dotted black line).  This indicates that individual episodes of jet activity can, in fact, distribute their energy rather isotropically into the cluster.  More detailed simulations including radiative cooling and multiple jet episodes will be necessary to develop a detailed picture of how heating and cooling in realistic cluster atmospheres proceed.   We have presented simulations of jet-cluster feedback that integrate a correct representation of cluster dynamics, including dark matter, and the collimated input of energy and momentum from an AGN.  The simulations accurately reproduce the observed morphological appearance of powerful radio sources like Cyg A, if a jitter is imposed on the jet to account for unresolved dynamical instabilities of the jet commonly referred to as the ``dentist drill effect''.  We find that the dynamic structure of the cluster in the form of rotation and shear in the velocity field and the inhomogeneity and anisotropy of the cluster have significant impact on the evolution of the jets and the radio lobes they inflate.  At late times, we find that the excess energy deposited into the cluster by the action of the jets is distributed well away from the axis, consistent with roughly spherically symmetric energy input.  After long times (>160 Myrs), the inner cluster has been sufficiently rearranged to essentially erase the low density channel blasted out by the jet and move enough material into the way of subsequent jet outbursts to couple efficiently with the inner cluster.  This solves the problem of low efficiency feedback found in simulations of spherically symmetric, static atmospheres \\citep{vernaleo:05}.  \\thanks{We thank Volker Springel for providing us with a set of Gadget simulated clusters.  We thank Mateusz Ruszkowski, Chris Reynolds, Mitch Begelman, and Paul Nulsen for helpful discussions.  SH acknowledges support by the National Aeronautics and Space Administration through Chandra Postdoctoral Fellowship Award Number PF3-40026 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics Space Administration under contract NAS8-39073.  MB acknowledges support by DFG grant BR 2026/2 and the supercomputing grant NIC 1658 at the John von-Neumann centre for computing at the Forschungszentrum J\\\"ulich. The software used in this work was in part developed by the DOE-supported ASCI/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.}"
    },
    "0606/astro-ph0606387_arXiv.txt": {
        "abstract": "We have analyzed \\chandra\\ calibration observations of \\bet\\ ($\\alpha$~Ori, M2\\,Iab, m$_V=0.58$, 131~pc) obtained at the aimpoint locations of the HRC-I (8 ks), HRC-S (8 ks), and ACIS-I (5 ks).  \\bet\\ is undetected in all the individual observations as well as cumulatively.  We derive upper limits to the X-ray count rates and compute the corresponding X-ray flux and luminosity upper limits for coronal plasma that may potentially exist in the atmosphere of \\bet\\ over a range of temperatures, $T=0.3-10$~MK. We place a flux limit at the telescope of $\\fx\\approx4\\times10^{-15}$~ergs~s$^{-1}$~cm$^{-2}$ at $T=1$~MK.  The upper limit is lowered by a factor of $\\approx3$ at higher temperatures, roughly an order of magnitude lower than that obtained previously.  Assuming that the entire stellar surface is active, these fluxes correspond to a surface flux limit that ranges from $30-7000$~ergs~s$^{-1}$~cm$^{-2}$ at $T=1$~MK, to $\\approx 1$~ergs~s$^{-1}$~cm$^{-2}$ at higher temperatures, five orders of magnitude below the quiet Sun X-ray surface flux.  We discuss the implications of our analysis in the context of models of a buried corona and a pervasive magnetic carpet.  We rule out the existence of a solar-like corona on \\bet, but cannot rule out the presence of low-level emission on the scale of coronal holes. ",
        "introduction": "\\label{s:intro}  \\begin{deluxetable}{lll} \\tablecolumns{2} \\tabletypesize{\\small} \\tablecaption{Stellar parameters for \\bet\\ \\label{t:stelpar}} \\tablewidth{0pt} \\startdata \\hline \\hline Other Names & \\multicolumn{2}{l}{$\\alpha$~Ori / 58 Ori / HD~39801 / HR~2061 / SAO 113271 / HIP 27989 } \\\\ (R.A., Dec) & (05:55:10.3053,~+07:24:25.426) & ICRS 2000.0 \\\\ ($l_{\\rm II}, b_{\\rm II}$) & ($199^{\\circ}.79, -8^{\\circ}.96$) & SIMBAD \\\\ Spectral Type & M2~Iab & SIMBAD \\\\ $m_{V}$ & 0$^m$.58 & SIMBAD \\\\ $B - V$ & 1$^m$.77 & SIMBAD \\\\ T$_{\\rm eff}$ & 3650~K & Levesque et al.\\ (2005) \\\\ $B.C.$ & -1.6 & Levesque et al.\\ (2005) \\\\ Distance & $131 \\pm 28$ pc & Hipparcos (Perryman et al.\\ 1997) \\\\ $\\lbol$ & $1.2~\\times~10^{38}$~ergs~s$^{-1}$ & \\hfil \\\\ Angular Diameter & $44.6 \\pm 0.2$ mas & CHARM (Richichi \\& Percheron 2002) \\\\ Radius & $631 \\pm 134$ R$_{\\odot}$ & \\hfil \\\\ Mass & $14$ M$_{\\odot}$ & Figure~\\ref{f:evoltrack} \\\\ Age & $10$ Myr & Figure~\\ref{f:evoltrack} \\\\ Gravity & $1\\pm0.4$~cm~s$^{-2}$ & \\hfil \\\\ Circumstellar Column & $6.2 \\times 10^{21}$~cm$^{-2}$ & Hagen (1978) \\\\ \\enddata \\end{deluxetable}  \\bet\\ (see Table~\\ref{t:stelpar}) is a nearby (131~pc), bright ($V=0.58^m$) evolved red supergiant star (M2\\,Iab, $B-V=1.77$). It has been monitored extensively in the optical (e.g., Wilson et al. 1997, Burns et al. 1997), and displays irregular brightness variations ($V=0.3-0.8$, Gray 2000) that have been interpreted as large-scale surface structures or activity (e.g., Lim et al.\\ 1998 and Gray 2000). A definitive estimate of its age does not exist since it cannot be identified with nearby Orion associations (Lesh 1968), and estimates of its mass vary from $5$~M$_{\\odot}$ (Dorch 2004) to $10-30$~M$_{\\odot}$ (Gray 2000; also Lambert et al.\\ 1984 and references therein). However, reasonable estimates can be obtained by comparing the evolutionary tracks of high-mass stars with the location of \\bet\\ on a color-magnitude diagram.  We show such a comparison using the Geneva stellar evolutionary tracks (Lejeune \\& Schaerer 2001) in Figure~\\ref{f:evoltrack}. The stellar models vary in initial mass, metallicity, and mass-loss rates, and no rotation effects are included.  While the systematic model uncertainties prevent a definitive assessment of the evolutionary history of \\bet, note nevertheless that most of the models predict similar values of age and mass for the star, and therefore an approximate estimate of the gross properties of \\bet\\ is possible. Based on Figure~\\ref{f:evoltrack}, we adopt a current mass of $\\approx14$~M$_{\\odot}$ and an age of $\\approx10$~Myr for \\bet.  The uncertainty in these parameters do not affect our conclusions.  \\begin{figure}[htb!] \\begin{center} {\\includegraphics{f1.eps}} \\caption{Evolutionary state of \\bet.  The tracks from the Geneva stellar evolutionary models (Lejeune \\& Schaerer 2001) are overlaid on a color-magnitude diagram, with the position of \\bet\\ represented by an asterisk.  The lines are solid until their closest approach to \\bet\\ in their evolutionary history, after which they are depicted as dotted lines (note that the magnitude range is different for the two axes).  The lines are labeled by the model used (the models differ in their masses, mass-loss, and metallicities; see Lejeune \\& Schaerer for a detailed description). The mass (in M$_{\\odot}$) and age (in Myr) of the star at the point of closest approach is also shown for each model. Inspection of the figure shows that while definitive measurements of the mass and age of \\bet\\ are not feasible due to the large systematic model uncertainties (e.g., masses ranging from 10 to 15 M$_{\\odot}$ are plausible) but the majority of the models correspond to a mass of $\\approx14$~M$_{\\odot}$ and an age of $\\approx10$~Myr, which we then adopt as the nominal mass and age of \\bet\\ throughout this paper. \\label{f:evoltrack} } \\end{center} \\end{figure}  \\bet\\ has never been detected in X-rays, and is in a region of the H-R diagram where a stable corona is not expected to exist (Ayres et al.\\ 1981, Linsky \\& Haisch 1979, Haisch, Schmitt, \\& Rosso 1991, Rosner et al.\\ 1995, H\\\"{u}nsch \\& Schr\\\"{o}der 1996, H\\\"{u}nsch et al.\\ 1998). Nevertheless, there is evidence from numerical MHD simulations that dynamo action which can produce large scale magnetic fields can exist on such stars (Dorch 2004).  Furthermore, chromospheric activity suggesting the presence of coronae has been detected in numerous late-type giant stars: coronal proxy lines such as Si\\,IV and C\\,IV have been detected in ostensibly non-coronal giants such as Arcturus and Aldebaran (Ayres et al.\\ 2003); and forbidden coronal lines have been reported in \\fuse\\ observations of $\\beta$ Cet by Redfield et al.\\ (2003).  Thus, while unlikely, it is plausible that hot coronal plasma may exist on supergiants like \\bet.  Here we present an analysis of X-ray observations of \\bet\\ obtained as part of the science instrument calibration program of the \\chandra\\ X-Ray Observatory (\\chandra). The data are described in \\S\\ref{s:data}. \\bet\\ is not detected in X-rays in any of our individual observations or in co-added data, and in \\S\\ref{s:ul} we set the most stringent upper limits to the X-ray flux determined thus far, improving upon limits obtained from the \\rosat\\ All-Sky Survey (RASS) by two orders of magnitude.  In \\S\\ref{s:discuss}, we discuss our results in the context of models of coronae on late-type stars. We summarize in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:discuss}  \\subsection{Magnetic Fields} \\label{s:magnetic_fields}  If any coronal plasma exists on the surface of \\bet, it is likely confined in place by magnetic fields.  From numerical studies (see \\S\\ref{s:magnetic_carpet} below), we expect fields of strength approaching 500~G.  How does it compare to the observational limit?  Here we estimate the magnetic field strength required to confine plasma, assuming that there does exist coronal plasma emitting at a level just below the derived flux upper limit (\\S\\ref{s:fluxul}), and further assuming equipartition of thermal and magnetic energy densities in the corona, i.e., $\\frac{B^2}{8\\pi} = 2 n_e k_B T$, and hence \\begin{equation} B = \\left( \\frac {16 \\pi \\mu^2 k_B^2 T^2 \\lx} {h R_*^2 f \\Lambda(T)} \\right)^{1/4} \\end{equation} where $T$ is the assumed plasma temperature, $\\Lambda(T)$ is the power emitted from a unit volume of the plasma, $\\mu$ is the effective mass of an average particle ($\\mu\\approx\\frac{1}{2}$ for fully ionized plasma), $R_*$ is the stellar radius, $k_B$ is Boltzmann's constant, $f$ is the filling fraction of the X-ray active region on the surface of the star, and $h$ is the height of the corona.  Note that $B~\\propto~(\\lx/f)^{1/4}$, i.e., it is only weakly dependent on the X-ray luminosity and the filling fraction.  \\begin{figure}[htb!] \\begin{center} {\\includegraphics{f4.eps}} \\caption{Magnetic field strength $B$ required to confine coronal plasma on \\bet, as a function of the plasma temperature.  The shaded regions represent the variation expected due to a patchy coverage of the surface by the magnetic field, with the lower edges corresponding to a filling fraction $f=1$ and the upper edges to $f=0.01$.  The lower shaded regions represent a corona with the same height as the scale height of the chromosphere, and for comparison, the upper shaded regions represent a corona with a similar height as that of the solar corona.  The two plots are calculated assuming different column densities of absorption: $N_H=6.2\\times10^{21}$ (left) and $10^{22}$~cm$^{-2}$ (right). \\label{f:BTf} } \\end{center} \\end{figure}  In Figure~\\ref{f:BTf}, we show the required field strength $B$, for a range of filling fraction $f$, for different estimates of the height of the corona.  We obtain estimates both for a corona that is approximately the same height as the solar corona ($h=10^{10}$~cm) and for one that extends as high as the chromospheric scale height ($h={kT_{\\rm chrom}R_{*}^2}/{\\mu~m_{\\rm H}GM_{*}}\\approx2\\times10^{12}$~cm). The latter is relevant in the context of a ``buried corona'' scenario expounded by Ayres et al.\\ (2003). We find that the required $B$ ranges from $\\lesssim0.1$~G up to $\\approx5$~G, considerably less than the field strength expected from modeling (Dorch 2004).  Weaker fields are sufficient for lower levels of X-ray emission. (Note that our calculation of $B$ is not a limit, and we cannot rule out the presence of a stronger magnetic field. However, we can say that a stronger magnetic field is {\\sl not required}.) We thus conclude that sustaining a weak corona on \\bet\\ is feasible, and such coronae cannot yet be ruled out based on the X-ray flux limits.  \\subsection{Magnetic Carpet} \\label{s:magnetic_carpet}  Numerical MHD simulations of stars such as \\bet\\ have shown that a highly structured magnetic dynamo may operate on them (Dorch 2004, Freytag et al.\\ 2002), with field strength as high as $B_*\\approx500$~G. The energy spectra of the magnetic energy density peaks cascade to small scales, with a preponderately high contribution at scales $\\sim \\frac{1}{10}R_*$.  This is very similar to the magnetic activity in the quiet Sun, where for comparison, field strengths reach up to $\\sim2000$~G (Cerde\\~{n}a et al.\\ 2006), and the magnetic structures of mixed polarity emerge uniformly over the surface on timescales approximately the same as the convective cells.  This is the so-called `magnetic carpet' (Schrijver et al.\\ 1997, Title 2000) which pervades the surface of the Sun away from active regions.  If red supergiants are magnetically active in a similar manner, the total energy dumped into the corona, and the resulting plasma temperatures and X-ray luminosities, are similar in scope to that seen in the quiet Sun and in coronal holes, scaled only by the energy densities and volumes involved. We assess the feasibility of this scenario by computing the expected surface flux upper limit $\\surffx$ for \\bet.  Figure~\\ref{f:bet_ul_nH_err} shows bands corresponding to the surface flux limits, calculated assuming a variety of plausible absorbing column densities, $N_H=(4,6.2,10)\\times10^{21}$~cm$^{-2}$ (see \\S\\ref{s:fluxul}). Here we assume that the entire surface is covered by magnetic structures that contribute to the X-ray emission, i.e., that the filling fraction $f=1$.  However, supergiant stars like \\bet\\ possess very large convection cells (Lim et al. 1998) which suggests that $f << 1$.  Smaller values of $f$ imply that the surface flux upper limit will be correspondingly higher.  We show the effects of smaller filling fractions in Figure~\\ref{f:BTf}.  The X-ray luminosity of the quiet Sun is $\\approx10^{27}$~ergs~s$^{-1}$ (Golub \\& Pasachoff 1997) corresponding to a surface flux of $1.64\\times10^{4}$~ergs~s$^{-1}$~cm$^{-2}$.  If similar processes to those on the Sun operate on \\bet, it is reasonable to expect that the energy flux deposited into the corona will be similar, scaled by the available magnetic energy density.  Thus, we expect a surface X-ray flux of $\\approx10^3$~ergs~s$^{-1}$~cm$^{-2}$ on \\bet.  Note however that the surface flux on \\bet\\ could be lower if it were dominated by features similar to the solar coronal holes. The X-ray flux from coronal holes on the Sun is much lower, $\\approx3\\times10^{3}$~ergs~s$^{-1}$~cm$^{-2}$ (Vernazza \\& Smith 1977, Schrijver et al., 2004), and again scaling it to the expected magnetic field energy density present on \\bet, we obtain a possible X-ray surface flux of $\\approx200$~ergs~s$^{-1}$~cm$^{-2}$ (see Figure~\\ref{f:bet_ul_nH_err}).  The upper limit on $\\surffx$ is strongly dependent on the temperature of the plasma because of the sensitive dependence of the observable spectrum on the absorption column.  We find that the limit is as low as $1$~ergs~s$^{-1}$~cm$^{-2}$ ($\\lx/\\lbol \\approx 10^{-10}$) at high temperatures, but cannot be reduced below $\\approx10^2-10^3$~ergs~s$^{-1}$~cm$^{-2}$ for $T<1$~MK ($\\lx/\\lbol \\approx 10^{-8}-10^{-7}$) for the sensitivity achieved thus far with \\chandra.  While the upper limit at low temperatures still lies above the flux expected from the solar analogy, this calculation does rule out the existence of pervasive quiet Sun type emission at high temperatures, since a surface flux of $\\approx10^2-10^3$~ergs~s$^{-1}$~cm$^{-2}$ from plasma at $T>2$~MK would have been easily detected.  Note however that this does not preclude patchy and highly localized regions of magnetic activity that produces high temperature plasma, or more pervasive low-temperature plasma emission arising from this mechanism: plasma at $T\\lesssim1$~MK will remain undetected at the sensitivity limit of our observations. Since the quiet Sun and coronal hole plasma is at $T\\gtrsim1$~MK, this further suggests that even if hot plasma at lower temperatures is present on \\bet, it will bear little resemblance to the solar case.  \\subsection{Coronal Proxies} \\label{s:coronal_proxies}  Numerous UV and FUV chromospheric lines have been identified as proxies for coronal activity in normal and giant stars.  Note that the physical relationship between the mechanisms that generate coronal and chromospheric line emission is poorly understood, and the proxy lines are often formed at temperatures very different from those that characterize coronae.  However, assuming that similar processes occur on \\bet\\ as on coronally active giants and main sequence stars, we investigate whether our derived flux upper limits (see \\S\\ref{s:fluxul}) are consistent with observations of coronal proxy lines on \\bet.  Based on known correlations between X-ray flux and proxy lines such as C\\,IV, Si\\,IV, etc., we can estimate the X-ray luminosity that can be expected from \\bet.  For main sequence stars, Redfield et al.\\ (2003) find a strong correlation between the soft X-ray flux and the Fe\\,XVIII\\,$\\lambda974$ flux (see their Figure 7) from \\fuse\\ (Far Ultraviolet Spectroscopic Explorer) spectra of main sequence and late-type stars.   Fe\\,XVIII\\, forms above $T=2$~MK, and has a peak response at $T=6$~MK, and is thus sensitive to plasma temperatures similar to that on active binaries such as Capella. Redfield et al.\\ place an upper limit of $8\\times10^{-13}$~ergs~s$^{-1}$~cm$^{-2}$ on the Fe\\,XVIII flux from \\bet\\ ($L_{\\rm FeXVIII}/\\lbol=7.5\\times10^{-11}$) which implies a value for the unabsorbed X-ray luminosity $\\lx/\\lbol < 10^{-7}$ if supergiants such as \\bet\\ follow the correlation seen in their  Figure 7. The flux upper limit we calculate (\\S\\ref{s:fluxul}, Figure~\\ref{f:bet_ul_nH_err}) corresponds to $\\lx/\\lbol<10^{-9}$ for $T>1$~MK, well below the value predicted by the main sequence correlation. Conversely, assuming that the correlation is valid for supergiants, our X-ray upper limit implies a stronger constraint on the Fe\\,XVIII emission, $L_{\\rm FeXVIII}<7.5\\times10^{-13}~\\lbol$.  Similarly, there exists a clear correlation between C\\,IV and X-ray luminosities for giant stars (see, e.g., Ayres et al.\\ 1997, especially their Figure 2).  Even though there is evidence for considerable scatter for different types of stars, one can establish a general correspondence that is valid to within an order of magnitude.  Based on \\textit{IUE} ultraviolet spectra, Basri et al.\\ (1981) place a $3\\sigma$ upper limit of $1.5\\times10^{-13}$~ergs~s$^{-1}$~cm$^{-2}$ on the C\\,IV\\,$\\lambda1549$ flux from \\bet, and this translates to an unabsorbed X-ray luminosity limit of $\\lx/\\lbol \\lesssim 10^{-9}$.  In the temperature range $T=1-10$~MK (which matches the sensitivity range of X-ray measurements from {\\sl Einstein} and \\rosat, on which this correlation is based), this limit is comparable to our observational limit determined above.  Lastly, from \\fuse\\ spectra of \\bet\\ published by Dupree et al.\\ (2005; see their Figure 4), we estimate an upper limit of $3\\times10^{-14}$~ergs~s$^{-1}$~cm$^{-2}$ on its Si\\,IV flux. If the correlation found by Ayres et al. (2003, see their Figure 4) for giant stars applies to supergiants like \\bet, then the limit of $f_{\\rm Si\\,IV} / f_{\\rm bol} < 5\\times10^{-10}$ corresponds to an expected X-ray unabsorbed flux of $\\fx / f_{\\rm bol} \\lesssim 10^{-9}$ which is again comparable to the observed limit we obtain.  Note that previous studies of \\bet\\ in the optical and UV have seen no evidence for chromospheric temperatures above 6000 K (e.g. Lobel et al. 2000 \\& 2001 and Carpenter et al. 1994) and our observations are consistent with those findings.  From the comparisons above, we conclude that we have now achieved a sensitivity in the X-ray regime that is comparable to the sensitivity achieved in the far UV with coronal proxy lines, especially at temperatures $T>3$~MK. In combination with the surface flux limit arguments made above (\\S\\ref{s:magnetic_carpet}) and the stringent upper limit on $\\lx/\\lbol$ obtained here, we conclude that any high-temperature plasma will have to arise from a mechanism other than that which normally operates on the Sun.  Finally, note also that if the hot plasma is ``buried'' in the highly extended chromospheric material (as suggested by Ayres et al.\\ 2003), then both the chromospheric proxies and the coronal X-ray flux will be subject to significant absorption, and our placement of \\bet\\ on these flux-flux correlation diagrams will be systematically low.  In such a case, the limit on $\\lx/\\lbol$ will be less restrictive, but will not affect our conclusions.    We have carefully analyzed over 20~ks of \\chandra\\ observations of \\bet\\ in an effort to detect X-ray emission from the massive red supergiant.  However, \\bet\\ remains undetected, and we derive an upper limit to the X-ray count rate by calculating the rate that would have resulted in a detection given the extant background.  We have converted this count rate limit to a flux limit at the telescope by computing the response of \\rosat\\ and \\chandra\\ instruments to isothermal plasma producing optically-thin thermal emission, and thereby derive the most stringent upper limits to the X-ray flux from \\bet\\ obtained thus far.  We find a limit for the flux from \\bet\\ at the telescope of $\\fx < 4\\times10^{-15}$~ergs~s$^{-1}$~cm$^{-2}$ for temperatures $T>1$~MK.  At lower temperatures, we place a limit of $\\fx < 3\\times10^{-14}$~ergs~s$^{-1}$~cm$^{-2}$ on the flux.  The flux limit at Earth can be converted to a stellar surface flux upper limit and to an $\\lx/\\lbol$ limit using the known distance and size of \\bet.  We compare the surface flux limits with the flux expected from a solar like emission mechanism, where a pervasive magnetic field maintains a low-level corona, as in the quiet Sun or solar coronal holes.  We rule out such emission at temperatures $>1$~MK, but such emission is still feasible at lower temperatures.  The minimum magnetic field necessary to maintain such a corona is $<10$~G, well within theoretical expectations.  We compare the $\\lx/\\lbol$ upper limit we derive with the limits obtained from non-detections of coronal tracer lines such as C\\,IV, Si\\,IV, and Fe\\,XVIII\\,$\\lambda974$ and find that we achieve sensitivities in the X-ray comparable to that in the coronal proxies.  These limits reinforce the conclusions arrived at above, that high-temperature plasma, even at levels expected in the presence of stellar coronal holes, is absent on \\bet, but the existence of low-temperature plasma cannot be ruled out."
    },
    "0606/astro-ph0606452_arXiv.txt": {
        "abstract": "The new three year WMAP data seem to confirm the presence of non-standard large scale features in the cosmic microwave anisotropy power spectrum. While these features may hint at uncorrected experimental systematics, it is also possible to generate, in a cosmological way, oscillations on large angular scales by introducing a sharp step in the inflaton potential. Using current cosmological data, we derive constraints on the position, magnitude and gradient of a possible step. We show that a step in the inflaton potential, while strongly constrained by current data, is still allowed and may provide an interesting explanation to the currently measured deviations from the standard featureless spectrum. Moreover, we show that inflationary oscillations in the primordial power spectrum can significantly bias parameter estimates from standard ruler methods involving measurements of baryon oscillations. ",
        "introduction": "The recent three year results from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite \\cite{wmap3cosm,wmap3pol,wmap3temp,wmap3beam} have further confirmed with an extraordinary precision the inflationary paradigm of structure formation in which primordial fluctuations are created from quantum fluctuations during an early period of superluminal expansion of the universe \\cite{Starobinsky:1979ty,muk81,bardeen83}.  Indeed, soon after the WMAP data release, a number of authors investigated the possibility to discriminate between several single-field inflationary models using this new, high quality, dataset \\cite{alabidi,PeirisEasther,Lewis:2006ma,Seljak:2006bg,Magueijo:2006we, Liddle06,kinney06,Martin:2006rs}. One of the main conclusions of these works is that some inflationary models, such as quartic chaotic models of the form $V(\\phi)\\sim \\lambda \\phi^4$, may be considered ruled out by the current data, while others, such as chaotic inflation with a quadratic potential $V(\\phi) \\sim m^2 \\phi^2$, are consistent with all data sets.  One important assumption in these analyses (apart for \\cite{Martin:2006rs}) is that the inflaton's potential is featureless, i.e., there is no preferred scale during inflation and the primordial power spectrum of density perturbations in Fourier $k$-space can be well approximated by a power law $k^n$, where the spectral index $n$ is almost scale independent. The main prediction of these models is that the anisotropy angular power spectrum should be ``smooth'' and not show features in addition to those provided by the baryon-photon plasma oscillations at decoupling within the framework of the standard $\\Lambda$CDM model of structure formation.  The current WMAP data is in very good agreement with this hypothesis: several non-standard features in the anisotropy angular power spectrum detected in the first year data have now disappeared thanks to the longer integration time of the observations and better control of systematics (see \\cite{wmap3cosm}).  However, features in the large scale anisotropy spectrum are still present in the new release. Moreover, some of the cosmological parameters derived from the new WMAP data, like, for instance, the low value of the variance of  fluctuations $\\sigma_8$, appear in tension with those derived by complementary data sets. It is therefore timely to investigate a larger set of inflationary models and to consider a cosmological origin of these unexpected features.  A departure from power law behavior of the primordial power spectrum could be caused by a change of the initial conditions, due to trans-planckian physics \\cite{Martin:2003kp,Easther:2002xe} or unusual initial field dynamics \\cite{Burgess:2002ub, Contaldi:2003zv} or by some brief violation of the slow roll conditions during inflation~\\cite{Kofman:1989ed,Starobinsky:1992ts}. We will investigate a model of the second type where features in the temperature and density power spectra arise due to a step-like change in the potential parameters, as proposed by~\\cite{ace}. A sharp step in the inflaton mass, caused, e.g., by a symmetry breaking phase transition, generates indeed $k$-dependent oscillations in the spectrum of primordial density perturbations.  The goal of our paper is to make use of the recent three year WMAP data (WMAP3) and other datasets to constrain the possibility of a step feature in the inflaton potential. For this purpose we adopt the phenomenological model proposed by Adams et al. \\cite{ace}, where a step feature is added to the chaotic inflationary potential in the following way: \\begin{equation} V(\\phi) = {1\\over 2} m^2 \\phi^2 \\, \\left( 1 + c \\tanh \\left( \\frac{\\phi-b}{d} \\right) \\right), \\label{tan-potential} \\end{equation} where $m$ is an overall normalization factor, $c$ determines the height of the step, $d$ its gradient and $b$ is the field value on which the step is centered. Previous phenomenological studies of the same~\\cite{Peiris:2003ff} or other oscillatory features~\\cite{Martin:2004yi,Kawasaki:2004pi} have been limited to the first year WMAP data, and in general a full analysis varying also all cosmological parameters is still missing and is the major result of this work.  The paper is organized as follows: In Sec.\\ II we briefly review step-inflation models. In Sec.\\ III we describe our analysis method. In Sec.\\ IV we present our results and, finally, in Sec.\\ V we derive our conclusions. ",
        "conclusions": "The new three year WMAP data seem to confirm the presence of non-standard large scale features on the cosmic microwave anisotropy power spectrum. While these features may hint at uncorrected experimental systematics, a possible cosmological way to generate large angular scale oscillations is to introduce a sharp step in the inflaton potential. By making use of current cosmological data we derive constraints on the position, magnitude and gradient of a possible step in the inflaton potential. Our conclusion is that such a step, while strongly constrained by current data, is still allowed and may provide an interesting explanation to the current measured deviations from the standard featureless spectrum at low $\\ell$.  Surprisingly though, the combination of all CMB data sets with the SDSS data seems to prefer a feature at small scales, which could mimic the effect of baryonic oscillations and reduces the best fit value of $\\Omega_b$. Note that for this it is sufficient to have a minute change, of order $0.1\\%$, in the inflaton mass parameter, but a relatively fast one.  It is an open question if such a sharp step can be realized in realistic inflationary models and if the effect is physical. In general we can exclude the presence of strong features with $c \\geq 0.003 $ in the observable range.  Future experiments like PLANCK will provide better measurements of the polarization and cross temperature-polarization spectra, providing an important check for possible non-standard features."
    },
    "0606/astro-ph0606178_arXiv.txt": {
        "abstract": "Soon after launch, the Advanced CCD Imaging Spectrometer (ACIS), one of the focal plane instruments on the Chandra X-ray Observatory, suffered radiation damage from exposure to soft protons during passages through the Earth's radiation belts. The primary effect of the damage was to increase the charge transfer inefficiency (CTI) of the eight front illuminated CCDs by more than two orders of magnitude. The ACIS instrument team is continuing to study the properties of the damage with an emphasis on developing techniques to mitigate CTI and spectral resolution degradation. We present the initial temperature dependence of ACIS CTI from --120 to --60 degrees Celsius and the current temperature dependence after more than six years of continuing slow radiation damage. We use the change of shape of the temperature dependence to speculate on the nature of the damaging particles. ",
        "introduction": "\\label{sect:intro}  The Chandra X-ray Observatory, the third of NASA's great observatories in space, was launched just past midnight on July 23, 1999, aboard the space shuttle {\\it Columbia}\\cite{cha2}.  After a series of orbital maneuvers Chandra reached its final, highly elliptical, orbit.  Chandra's orbit, with a perigee of 10,000~km, an apogee of 140,000~km and an initial inclination of 28.5$^\\circ$, transits a wide range of particle environments, from the radiation belts at closest approach through the magnetosphere and magnetopause and past the bow shock into the solar wind.  The Advanced CCD Imaging Spectrometer (ACIS), one of two focal plane science instruments on Chandra, utilizes frame-transfer charge-coupled devices (CCDs) of two types, front- and back-illuminated (FI and BI).  Soon after launch it was discovered that the FI CCDs had suffered radiation damage from exposure to soft protons scattered off the Observatory's grazing-incidence optics during passages through the Earth's radiation belts\\cite{gyp00,odell}.  Since mid-September 1999, ACIS has been protected during radiation belt passages and there is an ongoing effort to prevent further damage and to develop hardware and software strategies to mitigate the effects of charge transfer inefficiency on data analysis\\cite{cticorr}.  One symptom of radiation damage in CCDs is an increase in the number of charge traps.  When charge is transfered across the CCD to the readout, some portion can be captured by the traps and gradually re-emitted.  If the original charge packet has been transfered away before the traps re-emit, the captured charge is ``lost'' to the charge packet. The pulseheight read out from the instrument which corresponds to a given energy decreases with increasing transfer distance.  This process is quantified as charge transfer inefficiency (CTI), the fractional charge loss per pixel and is calculated from a linear fit to the pulseheight versus row number; CTI = (slope/intercept).  Damage can exist in the imaging or framestore array, causing parallel CTI in the column direction, or in the serial register, causing serial CTI along rows.  The ACIS CCDs come in two flavors, front- and back-illuminated, which have different manifestations of CTI.  The distribution of re-emission time constants of the electron traps which cause CTI vary depending on the type of damage.  In addition, the pixel to pixel transfer time in the imaging, framestore and serial arrays differs so that the same species of electron trap can produce different CTI results in different locations.  The eight front-illuminated CCDs had essentially no CTI before launch, but are strongly sensitive to radiation damage from low energy protons ($\\sim$100~keV) which preferentially create traps in the buried transfer channel.  The framestore covers are thick enough to stop this radiation, so the initial damage was limited to the imaging area of the FI CCDs.  Radiation damage from low-energy protons is now minimized by moving the ACIS detector away from the aimpoint of the observatory during passages through the Earth's particle belts.  Continuing exposure to both low and high energy particles over the lifetime of the mission slowly degrades the CTI further.\\cite{odell,ctitrend}  As of January 2000, the parallel CTI at 5.9~keV of the ACIS FI CCDs varied across the focal plane from $1 - 2 \\times 10^{-4}$ at the nominal operating temperature of --120$^\\circ$~C with a rate of increase of roughly $3 \\times 10^{-6}$/year.  Parallel CTI in the framestore array and serial CTI were not affected by the initial radiation damage and remain negligible with upper limits of $< 10^{-6}$ (framestore array) and $< 2 \\times 10^{-5}$ (serial array).  The two back-illuminated CCDs (ACIS-S1,S3) suffered damage during the manufacturing process and exhibit CTI in both the imaging and framestore areas and the serial transfer array, but are less sensitive to the low energy particles which damage the FI CCDs because they cannot reach the transfer channel.  The parallel CTI at 5.9~keV of the S3 BI CCD was  $\\sim 1.6 \\times 10^{-5}$ at a temperature of --120$^\\circ$~C at the beginning of the mission with a strong non-linear flattening of pulseheight at low row numbers due to CTI in the framestore array.  The serial CTI is much larger, $\\sim 8 \\times 10^{-5}$. BI CCD parallel CTI has been increasing at a rate of $1 \\times 10^{-6}$.  Measured CTI is a function of temperature.  Detrapping time constants decrease as the temperature increases so that different populations of traps can become more or less important.  If the detrapping time constant drops below the pixel transfer time or becomes much longer than the typical distance between charge packets, charge is no longer lost to the trap.  The distribution of trap time constants at a particular temperature determines the CTI, so temperature can positively or negatively correlate with CTI.  Conversely, the temperature dependence of CTI reflects the particular blend of electron traps.  The temperature dependence of ACIS CTI was first measured in 1999 \\cite{gyp00} and again in 2005.  In this paper, we present a comparison of the CTI temperature dependence in both time periods and use this to speculate as to the nature of the intervening particle damage.  This paper begins by describing the calibration data taken during the temperature tests in Section~\\ref{sect:data} and the CTI data reduction in Section~\\ref{sect:reduc}.  The CTI-temperature dependence for both time periods is presented and discussed in Section~\\ref{sect:results}. ",
        "conclusions": "The CTI-temperature dependence of the Chandra ACIS CCDs has been measured twice in orbit; soon after launch in 1999 and again in 2005.  We presented analysis of this data for back-illuminated and front-illuminated CCDs and discussed the differences between the two CCDs and between the two sets of observations.  The temperature dependence can be used to characterize the particular electron trap distribution.  In the case of ACIS, the initial temperature dependence of the FI and BI CCDs in 1999 was very different, due to the different causes of the damage.  The change in the temperature dependence between 1999 and 2005 was similar for both the FI and BI CCDs indicating that the spectrum of particle energies must be hard enough to reach the BI CCD buried channel.  Small differences between the change in temperature dependence in 1999 and 2005 and the initial FI CCD temperature dependence may indicate that the traps induced by the low energy protons in the radiation belts are slightly different than the traps created by the harder continuing particle damage."
    },
    "0606/astro-ph0606722_arXiv.txt": {
        "abstract": "{We report the results of  XMM-Newton and BeppoSAX observations of the radio and X--ray emitting star \\src, likely associated with the gamma-ray source 2CG 135+01 and recently detected also at TeV energies. The data include a long XMM-Newton pointing carried out in January 2005, which provides the deepest look ever obtained for this object in the 0.3--12 keV range.  During this observation the source flux decreased from a high level of $\\sim$13$\\times$10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$ to 4$\\times$10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$ within 2-3 hours. This flux range is the same seen in shorter and less sensitive observations carried out in the past, but the new data show for the first time that transitions between the two levels can occur on short time scales. The flux decrease was accompanied by a significant softening of the spectrum, which is well described by a power law with photon index changing from 1.62$\\pm{0.01}$ to 1.83$\\pm{0.01}$. A correlation between hardness and intensity is also found when comparing different short observations spanning almost 10 years and covering various orbital phases. \\src\\ was  detected in the 15--70 keV range with the PDS instrument in one of the BeppoSAX observations, providing evidence for variability also in the hard X--ray range. The X--ray spectra, discussed in the context of multiwavelength observations, place some interesting constraints on the properties and location of the high-energy emitting region. ",
        "introduction": "The peculiar radio source GT 0236+610, unambiguously associated with the B0 Ve star \\src, is unique in its high variability and periodic radio emission (Gregory \\& Taylor 1978). The radio outbursts show a periodicity of about 26.5 days (Taylor \\& Gregory 1982; Gregory et al., 1999) and a further modulation of both the outburst phase and outburst peak flux with a period of $\\sim$1600 days (Gregory et al. 1999), also displayed in the H$_{\\alpha}$ emission line (Zamanovet al. 1999).  A faint X--ray counterpart (F$_{\\rm 2-10 keV}$, 6$\\times$10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$) was identified by Bignami et al. (1981) with the Einstein satellite, and later observed with ROSAT (Goldoni \\& Mereghetti 1995; Taylor et al. 1996), ASCA (Leahy et al. 1997), RXTE (Harrison et al. 2000; Greiner \\& Rau 2001).  The hard power law X-ray spectrum extending up to 10~keV without breaks (photon index of $\\sim$1.7, absorbing column density, N$_{H}$, of $\\sim$5$\\times$10$^{21}$~cm$^{-2}$) is clearly inconsistent with low temperature plasma emission, thus excluding the Be-star as a main contribution to the X-ray emission.  The 26.5 days periodicity, reflecting the orbital motion of the system (Gregory \\& Taylor 1978, Taylor \\& Gregory 1982), has also been observed in the optical band (Hutchings \\& Crampton 1981; Mendelson \\& Mazeh 1989), in the infrared (Paredes et al. 1994), in soft X-rays (Paredes et al. 1997) and in the H$_{\\alpha}$ emission line (Zamanov et al. 1999).   Spectral line observations of the radio source give a distance of 2.0$\\pm{0.2}$~kpc (Frail \\& Hjellming 1991), implying an X--ray luminosity of 10$^{33}$~erg~s$^{-1}$. No periodic pulsations have been detected in its X--ray and radio emission. The low X-ray luminosity and the lack of iron line emission indicate that \\src\\ is not a classical accreting X-ray pulsar. Although this is not surprising in view of the strong radio emission (never seen from an X--ray pulsar, see e.g. Fender 2001), the low X-ray luminosity is remarkable in a source that can only be explained with the presence of a compact object.  \\src\\ is very likely associated with one of the most interesting unidentified gamma-ray sources near the plane of the Galaxy, the COS-B source 2CG~135+01 (Hermsen etal. 1977). Although the association between 2CG 135+01 and \\src\\ was initially weakened by the presence of another plausible counterpart in the COS-B error region (the quasar QSO 0241+622), the reduced error box determined by EGRET (Kniffen et al. 1997) is  compatible only with the position of \\src.   A recent detection of variable and likely periodic very high energy gamma-ray emission above 100~GeV has been reported by the MAGIC collaboration (Albert et al. 2006).   The strong radio outbursts suggest the presence of a compact star, although there is no direct evidence for the existence of a neutron star (no pulsations nor X-ray bursts). Two main classes of models have been proposed for \\src, involving a highly eccentric binary system with an orbital period of 26.5 days, composed of a compact star and the Be companion.  The first class of models suggests that the radio outbursts are produced by streams of relativistic particles originating in episodes of super-Eddington accretion onto the compact object (e.g., Taylor et al., 1992). The X--ray luminosity of \\src\\ is orders of magnitude lower than the Eddington limit. The association with the gamma-ray source 2CG 135+01 suggests that the bulk of the energy output is shifted from X--ray to $\\gamma$--ray wavelengths, but the actual mechanism responsible for this is not well understood in the context of the supercritical accretion model.   The second class of models assumes that \\src\\ contains a non-accreting young rapidly rotating neutron star (e.g., Maraschi \\& Treves, 1981; Tavani, 1994). The radio outbursts are produced by energetic electrons accelerated in the shock boundary between the relativistic wind of the young pulsar and the wind of the Be star. This scenario is supported by the shock-powered emission observed by ASCA and CGRO near the periastron in a possibly similar system, the binary PSR B1259-63 (Tavani et al., 1995), which is composed by a radio pulsar (with a period of 47~ms) orbiting a Be star (orbital period of 3.4~yr). In this case,the high-energy emission allows a sensible diagnostics of the shock emission region where the pulsar wind interacts with the circumstellar material originating from the surface of the massive star (Tavani, 1994). The modulation of the radio emission might be due to the time variable geometry of a `pulsar cavity' as a function of the orbital phase.  A third way of producing X-rays has been suggested by Campana et al. (1995), as accretion onto the pulsar magnetosphere when the stellar wind ram pressure exceeds the pulsar wind pressure, but is not large enough to penetrate the pulsar magnetospheric boundary. This situation can occur in \\src\\ for reasonable values of the neutron star magnetic field and spin period.  An object similar to \\src\\ has been discovered, the microquasar LS~5039 (Paredes et al. 2000, 2005) which is subluminous in the X--ray range (even more than \\src) and, if the proposed association with the EGRET source 3EG J1824-1514 is correct, also shows the same puzzling behavior,having L$_{\\gamma}$ $>$ L$_{\\rm X}$.  Therefore, \\src\\ and LS~5039 could be the first two examples of a new class of X--ray binaries with powerful ${\\gamma}$-ray emission (see also Aharonian et al. 2005 for a HESS detection of LS~5039 at energies above 250~GeV). In the case of LS~5039, an ejection process, probably fed by an accretion disk, is clearly proved by a VLBA map which shows bipolar jets emerging from a central core. Its high L$_{\\gamma}$ is tentatively explained by inverse Compton scattering. In \\src, VLBI observations (at 5~GHz) show a possible jet-like elongation at milliarcseconds scales, suggesting the presence of a one-sided collimated radio jet, as found in several other X--ray binaries. A lower limit of 0.4~c for the intrinsic velocity of the radio jet has been derived (Massi et al. 2001).  We report here the results of a deep XMM-Newton observation performed in January 2005, together with spectral analysis of archival \\sax\\ observations never reported in literature, and for completeness we studied also other five short archival \\xmm\\ observations recently analysed by Chernyakova et al. (2006). ",
        "conclusions": "We have reported here the deepest X--ray observation of \\src\\ ever carried out. A systematic analysis of other five \\xmm\\ and two \\sax\\ observations has also been performed, in order to investigate possible changes of the spectral parameters along the orbital phase (see Figs.~\\ref{fig:sum1} and ~\\ref{fig:sum2}) and to get a global picture at X-rays.  During the longest \\xmm\\ observation for the first time we found  evidence for a rapid (on timescales of few hours) change in flux and hardness ratio: the source flux decreased by a factor of $\\sim$3, together with a drop in hardness ratio (see Fig.~\\ref{fig:lc}) within about 1000~s.  This kind of variability (the source is harder when is brighter) is also visible when comparing different observations performed in different times and with different satellites (Fig~\\ref{fig:sum2}). Another interesting variability feature is the presence of a sort of ``flaring'' behaviour at the beginning of one of the \\sax\\ observations, but with no evidence for simultaneous hardness variations, mainly because of the low statistics.  The evolution of the source X--ray flux with the orbital phase shows X--ray emission along all the phases, although with a difference of a factor 3 in the flux level between two states, with a ``high'' state preferentially found in the phase range 0.4-1.0. We found no obvious correlation of any of the X--ray spectral parameters with the superorbital period.  The analysis of archival \\sax\\ observations allowed the first detection of the source with the PDS instrument (in the energy range 15--70 keV), with evidence for a change in flux between a high and low state also in the hard energy band.  The observations reported here demonstrate that the source is variable both at soft and hard X--rays, also at short timescales. Thus, a detailed study of the overall spectrum should be in principle performed with simultaneous observations at all wavelengths. We try here to discuss the broad band energy spectrum distinguishing at least two source states, a high and a low state, in order to get informations on the source energetics.  Figure~\\ref{fig:sed} shows our \\xmm\\ and \\sax\\ results in the multiwavelength spectral energy distribution (SED) of \\src. For radio, \\xmm\\ and \\sax\\ data we show both high and low state spectra. CGRO and MAGIC observations show low significance flux variations (peaking in the phase interval 0.4-0.7) correlated with those at X-ray energies. For these instruments we plot only average measurements of the datasets with positive detections. During the high state, the keV-MeV spectrum can be fit with a hardening power law ($\\Gamma$=1.6-1.5) up to the break in the COMPTEL energy range ($\\sim$1-10 MeV). Then the spectrum shows a shallow roll-off  (EGRET: $\\Gamma$=2.1; MAGIC: $\\Gamma$=2.6) extending up to TeV energies.  Two main models to account for the SED of \\src\\ have been proposed. The ``micro-quasar'' model involves streams of relativistic particles originating in episodes of super-Eddington accretion onto a compact star embedded in the mass outflow of the B-star (Taylor \\& Gregory 1984; Massi et al. 2004a, Massi 2004b). Alternatively \\src\\ might contain a non-accreting young pulsar in orbit around the mass-losing B-star (Maraschi \\& Treves 1981; Dubus 2006). In both cases, the high energy emission offers a diagnostics of the shock emission parameters in the jet or in the region where the hypothetical pulsar wind would interact with the circumstellar material.  The spectral index and break of the keV-MeV emission is consistent with that of diffusive shock acceleration models in the ``slow cooling\" regime (i.e. far from energetic equilibrium between particle injection and cooling, see e.g. Chevalier 2000). The shallow roll-off in the MeV-TeV energy range requires further model components in addition to the synchtrotron emission from shock-accelerated particles. However, the large error bars of the observations in this energy range can hardly constrain the several possibile contributions based both on leptonic mechanisms as Compton-scattered stellar radiation, synchrotron self-Compton, pair cascades (Dubus 2006; Dermer \\& B\\\"{o}ttcher 2006; Bosch-Ramon et al. 2006; Bednarek 2006), and on hadronic mechanisms as proton interactions, producing gamma-rays via neutral pion decay (Romero et al. 2003, 2005).  A clear signature of the SED is instead the spectral break in the 1-10 MeV range, possibly related to the maximum energy achievable from electrons in the shock acceleration accounting for synchrotron and Compton-scattered stellar radiation energy losses, the two major electron cooling mechanisms in the dense \\src\\ environment. The maximum Lorentz factor of shocked particle is obtained by solving the equation:  \\begin{equation} \\frac{1}{t_{\\rm acc}}=\\frac{1}{t_{\\rm synch}}+\\frac{1}{t_{\\rm comp}} \\end{equation} where $t_{\\rm acc}=\\xi t_{\\rm gyr}=m_{\\rm e} c \\gamma \\xi / e B $ is the electron acceleration time to reach a Lorentz factor $\\gamma$ in a magnetic field with mean intensity $B$, and $\\xi \\goe 1$ because the acceleration time in Fermi processes cannot be smaller than the gyration time, i.e. an electron gains at most a fraction of its energy when executing a single gyration in first-order shock or second-order stochastic Fermi processes (see e.g. Dermer \\& B\\\"{o}ttcher 2006). $t_{\\rm synch}=6 \\pi m_{\\rm e} c / \\sigma_{\\rm T} B^{2} \\gamma $ is the synchrotron cooling time for an electron with Lorentz factor $\\gamma$, and $t_{\\rm comp}$ accounts for inverse Compton energy losses. The solution of equation (1) is:  \\begin{equation} \\begin{array}{rll} \\gamma_{\\rm max}=\\sqrt{\\frac{6 \\pi e}{\\sigma_{\\rm T} \\xi B (1+ {t_{\\rm synch}}/{t_{\\rm comp}})}} \\\\ \\\\ =\\frac{1.2 \\times 10^{8}}{\\sqrt{\\xi B[{\\rm G}] (1+ {t_{\\rm synch}}/{t_{\\rm comp}})}} \\end{array} \\end{equation}  If $t_{\\rm synch} \\ll t_{\\rm comp}$ then the maximum synchrotron photon energy does not depend on the magnetic field B: $h \\nu_{\\rm max}=\\hbar e B \\delta \\gamma_{\\rm max}^{2} / m_{\\rm e} c \\simeq 167 \\delta / \\xi$ MeV, where $\\delta$ is the Doppler factor to be included for jet scenarios ($1<\\delta<2$, Chernyakova et al., 2006). Thus, neglecting inverse Compton losses, the spectral break frequency would be more than one order of magnitude higher than that observed.  Since photons from the Be star have a mean energy of $2.7k_{\\rm B}T \\simeq 6.5$ eV ($T_{\\rm s}=2.8 \\times 10^{4}$ K), the turning point between Thompson and Klein-Nishina inverse Compton regime is for $\\gamma_{\\rm KN}=m_{\\rm e}c^{2}/h \\nu \\simeq 8 \\times 10^{4}$. Even in case of severe inverse Compton energy losses, eq. 2 shows that $\\gamma_{\\rm max}$ cannot be lower than $\\gamma_{\\rm KN}$ for typical magnetic fields values expected in microquasar jets or pulsar-stellar winds shocks (up to several Gauss). Thus $t_{\\rm comp}=t_{\\rm KN}$ and considering electrons with Lorentz factor $\\gamma$ interacting with a blackbody surface radiation field (Blumenthal \\& Gould 1970): \\begin{equation} \\begin{array}{rll} t_{\\rm KN}=\\gamma \\left(\\frac{r}{R_{\\rm s}}\\right)^{2} \\frac{64 \\lambda_{\\rm C}^{3}}{8 \\pi^{3} c \\sigma_{\\rm T} \\Delta^{2}} \\left(\\frac{m_{\\rm e} c^{2}}{k_{\\rm B}T_{\\rm s}}\\right)^{2} \\left[\\ln\\left(0.552 \\Delta \\gamma \\frac{k_{\\rm B}T_{\\rm s}}{m_{\\rm e} c^{2}}\\right)\\right]^{-1} \\\\ \\\\ \\approx 10^{-30} \\gamma r[cm]^{2} ~~{\\rm s} \\end{array} \\end{equation} where $R_{\\rm s}=13.4R_{\\odot}$ is the Be star radius, $r$ is the distance between the star and the shock region and $\\Delta\\loe1$ accounts for the Lorentz boosting of the stellar radiation field into the co-moving frame of the emission region (in the microquasar scenario).  Substituting eq. 3 in eq.2 we obtain a general expression for the maximum Lorentz factor of the electrons accounting both for synchrotron and inverse Compton losses in the KN regime:   \\begin{equation} \\gamma_{\\rm max} \\simeq \\sqrt{\\frac{1.44 \\times 10^{16} B[{\\rm G}] - 8 \\times 10^{38} \\xi r[{\\rm cm}]^{-2}}{\\xi B[{\\rm G}]^{2}}} \\end{equation} which corresponds to a break energy of  \\begin{equation} h\\nu_{\\rm max}=\\delta \\left(\\frac{167}{\\xi}-\\frac{10^{25}}{B[{\\rm G}]r[{\\rm cm}]^{2}}\\right) ~~{\\rm MeV} \\end{equation}  The observed break energy at 1-10 MeV constrains the value of $Br^{2} \\approx 10^{23}$ G cm$^{2}$. Shock models associated to pulsar wind interacting with the circumstellar environments predict a standoff distance of the order of $r \\approx 10^{11}$ cm and magnetic fields of a few Gauss (Dubus 2006) in good agreement with the above constraint.  On the other hand, in accreting models, a shock originated in an extended jet ($r \\approx 10^{14}-10^{15}$ cm; Massi et al. 2001, 2004a) would imply unrealistic low values for the magnetic fields, unless other cooling processes apart from inverse Compton on stellar radiation could play a major role (e.g. SSC, Gupta \\& B\\\"{o}ttcher 2006). Shocks originating in the inner jet at a distance comparable with the compact object-star separation ($r \\approx 10^{12}$ cm) would imply more realistic values for the magnetic fields.  Future deeper and simultaneous observations in the MeV--GeV energy range with AGILE and GLAST satellites and in VHE gamma-rays with MAGIC and HESS telescopes will better assess the IC parameters (and/or hadronic mechanisms) and thus the actual emission region properties.      \\begin{figure*} \\hbox{\\hspace{-0.2cm} \\includegraphics[height=6.5cm,angle=0]{5933fig5a.ps} \\hspace{.0cm} \\includegraphics[height=6.5cm,angle=0]{5933fig5b.ps}} \\vspace{-0.2cm} \\caption[]{Summary of the dependence with the orbital phase of the X--ray spectral parameters of all the observations analysed here (six \\xmm\\ and two \\sax\\ observations) together with the ASCA results (published by Leahy et al. 1997) for completeness. The meaning of the symbols is as follows: {\\em thick circles} represent the longest \\xmm\\ observation, splitted into two (the ``first'' and the ``second'' spectra, with significantly different hardness ratio and flux); the other 5 \\xmm\\ observations are marked with  {\\em thick stars}, the 2 \\sax\\ observations are indicated with {\\em thick triangles} and the 2 ASCA observations (from Leahy et al. 1997)  with {\\em thin crosses}. Fluxes are corrected for the absorption and in units of 10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$. } \\label{fig:sum1} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[angle=0,height=6.5cm]{5933fig6.ps} \\caption{Unabsorbed flux (2--10 keV; in units of 10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$) versus power-law photon index, collecting all X--ray observations analysed here, together with the ASCA results (published by Leahy et al. 1997). The meaning of the symbols is the same as in Fig.~\\ref{fig:sum1}. The source appears to be softer when it is fainter. } \\label{fig:sum2} \\end{figure*} \\vspace{2cm}   \\begin{figure*} \\centering \\includegraphics[angle=90,height=12.5cm]{5933fig7.ps} \\caption{Broad band \\src\\ spectrum from radio wavelengths to TeV energies. Radio data for high and low states (open triangles), IR data (open stars) and OSSE data (filled triangles) are from Strickman et al. 1998. XMM-EPIC data (crosses) and SAX-PDS data (filled circles) both high and low states are from this work. COMPTEL data (filled squares) are from van Dijk et al. 1996. EGRET data (upside-down filled triangles) are from Kniffen et al. 1997. MAGIC data (filled stars) are from Albert et al. 2006. } \\label{fig:sed} \\end{figure*} \\vspace{2cm}"
    },
    "0606/astro-ph0606208_arXiv.txt": {
        "abstract": "I discuss the prospects of detecting the smallest dark matter bound structures present in the Milky Way by searching for the proper motion of $\\gamma$-ray sources in the upcoming GLAST all sky map. I show that for dark matter particle candidates that couple to photons the detection of at least one $\\gamma$-ray microhalo source with proper motion places a constraint on the couplings and mass of the dark matter particle. For SUSY dark matter, proper motion detection implies that the mass of the particle is less than 500 GeV and the kinetic decoupling temperature is in the range of [4-100] MeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606514_arXiv.txt": {
        "abstract": "We report the measurement results and compensation of the antenna elevation angle dependences of the Sub-millimeter Array (SMA) antenna characteristics. Without optimizing the subreflector (focus) positions as a function of the antenna elevation angle, antenna beam patterns show lopsided sidelobes, and antenna efficiencies show degradations. The sidelobe level increases and the antenna efficiencies decrease about $1\\%$ and a few \\%, respectively, for every $10^{\\circ}$ change in the elevation angle at the measured frequency of 237~GHz. We therefore obtained the optimized subreflector positions for X (azimuth), Y (elevation), and Z (radio optics) focus axes at various elevation angles for all the eight SMA antennas. The X axis position does not depend on the elevation angle. The Y and Z axes positions depend on the elevation angles, and are well fitted with a simple function for each axis with including a gravity term (cosine and sine of elevation, respectively). In the optimized subreflector positions, the antenna beam patterns show low level symmetric sidelobe of at most a few\\%, and the antenna efficiencies stay constant at any antenna elevation angles. Using one set of fitted functions for all antennas, the SMA is now operating with real-time focusing, and showing constant antenna characteristics at any given elevation angle. ",
        "introduction": "\\label{sect:intro}  %  The Submillimeter Array (SMA; left figure of Fig.~\\ref{fig:sma}) is the world's first dedicated submillimeter interferometer, and consists of eight 6~m antennas. It will cover the frequency range of 180--900~GHz with a 2~GHz bandwidth in both upper and lower side bands. An SMA antenna (right figure of Fig.~\\ref{fig:sma}) has a 6~m diameter main reflector, and its surface consists of 72 machined cast aluminum panels. These panels are supported by an open backup structure consists of carbon fiber tubes and steel nodes \\cite{ho04}. The secondary reflector (subreflector) is supported by a quadrupod. The main reflector and the quadrupod are deformed by gravity as the elevation angle of the telescope change \\cite{raf91}.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=7cm]{sma2-1crop.eps} \\includegraphics[height=7cm]{sma1-2.eps} \\end{tabular} \\end{center} \\caption[sma] {\\label{fig:sma} Pictures of the Submillimeter Array (SMA; left) and an SMA antenna (right). The SMA consists of eight 6~m diameter antennas. The main reflector consists of 72 aluminum panels supported by carbon fiber tubes and steel nodes.  The backup structure is usually covered, but in this figure, the bottom part is not covered yet, so that the backup structure is visible. The subreflector is supported by a quadrupod.} \\end{figure}  The gravitational deformation of the main reflector and the position change of the subreflector cause changes of the antenna beam shapes (beam patterns) and the focus positions, and affect the quality of the observational data, especially high frequency observations, extended object imagings, and mosaicing. It is therefore important to figure out the effect of the gravity on the SMA antenna reflectors, and to find out the ways to avoid or compensate these effects. We performed antenna beam pattern and efficiency measurements along various telescope elevation angles (Sect.~\\ref{sect:measfix}). We then measured the optimized focus positions at various telescope elevation angles, and modeled the focus curves with a simple function (Sect.~\\ref{sect:opt}). The optimized focus positions are verified with the measurements of beam pattern and efficiencies using the derived focus curves (Sect.~\\ref{sect:measopt}). Finally, we apply these results to the real-time subreflector position optimization (Sect.~\\ref{sect:real}). ",
        "conclusions": ""
    },
    "0606/astro-ph0606272_arXiv.txt": {
        "abstract": "{The galactic superluminal microquasar GRO~J1655-40 started a new outburst in February 2005, after seven years in quiescence, rising to a high/soft state in March 2005. In this paper we study the X-ray spectra during this rise. We observed GRO~J1655-40 with XMM-Newton, on 27 February 2005, in the low/hard state, and on three consecutive days in March 2005, during the rise of the source to its high/soft state. The EPIC-pn camera was used in the fast-read Burst mode to avoid photon pile-up. First, we contributed to the improvement of the calibration of the EPIC-pn, since the high flux received from the source required some refinements in the correction of the Charge Transfer Efficiency of the camera. Second, we find that the X-ray spectrum of GRO~J1655-40 is dominated in the high/soft state by the thermal emission from the accretion disk, with an inner radius of 13--14$\\left[\\rm{D}/3.2~\\rm{kpc}\\right]$~km and a maximum temperature of 1.3~keV. Two absorption lines are detected in the EPIC-pn spectra, at 6.7--6.8 and 7.8--8.0~keV, which can be identified either as blended Fe~XXV and Fe~XXVI K$\\alpha$ and K$\\beta$ lines, or as blueshifted Fe~XXV. We find no orbital dependence on the X-ray properties, which provides an upper limit for the inclination of the system of 73$\\degr$. The RGS spectrometers reveal interstellar absorption features at 17.2~$\\rm\\AA$, 17.5~$\\rm\\AA$ (Fe L edges) and 23.54~$\\rm\\AA$ (OI K$\\alpha$). Finally, while checking the interstellar origin of the OI line,\twe find a general correlation of the OI K$\\alpha$ line equivalent width with the hydrogen column density using several sources available in the literature. ",
        "introduction": "Galactic microquasars are accreting binary systems ejecting jets at relativistic velocities. Both black holes as well as neutron stars have been identified as the compact, accreting object. The analogy of these systems to extragalactic quasars and Active Galactic Nuclei (AGN) make them excellent laboratories for the study of the physics involved in accretion disks and ejection of relativistic jets associated with accreting black holes (Mirabel~et~al.~\\cite{mir92}). The physics ruling black hole systems is essentially the same in galactic microquasars, hosting stellar black holes, and in AGN, with supermassive black holes, but the differences in time scales makes the study of some aspects in microquasars easier than in AGN. The characteristic timescale for the flow of matter onto a black hole is proportional to its mass, being of order of minutes in a microquasar of a few solar masses, but of thousands of years in a massive black hole of $10^{9}\\rm M_{\\odot}$. In addition,  thanks to their proximity, microquasar jets have proper motions in the plane of sky with velocities about a thousand times faster than AGN, and two-sided jets can be observed.   The microquasar GRO\\,J1655-40 (X-ray nova Sco\\,1994, Zhang~et~al.~\\cite{zha94}) was the second superluminal source discovered in our Galaxy, after GRS\\,1915+105 (Mirabel~\\&~Rodriguez~\\cite{mir94}). The two sources may also be peculiar in the sense that they show evidence suggesting that both systems contain a maximally spinning black hole (Zhang~et~al.~\\cite{zha97}). Radio images of GRO\\,J1655-40 showed twin jets with apparent superluminal motion (Tingay~et~al.~\\cite{tin95}) moving in opposite directions at 0.92c, and the distance was determined to be 3.2$\\pm$0.2~kpc (Hjellming~\\&~Rupen~\\cite{hje95}). Using the dust scattering halo observed by ROSAT, Greiner~et~al.~(\\cite{gre95}) determined a distance of 3~kpc, compatible with the previous determination. Optical observations in 1996 provided an inclination angle of 69.5$\\pm$0.08$\\degr$, a radial velocity semiamplitude of 228.2$\\pm$2.2~km~s$^{-1}$ and a dynamical mass of the primary component of 7.02$\\pm$0.22~M$_{\\odot}$ (Orosz~\\&~Bailyn~\\cite{oro97}), indicating that it is a black hole. Hubble Space Telescope observations showed that the system moves in an eccentric orbit with a runaway space velocity (i.e., with respect to the Galactic rotation corresponding to its position in the Galactic plane) of 112$\\pm$18~km~s$^{-1}$ (Mirabel~et~al.~\\cite{mir02}), which together with abundance anomalies found in the atmosphere of the donor star (Israelian~et~al.~\\cite{isr99}) provide evidence for a supernova origin of the black hole in GRO~J1655-40.  ASCA observations of GRO~J1655-40 in August 1994 and August 1995 provided the first detection of absorption lines in an accretion powered source (Ueda~et~al.~\\cite{ued98}). The energy of the lines was found to depend on the X-ray intensity, being 6.95~keV (Fe~XXVI K$\\alpha$) at 2.2 Crab, and 6.63 and 7.66~keV (Fe~XXV K$\\alpha$ and K$\\beta$) at 0.27--0.57 Crab, revealing the presence of a highly ionized absorber. Similar absorption features were also detected for GRS 1915+105 (Kotani~et~al.~\\cite{kot00}). During the outburst in 1997, ASCA observed GRO~J1655-40 between 25 and 28 February, when it was at 1.1~Crab, and revealed an absorption feature at 6.8~keV, interpreted as the blend of the two resonance K$\\alpha$ lines of Fe~XXVI and Fe~XXV (Yamaoka~et~al.~\\cite{yam01}). Simultaneously, RXTE observations performed on 26 February 1997 showed emission lines at 5.85 and 7.32~keV, interpreted as the first red- and blueshifted Fe~K$\\alpha$ disk line in a Galactic source (Balucinska-Church~\\&~Church~\\cite{bal00}), indicating the presence of significantly ionized material in a region of the disk at $\\sim10$ Schwarzschild radii.   The first observed X-ray outbursts of GRO~J1655-40 occurred in 1994--1995 (Zhang~et~al.~\\cite{zha94}, Harmon~et~al.~\\cite{har95}), and a 16-month-long outburst began on 25 April 1996 (Sobczak~et~al.~\\cite{sob99}). After 7 years of inactivity, GRO~J1655-40 left quiescence again on 17 February 2005 (Markwardt~\\&~Swank~\\cite{mar05}). The X-ray evolution was followed with RXTE/ASM (Homan~et~al.~\\cite{hometal05}) and Swift (Brocksopp~et~al.~\\cite{bro06}). GRO~J1655-40 entered in February a low/hard state (Homan~\\cite{hom05}) until it experienced a first outburst in March, moving into a high/soft state and reaching $\\sim2$ Crab (Fig.~\\ref{fig_lc}). The decay of the first outburst was followed by a month and a half of increasing X-ray flux and finally a strong outburst in very high state in May 2005, when the source reached more than 4 Crab.  Here we present XMM-Newton observations performed in February-March 2005, aimed at obtaining detailed spectroscopy during the rise to the high/soft state. We use the canonical distance of 3.2~kpc throughout (see note added in proof, however).      \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1_xte_lc.eps}} \\caption{RXTE/ASM light curve of GRO~J1655-40, with dates of our XMM-Newton observations, and ASM hardness ratios: HR1, ratio of the ASM count rates (3--5~keV)/(1.5--3~keV), and HR2, (5--12~keV)/(3--5~keV). Data from the quick-look results provided by the ASM/RXTE team. \\label{fig_lc}} \\end{figure} ",
        "conclusions": "\\subsection{The dust halo}  The effect of the scattering dust halo is clearly detected in the soft X-ray spectra. The scattering fraction at 1~keV is changing during the three March observations, decreasing from 18.6($\\pm$0.2)\\% to 16.9($\\pm$0.3)\\% in the first 24 hours, and then increasing again to 18.4($\\pm$0.3)\\% during the next day. This temporal variation is consistent with a constant dust halo but a variable X-ray source flux: the time needed for the soft photons originating on the disk to travel to the scattering site far away from the source causes a delay in the flux variations of the dust halo with respect to the variations of the source. As a result, just after an increase of the disk temperature and thus of the source soft spectrum, the fraction of scattered photons with respect to the source is temporally smaller (during the second March observation) until the increasing source photons reach the outer part of the halo and their scattered photons also reach the observer, increasing again the scattering fraction (third March observation). This effect has been known for some time. It was also observed by ROSAT and used to estimate the distance to GRO~J1655-40 (Greiner~et~al.~\\cite{gre95}).    \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig3_pn_new.eps}} \\caption{EPIC pn residuals of GRO~J1655-40 spectra of our three March 2005 observations at zero velocity, after fitting an absorbed multi-temperature disk model.} \\label{fig_pn} \\end{figure}    \\subsection{Continuum}  The maximum disk temperature observed in the March observations, increasing from 1.25 to 1.35~keV, is consistent with the values commonly observed in black hole binaries in high/soft state (see for instance McClintock~\\&~Remillard~\\cite{mcc06}), and with the Swift observations of the same period (Brocksopp~et~al.~\\cite{bro06}). The inner radius, \\mbox{$\\rm R_{in}\\sim13-14\\left[\\rm{D}/3.2\\rm{kpc}\\right]$~km}, is only a little larger than the gravitational radius of the black hole, R$_{\\rm g}=\\rm G\\rm M/ \\rm c^2=10$~km, consistent with a rapidly spinning black hole (Zhang et al. \\cite{zha97}).  From the fit parameters of the MCD model the luminosity of the disk can be calculated as $\\rm L_x=4\\pi \\rm R_{in}^2\\sigma \\rm T_{in}^ 4$ (Makishima~et~al.~\\cite{mak86}). With our values, the accretion disk luminosity is \\mbox{$6-7\\times10^{37}\\left[\\rm{D}/3.2\\rm{kpc}\\right]^2$~erg~$\\rm s^{-1}$}.  With the accretion disk luminosity, the accretion rate can be calculated as $2\\rm L_x \\rm R_{in}/\\rm G \\rm M \\rm g^2$  (Makishima~et~al.~\\cite{mak86}), where $\\rm g=(1-\\rm R_{g}/\\rm R_{in})^{1/2}$ is a correction for the general relativity, R$_{\\rm g}$ is the gravitational radius and M is the mass of the black hole. From the accretion disk luminosity, the disk inner radius and the black hole mass of GRO~J1655-40, we obtain an accretion rate during the high/soft state of $\\sim10^{-8}~\\rm M_{\\odot}~\\rm{yr}^{-1}$.  The rise of the disk blackbody component during the high/soft state in March is reflected in the power law index of the corona, which increases from 1.48$\\pm$0.01 in February 27 to values higher than 2 in the March observations. With the rise of the thermal disk emission, the soft luminosity increases and the electrons in the corona are cooled more efficiently, resulting in a softer spectrum of the comptonized photons, i.e., a steeper power law.   \\subsection{Interstellar absorption}   The interstellar OI~K$\\alpha$ line is clearly detected at 23.5~$\\rm \\AA$ in our three March 2005 RGS spectra. This line was also found in the RGS spectra of other sources, and distinguished from the instrumental components around the interstellar oxygen edge, at 23.05 and 23.35 $\\rm \\AA$ (de Vries~et~al.~\\cite{cor03}). Juett~et~al.~(\\cite{jue04}) detected the same line in the high resolution spectra of seven X-ray binaries using the Chandra/HETGS, as part of their study of the structure of the oxygen absorption edge caused by the interstellar medium. More recently, it has also been detected in the Cyg~X-2 spectrum (Costantini~et~al.~\\cite{cos05}), in the observations that have allowed the first spatially resolved spectroscopic study of a scattering dust halo, as well as in the absorption towards LMC~X-3 (Wang~et~al.~\\cite{wan05}) and the recently discovered black hole candidate XTE~J1817-330 (Sala~\\&~Greiner~\\cite{sal06}).  To check the interstellar origin of the OI~K$\\alpha$ line observed in the RGS spectra of GRO~J1655-40, we have looked for a correlation of the equivalent width of the line and the hydrogen column density, using the above mentioned works (Fig.~\\ref{fig_oh}). We find that the equivalent width of the OI~K$\\alpha$ line, EW, is indeed correlated with the column density, N$_{\\rm H}$, as EW(eV)=$0.5(\\pm0.1)+2.8(\\pm0.3)\\times10^{22}\\rm N_{\\rm H}(\\rm{cm}^{-2})$. The values observed for GRO~J1655-40 fall well within this correlation for the ISM. We find that an alternative correlation with zero regression constant could be EW(eV)=$3.6(\\pm1.6)\\times10^{22}\\rm N_{\\rm H}(\\rm{cm}^{-2})$, but this would underestimate the equivalent width for the sources with hydrogen column density less than $5\\times10^{21}\\rm{cm}^{-2}$.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig5_new.eps}} \\caption{RGS1 (black) and RGS2 (grey, contributing only below 20$\\rm \\AA$ and above 24$\\rm \\AA$.) around the oxygen edge in the first order spectra of the 14th March 2005 observation, fitted with the new TBabs model.} \\label{fig_oxygen} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig4_new.eps}} \\caption{First and second order (with lower count rate) spectra of GRO~J1655-40 in the range 16--19$\\rm \\AA$, obtained with RGS1 (black) and RGS2 (grey) on 14 March 2005, showing the ISM Fe L edges fitted with the new TBabs model.} \\label{fig_rgs} \\end{figure}     \\subsection{Orbital dependence and inclination of the binary system}  Since the orbital period of GRO~J1655-40 is $\\sim2.6$ days and our three March observations were taken with one day intervals, we are covering approximately one whole orbital cycle. We can extrapolate the Orosz~\\&~Bailyn~(\\cite{oro97}) ephemeris, to obtain the orbital phase of our observations (error is 0.08 in all cases): 0.31--0.40 (27 February), 0.17--0.24 (14 March), 0.55--0.62 (15 March) and 0.93--0.00 (16 March).  Given the inclination of the system (between $\\sim70\\degr$, Orosz~\\&~Bailyn~\\cite{oro97}, van~der~Hooft~et~al.~\\cite{van98}; and $\\sim85\\degr$, Hjellming~\\&~Rupen~\\cite{hje95}), an orbital phase close to zero corresponds to having the donor star situated closer to the observer, while in phase 0.5 the disk would be in front of the secondary star. This means that if any of the observed spectral features were arising from the illuminated face of the secondary star, it would have its maximum at phase 0.5 and would not be present at phase zero. This is not the case in our observations. Another possible orbital effect would be that the disk emission were absorbed by the stellar wind of the secondary star, which would produce increased absorptions close to phase zero, i.e. on 16 March, which is also not the case.  Finally, neither dips nor eclipses were observed in the disk thermal emission at any of the phases. Given the size of the donor star, $\\sim5\\rm{R}_{\\odot}$ (Orosz~\\&~Bailyn~\\cite{oro97}), the binary separation ($1.17\\times10^{12}$cm), and assuming that the soft X-ray emission originates in the central 200\\,000 km of the disk (see below, section~\\ref{sect_fexxv}), the duration of a possible eclipse of the soft X-ray emission by the donor star would last more than 2.5 hours (i.e., a change in the orbital phase of 0.04) close to phase zero. This should have been clearly visible during our 16 March observation.  The fact that no orbital modulation is observed in the X-ray spectra, provides an upper limit for the inclination of the system. With the parameters mentioned above for the sizes of the system, the inclination must be smaller than 73$\\degr$ (for the innermost 200\\,000~km of the disk surface to be visible in all orbital phases). This limit is in agreement with the inclination determined from the optical light-curve by Orosz~\\&~Bailyn~(\\cite{oro97}), $69.\\!\\!^\\circ5\\pm0.\\!\\!^\\circ8$, and by van~der~Hooft~et~al.~(\\cite{van98}), $67.\\!\\!^\\circ2\\pm3.\\!\\!^\\circ5$, but incompatible with the inclination inferred from the radio jets, $\\sim84\\degr$ (Hjellming~\\&~Rupen~\\cite{hje95}). This points, as suggested by Orosz~\\&~Bailyn~(\\cite{oro97}), to an inclination of the jet axis of about 15$\\degr$ with respect to the normal to the orbital plane. Alternatively, non-symmetric jet ejections could have lead to a wrong inclination inference (Hjellming~\\&~Rupen~\\cite{hje95}).   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{fig6_oh.eps}} \\caption{OI K-$\\alpha$ line equivalent width versus hydrogen column density for several galactic sources. The dashed line indicates an approximate linear correlation, EW(eV)=$0.5+2.8\\times10^{-22}\\rm N_{\\rm H}(\\rm{cm}^{-2})$.} \\label{fig_oh} \\end{figure}    \\subsection{\\label{sect_fexxv}The Fe~XXV/Fe~XXVI absorber}  Clear highly ionized Fe absorption lines are detected in the EPIC-pn spectra of the three March observations in the high/soft state. It is worth noting that Ueda~et~al.~(\\cite{ued98}) found simultaneously Fe~XXV absorption features and an iron edge, while they detected no edge when absorption in the iron K band corresponded to Fe~XXVI.  Assuming that the absorbing plasma is photoionized, the ionization state of the observed elements provides information on the conditions and location of the absorber. The presence of He-like Fe ions indicates an ionization parameter $\\xi=$L/nr$^{2}\\sim10^{3}$~erg~cm~s$^{-1}$ (Kallman~et~al.~\\cite{kal96}). The lack of lower ionization absorption implies that the absorber cannot have a large extent.  The equivalent width of the Fe lines is increasing between the first and the last March observations. This may be pointing out either an increase in the total column density, or in the ionization state of the gas.   Ueda~et~al.~(\\cite{ued98}) determined the iron column density of the plasma to be $10^{19}-10^{20}$~cm$^{-2}$ from the observed equivalent width and the curve of growth of the Fe~XXV~K$\\alpha$ line, which relates the expected equivalent width to the iron column density. Using their curve of growth and assuming the detected features are only Fe~XXV, our equivalent width of the Fe~XXV~K$\\alpha$ line (between 50 and 160~eV) corresponds to a column density of $10^{19}$ and $5\\times10^{20}$~cm$^{-2}$, which is similar to the values found by Ueda~et~al.~(\\cite{ued98}). Assuming cosmic abundances, this corresponds to a hydrogen column density in the range $2\\times10^{23}-10^{25}$ cm$^{-2}$. Since the observed absorption lines may be a blend of Fe~XXV and Fe~XXVI, these must be taken as upper limits. In addition, for a hydrogen column density larger than 10$^{24}$ cm$^{-2}$, the line absorber would be optically thick to Thomson scattering (Ueda~et~al. 1998). We consider this to be unlikely and, taking into account the errors in the equivalent widths, the column density could be then in the range $2\\times10^{23}- 10^{24}$~cm$^{-2}$ in all three observations.  With the photoionization parameter, $\\xi\\sim10^{3}$~cm~s$^{-1}$, and the flux detected above $\\sim9$~keV (X-rays photoionizing He-like iron ions), $\\rm L_{\\ge 9\\rm{keV}}\\sim5\\times10^{36}$~erg~s$^{-1}$ (for a distance of 3.2~kpc), the column density indicates a distance to the central source between 50\\,000 and 200\\,000~km. Assuming that the disk radius is 70\\% of the Roche lobe radius, i.e., $4\\times10^{6}$~km, the Fe~XXV absorber extends to less than 5\\% of the disk surface.  In the case that the features correspond to blueshifted Fe~XXV lines, the error in the energy determination is too large as to reach any strong conclusion about its blueshift. Nevertheless, assuming the wind velocity is constant, the common range of blueshifts of the three March observations is reduced to 2600--4500~km~s$^{-1}$. From the parameters derived above for the Fe~XXV absorber, and assuming constant density and spherical symmetry for the expanding wind, the mass loss rate $4\\pi \\rm r^2 \\rho\\rm v \\rm b$ would be in the range $(2-13)\\times10^{-7}\\rm b ~\\rm M_{\\odot} \\rm {yr}^{-1}$, with b being the filling factor. This upper limit is a factor 20--130 larger than the accretion rate, which would indicate a highly non-stationary situation. However, there is considerable uncertainty in the wind mass loss, since both the wind density and wind velocity are likely a function of radius and the wind may be conical rather than spherically symmetric. Exploring these effects goes beyond the scope of this paper."
    },
    "0606/astro-ph0606334_arXiv.txt": {
        "abstract": "Recent surveys have detected Ly$\\alpha$ emission from $z=4.5-6.5$ at luminosities as low as  $10^{41}$ erg s$^{-1}$. There is good evidence that low numbers of AGN are among observed faint Ly$\\alpha$ emitters. Combining these observations with an empirical relation between the intrinsic Ly$\\alpha$ and B-band luminosities of AGN, we obtain an upper limit on the number density of AGN with absolute magnitudes $M_B \\in [-16,-19]$ at $z=4.5-6.5$. These AGN are up to two orders of magnitude fainter than those discovered in the Chandra Deep Field, resulting in the faintest observational constraints to date at these redshifts. At $z=4.5$, the powerlaw slope of the very faint end of the luminosity function of AGN is shallower than the slope observed at lower redshifts, $\\beta_l <1.6$, at the 98\\% confidence level. In fact, we find marginal evidence that the luminosity function rises with luminosity, corresponding to a powerlaw slope $\\beta_l <0 $, at magnitudes fainter than $M_B\\sim -20$ (75\\% confidence level). These results suggest either that accretion onto lower mass black holes is less efficient than onto their more massive counterparts, or that the number of black holes powering AGN with $M_B\\gsim-20$ is lower than expected from the $M_{\\rm BH}-\\sigma$ relation by one-two orders of magnitude. Extrapolating from reverberation-mapping studies suggests that these black holes would have $M_{\\rm BH}=10^6-10^7 M_{\\odot}$. To facilitate the identification of AGN among observed Ly$\\alpha$ emitters, we derive observational properties of faint AGN in the Ly$\\alpha$ line, as well as in the X-ray and optical bands. ",
        "introduction": "\\label{sec:intro} The quasar luminosity function (QLF) describes the space density of Active Galactic Nuclei (AGN) as a function of luminosity and redshift. The QLF encodes information on quantities like the black hole number density per unit mass, and the gas accretion efficiency. It therefore constrains physical models of AGN and of super massive black hole formation. The optical (B-band) QLF at $z \\lsim 4$ has been determined accurately for luminosities\\footnote{One solar B-Band luminosity, denoted by $L_{B,\\odot}$, corresponds to $4 \\times 10^{32}$ ergs s$^{-1}$.} exceeding $L_B \\gsim 10^{11}L_{B,\\odot}$ from the 2dF quasar survey (which corresponds to absolute magnitudes $M_B<-22$, Boyle et al. 2000; Croom et al. 2004). At higher redshifts the optical luminosity function of luminous quasars has been determined from quasars in the Sloan Digital Sky Survey \\citep{Fan01} at $M_B\\sim-27$. In addition, deep {\\it Chandra} and {\\it XMM} imaging has constrained the X-ray QLF at X-Ray luminosities as low as $L_X=10^{42}-10^{44}$ ergs s$^{-1}$ \\citep{Barger03,Cowie03,Hasinger05}. Following \\citet{Haiman98}, \\citet{Wyithe03} have used the X-Ray QLF to constrain the B-band QLF down to fainter optical luminosities ($M_B \\sim -22$) at $z \\gsim 4$.  At these redshifts, no observational constraints exist on the optical QLF at fainter luminosities. However, the details of the existence and evolution of low luminosity AGN is crucial for our understanding of the growth of low mass black holes, and of their role in the formation of super massive black holes.  In this paper we demonstrate how existing Ly$\\alpha$ surveys may be used to constrain the B-Band QLF at absolute magnitudes as low as $M_B =-16$. Existing wide-field narrow-band surveys \\citep[e.g.][]{Rhoads00} are optimised to detect Ly$\\alpha$ line emission from high-redshift galaxies, and in order to maximise their detection rate, deeply image fields as large as 1 deg$^{2}$ on the sky \\citep[e.g.][]{Taniguchi05}. The combination of wide field and deep images, together with concentration on a strong emission line, rather than continuum, allows these wide-field narrow-band surveys to put stringent constraints on the number density of AGN with $M_B\\sim -19$. Constraints on even fainter AGN are derived from deep spectroscopic surveys of regions around intermediate redshift clusters of galaxies, which offer an ultra-deep view into the high redshift universe (z=4.5-6.7) through strong gravitational lensing (with magnification factors of 10-1000, Santos et al, 2004).  Throughout this paper, the terms AGN and quasar are interchangeable and refer to broad lined active galactic nuclei, also known as 'Type I' AGN. The outline of this paper is as follows: In \\S~\\ref{sec:lum} we relate the Ly$\\alpha$ luminosity of a quasar to its B-band luminosity.  In \\S~\\ref{sec:lya} we summarise constraints on the number density of Ly$\\alpha$ emitters at high redshift, and discuss the abundance of quasars among these sources. We show how this existing data constrains the faint end of the quasar luminosity function.  In \\S~\\ref{sec:phys} we calculate the observable Ly$\\alpha$ properties of AGN and use these to physically interpret our results. We also model the appearance of faint AGN in the optical and X-ray bands.  Finally, in \\S~\\ref{sec:discuss} we discuss the possible cosmological implications of our work, before presenting our conclusions in \\S~\\ref{sec:conclusion}. We use the {\\it WMAP} cosmological parameters: $(\\Omega_m, \\Omega_{\\Lambda}, \\Omega_b, h, Y_{\\rm He})$ =$(0.3,0.7,0.044,0.7,0.24)$ \\citep{Spergel03} throughout the paper. ",
        "conclusions": "\\label{sec:conclusion}  Recent Ly$\\alpha$ surveys have detected Ly$\\alpha$ emitting objects from redshifts as high as $z=6.5$, and at luminosities as low as $10^{41}$ erg s$^{-1}$ \\citep[e.g.][]{Santos04lens}. No evidence of AGN activity exists among these several hundred Ly$\\alpha$ emitters \\citep{Dawson04,Wang04,Taniguchi05}. Wide field Ly$\\alpha$ surveys are designed to deeply image wide fields on the sky, yielding survey volumes in a narrow shell of redshift space as large as $10^5-10^6$ Mpc$^3$, while deep spectroscopic surveys of gravitationally lensed regions probe deeper into smaller volumes. The absence of AGN within these fields can place a tight upper limit on the number density of AGN with Ly$\\alpha$ luminosities exceeding the surveys detection thresholds, $L_{{\\rm Ly}\\alpha,{\\rm c}}$.  In \\S~\\ref{sec:lum} we have shown empirically that the Ly$\\alpha$ luminosities of AGN equal their B-band luminosities to within a factor of a few (Eq.~\\ref{eq:key}). As a result, deep Ly$\\alpha$ surveys can be used to obtain upper limits on the number density of AGN with B-Band luminosities exceeding $L_{\\rm B,min}\\sim 1.4 L_{{\\rm Ly}\\alpha,{\\rm c}}$. When expressed in B-band solar luminosities, $L_{\\rm B,min}$ is $\\sim10^{8.5}L_{\\rm B,\\odot}$ (which corresponds to an absolute magnitude of $M_B=-16$). In \\S~\\ref{sec:lyaagn} we demonstrated that such AGN are expected to be powered by black holes with masses in the range $M_{\\rm BH}=10^4-10^7 M_{\\odot}$ at $z=4.5-6.5$.  We derive upper limits on AGN number densities and constrain the quasar B-band luminosity function $\\Psi(L_{\\rm B,min},z)$ at a luminosity $L_{\\rm B,min}$.  The non-detection of AGN among $z=4.5-6.5$ LAEs rules out the model predictions by Wyithe \\& Loeb (2003), which succesfully reproduce the brighter end of the observed quasar luminosity function at $z=2-6$, at $\\geq 95\\%$ confidence levels at all redshifts. At $z=4.5$, we find that $\\partial$log$\\Psi/\\partial$log$L_B \\geq -1.6$, the value observed at lower redshifts, for log$[L_B/L_{B,\\odot}] \\lsim 11$ at the $98\\%$ confidence level (Fig~\\ref{fig:dPdB}). We find marginal evidence that at these luminosities, the luminosity function rises with luminosity, corresponding to a powerlaw slope $>0$ (75\\% confidence level). In other words, the QLF may increase with $L_B$ at these faint luminosities, in contrast to observations of more luminous AGN. These results represent the faintest observational constraints on the quasar luminosity function at these redshifts to date.  We have found that models of the quasar luminosity function which are successful in reproducing the bright end of the quasar luminosity function predict more AGN to be present than are observed, by up to two orders of magnitude (Fig.~\\ref{fig:lumfunc}) at $z\\sim4.5$. These results imply either that accretion onto lower mass black holes is less efficient than onto their more massive counterparts, or that the number of black holes powering AGN with $M_B\\gsim-20$ is lower than expected from the $M_{\\rm BH}-\\sigma$ relation by one-two orders of magnitude. Extrapolating from reverberation-mapping studies suggests that these black holes would have $M_{\\rm BH}=10^6-10^7 M_{\\odot}$.  Our work has demonstrated the effectiveness of Ly$\\alpha$ surveys in constraining the faint end of the quasar B-band luminosity function. Deeper and larger surveys will allow for a better determination of its slope, and whether indeed the quasar luminosity function rises with luminosity for $M_B\\gsim-20$ at $z=4.5$, and at other redshifts. These constraints will offer new insights on the growth of low mass black holes and their relation to the known super massive black holes. To help identify AGN among observed Ly$\\alpha$ emitters, we have modeled the observable properties of the Ly$\\alpha$ line for high redshift, faint AGN. Using the empirical Kaspi-relations, we estimate that the observable Ly$\\alpha$ line widths (Half Width at Half Maximum) of faint AGN will be $\\sim 1500$ km s$^{-1}$. For AGN embedded in a neutral medium the peak of the Ly$\\alpha$ line is redshifted up to $2000$ km s$^{-1}$ relative to the true line center (\\S~\\ref{sec:lyaagn}). To facilitate the identification of these faint AGN, we have estimated their observable properties in the visible (\\S~\\ref{sec:broadagn}, Fig~\\ref{fig:mab}) and X-Ray bands (\\S~\\ref{sec:Xray}). We caution that selection criteria used in Ly$\\alpha$ surveys to select candidate high redshift Ly$\\alpha$ emitters currently introduce a bias against detecting AGN with log$[L_B/L_{B,\\odot}] \\gsim 11.3$ (\\S~\\ref{sec:broadagn}), corresponding to the faintest AGN identified in the Chandra Deep Fields.  {\\bf Acknowledgments} Our research is supported by the Australian Research Council. The authors would like to thank Colin Norman for helpful discussions and Zolt\\'an Haiman for useful comments on an earlier version of the manuscript.  \\newcommand{\\noopsort}[1]{}"
    },
    "0606/astro-ph0606428_arXiv.txt": {
        "abstract": "A variability study of the young cluster IC 348 at Van Vleck Observatory has been extended to a total of seven years. Twelve new periodic stars have been found in the last two years, bringing the total discovered by this program to 40. In addition, we confirm 16 of the periods reported by others and resolve some discrepancies. The total number of known rotation periods in the cluster, from all studies has now reached 70. This is sufficient to demonstrate that the parent population of K5-M2 stars is rotationally indistinguishable from that in the Orion Nebula Cluster even though their radii are 20\\% smaller and they would be expected to spin about twice as fast if angular momentum were conserved. The median radius and, therefore, inferred age of the IC 348 stars actually closely matches that of NGC 2264, but the stars spin significantly more slowly. This suggests that another factor besides mass and age plays a role in establishing the rotation properties within a cluster and we suggest that it is environment. If disk locking were to persist for longer times in less harsh environments, because the disks themselves persist for longer times, it could explain the generally slower rotation rates observed for stars in this cluster, whose earliest type star is of class B5. We have also obtained radial velocities, the first for PMS stars in IC348, and \\emph{v} sin \\emph{i} measurements for 30 cluster stars to assist in the study of rotation and as an independent check on stellar radii. Several unusual variable stars are discussed; in some or all cases their behavior may be linked to occultations by circumstellar material.  A strong correlation exists between the range of photometric variability and the slope of the spectral energy distribution in the infrared. Nineteen of the 21 stars with I ranges exceeding 0.4 mag show infrared evidence for circumstellar disks. ",
        "introduction": "IC 348 is a compact, young cluster in the Perseus clouds at a distance of about 300 pc \\citep{c93}. It contains $\\sim 200$ members, with spectral types ranging from B5 to late M and about two dozen probable brown dwarfs that are also likely to be cluster members \\citep{l03,llme05}. It is especially important to star formation studies because it is relatively close to us, compact on the sky, well studied by a variety of methods, and has a nearly complete census of members. It is also intermediate in nature between the denser, more massive clusters, such as the Orion Nebula Cluster (ONC), and the looser associations, such as the Taurus T association. For a recent summary of work on the cluster, see the article by \\citet{h06} in {\\it The Handbook of Low Mass Star Forming Regions}.  Photometric monitoring of young clusters has led to the determination of accurate rotation periods for about 2,000 pre-main sequence stars (see, for example, the review by \\citet{hems06}). The best-studied cases are the ONC and NGC 2264. \\citet{hm05} have recently reviewed the data and shown that the evolution of rotation for stars in the 0.4-1.2 solar mass range may be understood in terms of two phenomena. About half of the stars of Orion age (the half which are already the more rapid rotators) spin up with essentially no loss of angular momentum all the way to the ZAMS. The other half continue to be slowed by interaction with circumstellar disks for times of up to 5 Myr and end up as very slow rotators on the ZAMS. This picture is based on only five clusters, two containing pre-main sequence stars and three containing ZAMS stars. Obviously it will be important to test it by obtaining data for more clusters and, in particular, clusters with different properties (such as stellar density and environment). IC 348 is a particularly useful cluster in this regard for the reasons discussed in the first paragraph.  The cluster has been photometrically monitored at Van Vleck Observatory (VVO) on the campus of Wesleyan University for seven years. Two previous papers in this series have reported on the progress of this program after one and five years, respectively \\citep{hmw00, chw04}. \\citet{lit05} have recently published results from a short (17 day) but intensive photometric monitoring program done with a larger telescope and going deeper than our work, and \\citet{kkb05} have even more recently published the results of a six month monitoring program designed to search for variability in stars with X-ray counterparts. In this paper, we present data from an additional two years of photometric monitoring at VVO as well as a complementary study of  radial velocities and \\emph{v} sin \\emph{i} measurements for 30 cluster members. We combine the results of all of the studies to discuss the rotation period distribution in the cluster and its significance for our understanding of angular momentum evolution in solar-like stars. A very recent infrared study of the cluster with {\\it Spitzer} by \\citet{lm06} allows us to correlate rotation and variability properties with the presence or absence of evidence for a disk. In addition, we discuss some unusual variable stars. ",
        "conclusions": "This study has confirmed several beliefs about T Tauri star variability and rotation based on other clusters, identified some interesting stars worthy of continued observation and provided a challenge to the simplest picture of rotational evolution of low mass stars. In particular, the new data confirm that WTTS variability arises primarily from the rotation of a spotted photosphere and that the large amplitude (greater than 0.4-0.5 mag in I) irregular variability of CTTS is linked to the presence of an accretion disk. This study also verifies the accuracy, reliability and remarkable consistency of periods of pre-main sequence stars over time scales of many years and confirms the link between rotation and photometric period through comparison with $v \\sin i$ data.  The most unusual variable in IC 348 is HMW 15 and it is the subject of a separate, forthcoming contribution. Also notable is HMW 73, the only star to show large amplitude periodic behavior. This may be the result of rotation of a stable accretion hot spot or more unusual variable behavior.  Finally, we have shown here that the rotation period distribution of stars in the 0.4-1.2 solar mass range in IC 348 matches that of the ONC but contains a higher proportion of slow rotators than is seen in NGC 2264. Since the age of IC 348 is closer to NGC 2264 than the ONC, this suggests that age and mass alone are not sufficient to predict the rotation properties of stars in an extremely young cluster. We suggest that environment may also be a factor, as \\citet{sw05} have also proposed based on their v sin i study of h and $\\chi$ Persei. If disk-locking controls rotational evolution for stars in this mass and age range then the longer accretion disks can persist, the longer the stars can continue to resist spin-up due to contraction with conservation of angular momentum. Perhaps in the less harsh environment of IC 348, whose earliest type star is of B5 and not a prodigious source of ultraviolet photons, the disks can persist longer than in NGC 2264. If so, the stars may continue to reflect an ONC-like distribution for times scales of order 2 My.  Continued study of this, and similar clusters, will be needed to ascertain whether this environment argument has merit. Because the rotation period distribution is so broad (and mass dependent) for young clusters, one requires a large number of stars to determine it and to distinguish significant differences among clusters. Obviously, it will continue to be of interest to monitor the variations of peculiar young stars such as HMW 15 and HMW 73 since they undoubtedly have much to tell us about the processes of accretion and disk evolution in solar-like stars."
    },
    "0606/astro-ph0606116_arXiv.txt": {
        "abstract": "We have exploited the large area coverage of the combined UKIDSS Ultra Deep Survey (UDS) and Subaru/XMM-Newton Deep Survey (SXDS) to search for bright Lyman-break galaxies (LBGs) at $z\\geq5$. Using the available optical+near-infrared photometry to efficiently exclude low-redshift contaminants, we identify nine $z\\geq5$ LBG candidates brighter than \\zp$_{AB}=25$ within the 0.6 square~degree overlap region between the UDS early data release (EDR) and the optical coverage of the SXDS. Accounting for selection incompleteness, we estimate the corresponding surface density of $z\\geq5$ LBGs with \\zp$_{AB}\\leq25$ to be $0.005\\pm0.002$ per square~arcmin. Modelling of the optical+near-infrared photometry constrains the candidates' redshifts to lie in the range $5.1<z<5.9$, and provides estimates for their stellar masses. Although the stellar mass estimates are individually uncertain, a stacking analysis suggests that the typical stellar mass of the LBG candidates is $\\gtsim\\,5\\times10^{10}\\Msolar$ which, if confirmed, places them amongst the most massive galaxies currently known at $z\\geq5$. It is found that $\\Lambda$CDM structure formation can produce sufficient numbers of dark matter halos at $z\\geq5$ to accommodate our estimated number density of massive LBGs for plausible values of $\\sigma_{8}$ and the ratio of stellar to dark matter. Moreover, it is found that recent galaxy formation models can also account for the existence of such massive galaxies at $z\\geq5$. Finally, no evidence is found for the existence of LBGs with stellar masses in excess of $\\simeq3\\times 10^{11}\\Msolar$ at this epoch, despite the large co-moving volume surveyed. ",
        "introduction": "The increasing availability of survey fields with deep multi-wavelength data has led to so-called ``dropout'' techniques becoming established for photometrically identifying high-redshift galaxies. At redshifts of $z\\simeq6$ it is now standard practice to select galaxy candidates using the $i-$drop technique, a straightforward extension of the Lyman-break method pioneered at $z\\simeq3$ by Guhathakurta, Tyson \\&  Majewski (1990) and Steidel \\& Hamilton (1992). Within this context, the unique combination of depth and image quality provided by the Hubble Space Telescope (HST) has arguably made the largest contribution to our understanding of high-redshift galaxies. Analysis of the Hubble deep fields, the GOODS fields, and in particular the Hubble ultra-deep field (HUDF), has led to the identification of hundreds of $i-$drop galaxy candidates at faint ($z_{850}\\geq26$\\footnote{All magnitudes are quoted in the AB system (Oke \\& Gunn 1983)}) magnitudes (e.g. Bouwens et al. 2006; Malhotra et al. 2005; Bunker et al. 2004; Dickinson et al. 2004). As a consequence, we now have a vastly improved understanding of the galaxies which undoubtedly dominate the star-formation density of the Universe during the epoch immediately following reionisation.  However, although HST has greatly advanced our knowledge of high-redshift galaxies, the various HST deep fields suffer from one fundamental weakness. Crucially, their small area (e.g. the HUDF covers $\\simeq$13 square~arcmin) makes them subject to significant cosmic variance (Somerville et al. 2004) and completely unsuitable for investigating the most luminous, rarest objects.  Indeed, of the many hundreds of $i-$drop galaxies identified in the deep HST fields, only one has $z_{850}\\leq25$ (SBM03$\\#3$, spectroscopically confirmed at $z=5.78$; Bunker et al. 2003). This is an important shortcoming because accurately determining the number density of massive galaxies at high-redshift ($z\\geq5$) has the potential to place important new constraints on galaxy formation models, only 1 Gyr after the Big Bang. From an observational perspective, it is clear that what is required is a combination of both depth and large area coverage.  In this respect, the study of $z\\simeq6$ LBGs in the Subaru Deep Field (SDF; Maihara et al. 2001) by Shimasaku et al. (2005) is perhaps the most notable. Based on deep Subaru optical imaging covering 767 square~arcmin, Shimasaku et al. identified twelve bright LBGs with \\zp-band magnitudes in the range 25.4$<$\\zp$<$26.6,  photometrically selected to lie in the redshift range $5.6<z<6.2$. Moreover, by using two medium-band filters ($z_{B}$ and $z_{R}$) which divide the Subaru \\zp-prime filter in half, Shimasaku et al. gained crucial information on the UV slope of their high-redshift candidates. As a consequence, Shimasaku et al. were able to select a very clean sample, which should have negligible contamination from low-redshift interlopers. In addition, Ota et al. (2005) performed a search for \\zp$\\leq26$ LBGs over a sky area of $\\sim1$ square~degree in the SXDS field, using a primary selection criterion of $i^{\\prime}-z^{\\prime}>1.5$. However, Ota et al. were only able to strongly constrain the surface density of LBG candidates at $z^{\\prime}\\geq25$, because at brighter magnitudes their lack of near-infrared data made it impossible to exclude significant contamination from low-redshift interlopers, and cool galactic stars in particular. Finally, it should be noted that there is already evidence for a large population of $z\\simeq6$ galaxies at fainter $z^{\\prime}-$band magnitudes within the SXDS/UDS field. Using a combination of the broad $i^{\\prime}-$band SXDS imaging and a narrow band filter ($\\lambda_c=8150$\\AA), Ouchi et al. (2005) identified 515 potential Ly$\\alpha$ emitters (LAEs) within the redshift interval $5.65<z<5.75$, over an area of one square~degree. Further deep optical spectroscopy subsequently confirmed eight LAEs at $z\\simeq5.7$ with asymmetric \\Lya emission lines. However, the LAEs detected by Ouchi et al. are fainter ($z^{\\prime}\\geq25$) than the objects identified in this study (and presumably have lower mass).  In this paper we exploit two crucial advantages of the combined SXDS/UDS data-set to explore a new area of LBG parameter space at $z\\geq5$. Firstly, the overlap region between the SXDS and the UDS EDR is $\\geq2.5$ times larger than the SDF ($\\simeq 0.6$ square~degrees), offering the prospect of identifying the rarest $z\\geq5$ LBGs with $z^{\\prime}\\leq25$. Secondly, the near-infrared data from the UDS EDR allows us to effectively clean the sample of low-redshift contaminants, and to obtain stellar mass estimates for the LBG candidates. The structure of the paper is as follows. In Section 2 we give a brief review of the optical and near-infrared data currently available in the SXDS/UDS field. In Section 3 we outline the criteria adopted for selecting the $5<z<6$ galaxy candidates, and discuss the possible sources of contamination. In Section 4 we present the results of our modelling of the candidates' spectral energy distributions (SEDs). In Section 5 we estimate the number densities of massive LBGs at $z\\geq5$ and compare with the predictions of current galaxy formation models. In Section 6 we compare our results with those in the recent literature, before presenting our conclusions in Section 7. Throughout the paper we adopt the following cosmology: $H_{0}=70$ km\\,\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$. When referring to the $i$~and $z-$band filters, $i^{\\prime}$ and $z^{\\prime}$ specifically refer to the Subaru filters, whereas $i_{814}$ and $z_{850}$ are the equivalent HST filters. ",
        "conclusions": "In this paper we have presented the results of a study aimed at exploiting the large co-moving volume of the SXDS/UDS data-set to identify the most massive LBGs at $z\\geq5$. Our main conclusions can be summarised as follows:  \\begin{enumerate}  \\item{We have identified a robust sample of nine bright ($z^{\\prime}\\leq~25$) LBG candidates with photometric redshifts in the range $5.1<~z~<5.9$.}  \\item{Our corresponding estimate of the surface density for $z^{\\prime}\\leq~25$ LBGs at $z\\geq5$ is $0.005\\pm0.002$ per square~arcmin. This surface density estimate is in good agreement with, although much more robust than, a prediction based on the one previously known LBG candidate with $z_{850}\\leq25$ in a compilation of the deep HST fields; $0.004\\pm0.004$ per square arcmin (Bouwens et al. 2006).}  \\item{SED fits to the stacked photometry of the final nine LBG candidates suggests that their typical stellar mass is $\\gtsim\\,5\\times10^{10}\\Msolar$ (Salpeter IMF), comparable with the most massive galaxies currently known at $z\\geq5$.}  \\item{Based on our best-fitting stellar mass estimates, five of the nine LBG candidates have estimated stellar masses $\\gtsim\\,10^{11}\\Msolar$. Our corresponding estimate for the number density of $z\\geq5$ LBGs with stellar masses $\\gtsim\\,10^{11}\\Msolar$ is $\\log(\\phi/{\\rm Mpc}^{-3}) = -5.2^{+0.2}_{-0.3}$.}  \\item{This number density is found to be compatible with the density of suitable dark matter halos predicted by $\\Lambda$CDM structure formation models for plausible combinations of $\\sigma_{8}$ and the ratio of stellar to dark matter.}  \\item{It is found that recent galaxy formation models can account for the existence of such massive galaxies at high redshift, even if LBGs only account for $\\leq50\\%$ of the total stellar mass at $z\\geq5$.}  \\item{Despite the large co-moving volume surveyed, no evidence is found for the existence of LBGs with masses in excess of $3\\times~10^{11}\\Msolar$ at $z\\geq5$.}  \\end{enumerate}"
    },
    "0606/astro-ph0606320_arXiv.txt": {
        "abstract": "GRB X-ray light curves display rapid declines followed by a gradual steepening or plateau phase in $\\gtrsim 30$\\% of GRBs in the Swift sample. Treating the standard relativistic blastwave model in a uniform circumburst medium, it is shown that if GRBs accelerate ultra-high energy cosmic rays through a Fermi mechanism, then the hadronic component can be rapidly depleted by means of photopion processes on time scales $\\sim 10^2$ -- $10^4$ s after the GRB explosion.  While discharging the hadronic energy in the form of ultra-high energy cosmic ray neutrals and escaping cosmic-ray ions, the blast wave goes through a strongly radiative phase, causing the steep declines observed with Swift. Following the discharge, the blast wave recovers its adiabatic behavior, forming the observed plateaus or slow declines. These effects are illustrated by calculations of model bolometric light curves. The results show that steep X-ray declines and plateau features occur when GRB sources take place in rather dense media, with $n \\gtrsim 10^2$ cm$^{-3}$ out to $\\gtrsim 10^{17}$ cm. ",
        "introduction": "The Swift telescope is providing a new database of GRB light curves consisting of a BAT light curve in the 15 -- 150 keV range followed, after slewing within $\\approx 100 $ s, by a detailed 0.3 -- 10 keV XRT X-ray light curve and UVOT monitoring \\citep{geh04}. Extrapolating the BAT light curve to the XRT range gives $\\sim 1$ keV X-ray light curves since the trigger.  \\citet{obr06} present a catalog of the combined 0.3 -- 10 keV light curves of 40 GRBs, of which $ 14$ or so have measured redshifts, $\\approx 30$\\% display rapid X-ray declines, and an additional $\\approx 30$\\% display features unlike simple blast wave model predictions. In some GRBs, the 0.3 -- 10 keV flux can decrease by 4 or 5 orders of magnitude over a period of  $\\lesssim 10^2$ seconds within several minutes after the GRB trigger (e.g., GRB 050915B, GRB 050422, GRB 050819). About one-half the XRT sample shows X-ray flares or short timescale ($\\Delta t/t\\ll 1$) structure at $\\gtrsim 10^3$ s after the GRB trigger, and in some cases out to $\\gtrsim 10^5$ s (e.g., GRB 050904 at $z = 6.29$ and GRB 050730 at $z = 3.97$).  Following the rapid X-ray declines, a more gradual steepening commences in most XRT light curves, as in the cases of GRB 050219A or GRB 050607. In several GRBs, e.g., GRB050315 or GRB050822, a plateau phase with rising and decaying features within $0.1$ -- $ 1$ day after the GRB trigger are observed; this phase may also be found in other GRBs (GRB 050724 or GRB 050819), but could be contaminated or mimicked by a late-time X-ray flare. In a few GRBs (e.g., GRB 050401, GRB 050717), the X-ray decline is gentle and monotonic, but in general the XRT light curves, whether displaying overall convexity or concavity, reveal temporal structure and, oftentimes, X-ray flares (e.g., GRB 050712, GRB 050716, GRB 050726). A phenomenological model with two distinct components can fit the rapid decays and hardenings \\citep{wil06}, but lacks a physical basis. An acceptable physical model for GRB afterglows must be able to explain this diversity of behaviors.  A combined leptonic-hadronic GRB model is proposed in this paper as the cause of the rapid X-ray declines and plateaus discovered with Swift. The analysis presented here follows the standard blastwave model \\citep[e.g.,][]{mes06}. We show that if GRBs accelerate cosmic rays to ultra-high energies, then for certain classes of GRBs, the blastwave will become strongly radiative during the early afterglow, and consequently will exhibit rapid X-ray declines. This class of GRBs is defined by a range of blast wave and environmental parameters. One set of parameters that produces such declines has initial Lorentz factors $\\Gamma_0 \\sim 100$ -- 300, apparent total energy releases $\\gtrsim 10^{54}$ ergs (and absolute energy releases $\\lesssim 10^{52}$ ergs), and surrounding medium density $n \\gtrsim 10^2$ cm$^{-3}$, which is assumed to be proton dominated.  The blast wave microphysical parameters $\\e_e, \\e_B$ must both be $\\gtrsim 0.1$.  For these GRBs, Fermi processes in the blast wave are assumed to accelerate proton and ions, like they do electrons, to ultrarelativistic energies. By making reasonable approximations for the acceleration rate as a fraction of the Larmor rate, and particle escape through Bohm diffusion, we find that photohadronic losses and particle escape significantly deplete the internal energy of the blast wave, causing the blast wave dynamics to be strongly affected. Photopion interactions by the ultrarelativistic protons and secondary neutrons with the internal synchrotron photons make a source of escaping neutrons, neutrinos and cascade $\\gamma$ rays, in addition to a generally weaker proton synchrotron component.  In Section 2, a standard blast-wave physics analysis for emissions from a GRB external shock in the prompt and early afterglow phase is presented, and timescales for the various processes are calculated for an adiabatic blast wave that decelerates by sweeping up material from a uniform surrounding medium. Parameter sets that allow a large fraction of the internal energy to be radiatively discharged through hadronic processes are graphically examined in Section 3. In Section 4, the equations for blastwave evolution are solved in the case of internal energy that is promptly radiated or is exponentially depleted with time. Synthetic bolometric light curves are calculated. Although these cannot be directly compared to the Swift X-ray light curves without taking into account spectral effects and more complicated radiation physics, the bolometric light curves exhibit many of the features observed with Swift.  Discussion of multiwavelength and multichannel  $\\gamma$-ray, cosmic-ray, and neutrino predictions for this model is found in Section 5, including a comparison with other models for the X-ray declines. The study is summarized in Section 6. ",
        "conclusions": "By carrying through an analysis of a complete leptonic/hadronic GRB blast-wave model, a region of parameter space was found for GRB external shocks where, even in the simplest constant density blastwave model, photohadronic processes have important effects on blastwave evolution. This occurs for the observer a few hundred seconds after the start of the GRB for circumburst medium densities $n_0\\gtrsim 10^2$ cm$^{-3}$, $100 \\lesssim \\Gamma_0 \\lesssim 300$, $E_{54} \\sim 1$ -- 10, and $\\e_e,\\e_B\\gtrsim 10$\\%.  For these parameters, the time scales for photopion losses and direct escape become comparable to the available time in the early afterglow, with observable consequences.  The GRB blast wave is predicted to evolve toward the strong cooling regime in the early afterglow, leading to a strong discharge of energy through photohadronic processes and direct escape of UHECR ions, causing the declines in the X-ray light curves observed by Swift.  The change of $\\e_{pk}$ with time may indicate the evolution of the radiative regime during the prompt and early afterglow phases, though this component can be concealed in an internal/external shock model by internal shock emissions.  Hadronic radiations consisting of high-energy neutrinos, cascade $\\gamma$ rays, and escaping cosmic-ray neutrons could be important to $\\gtrsim 10^{4}$ s after the GRB, producing a component that varies, in general, differently from the X-ray lepton sychrotron emission. As a prediction for GLAST, this component should be bright at GeV energies just before the times that the rapid X-ray declines are observed. A related flux of $\\gtrsim 10^{17}$ eV neutrinos is made at this time.  The release of energy in the form of neutrons or escaping ions in the early afterglow phase means that there was much more energy in the GRB blast wave during the prompt phase. Thus the efficiency of $\\gamma$-ray production during the prompt phase is low, which ameliorates efficiency concerns in both external or internal shock models. Hadronic processes imprint GRB light curves and the UHECR spectrum with distinctive signatures. The rapid X-ray declines and plateau phases discovered with Swift are argued here to be just such signatures."
    },
    "0606/astro-ph0606099_arXiv.txt": {
        "abstract": "{} {We investigate the formation of the \\ion{H}{i} Lyman~$\\alpha$ and Lyman~$\\beta$ lines in an equatorial coronal streamer. Particular attention is paid to frequency redistribution for the scattering of the incident radiation. The properties of the spectral lines are studied.} {The coronal model is given by a global 2.5-D three fluid solar wind model with $\\alpha$ particles. The emergent intensities and line profiles are calculated from the solution of the statistical equilibrium and radiative transfer equations for an hydrogen atom with 11 energy levels under non local thermodynamic equilibrium. The formation of the lines results from radiative excitation, collisional excitation, and takes into account the coupling with all other transitions between the hydrogen energy levels.} {We present new estimates of the radiative and collisional contributions of the Lyman line intensities within the streamer. It is also shown that within the streamer, the full width at half-maximum (FWHM) of the Lyman~$\\beta$ line is a better indicator of the plasma temperature than that of Lyman~$\\alpha$. These results show that care should be taken when inferring the proton temperature from the Lyman~$\\alpha$ line profile as observed in coronal streamers, e.g. by the \\emph{Ultraviolet Coronagraph Spectrometer} or the \\emph{Solar Ultraviolet Measurements of Emitted Radiation} experiments on board the \\emph{Solar and Heliospheric Observatory}.} {} ",
        "introduction": "A large amount of our knowledge of the solar extreme ultraviolet (EUV) corona comes from the observations of the first two \\ion{H}{i} Lyman lines at 1215.67 \\AA\\ and 1025.72 \\AA. The Lyman~$\\alpha$ line is indeed the brightest line in this range of radiation. To infer the coronal plasma properties from observations it is necessary to get the line profiles with enough spectral resolution. {While chromospheric profiles show a self-reversal at line centre, the coronal profile of the scattered Lyman~$\\alpha$ line is close to a gaussian profile.}  From the knowledge of the line width one can derive an {equivalent} temperature {(ie. including all causes of line broadening),} and with further assumptions, the kinetic {(thermal)} temperature of the hydrogen atoms. If the coupling due to charge exchange between hydrogen atoms and ions is strong enough, then the proton temperature is equal to the hydrogen temperature. {\\cite{loaetal98,loaetal00} concluded from fast solar wind models that this coupling is strong up to 3 R$_{\\sun}$, while \\cite{olh94} studied slow wind models and found strong coupling up to 10 R$_{\\sun}$ in some conditions. The decoupling of hydrogen atoms and protons will occur at different densities, depending on which physical assumptions are made in the models. In our case, this coupling is strong enough at densities above 10$^6$ cm$^{-3}$. This condition is met within the whole streamer \\citep[see for further details][]{lll06}.} Thus we can obtain a reliable estimate of the proton temperature from the full width at half maximum (FWHM) of the Lyman line profiles \\citep[see also][]{marschetal99}.  In this paper we show that in the case of coronal streamers, radiative transfer calculations in non local thermodynamic equilibrium (NLTE) are useful to predict the properties of the Lyman~$\\alpha$ and Lyman~$\\beta$ lines. The temperature diagnostic from these two lines is examined. We describe the streamer model in section~\\ref{bomodel} and the radiative transfer calculations in section~\\ref{radtrans}. Results are discussed in section~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  Figure~\\ref{intensites} presents the resulting integrated intensities of the first two Lyman lines as a function of distance from sun centre. The intensities are calculated by summing over the computed line profiles. The decrease of the intensity with altitude is more pronounced for Lyman~$\\beta$ than for Lyman~$\\alpha$, a first indication that they relate to the plasma parameters in different ways. The Lyman~$\\alpha$ intensities compare well with the computed and observed values presented in \\cite{vbr03}. {We find that the variation of the Ly$\\alpha$ intensity with height is not related to the variation of the electron density, while the decrease of Ly$\\beta$ and H$\\alpha$ (not shown on Fig.~\\ref{intensites}) intensities follows the decrease of $n_e^2$ closely up to $\\sim 1.5$~R$_{\\sun}$, consistent with the fact that these two lines are mostly formed by collisional excitation in the inner corona. Then the fall-off in intensity is less rapid than that of the square of the electron density, owing to the growing importance of radiative excitation (see also Fig.~\\ref{compo}). In fact there is a coupling between Ly$\\beta$ and H$\\alpha$, which means that {an} H$\\alpha$ photon can be absorbed and subsequently lead to an emission of a Ly$\\beta$ photon. We obtain a nearly constant ratio between Ly$\\beta$ and H$\\alpha$ intensities, with $I(\\mathrm{Ly}\\beta)/I(\\mathrm{H}\\alpha) \\simeq 8$. This coupling leads to lowered coherency effects, as was illustrated by \\cite{hgv87}. The coherence coefficient $\\gamma$ (eq.~\\ref{defredi}) is close to 1 for Ly$\\alpha$ and 0.57 for Ly$\\beta$ at the centre of the LOS closer to the Sun (height of 1.05~R$_{\\sun}$).  We have computed the parameter $\\lambda = (A_{ji} / P_j) \\times \\gamma$ as in \\cite{hgv87} -- although we use a slightly different definition for $\\gamma$, where $A_{ji}$ and $P_j$ are the spontaneous emission coefficient in the $j\\to i$ transition and the total depopulation rate of level $j$, respectively. Close to the Sun, in the streamer base, our value of $\\lambda$ is around 0.99 for Ly$\\alpha$ and 0.3 for Ly$\\beta$ (at the centre of LOS).  This confirms that the coherency effects in Ly$\\beta$ are less important than in Ly$\\alpha$. } \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{5059fig3.eps}}% \\caption{Integrated intensities in photons s$^{-1}$ cm$^{-2}$ sr$^{-1}$ of Lyman~$\\alpha$ (triangles) and Lyman~$\\beta$ (squares) as a function of heliocentric distance.} \\label{intensites} \\end{figure}  Now we turn to the study of the Lyman line profiles. We fit each computed profile with a gaussian profile obtained from a non-linear least squares fit. {We find that the Ly$\\alpha$ wings are broader than our gaussian fits at all altitudes. The line centre is well reproduced by the gaussian fits at altitudes above $\\sim$1.6~R$_{\\sun}$. The Ly$\\beta$ line deviates from the gaussian fits as the altitude increases. The line centre is always well reproduced by the gaussian fit, but the wings get broader with height.} From our gaussian fits we obtain the full width at half maximum (FWHM) of the computed line profiles.  As we go further up in the corona, the line widths decrease. This reflects the decrease in temperature. However, care should be taken when inferring the temperature from the line width. {As the profile is not exactly gaussian, this introduces an error in the temperature derivation \\citep[for a more thorough discussion see, e.g.,][]{loaetal98}}. In this respect it seems more reliable to exploit the diagnostic possibilities of the Lyman~$\\beta$ line which is closer to a gaussian {than the Lyman~$\\alpha$ line}.  To infer the temperature of the neutrals from the line widths, one can relate the FWHM and the temperature $T_H$ with: \\begin{equation} T_H = \\frac{m}{2k} \\left[ \\mathrm{FWHM}^2 \\frac{c^2}{4 \\lambda^2\\ln 2}- \\xi^2 \\right] \\ . \\label{temp} \\end{equation} We arbitrarily chose to use $\\xi = 20$ km s$^{-1}$ for the non-thermal motions in the calculations of the line profiles at all altitudes. While it might not be accurate, the exact value for $\\xi$ is not important for our discussion on the temperatures inferred from the Lyman lines, as we assume that they both have the same non-thermal broadening. The resulting temperatures derived from the width of the Lyman~$\\alpha$ line and the  Lyman~$\\beta$ line are plotted in Fig.~\\ref{temperatures} as a function of height of the LOS, together with the mean temperature derived from the model input and the temperature at the centre of the LOS. The mean temperature is defined as $T_\\mathrm{mean} = \\int_0^M{T(m)\\mathrm{d}m}/M$, with $m$ the column mass along the LOS, and $M$ the total column mass of the LOS. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{5059fig4.eps}}% \\caption{Temperature derived from the width of Lyman~$\\alpha$ (triangles), Lyman~$\\beta$ (squares), temperature at centre of LOS (thick solid line), and  mean temperature from the model input (thick dotted line), as a function of heliocentric distance.} \\label{temperatures} \\end{figure} Inspection of figure~\\ref{temperatures} clearly shows that the width of the Lyman~$\\beta$ line is indeed a good indicator of the plasma temperature from 1.05~R$_{\\sun}$ up to $\\sim$2~R$_{\\sun}$. Above this height, the temperature derived from the Lyman~$\\beta$ FWHM is slightly lower than the mean plasma temperature $T_\\mathrm{mean}$ and than the central temperature, an indication that the plasma conditions have significantly changed. Indeed, from the streamer model the magnetic cusp is located at about 2~R$_{\\sun}$. Within that distance, outflow velocities are very small (a few km~s$^{-1}$). Above the cusp, the protons reach a velocity of about 100 km~s$^{-1}$ at 3~R$_{\\sun}$.  Close to the Sun, the temperature derived from the Ly$\\beta$ FWHM is in very good agreement with the temperature at the centre of the LOS. However, higher in the streamer, the temperature derived from the Ly$\\beta$ line width is in better agreement with the mean plasma temperature than with the temperature at the centre of the LOS. Due to the temperature and hydrogen density variations along the LOS (see Fig.~\\ref{streamer}), the contribution of collisional excitation in the formation of the Ly$\\beta$ line is more concentrated at the centre of the LOS when close to the Sun, and more smoothly distributed along the line of sight higher in the streamer. We note that when neglecting CRD in frequency redistribution in eq.~(\\ref{defredi}), the temperature derived from the Ly$\\beta$ width \\emph{exactly} matches the mean temperature of the models up to a height of 1.6~R$_{\\sun}$. It can also be seen from Fig.~\\ref{temperatures} that the temperature derived from the Lyman~$\\alpha$ line significantly underestimates the mean plasma temperature and the temperature at the centre of the LOS. This difference can be as much as $3.8\\times10^5$~K at $r=1.39$~R$_{\\sun}$ for the mean temperature, and $5.0\\times10^5$~K at $r=1.58$~R$_{\\sun}$ for the central temperature. The difference in behaviour between the two Lyman lines points to the different relative contributions of radiative and collisional excitation in the formation of the two lines. From eq.~(\\ref{source}) we identify the first term of the right-hand side with the collisional component and the second term with the radiative component. Figure~\\ref{compo} shows their variation with altitude for the Lyman~$\\alpha$ and Lyman~$\\beta$ lines. The radiative component is represented with a solid line, and the collisional component with a dotted line, while triangles stand for Lyman~$\\alpha$ and squares for Lyman~$\\beta$. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{5059fig5.eps}}% \\caption{Relative contribution of the radiative (solid line) and collisional (dotted line) components for Lyman~$\\alpha$ (triangles) and Lyman~$\\beta$ (squares) as a function of heliocentric distance.} \\label{compo} \\end{figure} This figure shows that, as expected, Lyman~$\\beta$ is mostly formed by collisional excitation in the streamer. However it is interesting to note that the collisional component of Lyman~$\\alpha$ is not negligible in the streamer base, contributing to nearly 10\\% of the line intensity close to the Sun. We should stress here that both {the scattered and collisional} components will include contributions resulting from exchanges with other transitions (lines and continua) involving the rest of the atomic states. The values reported in Fig.~\\ref{compo} can be compared with those given by \\citet{rayetal97} in the centre of a streamer observed by UVCS. These authors obtained a collisional contribution of 1.1\\% for Lyman~$\\alpha$ and 57\\% for Lyman~$\\beta$ at $\\log T=6.2$, while our values at that temperature (at a heliocentric height of 1.7~R$_{\\sun}$) lead to 3\\% and 77\\%, respectively. \\\\  In this paper NLTE radiative transfer calculations are performed to compute the properties of the Lyman lines of hydrogen in the solar corona. It is shown that the width of the Lyman~$\\beta$ line is a better indicator of the plasma temperature than the width of the Lyman~$\\alpha$ line, especially within the streamer. {It is due to the formation mechanisms of the lines, and the coupling of Ly$\\beta$ with H$\\alpha$. This work has been done using the approximation of isotropic frequency redistribution {in the laboratory frame}. We do not expect that the inclusion of angle-dependent redistribution functions in our calculations would change the main conclusion of this Letter, namely that the width of Lyman~$\\beta$ is a better proxy for the plasma temperature in the streamer. It is known that the consideration of angular redistribution with dipole scattering has the effect of narrowing the line profile compared to isotropic redistribution, due to non-90$\\degr$ scattering of photons \\citep[e.g.][]{witetal82,loaetal98,cranmer98}. Therefore we believe that the discrepancy found here between temperatures derived from Ly$\\alpha$ line profiles and model temperatures could be even greater by using angle-dependent redistribution functions. Furthermore, \\cite{cv78} showed that the effects of angle-dependent PRD will be more important when the incident lines show substantial center-to-limb variations. This is not the case of hydrogen Lyman lines, but this is the case for, e.g., the \\ion{O}{vi} lines at 1032 and 1038 \\AA. We will include this effect in a future study when we include \\ion{O}{vi} in our calculations.} We also obtain new estimates of the radiative and collisional contributions of the Lyman line intensities in a non-isothermal streamer. These new values may have some importance in the derivation of element abundances. {Element abundances relative to hydrogen can be inferred independently from the ratio of the resonantly scattered (or collisional) component of a spectral line to the resonantly scattered (or collisional) component of, say, \\ion{H}{i}~Ly$\\beta$ \\citep{witetal82,rayetal97}. Therefore the relative contributions of the two components of the Lyman lines to their total observed intensities must be known with good accuracy. Finally, } our results can be compared with observations by the SUMER and UVCS spectrometers on SOHO. The radiative transfer calculations can be enhanced by including other effects such as Doppler dimming to improve the modelling in regions of the corona where outflow velocities cannot be ignored. This has been presented in \\cite{corsoho17}."
    },
    "0606/astro-ph0606050_arXiv.txt": {
        "abstract": "{BL Lacertae has been the target of several observing campaigns by the Whole Earth Blazar Telescope (WEBT) collaboration and is one of the best studied blazars at all accessible wavelengths. A recent analysis of the optical and radio variability indicates that part of the radio variability is correlated with the optical light curve. Here we present an analysis of a huge VLBI data set including 108 images at 15, 22, and 43\\,GHz obtained between 1995 and 2004. The aim of this study is to identify the different components contributing to the single-dish radio light curves. We obtain separate radio light curves for the VLBI core and jet and show that the radio spectral index of single-dish observations can be used to trace the core variability. Cross-correlation of the radio spectral index with the optical light curve indicates that the optical variations lead the radio by about 100 days at 15\\,GHz. By fitting the radio time lags vs. frequency, we find that the power law is steeper than expected for a freely expanding conical jet in equipartition with energy density decreasing as the square of the distance down the jet as in the K\\\"onigl model. The analysis of the historical data back to 1968 reveals that during a time range of 16 years the optical variability was reduced and its correlation with the radio emission was suppressed. There is a section of the compact radio jet where the emission is weak such that flares propagating down the jet are bright first in the core region with a secondary increase in flux about 1.0 mas from the core. This illustrates the importance of direct imaging to the interpretation of multi-wavelength light curves that can be affected by several distinct components at any given time. We discuss how the complex behaviour of the light curves and correlations can be understood within the framework of a precessing helical jet model.} ",
        "introduction": "The term ``blazars'' identifies a family of radio-loud active galactic nuclei (AGN) showing strong variability at all wavelengths from radio to $\\gamma$-rays, high degrees of polarization, and apparent brightness temperatures exceeding the Compton limit (e.g., \\citealt{1999APh....11..159U}). During the last decades a rather general consensus on the global mechanism responsible for the emission has been achieved: a supermassive black hole surrounded by a massive accretion disk feeding a powerful jet closely aligned to the line of sight. Relativistic electrons in the jet plasma produce the soft synchrotron photons (from radio wavelengths to the UV band and sometimes even X-rays) through, while hard photons (from X-rays to $\\gamma$-rays) are likely produced by inverse Compton scattering (e.g., \\citealt{1996ApJ...463..444S} and references therein). Blazars are divided into two subclasses: flat-spectrum radio quasars and BL Lac objects, whose main difference is the lack or weakness of emission lines in BL Lac objects (strongest rest-frame emission line width $<5$\\,\\AA; e.g., \\citealt{1991ApJ...374..431S} and \\citealt{1995PASP..107..803U} for a review).  BL Lacertae, located in the nucleus of a moderately bright elliptical galaxy at a redshift of $z=0.069$ (\\citealt{1978ApJ...219L..85M}) is the prototype of the class of BL Lac objects. However, emission lines with equivalent width in excess of $5$\\,\\AA\\ have been detected several times (\\citealt{1995ApJ...452L...5V,1996MNRAS.281..737C,2000MNRAS.311..485C}), suggesting that BL Lac objects can also have a broad line region (BLR), but it would be usually outshone by the beamed synchrotron emission of the jet. BL Lac has been studied intensively since its discovery and was the target of several multi-wavelengths campaigns (\\citealt{1997ApJ...490L.145B,1999ApJ...515..140S,1999ApJ...521..145M,2002A&A...390..407V,2003ApJ...596..847B,2004A&A...424..497V,2004A&A...421..103V}). In particular, the data collected during the Whole Earth Blazar Telescope (WEBT)\\footnote{http://www.to.astro.it/blazars/webt} campaigns provide an unprecedented time sampling for radio and optical light curves from 1994 up to now (\\citealt{2002A&A...390..407V,2004A&A...424..497V,2004A&A...421..103V}). Cross-corelation analysis of these light curves revealed well correlated variability in the radio bands, where variations at higher frequencies lead the lower-frequency ones by several days to a few months depending on the frequency separation. The detection of a fair correlation between the optical $R$ band variations and the ratio between the 22\\,GHz and 5\\,GHz radio flux densities suggests that the optical emission and part of the radio emission arises from a common origin in the inner portion of the jet.  Very Long Baseline Interferometry (VLBI) observations of the parsec-scale structure of BL Lac reveal a bright compact radio core and a jet extending to the south (e.g., \\citealt{1990ApJ...352...81M}). The jet emerges at a position angle of P.A.~$\\approx 195^\\circ$ and displays a small bend at $\\sim 4$\\,mas distance from the core towards P.A.~$\\approx 160^\\circ $. A number of superluminal jet components have been observed displaying bent trajectories and speeds from 3\\,$c$ to $9\\,c$ (e.g., \\citealt{1990ApJ...352...81M,2000ApJS..129...61D,2003MNRAS.341..405S,2004ApJ...609..539K,2005AJ....130.1418J}). Evidence for precession of the jet nozzle arising from periodic variations of the electric vector position angle (EVPA) at millimetre wavelengths and the position angle of newly emerging jet features (\\citealt{2003MNRAS.341..405S}) are still under debate (\\citealt{2005ApJ...623...79M}). A comparison of the position angles of optical polarization vectors with the orientation of the polarization vectors in 5\\,GHz VLBI images suggests that the optical emission is more likely to arise from the core than from the jet (\\citealt{1994AJ....107..884G}).  Here we present an investigation of the flux density evolution of the parsec- and subparsec-scale structure of BL Lac at 15\\,GHz, 22\\,GHz, and 43\\,GHz using a dataset of 108 VLBI observations performed between 1995 and 2004. The aim of this study is to disentangle the different variability patterns that seem to be present in the single-dish radio light curves and to identify the region of the radio events that seem to correlate with the optical variability. Throughout this paper we will assume a flat universe model with the following parameters: Hubble constant $H_0=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, a pressure-less matter content $\\Omega_{\\rm m}=0.3$, and a cosmological constant $\\Omega_{\\rm \\lambda}=0.7$ (\\citealt{2003ApJS..148..175S}). At the redshift of BL Lac this corresponds to a luminosity distance of 308\\,Mpc and to an angular resolution of 1.3\\,pc/mas. Cosmology-dependent values quoted from other authors are scaled to these parameters.  The observations and data reduction procedures will be described in Sect.~\\ref{observations}. Our results will be presented in Sect.~\\ref{results} and discussed in Sect.~\\ref{discussion}. At the end we will give a summary with the conclusions in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{discussion}  \\subsection{Optical-radio time delays}\\label{sec:timetest}  The long observational history of BL\\,Lac led to several attempts to search for radio-optical correlations. Some of these found correlations with radio time delays of several hundred days (\\citealt{1976AJ.....81..489P,1992ApJ...386..473H,1994A&A...289..673T,1995AJ....110..529C,2002A&A...394...17H}). Sometimes simultaneous events were reported (\\citealt{Balonek1982,1994A&A...286...80T}). A common result is that there are periods where radio and optical variations correlate better and others where the correlation seems to disappear. The correlations found here (Sect.~\\ref{sec:cross}) are calculated for data from 1986 to 2003, thus excluding the large outburst in 1980. Although the sampling is not homogeneous, it seems that the optical emission underwent a period of suppressed activity from 1981 to early 1997 (Fig.~\\ref{fig:Rspix}), during which flares were less frequent and not so pronounced. Such a period is visible neither in the radio light curves nor in the radio spectral index (Fig.~\\ref{fig:spix}).  By constraining the DCF analysis on smaller time ranges we test how the time delays change with time. This was done exemplarily for the cross-correlation of the R band fluxes with  $\\alpha_{8/15}$. The data between 1986 and 1989 are not so well sampled, but although the DCF results become less significant, they confirm the previous findings. During the period before 1986 a broad peak at about 300 days is found, which is in good agreement with the results from the studies of data up to 1992 (\\citealt{1976AJ.....81..489P,1992ApJ...386..473H,1994A&A...289..673T,1995AJ....110..529C,2002A&A...394...17H}). During the period of weak optical variability (from 1986 to early 1997) a small peak around zero time delay indicates some simultaneous events (\\citealt{Balonek1982,1994A&A...286...80T}), but in general the DCF during this time shows no peak above 0.25. Finally the most recent data (1997 to 2003), characterized by the high optical activity, are dominated by a peak at around 100 days. The fact that both \\cite{2004A&A...424..497V} and our new analysis detected this 100-day time delay for data obtained from 1986 until now, demonstrates that the most recent correlation seems to dominate over that found in the previous period of weak variability. However, the time range used here is not dominated by a major outburst, but rather consists of several flares in the optical (6--7 major events after 1997 and about the same number of minor flares before 1997), and the variability of the spectral index displays nearly the same amplitude throughout the whole period. Therefore, the optical-radio delays found in Sect.~\\ref{sec:cross} seem to represent the result of a true correlation in the variability behaviour of BL\\,Lac over at least the last 10 years of our dataset.  \\subsection{Frequency dependence of the delay}  Frequency dependent time delays in radio flares are commonly related to a combination of optical depth and travel time along the jet. Evolving flares first appear in the inner jet region at the highest frequencies and, as the perturbations (VLBI components) travel down the jet, the flares become visible at successively lower frequencies where the optical depth of synchrotron self-absorption decreases to $\\tau_{\\rm s} \\sim 1$ at that frequency. For a given magnetic field strength ($B(r)$) and electron energy distribution scale factor ($N(r)$), the corresponding $\\tau_{\\rm s}$ is (e.g., \\citealt{1998A&A...330...79L}):  \\begin{equation} \\tau_{\\rm s}=C_2(\\alpha)N_1\\left(\\frac{eB_1}{2\\pi m_{\\rm e}}\\right)^\\epsilon \\frac{\\delta^\\epsilon \\phi_0}{r^{(em+n-1)}\\nu^{\\epsilon+1}} \\end{equation}  Here $m_{\\rm e}$ is the electron mass, $e$ is the electron charge, $\\epsilon=3/2-\\alpha$, $\\alpha$ is the optically thin synchrotron spectral index, and $m$ and $n$ are the power-law exponents of the radial distance dependence of the magnetic field and $N(r)$, respectively (e.g., \\citealt{1998A&A...330...79L}). The observed jet opening angle is $\\phi_0=\\phi \\csc \\theta$. The factor $C_2(\\alpha)$ is described in \\cite{1970RvMP...42..237B}, and $C_2(-0.5)=8.4\\cdot 10^{10}$ in cgs units. Setting $\\tau_{\\rm s}$ to unity gives us an estimate of the distance from the convergence point of the jet to the core as observed at frequency $\\nu$.  Hence, we should be able to measure a frequency-dependent core position in VLBI images: for a conical jet geometry the shift is given by $r \\propto \\nu^{-1/k_{\\rm r}}$, where $k_{\\rm r}=[(3-2\\alpha)m+2n-2]/(5-2\\alpha)$ (e.g., \\citealt{1998A&A...330...79L}). For the case of equipartition between  the magnetic field and electron energy densities, the simplest choice of $m=1$ and $n=2$ (\\citealt{1981ApJ...243..700K}) leads to $k_{\\rm r}=1$ independent of the spectral index (\\citealt{1998A&A...330...79L}). Larger values are reached in regions with steeper pressure gradients than if, e.g., the electrons suffer adiabatic energy losses without reacceleration (\\citealt{1980ApJ...235..386M}). Thus, taking measurements of the core shift  between several frequencies allows to estimate $k_{\\rm r}$ along the jet (e.g., \\citealt{1998A&A...330...79L,2004A&A...426..481K}).  If the radio time lags are also due to opacity effects and the flare travels at a constant speed, then they should also be proportional to $\\nu^{-1/k_{\\rm r}}$. Recent kinematic studies on high frequency VLBA data of BL\\,Lac have consistently derived jet speeds with a bulk Lorentz factor of about 7 and an angle to the line of sight of  $\\sim12^\\circ$, with only minor deviations (\\citealt{2000ApJS..129...61D,2003NewAR..47..641R,2003MNRAS.341..405S,2004ApJ...609..539K,2005AJ....130.1418J}). We therefore assume that the Lorentz factor and angle to the line of sight are more or less constant, and that the time delays are proportional to the distance between the two emitting regions. E.g., a time delay,  $t'$, of 10 days would then correspond to a distance, $r$, of about 0.06\\,pc ($r=t\\,\\beta\\,c$, with $\\beta=\\sqrt{1-\\frac{1}{\\gamma^2}}$ and $t=\\gamma\\,t'$).  In Fig.~\\ref{fig:time lags} we plot the measured radio lags with respect to both the 37\\,GHz data and to the optical data {vs.} frequency. A power-law fit represents the data reasonably well in both cases (reduced $\\chi^2 < 1$) and there it seems that the assumption of a constant jet speed holds at least in the radio regime. It is important to note that the value of $k_{\\rm r}$ strongly depends on the absolute time lag, i.e. on the distance to the base of the jet, which might be not far from the central engine. Although the optical emission could be emitted close to the foot point of the jet, we should consider that this region is still at some considerable distance from the central engine. In this case $k_{\\rm r}$ would be larger. If, on the other hand, the jet speed between the radio and the optical emitting regions is not constant, the simple assumption of $\\tau \\propto \\nu^{-1/k_{\\rm r}}$ is invalid. However, it seems that at least in the radio regime the speed is constant and we can infer that $k_{\\rm r}$ is likely larger than 1.1. This implies that the gradients in jet parameters are steeper than in the case of equipartition with $m=1$. Our result is, on the other hand, consistent with a freely expanding jet with adiabatic energy losses, in which case $n=(2/3)(3-2\\alpha)$ (\\citealt{1980ApJ...235..386M}) and $k_{\\rm r}=(15-14\\alpha)/[3(5-2\\alpha)]$ for $m=1$, or $k_{\\rm r}\\approx 1.2$ for $\\alpha \\approx -0.5$.  \\begin{figure}[htbp] \\centering \\includegraphics[bb=  30 40 575 822,angle=-90,width=9cm,clip]{5235_f10.ps} \\caption{Measured time lags {vs.} frequency. The circles and dashed line represent the measured time lags and the corresponding fit with respect to the optical emission, while the squares and dotted line denote time lags and fit with respect to the 37\\,GHz light curve.}\\label{fig:time lags} \\end{figure}  \\subsection{Connection with jet components}  \\begin{figure*}[htbp] \\centering \\includegraphics[bb= 30 20 572 772,angle=-90,width=15.5cm,clip]{5235_f11.ps} \\caption{Top: Optical (grey squares) and radio (black circles) light curves in comparison with the ejection dates of new superluminal jet components, which are highlighted as grey vertical bars. The width of each bar corresponds to the uncertainty of the ejection date. Bottom: Ejection dates vs. radio spectral index (black circles and grey trend-line) and the 43\\,GHz core flux density (black squares and line). The ejection dates are taken from \\cite{2000ApJS..129...61D}, \\cite{2003MNRAS.341..405S} (S7 to S13), and \\cite{2004AIPC..714..167M} (M1 and M2). Only some of the new components seem to be connected with events in the light curves.}\\label{fig:spixzoom} \\end{figure*}  Since outbursts (especially at mm wavelengths) are often related to the ejection of superluminal jet components (e.g, \\citealt{1998A&A...334..489O,2002A&A...394..851S}), we have compared our radio and optical light curves and the spectral index, $\\alpha_{8/15}$, with the ejection dates of jet components reported in the literature. Figure~\\ref{fig:spixzoom} illustrates this comparison. There is no obvious connection among the events, but some new components are led by an optical flare or followed by an increase in the radio. We also tested whether we could find any evidence for a connection between the ejections and flares by calculating the variations in the flux density at different frequencies and the spectral indices for times shortly before and after the ejection, but no clear trend could be found. There is always a significant fraction of the components that does not show the expected behaviour. However, this kind of analysis might be affected by the relatively large errors in the dates of the earlier ejections (S7 to S10) and the fact that some components appear very close together (S9/S10 and S12/S13). Furthermore, the light curves are generally composites of multiple events, some of which are declining in flux while others are rising. This can render individual flares unnoticed in the total-flux light curves.  The best agreement can be seen for the two newest components M1 and M2. They appear accompanied by an optical flare and are followed by a flare and flattening of the spectrum at radio frequencies. Since previous works have also reported a good agreement between the appearance of flares and the ejection of components from S1 up to S7 (e.g., \\citealt{1990ApJ...352...81M,1999ASPC..159...31A}), it might be that the dates of S8 to S13 are not so well determined. Although S7, S11, and S12 can be identified with minor radio flares and a flattening of the spectrum, there are several large flares in the optical and radio light curves at 1997.6, 1998.4 and 1999.5 that seem to be unrelated to any new component. It will be interesting to see if the ongoing densely sampled VLBI monitoring will find corresponding components to the pronounced optical flares around 2001.5 and 2002.1.  As mentioned earlier, one should also expect to see a sign of the travelling jet components in the different parts of the VLBI light curves. After a flare in the VLBI core, a new component should brighten the jet after some time at every measured distance while it passes by. Unfortunately, the jet flux density is either too low or the time sampling is insufficient to reliably detect such events. In addition, the ejection of several components or the appearance of many flares in a short time can confuse the picture. A good example of the evolution of a new jet component is M1 (Fig.~\\ref{fig:spixzoom}) ejected in 1999.95 (JD=2451527). The corresponding flare is a nicely separated event and from Fig.~\\ref{fig:VLBAvstime} one can see that, some time after the peak in the core, the inner jet ($r\\approx 0.4$\\,mas) brightens at 43\\,GHz and 22\\,GHz, while later it appears at 2\\,mas and finally at 3\\,mas separation from the core. Since not all frequencies were observed simultaneously and in order to see better the evolution, we have converted the flux densities at 15\\,GHz and 43\\,GHz of the core and jet to 22\\,GHz according to their measured spectral index and show these together in Fig.~\\ref{fig:2000flare}. The plotted data points are three-point averages to smooth the residual differences, and the error bars correspond to their standard deviations. Parabolic fits to the data illustrate the appearance of the component at the different jet positions. The calculated speed of the component from the peaks in the light curves is about (2--4)\\,mas/yr or $\\beta_{\\rm app}\\approx$~(12--19)\\,$c$. This is somewhat higher than the speeds of 5\\,$c$ to $9\\,c$ from kinematic studies reported thus far (\\citealt{2000ApJS..129...61D,2003MNRAS.341..405S,2004ApJ...609..539K,2005AJ....130.1418J}), but since we have combined all frequencies in order to be able to fit the peak, our values might be more uncertain.  \\begin{figure}[htbp] \\centering \\includegraphics[bb= 50 50 554 775,angle=-90,width=9cm,clip]{5235_f12.ps} \\caption{VLBI light curves of the radio flare in early 2000. To better follow the flare evolution in both time and space, all fluxes are converted to 22\\,GHz according to their spectral index and combined in a three point average. One can follow the propagation of the core outburst along the jet. The ejection dates of new VLBI components are marked by arrows and correspond to M1 in 1999.95 and M2 in 2000.81 reported by \\cite{2003heba.conf..173M}.}\\label{fig:2000flare} \\end{figure}  Another jet component ejection (M2) is reported for 2000.81 (JD=2451841; \\citealt{2003heba.conf..173M}) and the related radio flare is also visible in Fig.~\\ref{fig:2000flare}. There is a sharp core peak at JD=2451920 and some time later a peak in the $r\\approx0.4$\\,mas light curve. Unfortunately, the dense VLBI monitoring ends in early 2001 and one cannot follow the flare further. However, both ejections are accompanied or led by a bright optical state, which suggests that the correlated optical-radio variability observed here is also related to the ejection of jet components.  During 1999 and 2000 BL\\,Lac was the target of several {\\it RXTE} (\\citealt{2004AIPC..714..167M}) and {\\it BeppoSAX} (\\citealt{2002A&A...383..763R,2003A&A...408..479R,2003ApJ...596..847B}) observations. Interestingly, the source was caught in very different states. The {\\it RXTE} light curve covers nearly the whole period of 1999 to 2001 with dense sampling. In particular, during the ejection in 2000.81 strong variability is observed together with a softening of the X-ray spectrum (\\citealt{2004AIPC..714..167M}). This is also confirmed by the {\\it BeppoSAX} observations from October 31 to November 2, 2000 (JD$\\sim$2451850; \\citealt{2003A&A...408..479R,2003ApJ...596..847B}), which show an unusually strong soft spectral component between 0.2\\,keV and about 10\\,keV. Unfortunately, the flare/ejection in 1999.95 was not covered by {\\it RXTE}, while the {\\it BeppoSAX} observations  of December 5--6 (JD$\\sim$2451518) obtained several days before the extrapolated ejection date showed no sign of a soft component, but a hard spectrum similar to that observed in July 2000. Another {\\it BeppoSAX} observation on June 5--7, 1999 (JD$\\sim$2451336) shows again a small soft component between 0.2\\,keV and about 3\\,keV (\\citealt{2002A&A...383..763R}), which seems also visible as a small enhancement in the {\\it RXTE} light curve and spectrum (\\citealt{2004AIPC..714..167M}). This small X-ray flare is also followed by optical and radio flares, but no new jet component has been reported. However, this might be due to confusion with the two earlier ejections of components S12 and S13, which were separated by only 0.3\\,yr.  Although further coordinated observations are needed to complete the picture, the very well sampled period from 1999 to 2001 already reveals some very interesting similarities in the variability behaviour from X-ray to radio wavelengths, which are also visible as changes in the source VLBI structure. The missing soft-spectrum component in the {\\it BeppoSAX} observations of 1999.92 might blear the claim of a correlation, but given that the typical uncertainty in the ejection dates is about 0.2\\,yr and the soft X-ray tail is potentially a short-lived phenomenon, more observations yielding better statistics are necessary.  \\subsection{A possible jet scenario}\\label{sec:jmod}  Due to the combination of the structural analysis using VLBI light curves and the single-dish observations, we are confident that the emission from the VLBI core, presumably close to the base of the jet, is responsible for the correlated variability observed at optical and radio wavelengths, and possibly also up to X-ray frequencies (\\citealt{2004AIPC..714..167M}). The mechanisms that can explain blazar variability are still the subject of discussion. Possibilities suggested in the literature include shocks in jets (\\citealt{1985ApJ...298..114M,1985ApJ...298..296A,1996etrg.conf...45M}), and flares introduced by the lighthouse effect from changes in the direction of forward beaming caused, e.g., by helical trajectories  of plasma elements (\\citealt{1992A&A...255...59C}), by a precessing binary black-hole system (\\citealt{1980Natur.287..307B,1988ApJ...325..628S}), or by the rotation of a helical jet (\\citealt{1998MNRAS.293L..13V,1999A&A...347...30V,2004A&A...419..913O}). Furthermore, variability caused by colliding relativistic plasma shells has been considered (\\citealt{2001MNRAS.325.1559S,2004A&A...421..877G}).  Regardless of the variability origin, the VLBI observations of BL\\,Lac indicate a bent path of the jet. The trajectories of several jet components have been well modelled and predicted by helical jet models (\\citealt{2000ApJS..129...61D,2003MNRAS.341..405S}). A bent jet path could also explain the region of low emission around 1\\,mas from the core noticed in Sect.~\\ref{sec:VLBIdata} (Fig.~\\ref{fig:slices}). One can see strong variability in both the core and the jet at core separations larger than 1\\,mas, but never between 0.7\\,mas and 1\\,mas. Since the spectral index does not show any evidence for absorption, a natural explanation would be reduced Doppler boosting either from a misalignment between the jet direction and our line of sight or from a change in the jet speed.  A possible scenario that could account for the observed variability behaviour, especially the changing variability pattern between optical and radio emission on timescales of several years, would be a helical jet structure that is precessing or rotating. Since the optical emission is most likely produced on small scales, a small misalignment could beam the emission less favourably, which could explain the low optical state with suppressed variability from 1981 to early 1997 (see Sect.~\\ref{sec:timetest}). On the contrary, the radio emission, coming from all along the jet, would be always strongly beamed in some part of the helical path and therefore would not suffer the less intense beaming of the optical component (e.g., \\citealt{2004ApJ...615L...5R,2005ChJAA...5S.305R} and references therein). The only beaming-suppressed region at radio frequencies would be the emission gap in the inner jet. In this case the gap might move or disappear with time and could be used by future VLBI studies to prove or disprove the rotating jet scenario."
    },
    "0606/gr-qc0606017_arXiv.txt": {
        "abstract": "The dilaton-gravity sector of the Two Measures Field Theory (TMT) is explored in detail in the context of cosmology.  The model possesses scale invariance which is spontaneously broken due to the intrinsic features of the TMT dynamics. The dilaton $\\phi$ dependence of the effective Lagrangian appears only as a result of the spontaneous breakdown of the scale invariance. If no fine tuning is made, the effective $\\phi$-Lagrangian $p(\\phi,X)$ depends quadratically upon the kinetic energy $X$. Hence TMT may represent an explicit example of {\\it the effective} $k$-{\\it essence resulting from first principles without any exotic term} in the fundamental action intended for obtaining this result. Depending of the choice of regions in the parameter space, TMT exhibits different possible outputs for cosmological dynamics: a) Possibility of a power law inflation driven by the field $\\phi$ which is followed by the late time evolution driven both by a small cosmological constant and the field $\\phi$ with a quintessence-like potential. TMT enables two ways for achieving small cosmological constant without fine tuning of dimensionfull parameters: either by a {\\it seesaw} type mechanism or due to {\\it a correspondence principle} between TMT and conventional field theories (i.e theories with only the measure of integration $\\sqrt{-g}$  in the action). b) Possibility of {\\it resolution of the old cosmological constant problem}. From the point of view of TMT, it becomes clear why the old cosmological constant problem cannot be solved (without fine tuning) in conventional field theories. c) The power law inflation without any fine tuning can end with damped oscillations of $\\phi$ around the state with zero cosmological constant. d) There is a broad range of the parameters such that: the equation-of-state in the late time universe $w=p/\\rho <-1$; $w$ {\\it asymptotically} (as $t\\rightarrow\\infty$) {\\it approaches} $-1$ {\\it from below}; $\\rho$ approaches a cosmological constant, the smallness of which does not require fine tuning of dimensionfull parameters. ",
        "introduction": "The cosmological constant (CC) problem \\cite{Weinberg1}-\\cite{CC}, the accelerated expansion of the late time universe\\cite{accel}, the cosmic coincidence \\cite{coinc} are challenges for the foundations of modern physics (see also reviews on dark energy\\cite{de-review}-\\cite{Copeland}, dark matter \\cite{dm-review} and references therein). Numerous models have been proposed with the aim to find answer to these puzzles, for example: the quintessence\\cite{quint},  coupled quintessence\\cite{Amendola}, $k$-essence\\cite{k-essence}, variable mass particles\\cite{vamp}, interacting quintessence\\cite{int-q}, Chaplygin gas\\cite{Chapl}, phantom field\\cite{phantom}, tachyon matter cosmology\\cite{tachyon}, brane world scenarios\\cite{brane}, etc.. These puzzles have also motivated an interest in modifications and even radical revisions of the standard gravitational theory (General Relativity (GR))\\cite{modified-gravity1},\\cite{modified-gravity2}. Although motivations for most of these models can be found in fundamental theories like for example in brane world\\cite{extra-dim},  the questions concerning the Einstein GR limit and relation to the regular particle physics, like standard model, still remain unclear.  It is very hard to imagine that  it is possible  to propose ideas which are able to solve  the above mentioned problems  keeping at the same time unchanged the basis of fundamental physics, i.e. gravity and particle field theory. This paper may be regarded as {\\it an attempt to find a way for satisfactory answers at least to a part of the fundamental questions on the basis of first principles}, i.e. without using semi-empirical models. In this paper we explore a toy model including gravity and a single scalar field $\\phi$ in the framework of the so called Two Measures Field Theory (TMT)\\cite{GK1}-\\cite{GK8}. In TMT, gravity (or more exactly, geometry) and particle field theory intertwined in a very non trivial manner, but the Einstein's GR is restored when the fermion matter energy density is much larger than the vacuum energy density\\cite{GK6},\\cite{GK7}.  Here we {\\it have no  purpose} of constructing a complete realistic cosmological scenario. Instead, we concentrate our attention on the possible role TMT may play in resolution of a number of the above mentioned problems. The field theory model we will study here is invariant under global scale transformation of the metric accompanied  with an appropriate shift of the dilaton field $\\phi$. This scale symmetry is spontaneously broken due to intrinsic features of the TMT dynamics. Except for a peculiar structure of the TMT action, the latter does not contain any exotic terms. The obtained dynamics represents an explicit example of $k$-essence resulting from first principles.  The parameter space of the model allows to find different regions where the following different effects can take place {\\it without fine tuning of the parameters and initial conditions}:  a) Possibility of a power law inflation driven by a scalar field $\\phi$ which is  followed by the late time evolution  driven both by a small cosmological constant and a scalar field $\\phi$ with a quintessence-like potential; smallness of the cosmological constant can be achieved {\\it without fine tuning of dimensionfull parameters}.  b) In another region of the parameters, there is a possibility of the power law inflation  ended with damped oscillations of $\\phi$ around the state with zero cosmological constant (we want to emphasize again that this is realized without fine tuning of the parameters and initial conditions). Thus this scenario includes {\\it an exit from inflation together with  resolution of the old CC problem}. This becomes possible because  the basic assumptions of the Weinberg's no-go theorem\\cite{Weinberg1} are generically violated in TMT models.  c) There is a region of the parameters where the model exhibits the possibility for a superacceleration phase of the universe: in the late time universe the dark energy density $\\rho$ increases asymptotically (as $t\\rightarrow\\infty$) approaching from below to a constant value; the equation-of-state $p/\\rho =w <-1$ and $w$ {\\it asymptotically approaches} $-1$ {\\it from below}.   The organization of the paper is the following: In Sec.II we present a  review of the basic ideas of TMT. Sec.III starts from formulation of a simple scale invariant model containing all the terms respecting the symmetry of the model but without any exotic terms. In Appendix A we present equations of motion in the original frame. Using results of Appendix A, the complete set of equations of motion in the Einstein frame is given in Sec.III as well. It is shown there that if no fine tuning of the parameters is made, the effective action of our model in the Einstein frame turns out to be a k-essence type action. We start in Sec.IV from a simple case with fine tuned parameters where the non-linear dependence of the kinetic term disappears. Then three different shapes of the effective potential are possible. For the spatially flat FRW universe we study some features of the cosmological dynamics for each of the shapes of the effective potential. For one of the shapes of the effective potential, the zero vacuum energy is achieved without any fine tuning.  In Sec.V we study the cosmological dynamics without tuning parameters. It is shown that there is the possibility of a superacceleration phase of the universe, and some details of the dynamics are explored. Sec.VI is devoted to studying in more details the CC problems: a) we discuss two mechanisms for resolution of the new CC problem; b) we analyze the difference between TMT and conventional field theories (where only the measure of integration $\\sqrt{-g}$ is used in the fundamental action) which allows to understand why from the point of view of TMT the conventional field theories failed to solve the old CC problem. In Appendix B we shortly discuss what kind of model one would obtain when choosing fine tuned couplings to the measures of integration in the action. Some additional remarks concerning the relation between the structure of TMT action and the CC problem are given in Appendix C. ",
        "conclusions": "\\subsection{Differences of TMT from the standard field theory in curved space-time} The main  idea of TMT is that the general form of the action $\\int L\\sqrt{-g}d^{4}x$ is not enough in order to account for some of the fundamental problems of particle physics and cosmology. The key difference of TMT from the conventional field theory in curved space-time consists in the hypothesis\\cite{GK2}-\\cite{GKatz} that in addition to the term in the action with the volume element $\\sqrt{-g}d^{4}x$ there should be one more term where the volume element is metric independent but rather it is determined either by four (in the 4-dimensional space-time) scalar fields $\\varphi_{a}$ or by a three index potential $A_{\\alpha\\beta\\gamma}$,  see Eqs.(\\ref{S})-(\\ref{Aabg}). We would like to emphasize that including in the action of TMT the coupling of the Lagrangian density $L_{1}$ with the measure $\\Phi$, we modify in general both the gravitational and matter sectors as compared with the standard field theory in curved space-time. Besides we made two more assumptions: the measure fields ($\\varphi_{a}$ or $A_{\\alpha\\beta\\gamma}$) appear only in the volume element; one should proceed in the first order formalism. {\\it These assumptions constitute all the modifications of the general structure of the theory we have made as compared with the conventional  field theory where only the measure of integration $\\sqrt{-g}$ is used in the action principle}. In fact, the Lagrangian densities $L_{1}$ and $L_{2}$ studied in the present paper, contain  only such terms which should be present in a conventional model with minimally coupled to gravity scalar field. In particular there is no need for the non-linear in kinetic energy terms as well as in the phantom type terms in the fundamental Lagrangian densities $L_{1}$ and $L_{2}$ in order to obtain a super-acceleration phase at the late time universe.  After making use of the variational principle and formulating the resulting equations in the Einstein frame, we have seen that the effective action (\\ref{k-eff}) represents a concrete realization of the $k$-essence\\cite{k-essence} obtained from first principles of TMT without any exotic terms in the Lagrangian densities.  \\subsection{Short summary of results}  \\subsubsection{The early universe inflation}  As $\\delta =0$, the dynamics of $\\phi$ can be analyzed by means of its effective potential (\\ref{Veffvac-delta=0}). As $\\phi\\ll -M_{p}$ the effective $\\phi$ potential has the exponential form and it is proportional to the integration constant $M^4$. In other words, the effective potential governing the dynamics of the early universe results from the spontaneous breakdown of the global scale invariance (\\ref{st}) caused by the intrinsic feature of TMT (see Eqs. (\\ref{varphi}) and (\\ref{app1})). We have seen that independently of the values of the parameters $V_{1}$, $V_{2}$ and under very general initial conditions, solutions rapidly approach a regime characterized by a power law inflation. If $\\delta\\neq 0$, we deal with the intrinsically $k$-essence dynamics. The numerical solutions in this case have showed that there is no qualitative difference from the  power law inflation obtained in the case with $\\delta =0$.  \\subsubsection{End of inflation}  In our toy model there are three regions of the parameters $V_{1}$ and $V_{2}$ and appropriate three shapes of the effective potentials, Fig.1. Therefore three different types of scenarios for exit from inflation can be realized:  a)  $V_{1}<0$ and $V_{2}<0$, \\,  Sec.IVB. In this case the power law inflation ends with damped oscillations of $\\phi$ approaching the point of the phase plane ($\\phi =\\phi_0, \\dot\\phi =0$) where the vacuum energy $V_{eff}^{(0)}(\\phi_0)=0$. This occurs without fine tuning  of the parameters $V_{1}$, $V_{2}$ and the initial conditions.  b) $V_{1}>0$ and $b_{g}V_{1}>2V_{2}$, \\, Sec.IVC. In this case the power law inflation monotonically transforms to the late time inflation asymptotically governed by the cosmological constant $\\Lambda_1$.  Qualitatively the same results are also obtained in  the cases a) and b) if $\\delta\\neq 0$.  c) $V_{2}<b_{g}V_{1}<2V_{2}$, \\, Sec.IVD. In this case the power law inflation ends without oscillations at the final value $\\phi_{min}$, corresponding to the (non zero) minimum of the effective potential.  The model we have studied in this paper may be extended by including the Higgs field, as well as gauge fields and fermions. It turns out that the scalar sector of such an extended model enables a scenario which resembles a hybrid inflation\\cite{hybrid}. These results will be presented in a future publication.   \\subsubsection{Cosmological constant problems}  1) {\\it The old cosmological constant problem}. In Sec.IVA we have seen in details that if $V_1<0$ then, for a broad range of other parameters, the vacuum energy turns out to be zero without fine tuning. This effect is a direct consequence of the TMT structure which yields the following results: a) the effective scalar sector potential generated in the Einstein frame is proportional to {\\it a perfect square} of two terms; b) one of those terms is proportional to the integration constant $\\pm M^4$ the appearance of which is also the intrinsic feature of TMT. Note that the spontaneously broken global scale invariance is not necessary to achieve this effect\\cite{GK3}. If such type of the structure for the scalar field potential in a conventional (non TMT) model would be chosen \"by hand\" it would be a sort of fine tuning.  In Sec.VI we have explained in details how this result avoids the well known no-go theorem by Weinberg\\cite{Weinberg1} stating that generically in field theory one cannot achieve zero value of the potential in the minimum without fine tuning. It is interesting that the resolution of the old CC problem in the context of TMT happens in the regime where $\\zeta\\rightarrow\\infty$. From the point of view of TMT, the latter is the answer to the question why the old cosmological constant problem cannot be solved (without fine tuning) in theories with only the measure of integration $\\sqrt{-g}$  in the action.  2) {\\it The new cosmological constant problem}. Interesting result following from the general structure of the scale invariant TMT model with $V_1>0$ is that the cosmological constant $\\Lambda$, Eq.(\\ref{lambda}), is a ratio of quantities constructed from pre-potentials $V_{1}$, $V_{2}$ and the dimensionless parameter $b_{g}$. Such structure of $\\Lambda$ allows to propose two  ways (see Sec.VI) for resolution of the problem of the smallness of $\\Lambda$ that should be $\\Lambda\\sim (10^{-3}eV)^{4}$:  \\quad a) The first way is a kind of a {\\it seesaw} mechanism\\cite{seesaw}. For instance, if  $V_{1}\\sim (10^{3}GeV)^{4}$ and $V_{2} \\sim (10^{18}GeV)^{4}$ then $\\Lambda_1\\sim (10^{-3}eV)^{4}$.  \\quad b) The second way is realized if the dimensionless parameters $b_{g}$, $b_\\phi$ and $V_2/V_1$ of the action (\\ref{totaction}) are huge numbers of the close orders of magnitude. For example, if $V_{1}\\sim (10^{3}GeV)^{4}$ then for getting $\\Lambda\\sim (10^{-3}eV)^{4}$ one should assume that $b_{g}\\sim 10^{60}$. Possibility of this idea means that the resolution of the new cosmological constant problem may have a certain relation to {\\it the correspondence principle} between TMT and conventional field theories (see details in Sec.VIA2).   \\subsubsection{Super-acceleration phase of the Universe. } If no fine tuning of the parameters is made in the fundamental action, namely if $b_g\\neq b_{\\phi}$, then our TMT model has big enough regions in the parameter space where the super-acceleration phase in the late time universe becomes possible. The appropriate phantom dark energy asymptotically approaches a cosmological constant. However  it is impossible to obtain {\\it a pure classical solution} which connects the early universe power law inflation with the late time super-acceleration. This problem  is apparently related with the toy character of the scenario where the role of the matter creation has been ignored: in TMT the fermionic matter generically contributes to the constraint equation for the scalar field $\\zeta$ and so can effect the field $\\phi$ dynamics as well.  \\subsection{What can we expect from quantization}   In this paper we have studied only classical TMT and its possible effects in the context of cosmology. However quantization of TMT as well as influence of quantum effects on the processes explored in this paper may have a crucial role. We summarize here some ideas and speculations which gives us a hope that quantum effects can keep the main results of this paper.  Recall first two fundamental facts of TMT as a classical field theory: (a) The measure degrees of freedom appear in the equations of motion only via the scalar $\\zeta$, Eq.(\\ref{zeta}); (b) The scalar $\\zeta$  is determined (as a function of matter fields, in our toy model - as a function of $\\phi$) by the constraint which is nothing but a consistency condition of the equations of motion (see Eqs.(\\ref{app1})-(\\ref{app3}) in Appendix A and Eq.(\\ref{constraint2})). Therefore the constraint plays a key role in TMT. Note however that if we were ignore the gravity from the very beginning in the action (\\ref{totaction}) then instead of the constraint (\\ref{app3}) we would obtain Eq.(\\ref{app1}) (where one has to put zero the scalar curvature). In such a case we would deal with a different theory. This notion shows that the gravity and matter intertwined in TMT in a much more complicated manner than in GR. Hence introducing the new measure of integration $\\Phi$ we have to expect that the quantization of TMT may be a complicated enough problem. Nevertheless we would like here to point out that in the light of the recently proposed idea of Ref.\\cite{Giddings}, the incorporation of four scalar fields $\\varphi_a$  together with the scalar density $\\Phi$, Eq.(\\ref{Phi}), (which in our case are the measure fields and the new measure of integration respectively), is  a possible way to define local observables in the local quantum field theory approach to quantum gravity. We regard this result as an indication that the effective gravity $+$ matter field theory has to contain the new measure of integration $\\Phi$ as it is in TMT.  The assumption formulated in item 2 in Sec.IIA, that the measure fields $\\varphi_a$ (or $A_{\\alpha\\beta\\gamma}$) appear in the action (\\ref{S}) {\\it only} via the measure of integration $\\Phi$, has a key role in the TMT results and in particular for the resolution of the old cosmological constant problem. In principle one can think of breakdown of such a structure by quantum corrections. However, TMT possesses an infinite dimensional symmetry mentioned in item 2 of Sec.II which, as we hope, is able to protect  the postulated structure of the action from a deformation caused by quantum corrections. Another effect of quantum corrections is the possible appearance of a nonminimal coupling of the dilaton field $\\phi$ to gravity in the form like for example $\\xi R\\phi^2$. Proceeding in the first order formalism of TMT one can show that the nonminimal coupling can affect the k-essence dynamics but the mechanism for resolution of the old CC problem exhibited in this paper remains unchanged. This conclusion together with expected effect of quantum corrections on the scale invariance (see our discussion in the paragraph after Eq.(\\ref{Veff=0})) allows us to hope that the exhibited resolution of the old CC problem holds in the quantized TMT as well.  Quantization of TMT being a constrained system requires developing the Hamiltonian formulation of TMT. Preliminary consideration shows that the Einstein frame appears in the canonical formalism in a very natural manner. A systematic exploration of TMT in the canonical formalism will be a subject of forthcoming research."
    },
    "0606/astro-ph0606266_arXiv.txt": {
        "abstract": "The recent three-year WMAP data have confirmed the anomaly concerning the low quadrupole amplitude compared to the best-fit $\\Lambda$CDM prediction. We show that, allowing the large-scale spatial geometry of our universe to be plane-symmetric with eccentricity at decoupling or order $10^{-2}$, the quadrupole amplitude can be drastically reduced without affecting higher multipoles of the angular power spectrum of the temperature anisotropy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606500_arXiv.txt": {
        "abstract": "We examine the practice of deriving interstellar medium (ISM) abundances from low-resolution spectroscopy of GRB afterglows. We argue that the multi-ion single-component curve-of-growth analysis technique systematically underestimates the column densities of the metal-line profiles commonly observed for GRB.  This systematic underestimate is accentuated by the fact that many GRB line-profiles (e.g.\\ GRB~050730, GRB~050820, GRB~051111) are comprised of `clouds' with a bi-modal distribution of column density.  Such line-profiles may be characteristic of a sightline which penetrates both a high density star-forming region and more distant, ambient ISM material. Our analysis suggests that the majority of abundances reported in the literature are systematically underestimates and that the reported errors are frequently over-optimistic. Further, we demonstrate that one cannot even report relative abundances with confidence. The implications are profound for our current understanding on the metallicity, dust-to-gas ratio, and chemical abundances of the ISM in GRB host galaxies.  For example, we argue that all but a few sightlines allow for the gas to have at least solar metallicity.  Finally, we suggests new approaches for constraining the abundances. ",
        "introduction": "With the launch of the {\\it Swift} satellite \\citep{gcg+04}, the rate of GRB detections has increased by more than an order of magnitude \\citep{gehrels06}.  In turn, one has seen roughly the same increase in the detection of bright afterglows. Though a few high-resolution spectra were obtained pre-Swift \\citep{cgh+03,fdl+05}, Swift has enabled the nearly routine acquisition of high signal-to-noise ratio (SNR), high-resolution spectra \\citep{cpb+05,pro_finegrb06}.  In addition,  there are now over a dozen low-resolution observations from pre-Swift GRB \\citep[e.g.][]{kdo+99,bsc+03} and the first year of Swift operation \\citep[e.g.][]{foley06,bpck+05,fynbo060206}. Observers are now progressing toward the examination of distributions of gas properties in the interstellar medium (ISM) of GRB host galaxies.  Because of the apparent faintness of many GRB afterglows and also instrument availability, the majority of GRB spectra are acquired with low-resolution spectrometers (FWHM$>1$\\AA) and at moderate to poor SNR ($<50$ per resolution element).  In these cases, abundance analysis must be performed on the equivalent width measurements of unresolved metal-line transitions. To date, the standard practice has been to perform a multi-ion single-component curve-of-growth (MISC-COG) analysis \\citep[e.g.][]{sff03} or single-component profile fits to unresolved data \\citep[e.g.][]{sf04,watson050401} which is fundamentally the same analysis. In this analysis, one derives an effective Doppler parameter $b_{eff}$ which principally characterizes the velocity width of strong, heavily saturated transitions.  The column density is then constrained with weaker, yet still potentially saturated transitions. In contrast, the \\ion{H}{1} column density can be determined from fits to the damping wings of the \\lya\\ profile even with low-resolution data \\citep[e.g.][]{phw05}.  The curve-of-growth technique dates far back \\cite[e.g.][]{unsold30,wilson39} and was primarily introduced to perform abundance analysis on lower spectral resolution data. The literature on the pitfalls of MISC-COG analysis is extensive \\citep{nh73,cru75,spitzer75}.  These include systematic differences in the results for ions of differing ionization state \\citep[e.g.\\ \\ion{Mg}{1} vs.\\ \\ion{Fe}{2};][]{spitzer75}, `hidden' saturation due to significant variations in the Doppler parameter of overlapping components \\citep{nh73}, and the dangers of mixing refractory and non-refractory elements \\citep{jss86}.  The Galactic ISM community is well aware of these issues; it has taken every effort to obtain high resolution spectroscopy and when limited to low resolution data, very cautiously interpret the observations. In short, few of the problems discussed in this paper are unique to GRB research, but we note that the discussion may translate to other abundance analyses with large, modern datasets of low resolution spectroscopy \\citep[e.g.][]{sga+04,trn+05}.  \\cite{jenkins86}  examined the accuracy of a single-component COG analysis on multi-component line-profiles for a range of reasonable distributions of column density and Doppler parameter. He showed that the COG results typically give results accurate to within 20\\% (i.e.\\ $<0.2$\\,dex) of the correct column density provided the $N$ and $b$ distributions of the individual components are well-behaved and so long as the peak optical depth $\\tau_0 < 5$. Jenkins also emphasized, however, that the results would likely significantly underestimate the column densities if the $N$ or $b$ values were bi-modally distributed \\citep[see also][]{nh73}.  We will find that this is frequently the case for GRB line-profiles.   In this paper, we will examine the main issues related to the ISM abundance analysis of GRB host galaxies as derived from low-resolution spectroscopy, i.e.\\ equivalent width analysis. We begin with an analysis of the GRB~051111 sightline where Keck/HIRES observations permit one to directly test the MISC-COG technique ($\\S$~\\ref{sec:051111}). The discouraging results are enlightening, namely the COG analysis systematically underestimates the column density measurements. We consider an additional example from the literature (GRB~020813) and raise similar concerns ($\\S$~\\ref{sec:020813}). Finally, we offer a set of guidelines to consider when analyzing low-resolution GRB spectra ($\\S$~\\ref{sec:proc}). In a future paper \\citep{pro_dataI06}, we will apply these guidelines to our datasets and previously published results with the intent of providing a uniform, database of measurements with realistic error estimates.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=3.5in]{f1.eps} \\caption{\\ion{Fe}{2} profiles from the ISM foreground to GRB~051111 in the original data (LHS) and smoothed to lower resolution (RHS; FWHM~$\\approx 5$\\AA) to mimic the line-profiles frequently observed. The velocity $v=0$ corresponds to $z=1.54948$. The rest equivalent width values labeled on the Figure correspond to the full line-profile. } \\label{fig:fe051111} \\end{center} \\end{figure} ",
        "conclusions": "Before concluding, we wish to briefly comment on the abundance measurements for the ISM of GRB presented in the literature. Table~\\ref{tab:summ} summarizes the results to date. The last two columns give the minimum and maximum metallicity values allowed by the data as reported in the literature\\footnote{We will ignore (for the time being) the fact that Zn is a trace element and that a large Zn/H ratio does not require a high gas metallicity.}.  One notes that solar (and even super-solar) abundances are allowed in nearly every case. It is evident from Table~\\ref{tab:summ} that the principal conclusion of \\cite{sff03} -- GRB sightlines have extreme metal column densities -- may be a remarkable understatement. We also caution that the common conclusion that GRB occur in a metal-poor ISM ([M/H]~$\\leq -1$) may need serious revision.   It is unfortunate that there are few transitions for mild or non-refractory elements that are weak enough to place meaningful upper limits to the abundances; the Zn and Si abundances, for example, are compromised by their relatively strong transitions. We expect, however, that an analysis of the lines listed in Table~\\ref{tab:wklin} may improve the constraints \\citep{pro_dataI06} and provide more realistic metallicity measurements for the gas within GRB host galaxies."
    },
    "0606/astro-ph0606736_arXiv.txt": {
        "abstract": "We report the detection of both the 67.9 and 78.4\\un{keV} \\ti{Sc} \\gammaray lines in Cassiopeia~A with the \\integ \\ibisgri instrument. Besides the robustness provided by spectro-imaging observations, the main improvements compared to previous measurements are a clear separation of the two \\ti{Sc} lines together with an improved significance of the detection of the hard \\xray continuum up to 100\\un{keV}. These allow us to refine the determination of the \\ti{Ti} yield and to constrain the nature of the nonthermal continuum emission. By combining COMPTEL, \\sax/PDS and \\isgri measurements, we find a line flux of (2.5 $\\pm$ 0.3) $\\times$ 10$^{-5}$ cm$^{-2}$ s$^{-1}$ leading to a synthesized \\ti{Ti} mass of 1.6 $^{+0.6}_{-0.3}$ $\\times$ 10$^{-4}$ M$_{\\odot}$. This high value suggests that Cas~A is peculiar in comparison to other young supernova remnants, from which so far no line emission from \\ti{Ti} decay has been unambiguously detected. ",
        "introduction": "\\label{s:intro}  Cassiopeia~A (hereafter, Cas~A) is the youngest known supernova remnant (SNR) in the Milky Way, located at a distance of 3.4$^{+0.3}_{-0.1}$\\un{kpc} \\cite{c:reed95}. The estimate of the supernova is A.D.~1671.3$\\pm$0.9, based on the proper motion of several ejecta knots \\cite{c:thorstensen01}. However, an event observed by Flamsteed (A.D.~1680) could be at the origin of the Cas~A remnant \\cite{c:ashworth80,c:stephenson02}. The large collection of data from observations in the radio, infra-red, optical, \\xray (see \\eg Hwang \\etal 2004) up to TeV \\gammarays \\cite{c:aharonian01} allows us to study its morphology, composition, cosmic-ray acceleration efficiency and secular evolution in details. Young SNRs are thought to be efficient particle accelerators and represent the main galactic production sites of heavy nuclei, some of them being radioactives. Soft \\gammaray observations, beyond the thermal \\xray emission ($\\geq$ 10\\un{keV}), can therefore provide invaluable information in both of these areas by studying the nonthermal continuum and the \\gammaray line emission. Cas~A then appears to be the best case for such investigations.  Few radioactive isotopes are accessible to \\gammaray astronomy for probing cosmic nucleosynthesis \\cite{c:diehl98}. Amongst them, \\ti{Ti} is a key isotope for the investigation of the inner regions of core-collapse SNe and their young remnants. This nucleus is thought to be exclusively created in SNe but with a large variation of yields depending on their type. Recent accurate measurements by several independent groups give a weighted-average \\ti{Ti} lifetime of 86.0 $\\pm$ 0.5 years \\cite{c:ahmad98,c:gorres98,c:norman98,c:wietfeldt99,c:hashimoto01}. The discovery of the 1157\\un{keV} \\ti{Ca} $\\gamma$-ray line emission from the decay chain of \\ti{Ti} ($^{44}$Ti$\\longrightarrow$$^{44}$Sc$\\longrightarrow$$^{44}$Ca) with \\gro/COMPTEL \\cite{c:iyudin94} was the first direct proof that this short-lived isotope is indeed produced in SNe. This has been strengthened by the \\sax/PDS detection of the two blended low energy \\ti{Sc} lines at 67.9\\un{keV} and 78.4\\un{keV} \\cite{c:vink01}. By combining both observations, Vink \\etal (2001) deduced a \\ti{Ti} yield of (1.5$\\pm$1.0) $\\times$ 10$^{-4}$ M$_{\\odot}$.  This high value compared to those predicted by \"standard\" models (\\eg Woosley \\& Weaver 1995b, WW95; Thielemann, Nomoto, \\& Hashimoto 1996, TNH96) as well as improved ones \\cite{c:rauscher02,c:limongi03} could be due to several effects. First of all, the explosion of Cas~A seems to have been intrinsically asymmetric since such asymmetries have recently been observed in the ejecta \\cite{c:vink04,c:hwang04}, and there are indications that its explosion energy was $\\sim$ 2 $\\times$ 10$^{51}$ erg \\cite{c:laming03}, higher than the canonical value of 10$^{51}$ erg. The sensitivity of the \\ti{Ti} production to the explosion energy and asymmetries may explain the high \\ti{Ti} yield compared to explosion models \\cite{c:nagataki98}.  It is generally accepted that Cas~A was formed by the explosion of a massive progenitor, from a 16 M$_{\\odot}$ single star \\cite{c:chevalier03} to a Wolf-Rayet (WR) remnant of a very massive ($<$ 60 M$_{\\odot}$) precursor \\cite{c:fesen91}. Type Ib explosions, originating from progenitors which have experienced strong mass loss (see Vink 2004, Vink 2005), should on average produce more \\ti{Ti} due to the lower fall back of material on the compact stellar remnant \\cite{c:wlw95a}. However, there is some debate on the detailed stellar evolution scenario that may have accounted for the low mass of the star prior to the explosion. The amount of oxygen present (1-2 M$_{\\odot}$, Vink \\etal 1996) suggests a main sequence mass of 20 M$_{\\odot}$. This may be too low to form a Type Ib progenitor by mass loss in a WR phase. Moreover, the high surrounding density is better explained if the shock wave is moving through the dense wind of a red supergiant rather than the more tenuous wind of a WR. Therefore, it has been recently suggested that the low mass of the progenitor is the result of a common envelope evolutionary phase in a binary system \\cite{c:young06}. The authors demonstrated that such a scenario of a 15-25 M$_{\\odot}$ progenitor which lost its hydrogen envelope due to a binary interaction can match the main observational constraints. In any case, the \\ti{Ti} production is highly sensitive to details of the explosion as well as nuclear reaction rates. It is of interest to point out that the major \\ti{Ti} production reaction \\ca{Ca}($\\alpha$,$\\gamma$)\\ti{Ti} has been revised \\cite{c:nassar06}, implying an increase of the \\ti{Ti} production by a factor of $\\sim$ 2.  In addition to the \\ti{Sc} \\gammaray lines, the hard \\xray spectrum is also of interest for its nonthermal continuum emission and because this underlying continuum is critical to properly measure the \\ti{Sc} line flux. Nevertheless, its nature is still under debate. The nonthermal hard \\xray continuum could be due to either synchrotron radiation of TeV electrons \\cite{c:allen97} or nonthermal bremsstrahlung from supra-thermal electrons which have been accelerated by internal shocks (Laming 2001a,b; Vink \\& Laming 2003). Both cases predict a gradual steepening at high energies and then, reliable continuum flux measurements beyond the two low energy \\ti{Sc} lines ($>$ 80\\un{keV}) are necessary, as initiated with \\gro/OSSE \\cite{c:the96}. Soft \\gammaray observations are therefore critical to better understand the nucleosynthesis and the particle acceleration processes in young SNRs such as Cas~A. \\ibis \\cite{c:ubertini03}, one of the two main coded mask aperture instruments onboard the \\integ satellite \\cite{c:winkler03}, is best suited to study both the hard \\xray continuum and the line emission thanks to its low energy (15\\un{keV} -- 1\\un{MeV}) camera \\isgri \\cite{c:lebrun03}. \\ibisgri provides spectro-imaging (13\\m~FWHM, 6\\un{keV} FWHM at 70\\un{keV}) over a large field of view (400 deg$^{2}$) in the energy range 15\\un{keV}-1\\un{MeV} with a milliCrab sensitivity at 70 keV (3 $\\sigma$, $\\Delta$E/E = 2, 10$^{6}$ s). The large field of view allows for long exposures devoted to the simultaneous observation of several sources. In this letter, we report the results of the spectro-imaging analysis of Cas~A based on \\ibisgri observations. ",
        "conclusions": "\\label{s:discuss}  The \\ibisgri observations confirm the presence of the two low energy \\ti{Sc} \\gammaray lines in Cas~A. By performing a weighted average of the three independent measurements of COMPTEL, \\sax/PDS \\cite{c:vink01} and \\isgri, we find a line flux of (2.5 $\\pm$ 0.3) $\\times$ 10$^{-5}$ cm$^{-2}$ s$^{-1}$. Taking into account uncertainties on its age \\cite{c:thorstensen01}, distance \\cite{c:reed95} and \\ti{Ti} lifetime \\cite{c:vink05}, this is translated into an initial synthesized \\ti{Ti} mass of (1.6$^{+0.6}_{-0.3}$) $\\times$ 10$^{-4}$ M$_{\\odot}$. This mass of ejected \\ti{Ti} is generally thought to be unusually large (or for few specific cases, marginally consistent) in comparison with spherical explosion models of WW95 and TNH96 \\cite{c:timmes96}. Moreover, in the standard frame where \\ti{Ti} and \\fe{Ni} are co-produced during the first stages of the explosion, Cas~A should have been a very bright, \\fe{Ni}-rich SN, in contrast with its non detection or with the Flamsteed's historical record. However, the large \\ti{Ti}/\\fe{Ni} ratio could be explained by the high degree of asymmetries \\cite{c:nagataki98}.The high \\ti{Ti} yield thus supports the idea that Cas~A is the result of an asymmetric and/or a relatively more energetic explosion, consistent with other observational evidence \\cite{c:vink04,c:hwang04}.  Anyway, the \\ti{Ti} production in core-collapse SNe is highly sensitive to the network used to compute nuclear reactions. With the recent revised \\ca{Ca}($\\alpha$,$\\gamma$)\\ti{Ti} reaction rate \\cite{c:nassar06}, theoretical models become more compatible with the \\ti{Ti} yield deduced from \\ibisgri and previous observations. However, this would make the lack of other Galactic \\ti{Ti} sources an even more serious problem: several \\gammaray line surveys \\cite{c:dupraz97,c:renaud04,c:the06} have highlighted the problem of the \"young, missing, and hidden\" Galactic SNe, those that should have occurred since Cas~A and are still not detected through the line emission from \\ti{Ti} decay. This would strengthen the idea that Cas~A is peculiar \\cite{c:young06}. On the other hand, the high \\ti{Ti} yield of both Cas~A and SN~1987A \\cite{c:fransson01} is more in accordance with the solar \\ti{Ca}/\\fe{Fe} ratio, whereas this ratio is underpredicted by current spherically symmetric explosive nucleosynthesis models \\cite{c:prantzos04,c:young06}.  Besides the robustness provided by these \\ibisgri spectro-imaging observations, the main improvements compared to previous observations \\cite{c:vink01,c:rothschild03} are the improved spectral resolution and the improved significance of the detection of the hard \\xray nonthermal continuum up to 100\\un{keV} well fitted by a single power-law. The latter gives more stringent constraints on both the line intensities and the underlying continuum. Therefore, the scenario of a synchrotron radiation by TeV electrons \\cite{c:allen97} as modeled by Reynolds \\& Keohane (1999) seems not appropriate in the case of Cas~A. On the other hand, the model developed by Laming (2001a,b) implying a nonthermal bremsstrahlung emission of supra-thermal electrons could be an alternative scenario. Based on this firm detection of the \\ti{Sc} lines with IBIS/ISGRI, the expected results with SPI, thanks to its fine spectral resolution ($\\Delta$E $\\sim$ 2\\un{keV} FWHM at 1\\un{MeV}), should help us for the first time to constrain the kinematics of the innermost layers of the explosion (Vink \\etal 2006, in preparation)."
    },
    "0606/astro-ph0606446_arXiv.txt": {
        "abstract": "This study investigates the expansion properties of nickel (Ni) bubbles in Type Ia supernovae (SNe Ia) due to radioactive heating from the $\\Ni \\rightarrow\\ \\co \\rightarrow\\ \\fe$ decay sequence, under the spherically symmetric approximation. An exponentially-declining medium is considered as the ejecta substrate, allowing for the density gradient expected in a Type Ia supernova (SN Ia). The heating gives rise to an inflated Ni bubble, which induces a forward shock that compresses the outer ambient gas into a shell. As the heating saturates, the flow tends toward a freely-expanding state with the structure frozen into the ejecta. The thickness of the shell is $\\sim 1\\%$ of the radius of the bubble, and the density contrast across the shell reaches $\\chi \\ga 100$ in a narrow region that is limited by numerical resolution.  In the adiabatic case (in which no radiative energy diffuses across the shell) the structure of the shell can be approximately described by a self-similar solution that is determined by the expansion rate and ambient density gradient of the shell. In the radiative case, the shell expansion weakens but remains comparable to the adiabatic solution. The density contrast of the inferred ejecta clumps (shell components) is greater than that given by the model that uses a uniform ejecta substrate, while the interaction of the clumps with the remnant is delayed to a more advanced evolutionary stage. The explosion parameters of the SN are varied to examine whether the created clump characteristics are consistent with those of the ejecta knots that are present at the edges of Tycho's remnant. Explosion conditions similar to those of successful explosion models are found, including the deflagration W7 and the delayed detonation yield with sufficient ejection velocities as well as timely initiations for the clump-remnant interaction, whereas the luminous helium (He) detonations and the low-energy Chandrasekhar-mass explosions are unfavorable. The clump speed can be increased as the initial density contrast of $\\Ni$, $0.5 \\la \\omega < 1$, is reduced, as determined by a realistic elemental distribution. In all cases, the occurrence of the reverse shock (RS) impact of the clump with the remnant is expected in under 2000 yrs after the supernova explodes. ",
        "introduction": "Observations of SN 1987A have shown that the distribution of Fe in the ejecta is not what would be expected in the simplest models; it has higher velocities than expected and a large filling factor for its mass of $0.07\\Msun$, as determined from the supernova light curve (McCray 1993; Li {\\it et al.} 1993). A plausible mechanism for the large filling factor is the Ni bubble effect, in which the radioactive progenitors of the Fe expand relative to their surroundings because of the deposition of radioactive power (Woosley 1988; Li {\\it et al.} 1993; Basko 1994). This effect is important during the first $\\sim$ ten days after the supernova is formed, when the radioactive power is significant and the diffusion of energy has not yet become important.  An expected effect of the Ni bubble expansion is to create clumps in the nonradioactive gas of intermediate elements. Widespread evidence indicates the ejecta of core-collapse supernovae are clumpy. The oxygen line profiles in the nearby Type II supernovae SN 1987A (Stathakis {\\it et al.} 1991) and SN 1993J (Spyromilio 1994; Matheson {\\it et al.} 2000) showed evidence of the structure, implying that the gas is clumped. The velocity range of the emission extends to $1,500\\kms$ in SN 1987A and $4,000\\kms$ in SN 1993J. Clumping was also evident in the prototypes of the oxygen-rich supernova remnants (SNRs) Cassiopeia A (Cas A) (Hughes {\\it et al.} 2000; Hwang {\\it et al.} 2000) and Puppis A (Winkler {\\it et al.} 1988), and in the Type Ib SN 1985F (Filippenko \\& Sargent 1989). As for SNe Ia, observations of SN1006 suggested that Si freely expands in the velocity range 5,600 - 7,000 $\\kms$, while its Fe mass is only estimated as $0.075-0.16 \\Msun$ (Hamilton {\\it et al.} 1997). Moreover, in Tycho's SNR (SN 1572), X-ray observations revealed two knots, one enriched in Si and the other in Fe, % protruding non-deceleratingly from the SE edge of the remnant at a velocity of $8,300 (D/2.5\\ \\rm kpc) \\kms$, % where $D$ is the distance (Vancura {\\it et al.} 1995; Hwang \\& Gotthelf 1997; Hughes 1997). X-ray spectra of Tycho's remnant % further suggested that the Fe line emitting gas is in general at a higher temperature and lower density than than the Si line emitting gas (Hwang {\\it et al.} 1998); % this result is also expected from the Ni bubble effect.   Wang \\& Chevalier (2001 and 2002, hereafter WC01 and WC02) studied the role of ejecta clumps in the evolution of SNRs. The ejecta knots associated with Tycho's remnant require a free expansion velocity $v \\sim 7,000 \\kms$ and a density contrast $\\chi\\ga 100$ relative to the surrounding ejecta, to survive crushing and remain distinct features in the remnant in the present epoch (WC01). The remarkable protrusions (Aschenbach {\\it et al.} 1995) on the periphery of the Vela SNR are also likely to be caused by ejecta clumps, and we estimated that $\\chi \\sim 1000$ and $v \\sim 3,000$ are needed to create the structure (WC02). My previous work (Wang 2005, hereafter W05) further examined the process of $\\Ni$ radioactive heating % as a plausible mechanism for forming such ejecta clumps, taking into account the effect of radiative diffusion. In the case of SNe Ia, the created clump properties seem compatible with those needed by Tycho's knots. This result nevertheless contrasts with the case of core-collapse SNe like Vela's bullets, whose required high compression is not anticipated in the simplest Ni bubble scenario.    Evidence may exist that for SNe Ia, ejecta clumping occurs around one large central Ni bubble. In radio observations of Tycho's remnant, Reynoso {\\it et al.} (1997) and Velazquez {\\it et al.} (1998) found a wide and regularly-spaced emission in the NW edge of the remnant, which extends very close to the blast wave. The large radius of the structure is not attributable to the intershock instabilities that grow from small perturbations, but is effectively explained by the expansion of a spherical shell of clumpy ejecta into the intershock region of the remnant (WC01). In fact, compared with various extensively observed Galactic Type II SNRs which in general exhibit quite complex abundance structures (such as in Hughes {\\it et al.} 2000 and Hwang, {\\it et al.} 2000 for Cas A; and Park {\\it et al.} 2002 for G292.0+1.8), both observations and modeling of X-ray emissions of Type Ia SNRs revealed a simple and stratified composition % (Badenes {\\it et al.} 2003, 2005; Decourchelle {\\it et al.} 2001 for Tycho; and Lewis {\\it et al.} 2003 for N103B). The stratification of ejecta of SNe Ia seems to be a natural consequence of the lack of neutrino-driven convection (Kifonidis {\\it et al.} 2000) and pulsar winds (Blondin {\\it et al.} 2001b) in them, which two mechanisms can mix elements and enhance compression. Conversely, the disturbance that is expected in core-collapse SNe may greatly complicate the evolution of the Ni bubble on a large scale.   Dwarkadas \\& Chevalier (1998, hereafter DC98) reviewed the solutions for the density structure of SNe Ia obtained by various 1-D explosion models; they found that an exponential profile, which effectively evolves from a steep power law profile to a flatter one, is the best overall approximation to the density profile that is obtained by summing over individual chemical elements. However, in W05, a uniform density distribution (comparable to the inner flat component of an $n=8$ power-law ejecta density model with the density $\\rho \\propto r^{-n}$ in the outer parts) was assumed, so the results may not be suitable for SNe Ia. Furthermore, the W05 models of SNe Ia show that the propagation of the Ni bubble is only slightly faster than the free motion of the ejecta, suggesting that the case of SNe Ia is more susceptible than the case of core-collapse SNe to the ejecta substrate structure. However, how the actual exponential decline in density affects the solution is unclear.   The aim of this work is to determine the properties of the Ni bubble-driven shell as ejecta clumps associated with SNe Ia, using a realistic ejecta structure. A simple, adiabatic and spherically symmetric scenario is first considered to determine how an evolving exponential density gradient influences the predicted clump properties. Radiative-transport radiation hydrodynamics (RHD) are then adopted to include the effect of the diffusion of radiative energy across the Ni bubble structure to show how the radiative solution differs. The explosion parameters used in our calculations are similar to those of the 1-D explosion models of Hoflich \\& Khokhlov (1996, hereafter HK96) and Hoflich {\\it et al.} (1998, hereafter HWK98). We note that the radiative transport process due the $\\Ni$ radioactive decay has not before been incorporated in the hydrodynamic (HD) phase of the explosion models, and such a proper treatment is particularly difficult for numerical convergence in the case of SNe Ia . \\S~2 elucidates our computational setup and methods. \\S~3 presents the evolutionary properties of the $\\nni$ bubble shell. \\S~4 draws on the self-similar solution for a pulsar bubble, to provide insight into the shell structure. \\S~5 compares the inferred ejecta clump properties of various exponential models. \\S~6 draws conclusions.     \\section {DENSITY STRUCTURE IN TYPE Ia SUPERNOVAE}  SNe Ia are widely regarded as being thermonuclear explosions of carbon-oxygen (C-O) white dwarfs (WDs). Soon after the explosion, the ejecta freely expand so that each gas element moves with a constant velocity $v=r/t$ and its density drops to $t^{-3}$, as in a spherical expansion. DC98 showed that, in this phase, the density distribution of valid SN Ia explosion models can generally be described by an exponential profile, \\begin{equation} \\rho_{SN} = A \\exp(-v/v_e) \\  t^{-3}, \\end{equation} where $A$ is a constant and $v_e$ is another constant called the velocity scale height, which are determined by the total mass $M$ and kinetic energy $E$ of the ejecta. The two constants can be obtained by integrating the mass density and the kinetic energy density over space: $M = 8 \\pi A v_e^3$ and $E=48 \\pi A v_e^5$, or \\begin{equation} v_e = (E/6M)^{1/2} = 2.44 \\times 10^8 \\   {M_{1.4}^{-1/2}} \\  {E_{51}^{1/2}}  \\ \\  \\rm{cm} \\ \\rm{s^{-1}}, \\end{equation} \\begin{equation} A = {6^{3/2}\\over 8\\pi} {M^{5/2}\\over E^{3/2}} = 7.67 \\times 10^6  \\ {M_{1.4}^{5/2}} \\  E_{51}^{- {3/2}} \\ \\ \\rm{g} \\ \\rm{s^3}\\  \\rm{cm^{-3}}, \\end{equation} where $M_{1.4}$ is the explosion mass in terms of the Chandrasekhar mass, $1.4 \\msun$, and $E_{51}$ is the explosion energy in units of $10^{51} \\erg$. A comparison with the power-law model $\\rho \\propto v^{-n}t^{-3}$ (Chevalier \\& Liang, 1989) yield the approximate power index of the density, $n =  - d ln \\rho/ d ln r =  v/v_{e} = r/v_{e}t$. Hence, models with higher velocity scale (increasing with the $E/M$ ratio) yield a flatter density distribution at a given flow velocity, and the gradient of the density increases with radius and decreases over time.    The creation of an exponential profile for SNe Ia is probably related to the fact that the explosion energy is steadily released behind a burning front, unlike in the case of pure shock acceleration in core-collapse SNe. Three progenitor scenarios have been proposed to distinguish SNe Ia (HK96; HWK98; for summary see Hoflich {\\it et al.} 2003). The most favored is the explosion of a near-Chandrasekhar mass C-O WD that has accreted mass through Roche-lobe overflow from an evolved companion star, triggered by compressional heating near the WD center. In this scenario, ignition is initiated in the central high-density region of the WD. In the deflagration model, the burning front propagates subsonically through the white dwarf; a popular example is the W7 model of Nomoto {\\it et al.} (1984). In the other class of models, which includes delayed detonation (DD) (Khokhlov 1991) and pulsating delayed detonation (PDD), the initial subsonic burning front gradually makes a transition to a supersonic burning front (detonation). In the DD models, the detonation front erases the chemical structure that is left behind by the deflagration and creates a layered chemical structure, as observed. With a particularly narrow range of energy variation, the DD models reproduce the optical and infrared light curves of typical SNe Ia quite accurately. In current 3-D pure deflagration models, however, a significant fraction of the C/O WD remains unburned and mixes with the burned material at all velocities, in disagreement with observations.   In the second scenario, or DET2env models, combustion occurs in a rotating low-density WD that is surrounded by an extended C-O envelope. The configuration forms during the merger of two low-mass binary WDs which lose angular momentum to gravitational radiation. The combined mass of the system may exceed the Chandrasekhar limit, but neither WD is above the limit. The third scenario is a double detonation of a sub-Chandrasekhar mass C-O WD, or the He detonation (HeD). The explosion starts with the detonation of a helium layer that has accreted from a close companion, which subsequently triggers the detonation of the inner C-O core. The exponential model is a good simplification of the density structure obtained from various 1-D explosion models, as presented in Fig.1 of DC98. Both the W7 model and the delay detonation model DD200c yield a perfectly exponential decline in density distribution, and they predict similar properties (HKW98). The PDD and the HeD (PDD3 and PDD1c; HeD10 and HeD6) also exhibit more of an exponential than a power-law density distribution, although the gradient appears to drop in the innermost regions. DC98 commented that the flattening may have been exaggerated by the numerical effect of inadequate zoning near the center. In the PDD and some of the HeD models, the explosion energy is below $10^{51} \\erg$. As low-energy models tend to be associated with a high density gradient, the density evolution of lower-energy models is more sensitive to numerical resolution. The merger scenario is believed to be responsible for only a small fraction of the SNe Ia population because the predicted large amount of unburned C/O at the outer layers is inconsistent with the IR spectra. Though not considered in DC98, this scenario may likewise bear an exponential density decrease, since a pure thermonuclear runaway reaction (as opposed to a subsonic burning) would incinerate almost all of the WD to form $\\Ni$, which process would be inconsistent with the observations of SNe Ia that only $\\sim 0.6 \\msun$ of $\\Ni$ is generated (Hoflich 2003).    The chemical structure of SNe Ia varies with the explosion mechanism (HK96; DC98; HWK98; Hoflich 2003). In the accretion scenario, thermonuclear runaway takes place near the center, such that neutron-rich isotopes such as Fe, rather than $\\Ni$, are synthesized in the core. However, for a model such as DD200c (HWK98), the central void of $\\Ni$ amounts to only $\\sim 10\\%$ of the total synthesized mass. The W7 model, whose density structure can be used as a prototype for an exponential model, shows that the amount of Si substantially exceeds that of Fe in the denser parts of the shocked interaction region, while the Fe dominates close to the reverse shock, for the age and ambient density of Tycho's remnant (DC98). A similar composition is produced in delayed detonation or pulsating delayed detonation models for Tycho's remnant, except that the densest parts of the shocked ejecta appear to consist mainly of C and/or O. In the merger scenario, by contrast, burning occurs in the low-density regions, causing little neutronization; therefore, $\\Ni$ is synthesized in the innermost region. In the HeD scenario,  two separate layers of $\\Ni$ are generated, The outer one is formed by burning in the He envelope above the velocity range 11,000-14,000 $\\kms$. The mass of the inner $\\Ni$ can be as little as that of the outer $\\Ni$, or it can be considerable and hence represent a bright SN Ia such as, for example, model HeD10 of HK96. In the models HeD6 and HeD10, the outermost dense regions of the shocked gas comprise mainly Fe, with O dominating in a dense shell inside Fe. The HeD scenario was advocated in view of the large population of low-mass stellar progenitors. However, Hoflich \\& Khokhlov (1997) maintained that its spectra are generally too blue to reproduce the observed light curves and late time spectra.    Subluminous models giving acceptable fits to the observed SNe Ia; for example, those in HK96 using a sample of 26 SNe Ia tend to eject similar amounts of Fe along with the $\\Ni$. In such models, the initial mass of $\\Ni$ ranges from 0.1 to 0.18 $\\msun$, while in normal models, it ranges from 0.49 to 0.83 $\\msun$. The properties of the Ni bubble expansion are dominated by the initial velocity distribution of $\\Ni$ and the effect of thermal pressure is negligible (\\S~3). Therefore, including a realistic ejecta elemental distribution in our models equivalently reduces the initial density contrast of $\\Ni$, which is assumed to be distributed over the whole SN interior, with a mass fraction equal to $\\omega$ everywhere within the bubble, such that $\\omega < 1$. The range $0.5 \\la \\omega < 1$ is presumed. Such a simple parameterization of the $\\Ni$ distribution may sufficiently reproduce results based on a more detailed $\\Ni$ distribution (\\S~3.4).  Figure~1 of DC98 shows that the explosion models significantly diverge from the exponential models as the velocity exceeds $\\ga 10,000 \\kms$, which transition presumably marks the burning interface between the core and the unburned envelope. Yet, inside the interface, the drop in density appears very smooth, and the results suggest that the Ni bubble is not capable of expanding close to this region.     To describe the evolution of an SNR in an exponential model, a set of scaling parameters, $R'$, $V'$ and $T'$, are used following DC98 and W01: \\begin{equation} R' = \\left( {3M \\over 4 \\pi \\rho_{am}}\\right)^{1/3}  \\approx 2.19 \\ \\left({M\\over M_{ch}}\\right)^{1/3} \\  n_{am} ^{-1/3} \\ \\ {\\rm pc}, \\end{equation} \\begin{equation} V' = \\left({2E\\over M}\\right)^{1/2}  \\approx 8.45 \\times 10^3 \\ \\left({E_ {51} \\over {M/M_{ch}}}\\right)^{1/2} \\ \\  {\\rm km \\  s^{-1}}, \\end{equation} \\begin{equation} T' = {R' \\over V'} \\approx 248 \\ E_{51}^{-1/2} \\ \\left({M\\over M_{ch}}\\right )^{5/6} \\ n_{am}^{-1/3} \\ \\ {\\rm yr}, \\end{equation} where $\\rho_{am}$ is the ambient ISM density of the SNR and $n_{am} = \\rho_{am}/(2.34 \\times 10^{-24})$ gm cm$^{-3}$, suitable for an ISM with an H/He ratio of 10/1 by number. The dimensional variables $r$, $v$ and $t$ can be expressed in terms of the non-dimensional quantities $r'=r/R'$, $t'=t/T'$ and $v'=v/V'$. Since the deceleration of the SNR is caused by the ambient ISM, a higher surrounding density is responsible for a later evolutionary phase of the remnant (larger $T'$). Once the non-dimensional solutions are obtained, they can be re-scaled to the dimensional solutions for a different set of $M$, $E$ and $\\rho_{am}$. Thus, one evolutionary sequence in the non-dimensional variables represents virtually all possible dimensional solutions.  Figure~1 plots the evolution of the characteristic velocities of an SNR that arises from the interaction of an exponential ejecta with a constant ambient medium, corresponding to Fig.~1 of WC01. As the ejecta density profile becomes flatter, the reverse shock, initially moving outward in the stellar frame, begins to turn inward at $t' \\approx 2.5 $. It reaches the stellar center at $t' \\approx 8$. Unlike in the power-law model in which the evolution is self-similar (Chevalier {\\it et al.} 1992), the exponential model gives rise to different deceleration rates as the remnant evolves. Nevertheless, the flow velocities immediately below the reverse shock front do not substantially deviate from those in the power-law model (Fig.~12 of W05) within the epoch $t' \\approx 0.1-2'$, when the clump-remnant interaction is likely to begin.   \\section {NICKEL BUBBLE SHELL STRUCTURE AND EVOLUTION} \\label{sec:evo}  \\subsection {Methods} \\label{subsec:met}   The computational methods used in this paper are similar to those adopted in the author's previous study W05. HD and RHD simulations were performed based on the two-dimensional finite-difference code ZEUS2D (Stone \\& Norman 1992; Stone {\\it et al.} 1992). Heating due to a central core of $\\Ni$ that has a density contrast $\\omega=1$ relative to the ambient ejecta was considered. The ejecta were initially freely expanding with an exponential density profile that is characterized by an explosion mass $M$ and energy $E$. The radioactive energy was deposited continually to the gas thermal energy (HD case) within the bubble as the bubble-shell interface was tracked in a uniformly-expanding grid in spherical polar coordinates, starting at $t_{0}=100 \\seco$ after the explosion. The inner boundary of the grid had a simple reflecting condition. At the outer boundary, a non-zero gradient outflow condition was applied to preserve the ambient exponential profile and eliminate spurious shocks that were caused by the grid expansion. Basko (1994) and W05 described the radioactive power input. Since the estimated mean life of $\\Ni$ is larger than that suggested in Firestone (1996), $\\la 2\\%$ less radioactive energy than indicated by Firestone was accumulated in the radiation-streaming epoch $\\sim 10^6 \\seco$. A $\\Ni$ abundance of $0.5 \\msun$ yields an effective energy input of $\\sim 6\\times10^{49} \\ \\rm erg$, or 6\\% of the initial kinetic energy. Therefore, the increase in the velocity from local free expansion is typically only around $\\la 3\\%$.    The initial ejecta velocity at the bubble edge was determined from the initial $\\Ni$ density contrast and age: $U_{0} = {R_{0} / t_{0}} \\sim \\omega^{-1/3}$. The background thermal pressure was distributed according to $p=\\kappa \\rho^\\gamma$, where $\\gamma = 4/3$ in the HD case, and $\\kappa = 6.1 \\times 10^{14}$ (cgs units) (derived from the change in entropy during nuclear burning of C-O to $\\Ni$, W05). The late-time radioactive power input prior to radiation streaming dominated the pressure in the bubble. In the RHD simulations, the radioactive power was deposited into the radiative energy % and $\\gamma=5/3$ for matter. In supernovae, since the dynamic time required for the ejecta to become freely-expanding is much longer than that required for the radiation field to be thermalized with electrons and for collisions between electrons and ions to occur, a black-body radiation field is well established (Shigeyama \\& Nomoto 1990). To ensure proper evolution of the flow toward the optically-thin regime, full radiative-transport RHD is invoked, under the assumptions of local thermal equilibrium (LTE) and gray opacity. The opacity of the ejecta is dominated by electron scattering; absorption is negligible (Shigeyama \\& Nomoto 1990). A Thomson scattering cross section of $0.2 \\ \\rm cm^2/gm$ was used to calculate the scattering opacity, which is suitable for a gas of completely ionized heavy elements. The time for radiative energy to diffuse out of the bubble depends on the extent of electron scattering and the column density of the shell. Notably, full-power radioactive heating can not be achieved in the radiative case because radiative diffusion limits shell expansion.    \\subsection {Hydrodynamic Simulations} \\label{subsec:hd}  Figure~2 and 3 show the ejecta structure and the evolution of the dynamic properties of the Ni bubble shell in the adiabatic case, with parameters $M=1.4 \\Msun$, $E=10^{51} \\erg$, $M_{Ni}=0.5\\Msun$ and $\\omega=1$, in what is referred to herein as the standard model. At the commencement of the simulation $t_{0} = 100 \\seco$, the bubble has the radius $R_0=5.18 \\times 10^{10} \\rm \\ cm$, velocity $U_{0} = 5.18 \\times 10^{8} \\cms$ and ambient density $0.92 \\gcm$ at its edge. The initial background thermal energy density is $2.77 \\times 10^{16} \\ergcc$ in the center. By this stage, a total radioactive energy of $4.85 \\times 10^{45} \\erg$ has been deposited since the explosion, % yielding an averaged radioactive energy density of $6.89 \\times 10^{19} \\ergcc$ in the bubble. % % % % % % % %   The inflation of the bubble induces a strong forward shock behind which the ambient gas is compressed into a dense shell (Fig.~2). A notable feature of the structure is the drop in density in the ejecta substrate. The density contrast across the shock front is seven, as for a radiation dominated ($\\gamma=4/3$) strong shock. The inner edge of the shell is a contact discontinuity  % where the gas has been shocked and cooled for the longest time and so has the highest density. To describe the shell acceleration, the shock radius is approximated as a power-law function of time, $R_{sh} \\propto t^a$, where $a$ is the expansion rate at the shock front, which is equivalent to the velocity contrast between the shock and the ambient freely-expanding ejecta, $(dR_{sh}/dt)/(R_{sh}/t)$. The expansion rate rises to the maximum $a \\la 1.05$ around $10^{7} \\seco$ and then falls to 1.0. That is, the shock front first accelerates and then freezes in the local free expansion of the flow (Figs.~3(a) and (b)). The shell cannot be slower than the free expansion.   The shell is very thin; at $10^{7} \\seco$ it reaches a maximum thickness ratio of $\\beta \\equiv h_{sh}/R_{sh} \\la 0.004$, where $h_{sh}$ is the thickness of the shell (Fig.~3(c)). A linear relationship holds between the thickness ratio and the expansion rate, $\\beta \\approx (a-1)/10$, because the expansion rate is low. The Mach number of the shock, $      {\\cal M}=\\left({dR_{sh}/ dt} -{R_{sh}/ t}\\right) \\left(\\rho_0 / \\gamma p_0\\right)^{1/2} = \\ (a-1){R_{sh}/ t}  \\left(\\rho_0 / \\gamma p_0\\right)^{1/2} $, is $\\cal{M}$$\\sim 87$ at $10^6 \\seco$, where $\\rho_0$ and $p_{0}$ are the preshock density and pressure, respectively. The expansion proceeds until pressure equilibrium is reached, or in the radiative case until the shell becomes transparent to the $\\gamma$-rays. The shell's compression ratio to the preshock ejecta is $\\chi \\ga 10$ (Fig.~3(d)), and is limited by numerical resolution. The compression in the bubble interior does not drop with time, as it does in the uniform-density model; rather it increases to $\\sim 0.9$ (Fig.~3(d)). % %   At $\\sim 10^6 \\seco$, the shell has a thickness ratio $\\beta \\ga 0.002$, a velocity $\\vf \\sim 6,000\\kms$, and a swept-up mass $\\Ms \\sim 0.09 \\Msun$. At this time, the surface density of the shell begins to drop below the mean free path for the $\\sim 1 MeV$ $\\gamma$-ray photons, $\\sim \\ 33 \\rm \\ g/cm^2$, determined by Monte Carlo simulations of X-ray transfer in SNe Ia (corresponding to an effective opacity of $0.03 \\cmsg$; Sutherland \\& Wheeler 1984, Shigeyama \\& Nomoto 1990). The gas is then transparent to the radiation. Although $75\\%$ of the total radioactive energy is yet to be released (Fig.1 of W05), further deposition of the radioactive energy does not affect the dynamics of the flow, and the flow evolves to a free expansion.   \\subsection{Variation of Initial Parameters - Adiabatic Cases}   The effect of explosion parameters and other initial conditions on the solutions is now examined. A wide range of explosion parameters, such as those used in the 1-D explosion models of HK96 and HWK98, have been studied, under both luminous and underluminous conditions (Tab. 1). The clump strength is conservatively evaluated by approximating the shell's frozen-in velocity by the flow velocity of the contact discontinuity when the shell's surface density starts to drop below $\\sim 33 \\ \\rm \\cmsg$.   A comparison with the uniform-density model (model 'flat') indicates that in the standard exponential model (model 'std'), the initial amount of $\\Ni$ was distributed at lower velocities, such that radioactive pressure within the bubble is larger. The shell reaches a higher expansion rate, and the density contrast and thickness ratio are therefore greater. Nevertheless, the expansion velocity and swept-up mass of the shell are lower (Figs.~3(e) and (f)). In the adiabatic uniform model, the frozen-in velocity is $\\vf \\sim 7,500 \\kms$ and the swept-up mass is $M_{s} \\sim 0.12 \\Msun$, while in the exponential model, $\\vf \\sim 5,800\\kms$ and $M_{s} \\sim 0.09 \\Msun$. The difference between the velocities of the radiative counterparts of the above two models should not exceed the difference between their maximum velocities, $\\sim 8,800 \\kms$ and $\\sim 7,500\\kms$.    For a lower initial $\\Ni$ density contrast $\\omega=0.5$, the frozen-in age is slightly advanced, with a lower expansion rate, a lesser shell mass $\\Ms \\sim 0.03 \\Msun$ and a higher frozen-in velocity $\\vf \\sim 9,300 \\kms$. The drop in the expansion rate and the swept-up mass follow from the lower radioactive pressure within the bubble. The frozen-in velocity is higher because the $\\Ni$ initially extends to higher velocities. In the exponential model, the radioactive pressure dominated the thermal pressure since the initial stage of heating; the thermal pressure thus has a negligible impact. In the uniform model, reducing $\\omega$ results in a greater swept-up mass (Fig.~3(f)), because our assumed thermal pressure equilibrium across the bubble causes the initial thermal energy in the bubble to be higher when $\\omega$ is lower. Notably, the $\\omega=0.5$ standard model yields a higher shell velocity than its uniform-density counterpart, because of the initially higher velocities of $\\Ni$.    Comparisons of various models show that models with a higher velocity scale height $v_e$, or a flatter density distribution, tend to yield a higher shell velocity, a lower expansion rate and a lower swept-up mass. For example, the HeD10 model (in which the density decline is the least among all of the models) yields a shell velocity that extends to $\\vf \\la 13,000 \\kms$ and a swept-up mass of $\\Ms \\la 0.02 \\msun$, whereas the PDD1c model (which has the greatest density gradient) gives $\\vf \\sim 2,500\\kms$ (or $\\sim 3,500\\kms$ at maximum) and $\\Ms \\sim 0.16 \\msun$. The shell velocity is correlated with the initial velocity distribution of $\\Ni$, or the velocity scale height, because the shell expansion rate barely exceeds the free expansion of the local ejecta. The expansion rate is generally expected to be under $a \\la 1.07$. The momentum of freely expanding ejecta that were swept up into the shell, obtained by integrating the momentum density over the velocity space from the initial fo final shell location, is found similar to the final momentum imparted to the shell. Thus, only a small fraction of the final shell momentum results from the bubble acceleration. Figure~2 shows that the density profile of the bubble interior deviates from an exponential model mostly in the regions that are close to the contact discontinuity, also because the expansion is weak and can be contrasted with that in the case of core-collapse SNe (Fig.~2 of W05). Figure~3(h) plots the time evolution of the approximate power index, $n \\equiv v/v_{e}$, at the contact discontinuity. The frozen-in shell velocity does not exceed about three times the velocity scale height of the model - except for the luminous HeD10 and low-$\\omega$ models.   The swept-up mass of the shell is not very sensitive to the initial abundance of $\\Ni$. In a Chandrasekhar-mass, low-energy model using $E_{51}=0.7$ and $M_{Ni}=0.49$, for example, the swept-up mass is comparable to that of a similar model with half the amount of $\\Ni$. It also is approximately the same across several models with intermediate explosion conditions, with a typical $M_{s}$ of $\\sim 0.1 \\msun$.   Solutions to the underluminous models are not sensitive to variations of the initial parameters including the initial $\\Ni$ density contrast, because their steep density distribution and low $\\Ni$ abundance allow little corresponding variations in velocity and radioactive pressure. A unique feature of the exponential models except for PDD1c is that the density contrast at the bubble interior increases over time, because of continual flattening of the exponential density distribution and a relatively low expansion rate.    In the HeD scenario, the outer layer of $\\Ni$ is present at high velocities of above 11,000-14,000 $\\kms$ (HK96). An off-center heating scenario was considered using the HeD6 model (model 'HeD6*'), in which the radioactive power from a half of the outer $\\Ni$ mass was gradually deposited in the space between the velocities $12,000 \\kms$ and $\\la 13,400 \\kms$, while the inner grid boundary continues to move at a fixed velocity of $12,000 \\kms$. A dense shell as in the central-heating model is swept up, but it exhibits very weak acceleration relative to the local fast ejecta. The onset of the optically-thin stage thus advances to $\\sim 10^5 \\seco$, leading to a 'low' frozen-in velocity at $\\la 14,000 \\kms$.    A flatter distribution of the thermal pressure, $p \\sim \\rho$, was assumed with a constant temperature, $kT \\sim 0.6 \\ MeV$ (derived from the nuclear burning process, W05). As the central background pressure is enhanced by a factor of $\\sim 500$, the radioactive pressure still dominates the thermal pressure, so the evolution of the Ni bubble remains largely unchanged. To account for the central $\\Ni$ void as in the accretion scenario, $\\Ni$ in model DD200c was redistributed to higher velocities, corresponding to a void mass of 1/10 of the synthesized mass. The result is indistinguishable from that of the central-heating model. Lastly, the simulated solutions that began at the age $t_0=1000 \\seco$ rapidly converged with those that began at $100 \\seco$. Thus, as long as $t_0 \\la 1000 \\seco$, the final solutions do not depend on the commencing ages of the simulations.   \\subsection {$\\Ni$ Mass Fraction in the W7 Model}   In our $\\omega = 1$ exponential W7 model, the shell velocities, 7558 $\\kms$, is actually lower than the minimum velocity $\\sim 10,000 \\kms$ of the Si-dominant layer, obtained using the W7 model of Nomoto {\\it et al.} (1984). In their original calculations, the velocity of the radioactive Ni was as high as 8300 km/s, and its fractional mass fell to 0.5 at 8800 km/s. To justify that a simple parametrization of the initial $\\Ni$ distribution in terms of $0.5 \\la \\omega < 1$ approximates the results that consider a realistic ejecta elemental distribution, the $\\Ni$ mass fraction given by the model of Nomoto {\\it et al.} ($M_{Ni}=0.58 \\msun$) was incorporated into the exponential model. As the $\\Ni$ mass fraction falls steeply above $8,800 \\kms$, the initial $\\Ni$ distribution was cut off at this velocity but the same radioactive power, proportional to the mass density within it, was deposited. Spatial variations in the background thermal pressure caused by the nonuniform elemental distribution were neglected. This modified model gives a similar initial outbound speed to that obtained in our $\\omega=0.7$ W7 model (Fig.~4). (The integrated density over all elements should follow an exponential profile as illustrated in DC98.) In both models, the final shell velocity reaches $\\sim 9000 \\kms$. In the modified model the swept-up mass is lower but does not differ significantly from that in the $\\omega=0.7$ case.   The reason why the use of $\\omega=0.7$ adequately reproduces  % the result based on a more detailed $\\Ni$ distribution in the W7 model is that it places the outer boundary of $\\Ni$ in the 8300-8800 velocity range, beyond which whose mass fraction drops down rapidly. However, in view of a finite sound speed propagation within the bubble and the presence of nonradioactive material in the ejecta, this result strongly depends on the assumed value of $\\omega$ and is merely an approximation. On the other hand, the clump speeds in our $\\omega=1$ cases stand only as the lower limits in more realistic Ni bubble heating scenarios.   \\subsection {Radiation-Hydrodynamic Solutions} \\label{subsec:rhd}  Here we present full-radiative transport RHD simulations to illustrate the effect of the diffusion of radiative energy across the bubble structure on the shell properties. % In Figs.~2, 4 and 5, the density structures of the radiative models are plotted over their adiabatic counterparts.  The shock has a widely extending radiative precursor, lowering the overall compression of the shell. When the radiative precursor is considered, the shell mass reaches $\\sim 90\\%$ of the adiabatic value (0.1 $\\msun$). Excluding the radiative precursor reduces the mass to $\\ga 0.04 \\msun$ (radiative case) in the standard model. The drop is typically $\\la 50\\%$.  Since the acceleration of the shell is weak, the shell formation processes can be regarded as perturbations to the initial free expansion, such that the final radius and velocity of the shell are only slightly reduced. The final shell velocity differs from the adiabatic solution by $\\la3\\%$ (Fig.~6). As the final shell velocity may not be an adequate measure of the radiative effects, the difference between the final shell velocity and the velocity of the initial free expansion at the bubble's edge is used to assess the importance of radiation diffusion effects. A $\\sim 25\\%$ difference is found between the adiabatic and radiative solutions. However, in the $\\omega=0.7$ W7 model, it is only 12\\%. Hence, the radiative effects are expected to be mild in a more realistic scenario.  The mass and number of clumps should drop as the acceleration and compression of the shell are decreased. Since the RHD simulations are affected by inadequate spatial resolution, the shell's density contrast is crudely estimated, based on the assumption that that the shell mass is 50\\% less and the shell thickness is double the size of that in the adiabatic solutions; a compression ratio of $\\sim 25\\%$ of the adiabatic value, or $\\chi \\sim 25$, is expected. Nonetheless, since ejecta clumps (shell fragments) most likely arise from the densest part of the CD, the true compression ratio may be higher. ",
        "conclusions": "\\label{sec:conclu}    We have investigated the structure and expansion properties of the Ni bubble shell in SNe Ia due to radioactive heating caused by $\\Ni \\rightarrow\\ \\co \\rightarrow\\ \\fe$ decay sequence, assuming an exponentially-declining density profile of the ejecta substrate. Based on the indication that ejecta clumps arise from the breakup of components of the Ni bubble shell, whether the inferred shell properties are compatible with those revealed for the ejecta knots in Tycho's remnant was studied. Adiabatic solutions were first obtained, and then radiative solutions, including the radiation diffusion process across the Ni bubble shell, up to the approximate frozen-in age $\\sim 10^6 \\seco$, were applied to estimate the clump properties. Since a realistic elemental distribution of $\\Ni$ results in a higher expansion velocity than a centralized distribution of $\\Ni$ used in our model, which effect is equivalent to the use of a lower initial $\\Ni$ density contrast in the range of $0.5 \\la \\omega < 1$, our $\\omega=1$ cases may only hint at the lower limit on the clump speeds. We believe that our inferred clump strength is at least not overestimated.   The expansion of the Ni bubble sweeps up a dense thin shell of $\\sim 0.1 \\Msun$ shocked gas in a typical Type Ia SNR. The density of the shell increases inwardly toward the bubble-shell interface, such that the highest computed density is limited by numerical resolution. Since the exponential profile is close to power-law profile that evolves from a steep power law to a flatter one, the adiabatic solutions of the shell were approximated as self-similar solutions for a pulsar bubble shell in power-law ejecta of power index $n$ with the shock propagated at an expansion rate $a$, which in our intermediate-class models falls in the range $n=2-3$ and $a\\la 1.07$. The shell is thicker than given by the flat ejecta model, representing $\\la 1\\%$ of the forward shock radius and the mean density contrast is enhanced by $50\\%$. In the radiative case, the shell is broader and less dense; the total swept-up mass of the shell is $\\sim 50\\%$ less without the radiative precursor, but the shell's expansion velocity is only $< 3\\%$ lower close to the frozen-in age. The radiative diffusion process does not substantially affect the clump speed or the initiative age of the clump-remnant interaction.   A wide range of explosion parameters that are similar to those used in valid 1-D explosion models was explored. The exponential model gives a higher shell expansion rate and thickness ratio than the uniform model, and a lower shell velocity and swept-up mass. The expansion of the shell is only slightly faster than the free expansion of the ambient ejecta substrate. Consequently, the frozen-in velocity of the shell often rises with the velocity scale height $v_e$ of the exponential profile -- which characterizes the free motion of the ejecta substrate -- or drops with the ambient ejecta density gradient. The velocity and swept-up mass of the shell appear to be inversely correlated: steeper models (with smaller velocity scale heights) tend to acquire a lower expansion velocity and a higher swept-up mass.    Under intermediate explosion conditions with an ambient ISM density in the range of $n_{am}=0.5-1\\cmc$, the clump-remnant impact is anticipated to occur at $\\sim 300$ yrs after the supernova is formed. Although the impact precedes Tycho's present state, the clump speed ($\\sim 6000 \\kms$ for $\\omega=1$) may not be sufficient. A higher clump speed and an earlier clump-remnant impact is expected when the model incorporates a more realistic elemental distribution. The explosion parameters obtained from the successful explosion models such as the deflagration W7 and the delay detonation DD200 yield the most favorable result for Tycho's knots. A sub-Chandrasekhar-mass, low-energy scenario such as the He detonation HeD6 may also satisfy Tycho's condition. A Chandrasekhar-mass and low-energy model such as the PDD is unfeasible, indicating that if Tycho's SN was intrinsically underluminous, then a near-Chandrasekhar-mass progenitor could not arise. Likewise, SN 1006 could not originate in a W7-like scenario since the currently-expected clump-forward shock impact is not observed. Recent studies nonetheless identified both Tycho's SN and SN 1006 as normal Type Ia supernovae.  The global approximations of Type Ia models with exponential models may not be sufficiently accurate enough close to the contact discontinuity, particularly for the PDD1c model, where the ejecta appears to diverge from the exponential model at low velocities. If the ejecta density profile deviates from the exponential or power-law models, a higher clump speed and an earlier clump-remnant impact is expected for a flatter ejecta substrate, and vice versa. However, significant deviations from the exponential model are not anticipated for the W7 and DD200c models. The clump-remnant impact is estimated to occur within $\\sim 2000$ yrs under all plausible explosion conditions of SNe Ia. As clumps are expected to be present only in a thin shell at the boundary of the Ni bubble, not throughout the bulk of the ejecta, the small clump size and a flat bubble density contrast should not induce vigorous mixing in the ensuing clump/bubble-remnant interactions, which such characteristic distinguishes the remnants of SNe Ia from those of the core-collapse SNe.   The author would like to thank Roger Chevalier, Dietrich Baade and Vikram Dwarkadas for useful discussions and correspondence, and the anonymous referee for valuable comments on the manuscript. This work was supported by the National Science Council of the Republic of China, Taiwan, and the National Center of High-Performance Computing of Taiwan. Ted Knoy is appreciated for his editorial assistance."
    },
    "0606/astro-ph0606670_arXiv.txt": {
        "abstract": "{% }{% We present a $8^\\circ\\times6^\\circ$, high resolution extinction map of the Pipe nebula using 4.5 million stars from the Two Micron All Sky Survey (2MASS) point source catalog. }{% The use of \\textsc{Nicer} (Lombardi \\& Alves 2001), a robust and optimal technique to map the dust column density, allows us to detect a $A_V = 0.5 \\mbox{ mag}$ extinction at a 3-$\\sigma$ level with a $1 \\mbox{ arcmin}$ resolution. }{% (\\textit{i\\/}) We find for the Pipe nebula a \\textit{normal\\/} reddening law, $E(J - H) = (1.85 \\pm 0.15) E(H - K)$. (\\textit{ii\\/}) We measure the cloud distance using Hipparchos and Tycho parallaxes, and obtain $130^{+24}_{-58} \\mbox{ pc}$.  This, together with the total estimated mass, $10^4 \\mbox{ M}_\\odot$, makes the Pipe the closest massive cloud complex to Earth. (\\textit{iii\\/}) We compare the \\textsc{Nicer} extinction map to the NANTEN ${}^{12}$CO observations and derive with unprecedented accuracy the relationship between the near-infrared extinction and the $^{12}$CO column density and hence (indirectly) the $^{12}$CO X-factor, that we estimate to be $2.91 \\times 10^{20} \\mbox{ cm}^{-2} \\mbox{ K}^{-1} \\mbox{ km}^{-1} \\mbox{ s}$ in the range $A_V \\in [0.9, 5.4] \\mbox{ mag}$. (\\textit{iv\\/}) We identify approximately $1\\,500$ OH/IR stars located within the Galactic bulge in the direction of the Pipe field. This represents a significant increase of the known numbers of such stars in the Galaxy. }{% Our analysis confirms the power and simplicity of the color excess technique to study molecular clouds.  The comparison with the NANTEN ${}^{12}$CO data corroborates the insensitivity of CO observations to low column densities (up to approximately $2 \\mbox{ mag}$ in $A_V$), and shows also an irreducible uncertainty in the dust-CO correlation of about $1 \\mbox{ mag}$ of visual extinction. } ",
        "introduction": "\\label{sec:introduction}  Nearby Galactic molecular clouds complexes represent our best chance to understand cloud formation and evolution and hence to learn how stars come to be.  But progress on the study of these objects has been slow.  Not only are molecular clouds the coldest objects known in the Universe, their main mass component ($\\mathrm{H}_2$) cannot be detected directly.  Most of all we know today about their physical properties has been derived through radio spectroscopy of $\\mathrm{H}_2$ surrogate molecules (CO, CS, $\\mathrm{NH}_3$; e.g., \\citealp{1999osps.conf....3B, 1999osps.conf...67M}) and more recently through thermal emission of the dust grains inside these clouds \\citep{2000prpl.conf...59A, 2000ApJ...545..327J}.  The results obtained using these techniques, especially the estimate of column densities, are not always straightforward to interpret and are plagued by several poorly constrained effects.  Moreover, although large scale maps of entire molecular cloud complexes are now available \\citep{1998ApJS..115..241H, 2001ApJ...551..747S}, maps with sufficient resolution and dynamic range to identify not only dense molecular cores but also their extended environment are still not existent. This wide view on molecular clouds is at present an obvious gap in our understanding of the relation between the dense Interstellar Medium (ISM) and star formation.  A straightforward and powerful technique to study molecular cloud structures, pioneered by \\citet{1994ApJ...429..694L} and known as the Near-Infrared Color Excess method (\\textsc{Nice}, \\citealp{1998ApJ...506..292A}) makes use of the most reliable tracer of $\\mathrm{H}_2$ in these clouds: extinction by pervasive dust grains in the gas.  The \\textsc{Nice} method relies on near-infrared (NIR) measurements of extinguished background starlight to derive accurate line-of-sight estimates of dust column densities.  Depending on the stellar richness and color properties of the background field this technique can produce column density maps with spatial resolutions down to $5 \\mbox{ arcsec}$ \\citep[e.g.][]{2001Natur.409..159A, 2002osp..conf...37A} and with dynamic ranges more then an order of magnitude larger than classical optical star count techniques.  This novel view on molecular clouds is providing not only new information on the physical structure of these object \\citep{2001Natur.409..159A} but also an insight into their chemical structure and the physical properties of the dust grains, when combined to molecular line and dust emission data \\citep{1998A&A...329L..33K, 1999A&A...342..257K, 1999ApJ...515..265A, 2001ApJ...557..209B, 2003ApJ...586..286L, 2003A&A...399L..43B, 2003A&A...399.1073K}.  In part to make use of the wealth of NIR data provided by the Two Micron All Sky Survey (2MASS; \\citealp{1994ExA.....3...65K}) the \\textsc{Nice} method was further developed into an optimized multi-band technique dubbed Near-Infrared Color Excess Revisited (\\textsc{Nicer}, \\citealp{2001A&A...377.1023L}, hereafter Paper~I). This generalization of \\textsc{Nice} can, nevertheless, be applied to any multi-band survey of molecular clouds.  Through use of optimal combinations of colors, \\textsc{Nicer} improves the noise variance of a map by a factor of two when compared to \\textsc{Nice}. This unique property of \\textsc{Nicer} makes it the ideal tool to trace large scale distributions of low column density molecular cloud material.  When applied to 2MASS data, \\textsc{Nicer} dust extinction maps trace not only the low column density regions ($A_V \\simeq 0.5 \\mbox{ mag}$) but have the dynamic range to identify dense molecular cores by reaching cloud depths of $\\sim 30 \\mbox{ mag}$ of extinction, corresponding to $8 \\times 10^{22} \\mbox{ protons} \\mbox{ cm}^{-2}$ \\citep{L05}.  In this paper we present an extinction map of the Pipe nebula covering $48 \\mbox{ sq deg}$, computed by applying the \\textsc{Nicer} technique on 4.5 millions JHK photometric measurements of stars from the 2MASS database.  The Pipe nebula is a poorly studied nearby complex of molecular clouds.  The only systematic analysis of this region is the one of \\citet{1999PASJ...51..871O} who present a $\\sim 27 \\mbox{ sq deg}$ map in the $J = 1-0$ line of $^{12}$CO observed on a $4 \\mbox{ arcmin}$ grid, and smaller maps of selected regions in the $J = 1-0$ lines of $^{13}$CO and C$^{18}$O.  They estimate the total $^{12}$CO mass to be $\\sim 10^4 \\mbox{ M}_\\odot$ (for a cloud distance of $160 \\mbox{ pc}$) and point out that star formation is only occurring on Barnard~59, where the C$^{18}$O column density is the highest and where they find a CO outflow.  Barnard~59 was also observed at $1300\\ \\mu\\mbox{m}$ by \\citet{1996A&A...314..258R} and a protostellar candidate B59-MMS1 was found.  Two pre-main sequence stars associated with Barnard~59 appear in the young binaries survey of \\citet{1993A&A...278...81R} and more recently in \\citet{2002AJ....124.1082K}.  This paper is organized as follows.  In Sect.~\\ref{sec:nicer-absorpt-map} we describe the technique used to map the dust in the Pipe nebula and we present the main results obtained.  In Sect.~\\ref{sec:distance} we address the determination of the cloud distance using Hipparcos data.  A statistical analysis and a discussion of the bias introduced by foreground stars is presented in Sect.~\\ref{sec:statistical-analysis}.  We compare the CO observations from \\citet{1999PASJ...51..871O} with our outcome in Sect.~\\ref{sec:comparison-with-co}.  Section~\\ref{sec:mass-estimate} is devoted to the mass estimate of the cloud complex.  Finally, we summarize the conclusions of this paper in Sect.~\\ref{sec:conclusions}.  In the following we will normally express column densities in terms of the 2MASS $K_\\mathrm{s}$ band extinction $A_K$.  When converting this quantity into the widely used visual extinction $A_V$, we will use the \\citet{1985ApJ...288..618R} reddening law converted into the 2MASS photometric system \\citep[see][]{2001AJ....121.2851C}, $A_K / A_V = 0.112$. ",
        "conclusions": "\\label{sec:conclusions}  The main results of this paper can be summarized as follows: \\begin{itemize} \\item We used 4.5 million stars from the 2MASS point source catalog to construct a $8^\\circ \\times 6^\\circ$ extinction map of the poorly studied Pipe nebula, a molecular cloud complex seen in projection towards the galactic bulge.  The map has a resolution of 1~arcmin and has a $3 \\sigma$ detection level of $0.5$ visual magnitudes. \\item We combined the 2MASS data with Tycho and Hipparcos parallaxes to obtain a distance of $130 \\pm 15 \\mbox{ pc}$ to this molecular complex. We estimated, for this distance, a total mass of $\\sim 10^4 \\mbox{ M}_\\odot$ for the Pipe complex.  This makes the Pipe nebula one the closest star forming region to the Earth of its type, closer than the Ophiuchus and the Taurus complexes. \\item We studied in details the statistical properties of the extinction map obtained through the \\textsc{Nicer} method.  We also confirmed the relation originally observed by \\citet{1994ApJ...429..694L} for IC~5146, suggesting that there is unresolved structure at the highest extinctions in our map. \\item We compared our near-infrared study with the CO observations of \\citet{1999PASJ...51..871O}, and derived fitting formulae to relate the $^{12}$CO column density with the $A_K$ extinction.  We also derived the dust-to-$^{12}$CO ratio, and showed that if a normal infrared reddening law is assumed, then the derived X-factor is as large as $2.91 \\times 10^{20} \\mbox{ cm}^{-2} \\mbox{ K}^{-1} \\mbox{ km}^{-1} \\mbox{ s}$ in the range $A_K \\in [0.1, 0.6] \\mbox{ mag}$. \\item We found that $^{12}$CO is only sensitive to about 65\\% of the total dust mass. About half of the missing mass is not traced by CO at column densities below $A_K < 0.25 \\mbox{ mag}$, and half is not traced at column densities above $A_K > 0.6 \\mbox{ mag}$, where the line begins to saturate.  There is an apparently irreducible uncertainty in the dust-CO correlation of about 1 magnitude of visual extinction. This uncertainty seems to be independent of the cloud or the interstellar radiation field. \\item We took advantage of the large number of background sources to accurately measure the NIR reddening law for this cloud, and obtained $E(J - H) = (1.85 \\pm 0.15) E(H - K)$, in very good agreement with the standard \\citet{1985ApJ...288..618R} reddening law. \\item Finally from analysis of the JHK color-color diagram for the Pipe region we identify a large population of red stars whose colors are distinct from those of typical reddened background giants. These stars are spatially distributed over the entire observed field, preferentially located at lower Galactic latitudes and are not associated with the molecular cloud.  These stars are the brightest stars detected in the field and have a very narrow distribution in magnitude space.  These properties are similar to those of Galactic OH/IR stars.  Our observations thus appear to have provided one of the largest samples of such stars yet discovered. \\end{itemize}"
    },
    "0606/astro-ph0606301_arXiv.txt": {
        "abstract": "Extended gas haloes around galaxies are a ubiquitous prediction of galaxy formation scenarios. However, the density profiles of this hot halo gas is virtually unknown, although various profiles have been suggested on theoretical grounds.  In order to quantitatively address the gas profile, we compare galaxies from direct cosmological simulations with analytical solutions of the underlying gas equations. We find remarkable agreement between simulations and theoretical predictions. We present an expression for this gas profile with a non-trivial dependence on the total mass profile. This expression is useful when setting up equilibrium galaxy models for numerical experiments.\\\\ ",
        "introduction": "Standard galaxy formation scenarios predict that the galaxies have an extended halo of hot gas.  This hot gas is gravitationally trapped by the dark matter potential and cools slowly due to thermal emission~\\citep{white91}.  These haloes of hot gas have been observed around massive elliptical galaxies through their X-ray emission (e.g. Matsushita 2001), and recently around normal quiescent spiral galaxies~\\citep{pedersen2006}. These observations provide strong support for the standard galaxy formation scenario.  One missing aspect is the qualitative understanding of the density profile of these hot haloes. We therefore investigate the density profiles extracted from virialized galaxies in cosmological simulations, as well as theoretical predictions based on the fundamental gas equations. We find remarkable agreement between simulations and theory. We present an analytical expression for the gas-profile, which depends non-trivially on the total mass distribution of both baryons and dark matter. ",
        "conclusions": "Standard galaxy formation scenarios predict that the central cold gas-disk and stars are surrounded by an extended halo of hot gas.  We investigate quantitatively the density profile of this gas-halo.  We find that results of cosmological $\\Lambda$CDM N-body/hydrodynamical simulations of the formation and evolution of galaxies, and analytical solutions to the fundamental gas equations are in remarkable agreement. We find that the gas density slope (the logarithmic derivative) is a non-trivial function of the slope of the total matter, expressed through eq.~(\\ref{eq:betagas}). This equation is useful when constructing realistic galaxies for controlled numerical experiments."
    },
    "0606/astro-ph0606137_arXiv.txt": {
        "abstract": "AAOmega is the new spectrograph for the 2dF fibre-positioning system on the Anglo-Australian Telescope. It is a bench-mounted, double-beamed design, using volume phase holographic (VPH) gratings and articulating cameras. It is fed by 392 fibres from either of the two 2dF field plates, or by the 512 fibre SPIRAL integral field unit (IFU) at Cassegrain focus. Wavelength coverage is 370 to 950nm and spectral resolution 1,000-8,000 in multi-Object mode, or 1,500-10,000 in IFU mode. Multi-object mode was commissioned in January 2006 and the IFU system will be commissioned in June 2006.  The spectrograph is located off the telescope in a thermally isolated room and the 2dF fibres have been replaced by new 38m broadband fibres. Despite the increased fibre length, we have achieved a large increase in throughput by use of VPH gratings, more efficient coatings and new detectors - amounting to a factor of at least 2 in the red. The number of spectral resolution elements and the maximum resolution are both more than doubled, and the stability is an order of magnitude better.  The spectrograph comprises: an f/3.15 Schmidt collimator, incorporating a dichroic beam-splitter; interchangeable VPH gratings; and articulating red and blue f/1.3 Schmidt cameras. Pupil size is 190mm, determined by the competing demands of cost, obstruction losses, and maximum resolution. A full suite of VPH gratings has been provided to cover resolutions 1,000 to 7,500, and up to 10,000 at particular wavelengths. ",
        "introduction": "\\label{sect:intro}  % The two degree field facility (2dF$^{\\rm{[1,2]}}$) of the Anglo-Australian Telescope (AAT), had its original spectrographs mounted on the telescope top end ring. This was mandated by the capabilities of fibres at that time, and the spectrographs were designed for the primary science goal - a large redshift survey of moderate brightness galaxies. However, they were of limited utility in obtaining spectra of fainter objects, or at higher signal-to-noise or resolution. The reasons for this were: \\begin{itemize} \\item Restricted size and weight, with implications for instrument stability and format \\item Flexure and temperature variations, compromising both spectral and spatial stability of the system \\item Imperfect optics giving variable point spread function (PSF), leading to systematic error in skyline subtraction that could not be reduced by increased integration. \\item Significant scattered light and halation problems, limiting the dynamic range within and between spectra. \\end{itemize}  Despite these limitations, both spectrographs performed admirably resulting in a multitude of high quality scientific results over their 10 year careers.  In August 2005 spectrographs 1 and 2 were retired from service to make way for installation of the next generation AAT multi-purpose optical spectrograph, \\aaomega.  The mechanical$^{\\rm{[3]}}$ and optical$^{\\rm{[4]}}$ design of \\aaomega\\ has been presented previously elsewhere.  In what follows we briefly describe the \\emph{as commissioned} capabilities of \\aaomega. The commissioning program is outlined in table \\ref{tab:time line}.  \\begin{table}[h] \\caption{\\aaomega\\ was commissioned November 2005 - January 2006 following the project time line outlined below.} \\label{tab:time line} \\begin{center} \\begin{tabular}{|l|l|} \\hline \\textbf{Date} & \\textbf{Status}\\\\ \\hline \\hline August 2005 & Decommission 2dF spectrographs\\\\ & Remove 2dF fibre retractors to install \\aaomega\\ fibres\\\\ \\hline November 2005 & First on sky commissioning (6 nights)\\\\ & 2dF positioner upgrade and new field plates\\\\ & Instrument and telescope control software integration\\\\ & AAOmega guide fibres\\\\ & Limited AAOmega slit units\\\\ & (2 Plates $\\times$ 3  Bundles $\\times$ 10 fibres)\\\\ & Blue camera only\\\\ \\hline December 2005 & Second commissioning run (7 nights)\\\\ & Red and Blue cameras\\\\ & Fully populated slit units\\\\ \\hline January 2006  & Final commissioning run (5 nights)\\\\ & Full system integration\\\\ \\hline \\hline January 2006  & Science verification (8 nights)\\\\ & 9 programs substantially completed\\\\ & Testing the full range of \\aaomega\\ capabilities.\\\\ \\hline February 2006  & Commence routine science operations\\\\ - May 2006   & 10 scheduled programs, 41 nights\\\\ & Active service program\\\\ \\hline \\hline June 2006     & SPIRAL IFU Scheduled for commissioning\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\subsection{Improvements over the original 2dF system} While maintaining the high multiplex (392 fibres) and low observational overhead (via configuration of one field plate during observation of the other) of 2dF, the key design goals for \\aaomega\\ have been to overcome many of the problems experienced with the original 2dF system:  \\begin{itemize} \\item Improved sensitivity at all wavelengths, even with the increased fibre run length. \\item Spectroscopic stability - removal of thermal and flexure effects \\item PSF uniformity - dramatically improved sky subtraction and higher velocity accuracy \\item Higher maximum resolution \\item Greater wavelength coverage at a given resolution \\item Minimisation of systematic effects, such as scattered light and light leaks \\item Improved support for \\emph{Nod-and-Shuffle} observations$^{\\rm{[5]}}$. \\end{itemize}  As an alternate to the 2dF MOS fibre feed, the SPIRAL Integral Field Unit (IFU) mounted at the Cassegrain focus, will also be available to \\aaomega.  SPIRAL will be commissioned in June 2006 and hence we defer discussion of the IFU feed until a later time. ",
        "conclusions": ""
    },
    "0606/astro-ph0606247_arXiv.txt": {
        "abstract": "A commonly used measure to summarize the nature of a photon spectrum is the so-called Hardness Ratio, which compares the number of counts observed in different passbands.  The hardness ratio is especially useful to distinguish between and categorize weak sources as a proxy for detailed spectral fitting.  However, in this regime classical methods of error propagation fail, and the estimates of spectral hardness become unreliable.  Here we develop a rigorous statistical treatment of hardness ratios that properly deals with detected photons as independent Poisson random variables and correctly deals with the non-Gaussian nature of the error propagation.  The method is Bayesian in nature, and thus can be generalized to carry out a multitude of source-population--based analyses.  We verify our method with simulation studies, and compare it with the classical method.  We apply this method to real world examples, such as the identification of candidate quiescent Low-mass X-ray binaries in globular clusters, and tracking the time evolution of a flare on a low-mass star. ",
        "introduction": "\\label{park:sec:intro}  The shape of the spectrum of an astronomical source is highly informative as to the physical processes at the source.  But often detailed spectral fitting is not feasible due to various constraints, such as the need to analyze a large and homogeneous sample for population studies, or there being insufficient counts to carry out a useful fit, etc.  In such cases, a hardness ratio, which requires the measurement of accumulated counts in two or more broad passbands, becomes a useful measure to quantify and characterize the source spectrum.\\footnote{ In the context of \\chandra\\ data, the passbands often used are called {\\sl soft} (0.3-0.9~keV), {\\sl medium} (0.9-2.5~keV), and {\\sl hard} (2.5-8.0~keV) bands. Sometimes the {\\sl soft} and {\\sl medium} bands are merged into a single band, also called the {\\sl soft} band (e.g., 0.5-2~keV). In all cases, note that the energies refer to the nominal detector PHA or PI channel boundaries and not to the intrinsic photon energies. The choice of energy or PI channel range is flexible and situational. } A hardness ratio is defined as either the ratio of the counts in two bands called the {\\sl soft} and {\\sl hard} bands, or a monotonic function of this ratio.  We consider three types of hardness ratios, \\begin{eqnarray} \\nonumber {\\rm Simple~Ratio,}~~&~\\R \\equiv& \\frac{S}{H} \\\\ \\nonumber {\\rm Color,}~~&~\\C \\equiv& \\log_{10}\\left(\\frac{S}{H}\\right) \\\\ {\\rm Fractional~Difference,}~~&\\HR \\equiv& \\frac{H-S}{H+S} \\label{e:RCHR} \\end{eqnarray} where $S$ and $H$ are the source counts in the two bands, called the {\\sl soft} and {\\sl hard} passbands.  The simple formulae above are modified for the existence of background counts and instrumental effective areas.  Spectral colors in optical astronomy, defined by the standard optical filters (e.g., UBVRIJK, U-B, R-I, etc), are well known and constitute the main observables in astronomy.  They have been widely applied to characterize populations of sources and their evolution. The hardness ratio concept was adopted in X-ray astronomy for early X-ray detectors (mainly proportional counter detectors) which had only a limited energy resolution. The first application of the X-ray hardness ratio was presented in the X-ray variability studies with SAS-3 observations by Bradt et al.\\ (1976). Zamorani et al.\\ (1981) investigated the X-ray hardness ratio for a sample of quasars observed with the {\\it Einstein} X-ray Observatory. They calculated each source intensity in two energy bands (1.2-3~keV and 0.5-1.2~keV) and discussed a possible evolution of the hardness ratio (the ratio of both intensities) with redshift for their X-ray sample of 27 quasars. Similar ratios have been used in surveys of stars (Vaiana et al.\\ 1981), galaxies (e.g., Kim, Fabbiano, \\& Trinchieri 1992) and X-ray binaries (Tuohy et al.\\ 1978) studied with the {\\it Einstein} and earlier observatories.  In the case of X-ray binaries in particular, they were used to define two different classes of sources (Z-track and atoll sources; Hasinger \\& van der Klis, 1989) depending on their time evolution on $\\HR$ v/s $\\HR$ diagrams.  Since then the concept of X-ray hardness ratio has been developed and broadly applied in a number of cases, most recently in the analysis of large data samples from the \\chandra\\ X-ray Observatory (Weisskopf et al.\\ 2000) and XMM-{\\sl Newton} (Jansen et al.\\ 2001).  Advanced X-ray missions such as \\chandra\\ and XMM-{\\sl Newton} allow for the detection of many very faint sources in deep X-ray surveys (see review by Brandt \\& Hasinger 2005). For example, in a typical observation of a galaxy, several tens and in some cases even hundreds of sources are detected, most of which have fewer than 50 counts, and many have only a few counts. Similar types of sources are detected in the ChaMP serendipitous survey (Kim et al.\\ 2005). They provide the best quality data to date for studying source populations and their evolution. The most interesting sources are the sources with the smallest number of detected counts, because they have never been observed before. Further, the number of faint sources increases with improved sensitivity limits, i.e., there are more faint sources in deeper observations. Because these faint sources have only a few counts, hardness ratios are commonly used to study properties of their X-ray emission and absorption (e.g., Alexander et al.\\ 2001, Brandt et al.\\ 2001a, Brandt et al.\\ 2001b, Giacconi et al.\\ 2001, Silverman et al.\\ 2005). With long observations the background counts increase, so the background contribution becomes significant and background subtraction fails to work. Still the background contribution must be taken into account to evaluate the source counts.  The sensitivity of X-ray telescopes is usually energy dependent, and thus the classical hardness ratio method can only be applied to observations made with the same instrument. Usually the measurements carry both detection errors and background contamination. In a typical hardness ratio calculation the background is subtracted from the data and only net counts are used. This background subtraction is not a good solution, especially in the low counts regime (see van Dyk et al.\\ 2001). In general the Gaussian assumption present in the classical method (see \\S\\S~\\ref{park:sec:classic}~and~\\ref{park:sec:verify}) is not appropriate for faint sources in the presence of a significant background. Therefore, the classical approach to calculating the hardness ratio and the errors is inappropriate for low counts.  Instead, adopting a Poisson distribution as we do here (see below), hardness ratios can be reliably computed for both low and high counts cases. The Bayesian approach allows us to include the information about the background, difference in collecting area, effective areas and exposure times between the source and background regions.  Our Bayesian model-based approach follows a pattern of ever more sophisticated statistical methods that are being developed and applied to solve outstanding quantitative challenges in empirical Astronomy and Astrophysics. Bayesian model-based methods for high-energy high-resolution spectral analysis can be found for example in Kashyap \\& Drake (1998), van Dyk et al. (2001), van Dyk and Hans (2002), Protassov et al.\\ (2002), van Dyk and Kang (2004), Gillessen \\& Harney (2005), and Park et al.\\ (2006, in preparation). More generally, the monographs on Statistical Challenges in Modern Astronomy (Feigelson and Babu, 1992, 2003, Babu and Feigelson, 1997) and the special issue of {\\it Statistical Science} devoted to Astronomy (May 2004) illustrate the wide breadth of statistical methods applied to an equal diversity of problems in astronomy and astrophysics.  Here, we discuss the classical method and present the fully Bayesian method for calculating the hardness ratio for counts data.\\footnote{ A Fortran and C-based program which calculates the hardness ratios following the methods described in this paper is available for download from \\url{{\\tt http://hea-www.harvard.edu/AstroStat/BEHR/}} }  \\subsection{The Classical Method} \\label{park:sec:classic}  A conventional hardness ratio is calculated as set out in Equations~\\ref{e:RCHR}, where $S$ and $H$ are the ``soft'' and ``hard'' counts accumulated in two non-overlapping passbands. In general, an observation will be contaminated by background counts, and this must be accounted for in the estimate of the hardness ratios. The background is usually estimated from an annular region surrounding the source of interest, or from a suitable representative region on the detector that is reliably devoid of sources. The difference in the exposure time and aperture area of source and background observations are summarized by a known constant $r$ for which the expected background counts in the source exposure area are adjusted. With the background counts in the soft band ($B_S$) and the hard band ($B_H$) collected in an area of $r$ times the source region, the hardness ratio is generalized to \\begin{eqnarray} \\nonumber \\R &=& \\frac{S-B_S/{r}}{H-B_H/{r}} \\\\ \\nonumber \\C &=& \\log_{10}\\left(\\frac{S-B_S/{r}}{H-B_H/{r}}\\right) \\\\ \\HR &=& \\frac{(H-B_H/{r})-(S-B_S/{r})}{(H-B_H/{r})+(S-B_S/{r})} \\label{e:bgRCHR} \\end{eqnarray} The adjusted counts in the background exposure area are directly subtracted from those in the source exposure area. The above equations can be further modified to account for variations of the detector effective areas by including them in the constant $r$, in which case the constants for the two bands will be different.\\footnote{ Gross differences in detector sensitivity, e.g., between different observations, or for sources at different off-axis angles, can be accounted for with additional multiplicative factors modifying the background-subtracted counts.  This is however not statistically meaningful and can lead to unreliable estimates and error bars in the small counts regime.  Instead, we apply such scaling factors directly to the source intensities (see Equation~\\ref{park:eq:SH}). } Errors are propagated assuming a Gaussian regime, i.e., \\begin{eqnarray} \\nonumber \\sig_{{\\R}} &=& \\ds\\frac{S-B_S/{r}}{H-B_H/{r}} \\sqrt{\\:\\frac{\\sig_{S}^2+\\sig_{B_S}^2/{r}^2}{(S-B_S/r)^2} + \\frac{\\sig_{H}^2+\\sig_{B_H}^2/{r}^2}{(H-B_H/r)^2}} \\\\ \\nonumber \\sig_{{\\C}} &=& \\ds\\frac{1}{\\ln(10)} \\ \\sqrt{\\:\\frac{\\sig_{S}^2+\\sig_{B_S}^2/{r}^2}{(S-B_S/r)^2} + \\frac{\\sig_{H}^2+\\sig_{B_H}^2/{r}^2}{(H-B_H/r)^2}} \\\\ \\sig_{{\\HR}} &=& \\ds\\frac{2\\: \\sqrt{(H-B_H/r)^2\\big(\\sig_{S}^2 +\\sig_{B_S}^2/{r}^2\\big) + (S-B_S/r)^2\\big(\\sig_{H}^2 +\\sig_{B_H}^2/{r}^2\\big)}}{\\big[(H - B_H/r)+(S - B_S/r)\\big]^2} \\label{e:sigbgRCHR} \\end{eqnarray} where $\\sig_{X_S}$, $\\sig_{X_H}$, $\\sig_{B_S}$, and $\\sig_{B_H}$ are typically approximated with the Gehrels prescription (Gehrels 1986) \\begin{equation} \\sig_{X} \\approx \\sqrt{X+0.75}+1 \\,, \\label{park:eq:sig2} \\end{equation} where $X$ are the observed counts, and the deviation from $\\sqrt{X}$ is due to identifying the 68\\% Gaussian (1$\\sigma$) deviation with the $16^{\\rm th}-84^{\\rm th}$ percentile range of the Poisson distribution. In addition to the approximation of a Poisson with a faux Gaussian distribution, classical error-propagation also fails on a fundamental level to properly describe the variance in the hardness ratios. Because $\\R$ is strictly positive, its probability distribution is skewed and its width cannot be described simply by $\\sig_{{\\R}}$. The fractional difference $\\HR$ is better behaved, but the range limits of $[-1,+1]$ cause the coverage rates (see \\S\\ref{park:sec:verify}) to be incorrect.  The color $\\C$ is the best behaved, but the error range will be asymmetrical when the errors are large.  Moreover, it has been demonstrated (van Dyk et al.\\ 2001) that using background subtracted source counts (which is the case in the method outlined above) gives a biased estimate of the source intensity, especially for very faint sources. Furthermore, although we can account for gross differences in the detector sensitivity between observations, this is not done in a statistically correct manner, but instead simply by rescaling the background-subtracted counts in each band using additional pre-computed correction factors. Finally, the classical hardness ratio method cannot be applied when a source is not detected in one of the two bands (which is quite common since CCD detectors are more sensitive in the soft band).  An alternative method  based on the median (or some quantile) of the distribution of the energies of the detected counts in a source was suggested by Hong et al.\\ (2004).  This method can be very powerful for classifying faint sources, but is not suited for comparisons of spectral properties between observations with different instruments. In this method, because the spectral parameter grids are highly non-linear and multi-valued in quantile-quantile space, it is also necessary that the spectral shape of the source being analyzed be known in order to interpret the numerical values.  We have therefore developed fully Bayesian approaches to appropriately compute hardness ratios and their errors, by properly accounting for the Poissonian nature of the observations.  The value of our method is heightened in several astrophysically meaningful cases. In \\S\\ref{park:sec:model}, we describe the model structure we adopt.  In \\S\\ref{park:sec:verify}, we carry out simulations to compare our method with the classical method.  In \\S\\ref{park:sec:apply}, we outline various applications of our method; we demonstrate its use in identifying candidate quiescent Low-mass X-ray binary % sources in a globular cluster and in studying the evolution of a stellar X-ray flare. In \\S\\ref{park:sec:discuss}, we discuss some nuances of our method such as the adopted prior distributions, as well as the advantages and limitations.  A summary of the method and its advantages is presented in \\S\\ref{park:sec:conclusion}. A detailed mathematical derivation of the posterior probability distributions of the hardness ratios and the computational methods we use to calculate them are described in Appendix~\\ref{sec:appA}, and a detailed comparison of the effects of priors is described in Appendix~\\ref{sec:appB}. ",
        "conclusions": "\\label{park:sec:discuss}  \\subsection{Informative and Non-informative Prior Distributions} \\label{park:sec:priors}  A major component of Bayesian analysis is the specification of a prior distribution, the probability distribution assigned {\\sl a priori} to the parameters defining the model.  This is often deemed controversial because of the apparent subjectivity in the assignment: if different people assign different priors and obtain different results, how is it possible to evaluate them?  The answer to this conundrum is to realize that all problems always include a bias brought to the analysis by the researcher, either in the choice of statistic, or even in the choice of analysis method, and the specification of a prior distribution codifies and makes explicit this bias.  Furthermore, if indeed prior information about the value or the range of the parameter is available, the prior allows a simple and principled method by which this information can be included in the analysis.  If no prior information is available, a so-called {\\sl non-informative} prior distribution must be chosen in order to not bias the result.  In our case, if there is a strong belief as to the value or range of the hardness ratio, we can incorporate this via an informative prior distribution, encoded as a $\\gamma$-distribution (see van Dyk et al.\\ 2001).  In addition, the analysis of previously acquired data will produce a posterior probability distribution that can be used as an informative prior on future observations of the same source.  When no prior information is available, we normally use a non-informative prior distribution. Since a Poisson intensity is necessarily positive, three types of non-informative prior distributions immediately present themselves: when $X|\\th\\sim\\Pois(\\th)$, \\begin{mathletters} \\begin{enumerate} \\item a non-informative prior distribution on the original scale, \\begin{equation} p(\\th)\\propto1 \\,, \\label{e:prior_standard} \\end{equation} \\item a Jeffrey's non-informative prior distribution, \\begin{equation} p(\\th)\\propto I_\\th^{1/2} \\,, ~{\\rm and} \\label{e:prior_jeffrey} \\end{equation} \\item a non-informative prior distribution under a log transformation, \\begin{equation} p(\\log\\th)\\propto1 ~{\\rm or~equivalently}~p(\\th)\\propto \\frac{1}{\\th} \\,, \\label{e:prior_log} \\end{equation} \\end{enumerate} \\end{mathletters} where $I_\\th=\\E[-\\partial^2\\log p(X|\\th)/\\partial\\th^2|\\th]$ is the expected Fisher information (Casella \\& Berger 2002). The first non-informative prior distribution (Equation~\\ref{e:prior_standard}) is flat from 0 to $\\infty$; the second (Equation~\\ref{e:prior_jeffrey}) is proportional to the square root of the Fisher information; and the third (Equation~\\ref{e:prior_log}) is flat on the whole real line under a log transformation. The functional forms of these prior distributions can be generalized into $p(\\th)\\propto\\th^{\\phi-1}$, i.e., $\\th\\sim\\gamma(\\phi,0)$ where $\\phi$ is an index that varies from $0$ to $1$: the first corresponds to $\\phi=1$, the second to $\\phi=1/2$; and the third to $\\phi=0^+$ (where this notation indicates that $\\phi>0$, but arbitrarily close to zero, for reasons described below).  We note that these non-informative prior distributions are {\\sl improper}, i.e., are not integrable.  If an improper prior distribution causes a posterior distribution to be improper as well, then no inferences can be made using such non-integrable distributions.  In our case, as long as $\\phi$ is strictly positive, a posterior distribution remains {\\sl proper}, i.e., integrable.  Hence, in the third case, we adopt values of $\\phi$ that are strictly positive but close in value to $0$, e.g., $\\phi=0.01$ or $\\phi=0.1$.  Note that while these prior distributions are non-informative, in the sense that in most cases they do not affect the calculated values of the hardness ratios (see Appendix~\\ref{sec:appB}), they do codify some specific assumptions regarding the range of values that it is believed possible.  For instance, if a large dynamic range is expected in the observations, a flat prior distribution in the log scale is more appropriate than a flat prior distribution in the original scale.  The choice of the prior distribution is dictated by the analysis and must be reported along with the results.  \\subsection{$\\R$ versus $\\C$ versus $\\HR$} \\label{park:sec:HR}  At low counts, the posterior distribution of the counts ratio, $\\R$, tends to be skewed because of the Poissonian nature of data; $\\R$ only takes positive values. The color, $\\C=\\log_{10}\\R$, is a log transformation of $\\R$, which makes the skewed distribution more symmetric. The fractional difference hardness ratio, $\\HR=(1-\\R)/(1+\\R)$, is a monotonically decreasing transformation of $\\R$, such that $\\HR\\rightarrow+1$ as $\\R\\rightarrow0$ (i.e., a source gets harder) and $\\HR\\rightarrow-1$ as $\\R\\rightarrow\\infty$ (i.e., a source gets softer). The monotonic transformation results in a bounded range of $[-1,+1]$ for $\\HR$, thereby reducing the asymmetry of the skewed posterior distribution. $\\R$ and $\\HR$ are bounded on one side or two sides, while $\\C$ is unbounded due to the log transformation.  The posterior distribution of any hardness ratio becomes more symmetric as both soft and hard source intensities increase.  Regardless of the size of the intensities, however, the color has the most symmetric posterior distribution among the popular definitions of a hardness ratio. Figure~\\ref{park:fig:dens} illustrates the effect of the magnitude of the source intensities on the symmetry of the posterior distribution of each hardness ratio; the posterior distribution of $\\C$ is confirmed to have the most symmetric posterior distribution. In the figure, we fix $\\R=2$ and the soft and hard intensities are determined by beginning with $\\lam_S=2$ and $\\lam_H=1$ (in units of counts~(source~area)$^{-1}$) and increasing the intensities by a factor of 5 in each subsequent column. We assume no background contamination in each simulation.  \\begin{figure}[htb!] \\centerline{ \\includegraphics[width=5.8in,angle=270]{f8.eps}} \\caption{The posterior distributions of hardness ratios calculated for $\\R$ (top row), $\\C$ (middle row), and $\\HR$ (bottom row), with different source intensities (low at left and high at right) and prior distributions. The different curves in each figure represent the posterior probability distributions computed with different non-informative prior distribution indices, i.e., $\\phi=0^+$ (solid), $\\phi=1/2$ (dashed), and $\\phi=1$ (dotted). At higher counts (large Poisson intensities; right column of figures), the posterior distributions tends to be symmetric and the effect of the prior distributions is minimal, i.e., the posterior distributions with different indices are nearly the same.  At small counts (small Poisson intensities; the left column of figures), the non-symmetric shape of the posterior distribution for each hardness ratio becomes clear as does the effect of the choice of a non-informative prior distribution.} \\label{park:fig:dens} \\end{figure}  Figure~\\ref{park:fig:dens} also illustrates the effect of the use of different indices for the non-informative prior distribution on the resulting posterior distribution of each hardness ratio. In the case of $\\R$, the posterior distribution becomes diffuse and more skewed toward $0$ as the value of $\\phi$ decreases; as expected, this is also the case for $\\C$.  And in the case of $\\HR$, the posterior distribution has more mass near the boundaries (i.e., $\\pm1$) as the value of $\\phi$ decreases, thereby exhibiting a U-shaped density in the case of very faint sources.  As the intensities increase, however, the choice of a prior distribution has less effect on the posterior distribution of any of the hardness ratios.  Finally, as pointed out in \\S\\ref{park:sec:verify}, for $\\R$ and $\\C$, the mode of the posterior probability distribution is a robust estimator, whereas for $\\HR$, the mean of the distribution is a better choice. (Notice that the posterior mode of $\\HR$ in the low count scenario is $-1$ when $\\phi=0^+$ in Figure~\\ref{park:fig:dens}.)  \\subsection{Advantages} \\label{park:sec:advan}  A significant improvement that our method provides over the classical way of computing hardness ratios is that we use the correct Poisson distribution throughout and do not make any limiting (high counts) Gaussian assumptions.  Thus, while the classical method works well only with high counts data and fails to give reliable results with low counts, our method is valid in all regimes.  Indeed, because the observed counts are non-negative integers, it is not appropriate to model low counts with Gaussian distributions which are defined on all the real numbers.  Because our methods are based on Poisson assumptions in a fully model based statistical approach, we need not rely on plug-in estimates of the hardness ratio. Instead we can evaluate the full posterior probability distribution of the hardness ratio, which provides reliable estimates and correct confidence limits even when either or both soft and hard counts are very low.  In particular, our method is not limited to ``detectable'' counts, and can properly calculate upper and lower bounds in cases where the source is detected in only one passband.  For high counts, our method gives more precise error bars for the hardness ratios than the classical method (as defined by the coverage; see Table~\\ref{park:tbl:coverage} and Appendix~\\ref{sec:appA}), although both methods yield similar results.  Further, {\\sl a priori} information can be embedded into our model and can be updated by data, thereby producing more accurate results as we collect more data.  \\subsection{Limitations} \\label{park:sec:limit}  Unlike the quantile-width method (Hong et al.\\ 2004), the calculation of hardness ratios is limited by necessity to pre-selected non-overlapping passbands.  In general, we must use different band definitions in order to achieve the maximum discriminating power for soft and hard sources.  {\\sl A priori}, with no knowledge of the character of the source populations, there is no way to ensure that the chosen passbands are the most efficient for splitting the spectrum and obtaining maximum leverage to distinguish spectral model parameters.  However, this restriction allows for a more uniform analysis and comparison of results across different instruments and telescopes.  Because the Bayesian method does not allow for a simple analytic solution similar to standard error propagation in the Gaussian case, the computational methods used to determine the probability distributions become important.  We have implemented a Monte Carlo integration method (the Gibbs sampler, see Appendix~\\ref{park:sec:gibbs}) and a method based on numerical integration (quadrature, see Appendix~\\ref{park:sec:quad}).  The Gibbs sampler is based on Monte Carlo simulation and hence causes our estimates to have simulation errors in addition to their true variability; to reduce the levels of simulation errors, the Markov chain must be run with large enough iterations. On the other hand, the quadrature precisely computes the posterior distribution as long as the number of bins is large enough; however, its computation becomes expensive as the counts become large because of the binomial expansion in Equation~\\ref{park:eq:binom-expan}. In general, the Gibbs sampler is much faster than the method based on quadrature, but care must be taken to ensure that the number of iterations is sufficient to determine the posterior mode and the required posterior interval with good precision. We recommend using the Gibbs sampler for high counts and the quadrature for low counts."
    },
    "0606/astro-ph0606592_arXiv.txt": {
        "abstract": "Particle acceleration at astrophysical shocks may be very efficient if magnetic scattering is self-generated by the same particles. This nonlinear process adds to the nonlinear modification of the shock due to the dynamical reaction of the accelerated particles on the shock. Building on a previous general solution of the problem of particle acceleration with arbitrary diffusion coefficients (\\cite{amato1}), we present here the first semi-analytical calculation of particle acceleration with both effects taken into account at the same time: charged particles are accelerated in the background of Alfv\\'en waves that they generate due to the streaming instability, and modify the dynamics of the plasma in the shock vicinity. ",
        "introduction": "\\label{sec:intro}  Soon after the pioneering papers by \\cite{krymskii,bo78,bell78a,bell78b}, introducing the test particle theory of particle acceleration at collisionless shocks, it became clear that the dynamical reaction of the accelerated particles on the plasmas involved in the shock formation may not be negligible. It is now clear that such reaction may in fact make shocks efficient accelerators and change quite drastically the predictions of the test particle theory. The main consequences of the shock modification induced by the accelerated particles can be summarized as follows: 1) a precursor, consisting in a gradual braking of the upstream fluid, is created; 2) particles with different momenta {\\it feel} different effective compression factors, which reflects in the fact that the spectrum of accelerated particles is no longer a power law, but rather a concave spectrum, as hard as $p^{-3.2}$ at high momenta; 3) the shock becomes less efficient in heating the background plasma, so that the temperature of the downstream gas is expected to be lower than predicted through the usual Rankine-Hugoniot relations at an unmodified shock front (see \\cite{drury83,be87,je91,maldru2001} for reviews on different aspects of the subject). The reaction of the accelerated particles has been calculated within different approaches: the so-called two-fluid models (\\cite{dr_v80,dr_v81}), kinetic models (\\cite{malkov1,malkov2,blasi1,blasi2}) and numerical approaches, both Monte Carlo and other simulation procedures (\\cite{je91,bell87,elli90,ebj95,ebj96,kj97,kj05,jones02}). In most of these calculations, the diffusion properties of the plasma upstream and downstream are provided as an input to the problem. This also results in fixing the value of the maximum momentum of the accelerated particles. However, one of the well known and most disturbing problems associated with the mechanism of particle acceleration at shock fronts is that a substantial amount of magnetic scattering of the particles is required (e.g. \\cite{lc83a,lc83b}). In the absence of it, the maximum energy of the accelerated particles is exceedingly low and uninteresting for astrophysical applications (e.g. \\cite{blasirev}). \\cite{bell78a} proposed that the streaming instability of cosmic rays could be responsible for the generation of perturbations in the magnetic field of an amplitude necessary to provide pitch angle scattering (and therefore spatial diffusion) of the accelerated particles. \\cite{lc83b} used this argument to estimate the maximum energy of particles accelerated at shocks in supernova remnants. In all previous works either the shock was considered unmodified, or the diffusion coefficients were fixed {\\it a priori}, because a comprehensive theory of particle acceleration was missing. Recently, \\cite{amato1} have found a general exact solution of the system of equations describing the diffusion-convection of accelerated particles, and the dynamics and thermodynamics of plasmas in the shock region, for an arbitrary choice of the spatial and momentum dependence of the diffusion coefficient. In the present paper we use the formalism proposed by \\cite{amato1} and combine it with calculations of the perturbations created through streaming instability, so that the diffusion coefficient, as a function of spatial location and momentum, is determined from the spectrum and spatial distribution of the accelerated particles. This provides the first combined description of the process of particle acceleration at collisionless shocks in the presence of particle reaction and wave generation. In this approach, the spectrum of accelerated particles, their distribution in the upstream plasma and the diffusion coefficient are outputs of the problem.  The paper is organized as follows: in Sec.~\\ref{sec:solution} we summarize the findings of \\cite{amato1}. In Sec.~\\ref{sec:diff} we illustrate our treatment of the streaming instability and determine a relation between the power spectrum of magnetic fluctuations and the diffusion coefficient upstream. In Sec.~\\ref{sec:full} we describe the results of our calculations. In Sec.~\\ref{sec:heating} we shortly discuss how the results presented in the previous section change when the effects of turbulent heating are taken into account. We conclude in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We described the mathematical theory of particle acceleration at non-relativistic shock fronts with dynamical reaction of accelerated particles and self-generated scattering waves. The diffusion coefficient itself is an output of the calculations, though within the limitations imposed by the usage of quasi-linear theory applied to the case of potentially strong magnetic field amplification. The scattering in the upstream plasma is generated through streaming instability, as discussed extensively in previous literature.  We determined the spectra of accelerated particles, their spatial distribution and the space dependence of the fluid velocity, pressure and temperature. The diffusion coefficient and the strength of the self-generated magnetic perturbations are also calculated, as a function of the distance from the shock front in the precursor. We confirm the general finding that the spectra of accelerated particles are concave, an effect which is particularly evident for strongly modified shocks, namely for large Mach numbers of the moving fluid. However, the shape of the concavity is somewhat affected by the self-determined diffusion coefficient, as visible in Fig. \\ref{fig:modspec}.  Having in mind the comparison between the predicted spectra at the sources and the observed cosmic ray spectrum at the Earth, it is worth reminding the reader that what can actually be measured is the combination of the diffusion time, the gas density along the trajectory (responsible for the spallation) and the injection spectrum. In order to infer some conclusions about the spectrum at the source, one has to make assumptions on the diffusion coefficient in the interstellar medium. In alternative it would be a precious step forward if we could measure unambiguously the spectrum of gamma rays generated by $\\pi^0$ decays close to the source itself, an evidence that unfortunately is still missing.  The asymptotic slope of the spectra for $p\\to p_{max}$ may be as flat as $\\sim 3.1-3.2$, but this conclusion is not strongly affected by the fact that the diffusion coefficient is calculated self-consistently.  The most striking new result of our calculations is the energy dependence of the diffusion coefficient and the strength of the amplified turbulent magnetic field. As could be expected, the diffusion coefficient is not Bohm-like, and the turbulent component of the magnetic field is amplified so efficiently that the diffusion coefficient is much smaller than the Bohm coefficient in the background magnetic field. This is especially true at the highest momenta, which leads to think that a full non-linear theory might predict higher values of the maximum momentum than expected on naive grounds. Unfortunately a full, self-consistent calculation of $p_{max}$ for a strongly modified shock has never been carried out, the main difficulty being in accounting for the spatial dependence of all the quantities involved.  When compared with the Bohm diffusion coefficient as calculated in the amplified magnetic field, our diffusion coefficient remains always larger, with the possible exception of a narrow momentum region close to $p_{max}$.  While the calculation presented here is fully self-consistent in the determination of the shock modification due to the reaction of the accelerated particles, the part related to the amplification of the background field suffers from all the limitations related to the usage of quasi-linear theory for the streaming instability. This approach, initially developed for weakly amplified magnetic fields, is widely applied in the literature to situations that violate this condition. Unfortunately at the present time this is the only way we have to achieve a (at least partially) self-consistent picture of the process of particle acceleration at cosmic ray modified shocks with self-generated turbulence. This problem is in fact even more serious for those approaches that predict levels of magnetic field amplification which are much higher than those found here (e.g. \\cite{lb01,bell04}).  The high acceleration efficiencies obtained in the context of all approaches to particle acceleration at shocks are known to be reduced by the effect of turbulent heating. Any non-adiabatic heating of the gas in the precursor leads to reducing the energy channelled into non-thermal particles at the shock. This is a serious problem, because the effect of turbulent heating depends dramatically on the type of mechanism that is responsible for the heating: Alfv\\'en heating, often used in the literature, is only one of these mechanisms, and not necessarily the most efficient. For instance, the instability induced by the propagation of acoustic waves in the precursor was shown to lead to the formation of weak shocks in the precursor, which in turn heat the upstream plasma (e.g. \\cite{drury}).  These non-linear effects can hardly be taken into account in a credible way. Most notably, the phenomenological expressions proposed in the literature and used also in the present paper, have originally been proposed as mechanisms to reduce the amount of magnetic field amplification and remain in the context of small perturbations of the background magnetic field. However, as shown in Fig. 5, even with the Alfv\\'en heating taken into account, the magnetic field can be amplified by a factor in excess of $\\sim 10$ with respect to the background field. This means that a fully non-linear theory of the turbulent heating is required in order to make fully reliable predictions.  From the phenomenological point of view, the best evidences for both magnetic field amplification and efficient particle acceleration come from observations of supernova remnants (see the reviews of \\cite{hillas} and \\cite{blasirev} and references therein). In fact, it has been argued that the amount of field amplification required to explain the thickness of the X-ray bright rims in several remnants is of the order of $\\sim 200-300\\mu G$ (\\cite{voelk05}). An important role in explaining this level of amplification could be played by different versions of the streaming instability (\\cite{lb01,bell04,bell05}), not requiring resonant interactions of particles and waves. A full non-linear theory including these effects will be described elsewhere (Amato and Blasi, in preparation)."
    },
    "0606/astro-ph0606071_arXiv.txt": {
        "abstract": "The interstellar medium is highly dynamic and turbulent.  However, little or no attention has been paid in the literature to the implications that this fact has on the validity of at least six common assumptions on the Virial Theorem (VT), which are: (i) the only role of turbulent motions within a cloud is to provide support against collapse, (ii) the surface terms are negligible compared to the volumetric ones, (iii) the gravitational term is a binding source for the clouds since it can be approximated by the gravitational energy, (iv) the sign of the second-time derivative of the moment of inertia determines whether the cloud is contracting ($\\ddot{I}<0$) or expanding ($\\ddot{I}>0$), (v) interstellar clouds are in Virial Equilibrium (VE), and (vi) Larson's (1981) relations (mean density-size and velocity dispersion-size) are the observational proof that clouds are in VE.  Turbulent, supersonic interstellar clouds cannot fulfill these assumptions, however, because turbulent fragmentation will induce flux of mass, moment and energy between the clouds and their environment, and will favor local collapse while may disrupt the clouds within a dynamical timescale. It is argued that, although the observational and numerical evidence suggests that interstellar clouds are not in VE, the so-called ``Virial Mass'' estimations, which actually should be called ``energy-equipartition mass'' estimations, are good order-of magnitude estimations of the actual mass of the clouds just because observational surveys will tend to detect interstellar clouds appearing to be close to energy equipartition.  Similarly, order of magnitude estimations of the energy content of the clouds are reasonable.  However, since clouds are actually out of VE, as suggested by asymmetrical line profiles, they should be transient entities.  These results are compatible with observationally-based estimations for rapid star formation, and call into question the models for the star formation efficiency based on clouds being in VE. ",
        "introduction": "\\label{sec:intro}  Interstellar clouds are thought to be turbulent and supersonic.  Their Mach numbers range from a 1 ($T \\sim$ 7000--8000~K, H~I clouds) through 10 ($T\\sim$~10~K, molecular clouds forming low-mass stars) to 50 ($T\\sim$~10--50~K, molecular clouds forming high-mass stars). Since supersonic turbulent motions carry mass and produce large-amplitude density fluctuations, turbulent fragmentation\\footnote{Turbulent fragmentation is defined as the process through which a chaotic velocity field produces a clumpy density structure in the gas within a few dynamical timescales \\cite[see e.g.,][]{VonWeiczacker51, Sasao73, Scalo88, Padoan95, BVS99, KHM00, HMK01, BP04M}.} is expected to occur in the interstellar medium.  A useful tool for describing the overall structure of interstellar clouds is the scalar Virial Theorem (VT).  It is obtained by dotting the momentum equation by the position vector and integrating over the volume of interest (\\S\\ref{sec:VT}).  Being directly derived from the momentum equation, the VT always holds for any parcel of fluid.  Virial Equilibrium (VE) is a restrictive condition of the VT, and it is defined by the condition that the parcel of fluid under study has a second time derivative of the moment of inertia equal to zero.  It has been invoked extensively to analyze the stability of interstellar clouds.  However, since in a trans- or super-sonic turbulent interstellar medium, clouds should be redistributing their mass as a consequence of their own turbulent motions, it already seems difficult to achieve VE in a supersonic, turbulent cloud.  Other simplifications, such as the assumption that the cloud is isolated, or that the surface terms in the VT are negligible, are frequently made in many (if not in most) astrophysical studies of the VT.  Those simplifications have been thought to be applicable to molecular clouds and their substructure for nearly three decades \\cite[e.g., the textbooks by][ and references therein]{Spitzer78, Shu91, Stahler_Palla05, Lequeux05}, maybe as a consequence of the old idea that in the ISM, ``all forces are in balance and the medium is motionless, with no net acceleration'' \\cite[][ Chap. 11]{Spitzer78}, in which ``the observational evidence'' seemed to be consistent with the expectation that interstellar ``clouds tend toward pressure equilibrium'' \\cite[][]{Spitzer78}.  The possible inapplicability of these assumptions has been mentioned in passing in some previous papers \\cite[e.g., ][]{BVS99, Shadmehri_etal02, BP04P} but no attention has been paid in general to its implications.  Thus, in the present paper I discuss in detail the applicability of those assumptions for molecular clouds and their cores.  In \\S\\ref{sec:VT} I write explicitly the VT for fluids in its Lagrangian and Eulerian forms.  In \\S\\ref{sec:myths} I discuss the six more common assumptions of the VT and their validity in a turbulent environment.  In \\S\\ref{sec:reloading} I explain why even though clouds are not in VE, they appear to be in energy equipartition, and argue that asymmetries in the line profiles are the evidence for clouds out of VE.  Finally, in \\S\\ref{sec:conclusions} I draw the main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  The present contribution has discussed the applicability of the six more common assumptions on the Virial Theorem.  Specifically,  \\begin{enumerate}  \\item{} It was shown that a decomposition of the velocity field into its vortical and compressible modes is necessary, since only modes satisfying the condition $\\nabla\\cdot u > 0$ provide support, while modes satisfying $\\nabla\\cdot u < 0$ foment collapse.  \\item{} It was argued that for a supersonic, turbulent ISM, surface terms should not be neglected.  \\item{} It was shown that the gravitational term can be decomposed into a contribution from the cloud itself, and a contribution from the outside.  The first part is the well-known gravitational energy. The second part is the sum of three terms: The gravitational pressure evaluated at the surface of the cloud plus three times the work done against the external mass to assemble the density distribution of the cloud, plus a term that depends on the gradient of the density distribution of the cloud.  These represent the tidal forces due to an external potential, as could be the case of a dense core within a giant molecular cloud, or a giant molecular cloud close to a spiral arm.  It is argued that this contribution can be as important as the self-gravitational energy.  \\item{} Using a simple counter-example, it was shown that the sign of the second-time derivative of the moment of inertia does not determine whether the cloud is contracting or expanding. An expanding cloud may very well satisfy the condition $\\ddot{I}<0$, and a contracting one may satisfy $\\ddot{I}>0$, contrary to the common belief.  In other words, $\\ddot{I}$ has been treated in the literature as if it were $\\dot{I}$.  \\item{} It was argued that interstellar clouds are not likely to satisfy the Virial Equilibrium (VE) condition $\\ddot{I}=0$.  \\item{} It was shown that Larson's (1981) relations are not observational proof for clouds being in VE.  \\item{} Clouds seem to be in energy equipartition because of either observational limitations, as well as because of the intrinsic definition of a cloud.  \\end{enumerate}   Turbulent fragmentation plays a crucial role for the inapplicability of the VT to interstellar clouds, since it will induce a flux of mass, moment and energy between the clouds and their environment, and will favor local collapse while disrupting the clouds within a dynamical timescale.  The common assumptions discussed in the present contribution drive our understanding of the dynamical state of the interstellar clouds toward a picture that favors a static ISM. However, they are highly difficult to fulfill if the ISM is highly turbulent, as it was found to be many years ago \\cite[e.g., ][]{McCray_Snow79}.  Inferences of the star formation efficiency for supersonic (Mach numbers $\\sim 20-40$) clouds in virial equilibrium living several dynamical timescales \\citep[e.g., ][]{Krumholz_McKee05, Tan_etal06} should be taken with caution.  The lack of observational evidence for clouds being in VE, and the identificaion of asymmetrical line profiles observed toward interstellar clouds using different tracers are the best evidence of clouds being out-of-equilibrium systems. These facts lead us to the conclusion that clouds should be transient structures, which exchange mass, momentum and energy with their environment.  This is precisely the opposite point of view to the old one in which clouds should be at rest, which still is present in textbooks and papers, but it is consistent with recent observationally-based works that favor a scenario of rapid cloud and star formation \\cite[see, e.g., the reviews by ][ and references therein.]{MK04, BKMV06}."
    },
    "0606/hep-ph0606014_arXiv.txt": {
        "abstract": "s{ Axions solve the Strong CP Problem and are a cold dark matter candidate.  The combined constraints from accelerator searches, stellar evolution limits and cosmology suggest that the axion mass is in the range $3 \\cdot 10^{-3} > m_a > 10^{-6}$ eV. The lower bound can, however, be relaxed in a number of ways. I discuss the constraint on axion models from the absence of isocurvature perturbations.   Dark matter axions can be searched for on Earth by stimulating their conversion to microwave photons in an electromagnetic cavity permeated by a magnetic field.  Using this technique, limits on the local halo density have been obtained by the Axion Dark Matter eXperiment.} ",
        "introduction": "The standard model of elementary particles has been an extraordinary breakthrough in our description of the physical world.  It agrees with experiment and observation disconcertingly well and surely provides the foundation for all further progress in our field.  It does however present us with a puzzle.  Indeed the action density includes, in general, a term \\begin{equation} {\\cal L}_{\\rm stand~mod} = ... ~+~{\\theta g^2\\over 32\\pi^2} G^a_{\\mu\\nu} \\tilde G^{a\\mu\\nu} \\label{ggdual} \\end{equation} where $G^a_{\\mu\\nu}$ are the QCD field strengths, $g$ is the QCD coupling constant and $\\theta$ is a parameter.  The dots represent all the other terms in the action density, i.e. the terms that lead to the numerous successes of the standard model.  Eq.~(\\ref{ggdual}) perversely shows the one term that isn't a success.  Using the statement of the chiral anomaly \\cite{abj}, one can show three things about that term.  First, that QCD physics depends on the value of the parameter $\\theta$ because in the absence of such dependence QCD would have a $U_A(1)$ symmetry in the chiral limit, and we know QCD has no such $U_A(1)$ symmetry \\cite{SW}.  Second, that $\\theta$ is cyclic, that is to say that physics at $\\theta$ is indistinguishable from physics at $\\theta + 2\\pi$.  Third, that an overall phase in the quark mass matrix $m_q$ can be removed by a redefinition of the quark fields only if, at the same time, $\\theta$ is shifted to $\\theta - \\arg\\det m_q$.  The combination of standard model parameters $\\bar{\\theta} \\equiv \\theta - \\arg \\det m_q$ is independent of quark field redefinitions.  Physics therefore depends on $\\theta$ solely through $\\bar{\\theta}$.  Since physics depends on $\\bar{\\theta}$, the value of $\\bar{\\theta}$ is determined by experiment.  The term shown in Eq.~(\\ref{ggdual}) violates P and CP.  This source of P and CP violation is incompatible with the experimental upper bound on the neutron electic dipole moment unless $|\\bar{\\theta}| < 10^{-10}$.  A new improved upper limit on the neutron electric dipole moment ($|d_n| < 3.0 \\cdot 10^{-26}~e~$cm) was reported at this conference by P. Geltenbort \\cite{ned}.  The puzzle aforementioned is why the value of $\\bar{\\theta}$ is so small.  It is usually referred to as the ``Strong CP Problem\".  If there were only strong interactions, a zero value of $\\bar{\\theta}$ could simply be a consequence of P and CP conservation.  That would not be much of a puzzle. But there are also weak interactions and they, and therefore the standard model as a whole, violate P and CP.  So these symmetries can not be invoked to set $\\bar{\\theta} = 0$.  More pointedly, P and CP violation are introduced in the standard model by letting the elements of the quark mass matrix $m_q$ be arbitrary complex numbers \\cite{KM}.  In that case, one expects $\\arg \\det m_q$, and hence $\\bar{\\theta}$, to be a random angle.  The puzzle is removed if the action density is instead \\begin{equation} {\\cal L}_{\\rm stand~mod~+~axion} =~... ~+~{1 \\over 2}\\partial_\\mu a \\partial^\\mu a ~+~{g^2\\over 32\\pi^2}~{a(x) \\over f_a}~G^a_{\\mu\\nu} \\tilde G^{a\\mu\\nu} \\label{ax} \\end{equation} where $a(x)$ is a new scalar field, and the dots represent the other terms of the standard model.  $f_a$ is a constant with dimension of energy.  The $a G \\cdot \\tilde G$ interaction in Eq.~(\\ref{ax}) is not renormalizable.  However, there is a recipe for constructing renormalizable theories whose low energy effective action density is of the form of Eq.~(\\ref{ax}).  The recipe is as follows: construct the theory in such a way that it has a $U(1)$ symmetry which (1) is a global symmetry of the classical action density, (2) is broken by the color anomaly, and (3) is spontaneously broken.  Such a symmetry is called Peccei-Quinn symmetry after its inventors \\cite{PQ}.  Weinberg and Wilczek \\cite{WW} pointed out that a theory with a $U_{\\rm PQ}(1)$ symmetry has a light pseudo-scalar particle, called the axion.  The axion field is $a(x)$.  $f_a$ is of order the expectation value that breaks $U_{\\rm PQ}(1)$, and is called the ``axion decay constant\".  In the theory defined by Eq.~(\\ref{ax}), $\\bar{\\theta} = {a(x) \\over f_a} - \\det\\arg m_q$ depends on the expectation value of $a(x)$.  That expectation value minimizes the effective potential.  The Strong CP Problem is then solved because the minimum of the QCD effective potential $V(\\bar{\\theta})$ occurs at $\\bar{\\theta} = 0$ \\cite{VW}. The weak interactions induce a small value for $\\bar{\\theta}$ \\cite{EG,GR}, of order $10^{-17}$, but this is consistent with experiment.  The notion of Peccei-Quinn (PQ) symmetry may seem contrived.  Why should there be a $U(1)$ symmetry which is broken at the quantum level but which is exact at the classical level?  However, the reason for PQ symmetry may be deeper than we know at present.  String theory contains many examples of symmetries which are exact classically but which are broken by quantum anomalies, including PQ symmetry \\cite{Wit,Kiw,Svr}.  Within field theory, there are examples of theories with {\\it automatic} PQ symmetry, i.e. where PQ symmetry is a consequence of just the particle content of the theory without adjustment of parameters to special values.  The first axion models had $f_a$ of order the weak interaction scale and it was thought that this was an unavoidable property of axion models.  However, it was soon pointed out \\cite{KSVZ,DFSZ} that the value of $f_a$ is really arbitrary, that it is possible to construct axion models with any value of $f_a$.  A value of $f_a$ far from any previously known scale need not lead to a hierarchy problem because PQ symmetry can be broken by the condensates of a new technicolor-like interaction \\cite{Kim}.  The properties of the axion can be derived using the methods of current algebra \\cite{curr}.  The axion mass is given in terms of $f_a$ by \\begin{equation} m_a\\simeq 6~eV~{10^6 GeV\\over f_a}\\, . \\label{ma} \\end{equation} All the axion couplings are inversely proportional to $f_a$. For example, the axion coupling to two photons is: \\begin{equation} {\\cal L}_{a\\gamma\\gamma} = -g_\\gamma {\\alpha\\over \\pi} {a(x)\\over f_a} \\vec E \\cdot\\vec B~~~\\ . \\label{aEB} \\end{equation} Here $\\vec E$ and $\\vec B$ are the electric and magnetic fields, $\\alpha$ is the fine structure constant, and $g_\\gamma$ is a model-dependent coefficient of order one.  $g_\\gamma=0.36$ in the DFSZ model \\cite{DFSZ} whereas $g_\\gamma=-0.97$ in the KSVZ model \\cite{KSVZ}.  The coupling of the axion to a spin 1/2 fermion $f$ has the form: \\begin{equation} {\\cal L}_{a \\overline f f} = i g_f {m_f \\over f_a} a \\overline f \\gamma_5 f \\label{cf} \\end{equation} where $g_f$ is a model-dependent coefficient of order one.  In the KSVZ model the coupling to electrons is zero at tree level.  Models with this property are called 'hadronic'.   The axion has been searched for in many places but not found \\cite{arev}. The resulting constraints may be summarized as follows.  Axion masses larger than about 50 keV are ruled out by particle physics experiments (beam dumps and rare decays) and nuclear physics experiments.  The next range of axion masses, in decreasing order, is ruled out by stellar evolution arguments.  The longevity of red giants rules out 200 keV $> m_a >$ 0.5 eV \\cite{astro,Raff87} in the case of hadronic axions, and 200 keV $> m_a > 10^{-2}$ eV \\cite{Schramm} in the case of axions with a large coupling to electrons [$g_e = 0(1)$ in Eq. \\ref{cf}]. The duration of the neutrino pulse from Supernova 1987a rules out 2 eV $> m_a > 3 \\cdot 10^{-3}$ eV \\cite{1987a}.  Finally, there is a lower limit, $m_a \\gtwid 10^{-6}$ eV, from cosmology which will be discussed in detail in the next section.  This leaves open an ``axion window\": $3 \\cdot 10^{-3} > m_a \\gtwid 10^{-6}$ eV.  We will see, however, that the lower edge of this window ($10^{-6}$ eV) is much softer than its upper edge. ",
        "conclusions": ""
    },
    "0606/astro-ph0606717_arXiv.txt": {
        "abstract": "Physical and wind properties of Galactic B supergiants are presented based upon non-LTE line blanketed model atmospheres, including Sher~25 toward the NGC~3603 cluster. We compare H$\\alpha$ derived wind densities with recent results for SMC B supergiants and generally confirm theoretical expectations for stronger winds amongst Galactic supergiants. Mid B supergiant winds are substantially weaker than predictions from current radiatively driven wind theory, a problem which is exacerbated if winds are already clumped in the H$\\alpha$ line forming region. We find that the so-called `bistability jump' at B1 (\\teff\\ $\\sim$ 21kK) from Lamers et al. is rather a more gradual downward trend.  CNO elemental abundances, including Sher~25, reveal partially processed material at their surfaces. In general, these are in good agreement with evolutionary predictions for blue supergiants evolving redward accounting for rotational mixing. A few cases, including HD~152236 ($\\zeta^{1}$ Sco), exhibit strongly processed material which is more typical of Luminous Blue Variables. Our  derived photospheric [N/O] ratio for Sher~25 agrees with that for its ring nebula, although a higher degree of CNO processing would be expected if the nebula originated during a red supergiant phase, as is suspected for the ring nebula ejected by the B supergiant progenitor of SN~1987A, Sk--69$^{\\circ}$ 202. Sher~25 has an inferred age of $\\sim$5Myr in contrast with $\\sim$2Myr for HD~97950, the ionizing cluster of NGC~3603, so it may be a foreground object or close binary evolution may be responsible for its unusual location in the H-R diagram. ",
        "introduction": "O stars dominate the energetics of young starbursts since their ionizing output is a very steep function of effective temperature (Leitherer et al. 1992). As such, numerous intensive studies of O stars in the Milky Way and Magellanic Clouds have been undertaken, with regard to determining their fundamental parameters and the empirical dependence of their wind strengths on metallicity (e.g. Mokiem et al. 2006ab). Unfortunately, O dwarfs are visually faint in external galaxies due to their large bolometric corrections and small radii, with typically $M_{\\rm V}$ = -5.0 mag.   O star abundances of CNO elements are difficult to derive (Crowther et al. 2002), yet these are sensitive to early rotational mixing (e.g. Meynet \\& Maeder 2000).  In contrast, B supergiants have received rather less attention since they are less relevant for energetics of young starbursts. Nevertheless, they are visually bright, typically $M_{\\rm V} = -7.0$ mag, possess larger radii and so are easy to observe individually in nearby galaxies, permitting robust tests of the metallicity dependence of radiatively driven wind theory, which is predicted to increase from early to mid B supergiants (Vink et al. 2000). In addition, CNO abundances are readily obtained from optical spectroscopy, for comparison to evolutionary model predictions. Notably, Trundle et al. (2004) and Trundle \\& Lennon (2005) have studied a large sample of SMC B supergiants using the line blanketed version of FASTWIND (Puls et al. 2005). To date, the most extensive study of Galactic B supergiants is by Kudritzki et al. (1999), who adopted the \\teff\\ scale from the plane-parallel unblanketed study of McErlean et al. (1999), plus mass-loss rates from the unblanketed version of FASTWIND.  Our present study of Galactic B supergiants employs CMFGEN (e.g.  Hillier et al. 2003) which also incorporates line blanketing and spherical geometry, and so permits a direct comparison with recent SMC results, plus CNO abundance determinations. Do the abundances of normal B supergiants match those of stars evolving towards red supergiants (RSG) or returning from RSG following convective dredge-up? Comparisons with Luminous Blue Variables (LBVs) are also made, since they exhibit spectral characteristics of early B supergiants at visual minimum and A supergiants at visual maximum (Crowther 1997). Notably, $\\zeta^{1}$ Sco (HD~152236) is an early B hypergiant and lies above the empirical Humphreys-Davidson limit (Humphreys \\& Davidson 1979). We also include new stellar and nebular studies of Sher~25 (Hendry et al. 2006), an apparently normal B supergiant in the giant HII region NGC~3603 yet possessing an ejecta nebula reminiscent of that associated with SN~1987A (Brandner et al.  1997b).  Finally, we consider the empirical bistability jump amongst early B supergiants, which was originally developed for the LBV P~Cygni by Pauldrach \\& Puls (1990), in the sense that slight variations in its effective temperature resulted in hydrogen being ionized or neutral in its outer stellar wind (Najarro et al.  1997).  Lamers et al. (1995) studied wind velocities of OB stars, and established a discontinuity in the ratio of \\vinf\\ to \\vesc\\ for B0.5 and B1 supergiants, which is assumed in current stellar wind models of Vink et al. (2000).  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.6\\columnwidth,clip,angle=-90]{t.eps} \\end{center} \\caption{Temperature scale of B supergiants from this Milky Way study (CMFGEN, circles) plus Trundle et al. (2004) and Trundle \\& Lennon (2005) for the SMC (FASTWIND, triangles) versus the McErlean et al. (1998, solid line) calibration.\\label{teff}} \\end{figure} ",
        "conclusions": ""
    },
    "0606/astro-ph0606521_arXiv.txt": {
        "abstract": "We examine the origins of the bimodality observed in the global properties of galaxies around a stellar mass of \\mbox{$3\\!\\times\\!10^{10}{\\rm M}_{\\sun}$} by comparing the environmental dependencies of star-formation for the giant and dwarf galaxy populations. The Sloan Digital Sky Survey DR4 spectroscopic dataset is used to produce a sample of galaxies in the vicinity of the supercluster centered on the cluster A2199 at $z=0.03$ that is \\mbox{$\\ga\\!9$0\\%} complete to a magnitude limit of \\mbox{${\\rm M}_{r}^{*}+3.3$}. From these we measure global trends with environment for both giant \\mbox{(M$_{r}\\!<\\!-20$ mag)} and dwarf \\mbox{($-19<{\\rm M}_{r}\\!<\\!-17.8$ mag)} subsamples using the luminosity-weighted mean stellar age and H$\\alpha$ emission as independent measures of star-formation history. The fraction of giant galaxies classed as old (\\mbox{$t>\\!7$\\,Gyr}) or passive (\\mbox{EW$[{\\rm H}\\alpha]\\le4$\\,\\AA}) falls gradually from $\\ga\\!8$0\\% in the cluster cores to \\mbox{$\\sim\\!4$0\\%} in field regions beyond \\mbox{3--4\\,$R_{vir}$}, as found in previous studies. In contrast, we find that the dwarf galaxy population shows a sharp transition at \\mbox{$\\sim\\!1\\,R_{\\rm vir}$}, from being predominantly old/passive within the cluster, to outside where virtually all galaxies are forming stars and old/passive galaxies are {\\em only} found as satellites to more massive galaxies. These results imply fundamental differences in the evolution of giant and dwarf galaxies: whereas the star-formation histories of giant galaxies are determined primarily by their merger history, star-formation in dwarf galaxies is much more resilient to the effects of major mergers. Instead dwarf galaxies become passive only once they become satellites within a more massive halo, by losing their halo gas reservoir to the host halo, or through other environment-related processes such as galaxy harassment and/or ram-pressure stripping. ",
        "introduction": "\\label{intro}  The star-formation histories, masses and structural properties of galaxies, are strongly dependent on their environment: massive, passively-evolving spheroids dominate cluster cores, whereas in field regions galaxies are typically low-mass, star-forming and disk-dominated \\citep[e.g.][hereafter K04]{blanton05,k04}. The Sloan Digital Sky Survey (SDSS) has allowed these environmental dependences to be studied in detail \\citep[e.g.][]{gomez,tanaka04}, showing that for massive galaxies at least \\mbox{($\\la\\!{\\rm M}^{*}+1$ mag)}, star-formation is most closely dependent on local density, independent of the richness of the nearest cluster or group, and is still systematically suppressed for regions as far as 3--4 virial radii  \\mbox{(${\\rm R}_{vir}$)} from the nearest cluster. Hence the evolution of massive galaxies is not primarily driven by mechanisms related to the cluster environment, and instead they are most likely to become passive {\\em before} encountering the cluster environment through galaxy mergers, which are most frequent in galaxy groups or cluster infall regions.   The environmental trends of fainter galaxies \\mbox{(M$_{r}\\!\\ga\\!{\\rm M}^{*}+1$} mag) have generally been examined using galaxy colors as a measure of their star-formation history. Whereas the color of massive galaxies becomes steadily redder with increasing density, a sharp break in the mean color of faint galaxies is observed at a critical density corresponding to $\\sim\\!{\\rm R}_{vir}$ \\citep{gray,tanaka}. In addition, the relative fraction of red and blue galaxies and the shape of the luminosity function of red galaxies over \\mbox{${\\rm M}^{*}\\!+1\\!\\la\\!{\\rm M}_{r}\\!\\la\\!{\\rm M}^{*}\\!+6$ mag} change dramatically with density {\\em inside} the virial radius \\citep{mercurio}. These results imply that the evolution of dwarf galaxies is primarily driven by mechanisms directly related to the structure to which the galaxy is bound, such as suffocation (whereby galaxies lose their halo gas reservoir to the host halo), galaxy harassment and/or ram-pressure stripping.   These differences in the environmental trends and evolution of high- and low-mass galaxies are most likely related to the observed strong bimodality in the properties of galaxies around a stellar mass of \\mbox{$\\sim\\!3\\times10^{10}{\\rm M}_{\\sun}$} \\mbox{($\\sim\\!{\\rm M}^{*}+1$\\,mag)}, with more massive galaxies predominately passive red spheroids, and less massive galaxies tending to be blue star-forming disks \\citep[][hereafter K03a,K03b]{k03a,k03b}. This implies fundamental differences in the formation and evolution of giant and dwarf galaxies, and it has been proposed \\citep[e.g.][]{dekel} that this is related to the way that gas from the halo cools and flows onto the galaxy, which affects its ability to maintain star-formation over many Gyr.  We examine the origins of this bimodality by comparing the environmental dependences of star-formation in giant and dwarf galaxies. The SDSS Data Release 4 \\citep[DR4;][]{sdssdr4} is used to construct a sample of galaxies over a \\mbox{$26\\times26\\,{\\rm Mpc}^{2}$} region containing the supercluster centered on the rich cluster A2199 at \\mbox{$z=0.0309$} \\citep{rines} that is \\mbox{$\\ga\\!9$0\\%} complete to \\mbox{M$^*+3.3$}, and for which we estimate the mean stellar ages and current star-formation rates (SFRs) of galaxies through their spectral indices. We adopt a cosmology with \\mbox{$\\Omega_{M}=0.27$}, \\mbox{$\\Omega_{\\Lambda}=0.73$} and \\mbox{H$_{0}=73$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$} \\citep{spergel}. ",
        "conclusions": "\\label{discussion} Using two independent indicators of the past and present star-formation in galaxies, we present a clear demonstration of the quite different relationships between star-formation and environment for giant \\mbox{($\\sim\\!L^*$)} and dwarf \\mbox{($\\la\\!0.1L^*$)} galaxies. As found in previous studies, giant galaxies \\mbox{(M$_{r}\\!<\\!-20$ mag)} show a gradual transition from \\mbox{$\\ga\\!8$0\\%} being classed as old \\mbox{($t\\!>\\!7$\\,Gyr)} or passive \\mbox{(EW$[{\\rm H}\\alpha]\\le4$\\,\\AA)} in the cluster cores, to field regions beyond \\mbox{3--4\\,$R_{vir}$} where still \\mbox{$\\sim\\!4$0\\%} are classed as passive. These results can be understood in the context of massive galaxies becoming passive through galaxy-galaxy interactions: the finding of both passive and star-forming galaxies in all environments reflects its stochastic nature and its independence from large-scale structure. The gradual overall trend reflects the increasing probability with density for a galaxy to have undergone a major merger in its lifetime.   In contrast the dwarf galaxy population (\\mbox{$-19\\!<\\!{\\rm M}_{r}\\!<\\!-17.8$} mag) shows a sharp transition at \\mbox{$\\sim\\!R_{\\rm vir}$}, from being predominantly old and passive within the cluster, to outside where virtually all galaxies are young and forming stars, and passive galaxies are {\\em only} found as satellites to more massive galaxies. These findings are supported by the observed dependencies of galaxy clustering on luminosity and color, in which the galaxies associated with the most overdense regions on $\\la\\!1$\\,Mpc scales are found to be both the bright \\mbox{(${\\rm M}_{r}\\!<\\!-22$ mag)} and faint \\mbox{(${\\rm M}_{r}\\!>\\!-19$ mag)} red galaxies \\citep{hogg,zehavi}. These dependencies were successfully reproduced by the smoothed particle hydrodynamics cosmological simulations of \\citet{berlind}, who identify the passive low-mass galaxies as satellites within more massive halos.  The observation that no passive dwarf galaxies are found in isolated (with respect to \\mbox{$\\ga\\!L^{*}$} galaxies) low-density regions implies that galaxy merging cannot be effective in completely terminating star-formation in low-mass galaxies, particularly as in this study these regions represent the infall regions of the supercluster where low-velocity encounters should be most frequent. Low-mass galaxies become passive only when they become satellites within a more massive DM halo and their halo gas reservoir is lost to that of the host, ``suffocating'' the galaxy \\citep{larson,bekki}, or through directly cluster-related mechanisms such as galaxy harassment and/or ram-pressure stripping.   These differences can be understood in the context of the hot and cold gas infall \\citep{dekel,keres} or AGN feedback \\citep[e.g.][]{croton,hopkins} models of galaxy evolution. When gas-rich galaxies merge, tidal forces trigger a starburst and fuel the rapid growth of the central black hole, until outflows from the AGN drive out the remaining cold gas from the galaxy, rapidly terminating the starburst. In massive galaxies, the gas in the halo is also heated by stable virial shocks, and is prevented from cooling by feedback from quiescent accretion of the hot gas onto the black hole, effectively shutting down star-formation \\citep{croton}. Because black hole growth is strongly dependent on galaxy mass, AGN feedback in low-mass galaxies is much less efficient at expelling cold gas or affecting star-formation \\citep{springel}. In addition, low-mass galaxies self-regulate their star-formation through supernova feedback, preventing starbursts that exhaust the cold gas, which in turn is constantly replenished by cold streams, so that their ability to maintain star-formation over many Gyr is much less affected by their merger history."
    },
    "0606/astro-ph0606698_arXiv.txt": {
        "abstract": "The luminosity distance - redshift relation is analytically given for generalized Randall-Sundrum type II brane-world models containing Weyl fluid either as dark radiation or as a radiation field from the brane. The derived expressions contain both elementary functions and elliptic integrals of the first and second kind. First we derive the relation for models with the Randall-Sundrum fine-tuning. Then we generalize the method for models with cosmological constant. The most interesting models contain small amounts of Weyl fluid, expected to be in good accordance with supernova data. The derived analytical results are suitable for testing brane-world models with Weyl fluid when future supernova data at higher redshifts will be available. ",
        "introduction": "At present the Universe is considered a general relativistic Friedmann space-time with flat spatial sections, containing more than $70\\%$ dark energy and at about $25\\%$ of dark matter. Dark energy could be simply a cosmological constant $\\Lambda$, or quintessence or something entirely different. There is no widely accepted explanations for the nature of any of the dark matter or dark energy (even the existence of the cosmological constant remains unexplained).  An alternative to introducing dark matter would be to modify the law of gravitation, like in MOND \\cite{MOND} and its relativistic generalization  \\cite{MONDrel}. These theories are compatible with the Large scale structure of the Universe \\cite{MOND2}. However in spite of the successes, certain problems were signaled on smaller scales \\cite{MOND3}.  Quite remarkably, supernova data, which in the traditional interpretation yield to the existence of dark energy, can be explained by certain f(R) \\cite{fR} or inverse curvature gravity models \\cite {modgrav}. However the parameter range, in which the latter is in goood agrement with the supernova data, also presents stability problems \\cite{excmodgrav}.  Modifications of the gravitational interaction could also occur by enriching the space-time with extra dimensions. Originally pioneered by Kaluza and Klein, such theories contained compact extra dimensions. The so-called brane-world models, motivated by string / M-theory, containing our observable 4-dimensional universe (the brane) as a hypersurface, were introduced in \\cite{ADD}, \\cite{RS1} and \\cite{RS2}, the latter model allowing for a non-compact extra dimension.  The curved generalizations of the model presented in \\cite{RS2} have evolved into a 5-dimensional alternative to general relativity, in which gravity has more degrees of freedom. In contrast with standard model fields, these evolve in the whole 5-dimensional bulk. In this generalized Randall-Sundrum type II (RS) theory, the brane has a tension $\\lambda $ and gravitational dynamics is governed by the 5-dimensional Einstein equation. Its projections to our observable 4-dimensional universe (the brane) are the twice contracted Gauss equation, the Codazzi equation and an effective Einstein equation, the latter being obtained by employing the junction conditions across the brane \\cite{Decomp}. The effective Einstein equation (for the case of symmetric embedding and no other contribution to the bulk-energy-momentum than a bulk cosmological constant) was first given in a covariant form in \\cite{SMS}. Supplementing this by the pull-back to the brane of the bulk energy momentum tensor $\\widetilde{\\Pi }_{ab}$, which is \\begin{equation} \\mathcal{P}_{ab}=\\frac{2\\widetilde{\\kappa }^{2}}{3}\\left( g_{a}^{c}g_{b}^{d} \\widetilde{\\Pi }_{cd}\\right) ^{TF} \\end{equation} (with $\\widetilde{\\kappa }^{2}$ the bulk coupling constant and $g_{ab}$ the induced metric on the brane) the effective Einstein equation reads \\cite{Decomp}: \\begin{equation} G_{ab}=-\\Lambda g_{ab}+\\kappa ^{2}T_{ab}+\\widetilde{\\kappa }^{4}S_{ab}- \\mathcal{E}_{ab}+\\mathcal{P}_{ab}\\ .  \\label{modE} \\end{equation} Here $\\kappa ^{2}$ is the brane coupling constant, related to the bulk coupling constant and the brane tension $\\lambda $ as $6\\kappa ^{2}=\\widetilde{\\kappa }^{4}\\lambda $, and \\begin{equation} \\Lambda =\\frac{\\kappa^{2}}{2} \\lambda -\\frac{\\widetilde{\\kappa }^{2}}{2} n^{c}n^{d} \\widetilde{\\Pi }_{cd} \\end{equation} represents a cosmological \"constant\" which possibly varies due to the normal projection of the bulk energy-momentum tensor (this includes the contribution $-\\widetilde{\\Lambda }g_{ab}$ due to the bulk cosmological constant $\\widetilde{\\Lambda }$). The source term $S_{ab}$ is quadratic in the brane energy-momentum tensor $T_{ab}$: \\begin{equation} S_{ab}=\\frac{1}{4}\\Biggl[-T_{ac}^{\\ }T_{b}^{c}+\\frac{1}{3}TT_{ab}-\\frac{ g_{ab}}{2}\\left( -T_{cd}^{\\ }T^{cd}+\\frac{1}{3}T^{2}\\right) \\Biggr]\\ , \\label{S} \\end{equation} and $\\mathcal{E}_{ab}$ is the electric part of the bulk Weyl tensor $ \\widetilde{C}_{abcd}$, given as  \\begin{equation} \\mathcal{E}_{ac}=\\widetilde{C}_{abcd}n^{b}n^{d}\\ .  \\label{calE} \\end{equation} In a cosmological context and suppressing any energy exchange between the brane and the bulk, this latter term generates the so-called dark radiation. Otherwise it can be called a Weyl fluid.  A review of many aspects related to the theories described by the effective Einstein equation (\\ref{modE}) can be found in \\cite{MaartensLR}. Both early cosmology \\cite{BDEL} and gravitational collapse \\cite{BGM}-\\cite{Pal} are essentially modified in these theories. There is also possible to replace dark matter with geometric effects in the interpretation of galactic rotation curves, weak lensing and galaxy cluster dynamics \\cite{Harko}.  The possible modifications of gravitational dynamics are even more versatile in the so-called induced gravity models. These can be regarded as brane-world models enhanced with the first quantum-correction arising from the interaction of the brane matter with bulk gravity. The induced gravity correction couples to the 5-dimensional Einstein-Hilbert action with the coupling constant $\\gamma \\widetilde{\\kappa }^{2}/\\kappa ^{2}$. The simplest of such models, the DGP model was introduced in \\cite{DGP}. This model however suffers from linear instabilities (ghost modes in the perturbations), as shown for de Sitter branes \\cite{ghost}. The ghost modes withstand even the introduction of a second brane \\cite{ghost2}. Generalizations of the DGP model are discussed covariantly in \\cite{SS} and  \\cite{MMT} when the embedding is symmetric, and in \\cite{Induced} when it is asymmetric. In these models the role of the effective Einstein equation (\\ref{modE}) is taken by a more complicated equation (see for example Eq. (29) of  \\cite{Induced}), which contains the square of the Einstein tensor $G_{ab}$. This implies that in certain sense the degree of nonlinearity of the theory is squared. In a cosmological setup the square root of this equation can be taken, leading to a set of modified Friedmann and Raychaudhuri equations, which however contain a sign ambiguity $\\varepsilon =\\pm 1$ due to the involved square root. These are called the BRANE1 [DGP(-)] branch for $\\varepsilon =-1$ and BRANE2 [DGP(+)] for $\\varepsilon =1$ in the terminology of \\cite{SS} [or \\cite{LMM}, respectively]. Both the original Randall-Sundrum type II model and the DGP model are contained as special subcases. Notably, the BRANE2 branch contains cosmological models which self-accelerate at late-times. We give in Fig \\ref{Fig1} a diagram containing a classification of these theories and how they emerge as different limits from each other.  \\begin{figure}[tbp] \\includegraphics[height=8cm]{Fig1.eps} \\caption{(Color online) A diagram presenting various brane-world models and their inter-relations. LWRS is the generalized Randall-Sundrum model with cosmological constant and a Weyl fluid reflecting a brane radiating into the bulk during nowadays or at least until recent cosmological times.} \\label{Fig1} \\end{figure}  In this paper we discuss analytically the luminosity distance - redshift relation in various generalized Randall-Sundrum type II brane-world models described by Eq. (\\ref{modE}). Our analytical approach can enhance the confrontation of these models with current and most notably, with future supernova observations. We note that recently analytical results have been given in Ref. \\cite{DabrowskiS} for a wide class of phantom Friedmann cosmologies too, in terms of elementary and Weierstrass elliptic functions.  In section 2 we review the notion of luminosity distance, its relation with the redshift and how these can be measured independently. This section was included mainly for didactical purposes.  In section 3 we review the modification of this relation in the Randall-Sundrum type II brane-world scenario. These include the introduction of the parameters $\\Omega _{\\lambda }$ and $\\Omega _{d}$ which can be traced back to the source terms $S_{ab}$ and $\\mathcal{E}_{ab}$ of the modified Einstein equation (\\ref{modE}). The other cosmological parameters are $\\Omega _{\\rho }$, representing (baryonic and dark) matter and $\\Omega _{\\Lambda }$. We do not include bulk sources in the analysis, with the notable exception of a bulk cosmological constant.  Section 4 contains the derivation of the analytic expression for the luminosity distance - redshift relation for the brane-worlds which are closest to the original Randall-Sundrum scenario \\cite{RS2}, thus with no cosmological constant (Randall-Sundrum fine-tuning). The generic expression (\\ref{solution1}) of the luminosity distance derived here is given in terms of elementary functions and elliptic integrals of the first and second kind. From this most generic case we take the subsequent limits: $\\Omega _{d}=0$ (subsection 4.2), $\\Omega _{\\lambda }=0$ (subsection 4.3); and both $\\Omega _{d}=\\Omega _{\\lambda }=0$, this being the general relativistic Einstein-de Sitter case (subsection 4.4).  Such models however could not allow for late-time acceleration, therefore in section 5 we discuss the luminosity distance - redshift relation for brane-worlds with $\\Lambda $. First we present in subsection 5.1 a class of models, for which the luminosity distance can be given in terms of elementary functions alone. These models are characterized by an extremely low value of the brane tension, thus are in conflict with various constraints on brane-world models.  Next, in subsection 5.2 we discuss brane-worlds for which the brane-characteristic contributions $\\Omega _{\\lambda }$ and $\\Omega _{d}$ represent small perturbations. This is a good assumption as observational evidences suggest that general relativity is a sufficiently accurate theory of the universe, and as such the deviations from it could not be very high, at least at late-times. We give analytical expressions in terms of both elementary functions and elliptic integrals of the first and second kind for the luminosity distance, to first order accuracy in the chosen small parameters of the model. Some of the most lengthy computations needed in order to achieve the result are presented in the Appendix.  Section 6 contains the concluding remarks.  Throughout the paper $c=1$ was employed. ",
        "conclusions": "The main purpose of this paper was to present the analytical formulation of the luminosity distance - redshift relation in the generalized Randall-Sundrum type II brane-world models containing a Weyl fluid either in the form of dark radiation or as radiation leaving the brane and feeding the bulk black holes. We have given the luminosity distance in terms of elementary functions and elliptical integrals of first and second type and we have also shown how the different limits arise from the generic result. Our results hold for:  (a) Models with Randall-Sundrum fine-tuning ($\\Lambda =0$), with or without dark radiation from the bulk and with or without considerable contribution from the energy-momentum squared source terms, discussed in section 4.  (b) The models discussed in subsection 5.1, obeying $\\Lambda =\\kappa ^{2}\\lambda /2$, integrable in terms of elementary functions and  (c) Models with a brane cosmological constant, discussed to first order accuracy in both the Weyl fluid and energy-momentum squared sources.  This last class of models, presented in subsection 5.2 in the latest times of the cosmological evolution are only slightly different from the $\\Lambda $ CDM model, as they have $\\Omega _{d}\\ll 1$ and $\\Omega _{\\lambda }\\ll 1.$ The derived modifications in the luminosity distance - redshift formula then represent corrections to the corresponding formula of the $\\Lambda $CDM model.  While the focus of the present paper is the integrability of the luminosity distance - redshift relation in various brane-world models, in a forthcoming paper \\cite{BraneLuminosityDistance2} we will discuss how well the presently available supernova data support the brane-world models with a small amount of Weyl fluid.  \\ack  This work was supported by OTKA grants no. T046939, 69036 and T042509. L \\'{A}G and GyMSz were further supported by the J\\'{a}nos Bolyai Grant of the Hungarian Academy of Sciences and GyMSz by the Magyary Zolt\\'{a}n Higher Educational Public Foundation.  \\appendix"
    },
    "0606/astro-ph0606651_arXiv.txt": {
        "abstract": "Neutrino emission caused by singlet Cooper pairing of baryons in neutron stars is recalculated by accurately taking into account for conservation of the vector weak currents. The neutrino emissivity via the vector weak currents is found to be several orders of magnitude smaller than that obtained before by different authors. This makes unimportant the neutrino radiation from singlet pairing of protons or hyperons. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606184_arXiv.txt": {
        "abstract": "We present a systematic investigation of X-ray thermal coronae in 157 early-type galaxies and 22 late-type galaxies from a survey of 25 hot ($kT >$ 3 keV), nearby ($z <$ 0.05) clusters, based on \\chandra\\ archival data. Cool galactic coronae ($kT$ = 0.5 - 1.1 keV generally) have been found to be very common, $>$ 60\\% in NIR selected galaxies that are more luminous than 2 \\Ls, and $>$ 40\\% in \\Ls\\ $< L_{\\rm Ks} <$ 2 \\Ls\\ galaxies. These embedded coronae in hot clusters are generally smaller (1.5-4 kpc radii), less luminous ($\\lsim 10^{41}$ erg s$^{-1}$), and less massive (10$^{6.5}$-10$^{8}$ M$_{\\odot}$) than coronae in poor environments, demonstrating the negative effects of hot cluster environments on galactic coronae. Nevertheless, these coronae still manage to survive ICM stripping, evaporation, rapid cooling, and powerful AGN outflows, making them a rich source of information about gas stripping, microscopic transport, and feedback processes in the cluster environment. Heat conduction across the boundary of the coronae has to be suppressed by a factor of $\\gsim$ 100, which implies the X-ray gas in early-type galaxies is magnetized and the magnetic field plays an important role in energy transfer. Stripping through transport processes (viscosity or turbulence) also needs to be suppressed by at least a factor of ten at the coronal boundary. The stellar mass loss inside the corona is key to maintaining the gas balance in coronae. The luminous, embedded coronae, with high central density (0.1 - 0.4 cm$^{-3}$), are mini-versions of group and cluster cooling cores. As the prevalence of coronae of massive galaxies implies a long lifetime ($\\gsim$ several Gyr), there must be a heat source inside coronae to offset cooling. While we argue that AGN heating may not generally be the heat source, we conclude that SN heating can be enough as long as the kinetic energy of SNe can be efficiently dissipated. We have also observed a connection between radiative cooling and the SMBH activity of their host galaxies as many coronae are associated with powerful radio galaxies. Cooling of the coronal gas may provide fuel for the central SMBH and nuclear star formation in environments where the amount of galactic cold gas is otherwise at a minimum. Diffuse thermal coronae have also been detected in at least 8 of 22 late-type (Sb or later) galaxies in our sample. Evidence for enhanced star formation triggered by the ICM pressure has been found in four late-type galaxies. The fraction of luminous X-ray AGN ($> 10^{41}$ ergs s$^{-1}$) is not small ($\\sim$ 5\\%) in our sample. ",
        "introduction": "One of the most important discoveries by the \\einstein\\ observatory was the ubiquity of X-ray coronae of early-type galaxies in the field and poor environments (Forman, Jones \\& Tucker 1985). This discovery changed our view of early-type galaxies: they are gas-rich instead of gas-poor, but the gas is in the X-ray phase with a temperature of $\\sim 10^{7}$ K. The consensus is that the X-ray gas of early-type galaxies originates as the stellar mass lost from evolved stars and planetary nebulae, accumulating in the galaxy without much escaping via galactic winds (e.g., Mathews 1990). Their origin is different from that of the cluster ICM, which should be mostly primordial. The gas ejected from evolved stars collides and passes through shocks. The gas temperature has been raised to the stellar kinetic temperature determined by the stellar velocity dispersion, and may further be raised by heating from supernovae (SNe). However, it has been known that the kinetic energy of SNe is not efficiently dissipated into the hot ISM gas, especially for less massive galaxies with shallow potentials, so galactic winds can form in less massive galaxies and make them X-ray faint (Mathews \\& Baker 1971; David, Forman \\& Jones 1991; Brown \\& Bregman 1998). Galactic cooling cores have generally formed in massive galaxies with deep potentials. Cooling of the X-ray coronal gas indeed happens in at least some galaxies (e.g., the O VI detections by {\\em FUSE}, Bregman et al. 2005).  It has been long known that the X-ray luminosities ($L_{\\rm X}$) of coronae correlates with the optical $B$ band luminosities ($L_{\\rm B}$) of their host galaxies, but the dispersion is large (Forman et al. 1985; Canizares, Fabbiano \\& Trinchieri 1987; Brown \\& Bregman 1998; O'Sullivan et al. 2001). The large dispersion has been a puzzle for almost two decades. The narrow color-magnitude relation and the narrow fundamental plane of early-type galaxies imply that local early-type galaxies are a homogeneous group. Nevertheless, early-type galaxies are not a one-parameter class, and the X-ray luminosity of coronae is also affected by properties besides the optical luminosity of the galaxy. Various factors have been proposed to explain the dispersion of the $L_{\\rm X}$ - $L_{\\rm B}$ relation for the X-ray coronae of early-type galaxies, including SN rate, galactic rotation, metallicity, dark matter halo and environments (see a review by Mathews \\& Brighenti 2003 and the references therein). However, no extra parameter has been found to significantly tighten the correlation, and no single scenario seems to be able to explain the large dispersion. It is likely that both internal effects (e.g., SN rate \\& dark matter halo) and external effects (environment) contribute to the dispersion.  Limited by its angular resolution (30$''$), the X-ray coronae detected by \\rosat\\ are almost all in the field, or poor environments, or cool clusters (e.g., Virgo). There were some attempts to study X-ray coronae of galaxies in hot clusters with the \\einstein\\ and \\rosat\\ data (e.g., Canizares \\& Blizzard 1991 for Coma; Dow \\& White 1995 for Coma; Sakelliou \\& Merrifield 1998 for A2634), but the results were generally non-detections or inconclusive. Furthermore, the \\rosat\\ results on galaxy coronae suffer greatly from the contamination by AGN, X-ray binaries and ICM emission. Prior to the launch of \\chandra, no 10$^{7}$ K galaxy coronae (not including cluster cooling cores) were firmly detected in hot ($T \\sim$ 10$^{8}$ K) clusters, nor were they expected, since evaporation and ram-pressure stripping by the hot, dense ICM should very efficiently remove gas from galaxies. The first direct evidence for the survival of galaxy coronae in hot clusters came from the \\chandra\\ observations of the Coma cluster. Vikhlinin et al. (2001, V01 hereafter) found small, but extended ($\\sim$ 2 kpc in radius), X-ray coronae ($kT \\sim$ 1-2 keV) in the cores of the two dominant Coma galaxies, NGC 4874 and NGC 4889, where the surrounding ICM has a temperature of 8 - 9 keV. More and more similar embedded coronae have been found since then. At present, eight more coronae, small but spatially resolved, were unambiguously revealed by \\chandra\\ in hot clusters ($>$ 3 keV) and investigated in detail: two lie in the 3.2 keV cluster A1060 (Yamasaki et al. 2002), four in a 5-6 keV region of A1367 (Sun et al. 2005, hereafter S05), one in a 6-7 keV region of Perseus (Sun, Jerius \\& Jones 2005, hereafter SJJ05), and one associated with the cD galaxy of the 4 keV cluster A2670 (Fujita, Sarazin \\& Sivakoff 2006). Thermal emission from cluster galaxies was also detected in several other cases (e.g., Smith 2003; Hardcastle, Sakelliou \\& Worrall 2005), although detailed analysis was not done because the X-ray sources are either unresolved or faint. Among all these detections, the corona associated with NGC~1265, the prototype of narrow-angle tail (NAT) radio galaxies in Perseus, is the most interesting one, as its corona survives gas stripping despite motion with a Mach number of $\\sim$ 3, ICM evaporation ($T_{\\rm ICM}/T_{\\rm ISM} \\sim$ 10), fast cooling ($t_{\\rm cooling} \\sim$ 30 Myr at the center) and powerful AGN outbursts. These galaxy coronae ($kT \\sim$ 1 keV) in rich environments (high ambient pressure) are pressure confined and small (1.5 - 4.5 kpc, or 3 - 10$''$ in radius at the distance of the Coma cluster) and require the superior \\chandra\\ angular resolution to resolve them. These small coronae should form before their host galaxies began to evolve in hot clusters. If they were once destroyed in hot clusters, the hot ICM should have filled the interstellar space. It is then difficult for galactic cooling cores to re-form against strong stripping and evaporation.  These embedded galaxy coronae are perfect targets to study the gas physics and microscopic transport processes involved in the ICM-corona interaction. Gas stripping and evaporation by the surrounding hot dense ICM should be very efficient to destroy coronae. However, V01 pointed out that the survival of two cool coronae in Coma requires that heat conduction must be suppressed by a factor of 30 - 100 at the ISM-ICM boundary. S05 and SJJ05 derived similar conclusions and the results also require viscosity to be suppressed. The galactic magnetic field should be responsible for the suppression, and so a better understanding of the galactic magnetic field evolution in rich environments is required. We shall emphasize that these embedded coronae in hot clusters are generally smaller than the mean free path of particles in the ICM without magnetic field, which is:  \\begin{eqnarray} \\lambda &=& 7.8 (\\frac{kT_{\\rm ICM}}{\\rm 5 keV})^{2} (\\frac{n_{\\rm e, ICM}}{10^{-3} {\\rm cm}^{-3}})^{-1} {\\rm kpc} \\end{eqnarray}  Therefore, it is questionable whether we can apply hydrodynamics to the ICM flow around the embedded coronae. The distribution function may not be Maxwellian so ideally kinetic theory or at least rarefied gas dynamics should be applied. This situation is different from coronae in cool groups (e.g., NGC~1404 in Fornax, Machacek et al. 2005), where $\\lambda$ is generally smaller than 1 kpc while coronae are larger (e.g., $\\sim$ 9 kpc for NGC~1404). The inclusion of magnetic field in the ICM may further complicates the problem, by making transport processes anisotropic. The details are certainly complicated, largely depending on the magnetic field configuration and evolution. We notice that simulations of the evolution of X-ray coronae in clusters and groups all use hydrodynamics (e.g., Toniazzo \\& Schindler 2001), while the detailed transport processes involving in the corona-ICM interaction were not incorporated.  S05 and SJJ05 also showed that many host galaxies of coronae are active in radio and sometimes are luminous (e.g., NGC~1265). The radio emission generally ``turns on'' after traversing the dense coronae (i.e. morphologically anti-correlated, e.g., NGC~3842, NGC~4874 and NGC~1265). In NGC~1265, since the jet power is so strong ($>$ 1000 times the thermal energy of the corona), the narrow jets must carry nearly all their energy away from the central SMBH and release the energy outside the corona to preserve the small corona. Thus, this fact implies that AGN heating may not always be able to heat the central few kiloparsecs of a cooling core significantly. For NGC~1265, we also found that the Bondi accretion luminosity of the detected coronal gas is similar to the jet power. As the inner $\\sim$ 2 kpc dense core of the corona can survive both high-speed stripping and powerful AGN feedback, the cooling of coronae can fuel the central SMBH in rich environments where the amount of cold galactic gas is at a minimum.  These embedded coronae are also ideal targets to examine the effects of hot cluster environments on the properties of coronae (e.g., the $L_{\\rm X} - L_{\\rm B}$ relation, the size and X-ray gas mass of coronae), especially through a comparison with a sample of coronae in the field or poor environments. These cool coronae may also be seeds for future cluster cooling cores. Motl et al. (2004) presented a scenario for the formation of cluster cooling cores via hierarchical mergers. Although the cool cores in their simulations are larger than what we observed, it is interesting to explore the evolution of these coronae and their frequency in clusters, and to further compare with simulations. In principle, the properties of the coronae also shed light on the evolution of their early-type host galaxies, as the evolution processes (e.g., mergers) can easily impact the small coronae.  Compared to the X-ray coronae of early-type galaxies, we know even less about the X-ray coronae of late-type galaxies in hot clusters. One of the first detailed studies with \\chandra\\ or \\xmm\\ is UGC~6697 in A1367 (Sun \\& Vikhlinin 2005). We suggested that the starburst in UGC~6697 was triggered by the interaction with the ICM. During the course of this project, we have found a spectacular X-ray tail behind a small starburst galaxy ESO~137-001 (SFR $\\sim$ 10 M$_{\\odot}$ / yr) near the center of a massive cluster A3627 (Sun et al. 2006). The X-ray tail is long (70 kpc) with a length-to-width ratio of $\\sim$ 10. We interpret this tail as the stripped ISM of ESO 137-001 mixed with the hot cluster medium, with this blue galaxy being converted into a gas-poor galaxy.  These remarkable examples of coronae in early-type and late-type galaxies are important to understand the galactic ISM in hot clusters. However, the results based on just a few detections can be biased. We do not know how common the embedded coronae are. We know nothing about the properties of the whole population (e.g., luminosity, temperature, size and gas mass). What are the environmental effects on the properties of embedded coronae? What processes mediate the energy balance and transfer inside coronae? What is the general connection between coronae and the AGN activity of their host galaxies. Only a systematic analysis based on a well-selected sample can answer these questions and characterize the general properties of the whole population. Only with the knowledge of the whole population, can we better understand the results from detailed investigations of individual coronae. In this paper, we present the first systematic study of X-ray coronae of cluster galaxies, in 25 nearby ($z <$ 0.05), hot ($kT >$ 3 keV) clusters using archival \\chandra\\ data. The cluster and galaxy samples are defined in $\\S$2. The data analysis is summarized in $\\S$3. The properties of the corona population are present in $\\S$4 (for early-type galaxies) and $\\S$5 (for late-type galaxies). We summarize the fraction of luminous X-ray AGN in our sample in $\\S$6. $\\S$7 is the discussion, while $\\S$8 is the conclusions. Notes of interesting coronae and galaxies in the sample are present in the appendix. ",
        "conclusions": "We have systematically searched for and investigated thermal coronae of 179 galaxies (both early and late types) in 25 nearby ($z <$ 0.05), hot ($kT >$ 3 keV) clusters, based on 68 \\chandra\\ archival pointings with a total exposure of 2.77 Msec and a sky coverage of 3.3 deg$^{2}$. The galaxy sample is complete for all NIR luminous ($>$ 0.74 \\Ls\\ in the $K_{\\rm s}$ band) or radio luminous ($L_{\\rm 1.4 GHz} > 10^{22.8}$ W Hz$^{-1}$) galaxies in the \\chandra\\ field. This work represents the first systematic study of X-ray thermal emission of galaxies in rich environments. The main observational results and conclusions of our study are:  1) We find a new population of embedded X-ray coronae of early-type galaxies in hot clusters. Despite the effects of ICM stripping, evaporation, rapid cooling, and powerful SMBH bursts, X-ray coronae of massive early-type galaxies (excluding cDs in cluster cooling cores) are very common ($>$ 60\\% of $>$ 2 \\Ls\\ galaxies in the $K_{\\rm s}$ band) in hot clusters, although their properties have been significantly modified by the dense ICM. Significant number of coronae have also been found in less massive galaxies ($>$ 40\\% of \\Ls\\ $<$ $L_{\\rm Ks} <$ 2 \\Ls\\ galaxies; $>$ 15\\% of $<$ \\Ls\\ galaxies), although the \\chandra\\ data of most clusters are not deep enough to unambiguously identify faint coronae ($L_{\\rm 0.5 - 2 keV} \\lsim 10^{40}$ ergs s$^{-1}$). These embedded coronae are smaller (generally 1.5 - 4 kpc in radius) and contain less gas (10$^{6.5} - 10^{8}$ M$_{\\odot}$) than their counterparts in poor environments. The negative effect on the coronal luminosity is also suggested by the data. These embedded coronae may correspond to the dense cores of coronae in poor environments. Therefore, our work has demonstrated the negative environmental effects of rich environments on galaxy coronae. The ubiquity of embedded coronae in massive galaxies implies a lifetime comparable to that of clusters (or at least several Gyr).  2) The temperatures of embedded coronae range from 0.3 to 1.7 keV. The gas temperature is generally higher than the stellar kinetic temperature of the galaxy ($\\beta_{\\rm spec}$ = 0.2 - 1.1, Fig. 8). Internal temperature profiles of embedded coronae generally show a decrease of gas temperature towards the center. The abundance of the coronal gas, constrained from the joint data of 20 coronae with best statistics, is $\\sim$ 0.8 solar, which implies a stellar origin of the coronal gas.  3) For the cool coronae of early-type galaxies to survive in the hot ICM, heat conduction across the boundary of the coronae has to be suppressed by a factor of $\\gsim$ 100. We argue that this fact implies the X-ray gas in early-type galaxies is magnetized. Magnetic field plays an important role in energy transfer. However, it is unclear how the galactic magnetic field can remain disconnected from the field in the ICM for the lifetime of embedded coronae.  4) The embedded coronae of early-type galaxies are subject to ICM stripping. The internal thermal pressure of coronae is generally high enough to overcome the ram pressure. However, the stripping through transport processes (viscosity or turbulence) at the coronal boundary is too strong and has to be suppressed by at least a factor of ten. The stellar mass loss in the boundary layer can balance the mass lost from the suppressed stripping and the balance is stable.  5) The embedded coronae of early-type galaxies have high gas densities ($\\sim$ 0.3 cm$^{-3}$ at the center of luminous coronae) so they are the mini version of big cluster cooling cores, with a boundary. As the prevalence of the coronae of massive galaxies implies a long lifetime ($\\gsim$ several Gyr), there must be a heat source inside these embedded mini cooling cores. While we argue that both AGN and the stellar ejecta cannot be the major heat source, SN heating inside coronae is a good candidate if its kinetic energy can be efficiently (20\\% - 100\\%) couple into the coronal gas.  6) We have observed a connection between these mini-cooling-cores and the radio activity of their host galaxies. Radiative cooling of the coronal gas may provide fuel for the central SMBH in environments where the amount of galactic cold gas is at a minimum. We have also found a general morphological anti-correlation of the radio jet emission and the X-ray coronae. The radio jets generally ``turn on'' after traversing the dense coronae. The embedded dense coronae may also provide cold gas for nuclear star formation, while the star formation in ``naked'' galaxies exposed to the hot ICM may have been truncated.  7) We have observed a variety of X-ray components associated with or centered on the cD galaxies in our sample, ranging from no cool component, to small cool coronae, to large X-ray cool cores, to big cluster cores. Most clusters in our sample either have a small cool corona associated with the cD galaxy, or a big cluster cooling core centered on the cD galaxy.  8) We also detected thermal coronae of at least 8 late-type galaxies (Sb or later) from 22 galaxies in our sample. Late-type galaxies with luminous X-ray thermal emission usually have substantial star formation activity. In four galaxies with the most X-ray luminous coronae in the sample, evidence for enhanced star formation triggered by the ICM pressure has been found.  9) Nine luminous X-ray AGN ($L_{0.3 - 8 keV} > 10^{41}$ ergs s$^{-1}$) are found from 163 galaxies (both early-types and late-types) brighter than 0.74 \\Ls\\ (in the $K_{s}$ band), which indicates a not small fraction ($\\sim$ 5\\%) of X-ray luminous AGN in local cluster galaxies. Fainter nuclear hard sources are also found in at least 20\\% of galaxies in our sample."
    },
    "0606/astro-ph0606467_arXiv.txt": {
        "abstract": "We derive the constraints on the mass ratio for a binary system to merge in a violent process. We find that the secondary to primary stellar mass ratio should be $0.003 \\la (M_2/M_1) \\la 0.15$. A more massive secondary star will keep the primary stellar envelope in synchronized rotation with the orbital motion until merger occurs. This implies a very small relative velocity between the secondary star and the primary stellar envelope at the moment of merger, and therefore very weak shock waves, and low flash luminosity. A too low mass secondary will release small amount of energy, and will expel small amount of mass, which is unable to form an inflated envelope. It can however produce a quite luminous but short flash when colliding with a low mass main sequence star.  Violent and luminous mergers, which we term {\\it mergebursts,} can be observed as V838 Monocerotis type events, where a star undergoes a fast brightening lasting days to months, with a peak luminosity of up to $\\sim 10^6 L_\\odot$ followed by a slow decline at very low effective temperatures. ",
        "introduction": "Stellar mergers have been recognized for a long time as events which can be important for evolution of binary systems. In discussions of globular clusters stellar mergers are usually considered as the most probable source of blue stragglers, e.g. De Marco et al. (2005, and earlier references therein; Sills et al. 2005). In some cases three or more stars might be involved (Knigge et al. 2006). However, the main interest in these cases has been directed toward understanding the nature of the final product of a merger in terms of its mass, chemical structure and farther evolution. Little attention, if any, has usually been paid to direct observational appearances of these events. This was obviously due to the common belief that the stellar mergers are very rare events and that, consequently, there is very little chance to observe them. However, the discovery of the eruption of V838~Mon in 2002 (Brown 2002) and subsequent studies of its observed evolution (Munari et~al. 2002; Kimeswenger et~al. 2002; Crause et~al. 2003; Kipper et~al. 2004; Tylenda 2005), as well as, of other similar objects, i.e. V4332 Sgr (Martini et~al. 1999; Tylenda et~al., 2005) and M31~RV (Mould et~al. 1990) have led to suggestions that these observed events were likely to be due to stellar mergers (Soker \\& Tylenda 2003, Tylenda \\& Soker 2006, hereafter TS06). Likewise an analysis done by Bally \\& Zinnecker (2005) shows that stellar mergers in cores of young clusters, which might be one of channels for producing very massive stars, can be source of luminous and spectacular observational events.  Different reasons can lead to stellar mergers. In dense stellar systems, as globular clusters or cores of young clusters, direct collisions of two stars or interactions of binaries with other cluster members can quite easily happen, often leading to a merger (for recent papers and more references see, e.g., Lombardi et al. 2003; Fregeau et al. 2004; Mapelli et al. 2004; Dale \\& Davies 2006; for a review of these merger possibilities see Bailyn 1995). In multiple star systems dynamical interactions between the components or encounters between the system and other stars can destabilize stellar trajectories so that two components collide and merge. A binary stellar system can lose angular momentum during its evolution, e.g. due to mass loss, so the separation of the components decreases, which may finally lead to a merger. In the latter case, the merger is probably often relatively gentle and does not lead to spectacular events. This happens when the system reaches and keeps synchronization until the very merger. The relative velocity between the matter elements from different components is then very low, there is no violent shock heating and the orbital energy is released on a very long time scale. However when the binary component mass ratio is low the secondary is unable to maintain the primary in synchronization so the so-called Darwin instability sets in and the merger takes place with a large difference between the orbital velocity of the secondary and the rotational velocity of the primary. In this case the merger is expected to be violent, at least in the initial phase when the large velocity differences are dissipated in shocks. In this paper we analyze this possibility in more detail and discuss observational appearances of violent mergers triggered by the Darwin instability in binaries.  We can schematically distinguish between three basic types of merger events. Of course, there is a continuous variation between the three types, but it is instructive to make these three ideal classes.  (1) The secondary is disrupted during the collision, and contributes most of the mass in the inflated envelope. This is likely to happen in a close to grazing collision with a not too-compact companion in a very eccentric orbit. We suggest this type of merger as explanation for V838 Mon (TS06).  (2) Merger in a binary system which reaches merger in an unsynchronized rotation (the spin of the primary star is not synchronized with the orbital angular frequency), and where the secondary survives the initial merger stages. The secondary then spirals-in inside the primary envelope. The inflated envelope comes mainly from the primary mass. The conditions for the occurrence the type of merger are studied in this paper.  The above two merger types are expected to be violent. The third type is a non-violent merger: \\newline (3) Merger between two stars having a synchronized rotation. The secondary is massive enough to maintain synchronized orbital motion until merger occurs. The secondary survives, and the 2 stars form a massive star in a relatively gentle process. {{{Although this process is termed non-violent by us, it might still evolve on a dynamical time scale at some phases, and it has many interesting properties. The spiraling-in process of the secondary deep inside the envelope will release a huge amount of orbital energy, which might result in highly distorted mass loss event (Morris \\& Podsiadlowski 2006). On a later time, the process can alter the evolution of the star on the HR diagram (e.g., Podsiadlowski et al. 1990). However, we don't expect a bright flash in these cases. }}}  The gravitational and kinetic energy of the merging binary system can result in the following observational events: \\begin{enumerate} \\item Flash of light. This flash is formed by emission from a strongly shocked gas, in the primary and/or secondary envelope, and will be observed as a flash lasting as long as the secondary is violently slowed-down in the outer regions of the primary star. This can be from few times the dynamical time of the system up to a very long time, depending on the condition of the merging system. For the flash to be bright, the duration should be short, which implies a large relative velocity between the secondary and rotating primary envelope. \\item Gravitational and kinetic energy of matter expelled to large distances, and even leaving the system. The matter that does not leave the system, falls back on a dynamical time scale at its maximum distance, and when it becomes optically thick it contracts on its Kelvin-Helmholtz time scale. The Kelvin-Helmholtz time scale of the inflated envelope is much shorter than the Kelvin-Helmholtz time scale of the primary star, due to the very high luminosity and very low mass of the inflated envelope. The large energy in the inflated envelope and its relatively short contraction time implies that the energy deposited in the inflated envelope results in a bright phase of the merging system, lasting much longer than the initial flash. \\item Gravitational energy of the expanding inner layers of the primary and/or secondary (even destroying the secondary). This will be largely so when the secondary penetrate the deep layers of the primary star. When the inner layers of the primary finally relax to equilibrium it will be on a very long Kelvin-Helmholtz time scale. \\end{enumerate} To form a bright transient event only the first two energy channels are relevant. These two channels require violent interaction between the secondary and primary star. In the present paper we study the conditions for a violent merger triggered by the Darwin instability in a binary system. ",
        "conclusions": "In Section 1 we distinguish between three basic types of merger events: (1) A violent event where the secondary is disrupted and supplies most of the ejected and inflated envelope mass. This occurs when the secondary star is loosely bound, e.g., a pre-main-sequence star, and there is a grazing (rather than a head-on) collision from an eccentric orbit. (2) A violent merger where the secondary survives to a deeper depth in the primary envelope, and most of the inflated envelope originates in the very outer layers of the primary stellar envelope. Such a merger is likely to occur in a binary system when the secondary cannot maintain the primary envelope in synchronization (corotation) with the orbital motion. (3) A non-violent merger, occurring when the secondary maintains synchronization until merger occurs.  The basic results of the paper is that the type-2 merger defined above can cause a very bright event when the primary is a main sequence star. The peak luminosity can reach $L_{\\rm peak} \\ga 10^5 L_\\odot$, and last up to a few months. Our claim might seem somewhat contrary to intuition, as one might expect the merger event to become more violent as the secondary mass increases. However, as the secondary stellar mass increases, the relative velocity between the secondary star and the primary stellar envelope at the moment of merger decreases, implying weaker shock waves in the very outer layers of the envelope.  Interestingly, we find that the type-2 merger can lead to a merger event similar to that observed in V838 Mon, although type-1 event seems to fit observations a little better (TS06; Soker \\& Tylenda 2007). The hugely inflated envelope in the type-1 merger event discussed by TS06 for V838 Mon, must have both a grazing collision and a pre-main-sequence secondary star. These ingredients are not required in the type-2 merger scenario. However, the type-2 merger scenario cannot lead to a massive inflated envelope. The total mass in the expelled shell and in the inflated envelope of V838 Mon is $M \\simeq 0.1-0.3 M_\\odot$ (Tylenda 2005; Soker \\& Tylenda 2007). This mass favors the type-1 merger for V838 Mon. However, if this mass will turn out to be $M \\la 0.05 M_\\odot$ the type-2 merger model might work as well.  One of the important observational aspects of the merger events is the decline phase. Mass loss from a merger is always much smaller than the mass disturbed in the event. Most of this disturbed matter will form a more or less inflated envelope of the primary. When the merger processes dissipating energy are over the only source of luminosity in the envelope is the gravitational energy released in its contraction. However, the thermal time scale of the envelope is comparable to or lower than its dynamical time scale, especially if the envelope outer radius is much larger than the thermal equilibrium radius of the primary. Therefore the envelope contraction phase will be proceeded by a rapid cooling of the envelope outer layers. This cooling can go down to the Hayashi limit. Therefore the merger remnant is expected to decline as a very cool star. This is the main observational aspect allowing to distinguish the merger events from thermonuclear runaway events, e.g. nova-type outbursts. In the latter case, as it is well known from theoretical models and observations, the object has to evolve to very high effective temperatures ($\\ga 10^5$\\,K) before the final decline.  {{{ The exact evolution of the merger remnant as it re-establishes equilibrium requires numerical calculations. Podsiadlowski (2003) has considered an instantaneous removal of mass from a subgiant, together with an instantaneous addition of energy to its remaining envelope. When the heating is sufficiently high, in a short time the star reaches high luminosity, and then it contracts more or less along the Hayashi line. However, when a mass is removed and the heating (energy addition) is not sufficient, the star becomes very underluminous (Podsiadlowski 2003). The less-heated star starts its life to the left of the Hayashi line on the HR diagram, and then contracts and fades, but as a less luminous and smaller star than a star of the same parameters on the Hayashi line. The reason is that after mass is instantaneously removed from the outer layers of the envelope deep layers have to expand outward. This requires a lot of energy. In the merger process mass is added and lots of orbital energy is released. Thus there is no need for the inner envelope's layers to expand. On the contrary, after the short outburst which makes the star luminous and inflates an envelope, layers in the extended envelope have to contract. We therefore expect the merger remnant to decline along the Hayashi line. A behavior like this was observed during the eruption and subsequent decline of V838 Mon (Tylenda 2005). }}}  Soker \\& Tylenda (2007) suggest to call the violent merger events {\\it mergeburst.} As it is clear from our present study these events can be easily observable not only in our Galaxy but also beyond it. What is the galactic {\\it mergeburst} rate? We can crudely estimate it as follows. We use the estimated formation rate of blue stragglers as done by Ciardullo et al. (2005) to explain bright planetary nebulae in elliptical galaxies. Ciardullo et al. (2005) argue that blue stragglers can account for the formation of bright planetary nebulae in elliptical galaxies with an old stellar population, and estimate the blue stragglers formation rate per solar luminosity in galaxies. Their estimates yields a galactic blue stragglers formation rate of $\\sim 0.1 \\yr^{-1}$. The blue stragglers in their scenario requires that the secondary mass be close to that of the primary, $0.7 M_1 \\la M_2 <M_1$, whereas in our scenario $0.003 M_1 \\la M_2 \\la 0.15 M_1$. Because the allowed range in $M_2$ is smaller in our proposed model, and the mass ratio is much smaller, we expect the type-2 {\\it mergeburst} rate to be less frequent, say $0.01-0.05 \\yr^{-1}$. On the other hand, the merger rate of young systems due to perturbation from the environment, as proposed for V838 Mon (TS06), can add to the number of bright transient {\\it mergeburst} events. Over all, we estimate that bright transient {\\it mergeburst} events, resulting from stellar merger, will occur in the galaxy once every $\\sim 10-50$~yr."
    },
    "0606/astro-ph0606159_arXiv.txt": {
        "abstract": "In this paper we discuss the evolution of gravitationally unstable pre-galactic discs that result from the collapse of haloes at high redshift $z \\approx 10$ or so, which have not yet been enriched by metals. In cases where molecular hydrogen formation is suppressed the discs are maintained at a temperature of a few thousand degrees Kelvin.  However, when molecular hydrogen is present cooling can proceed down to a few hundred degrees Kelvin. Analogous to the case of the larger scale proto-galactic discs, we assume that the evolution of these discs is mainly driven by angular momentum redistribution induced by the development of gravitational instabilities in the disc. We also properly take into account the possibility of disc fragmentation. We thus show that this simple model naturally predicts the formation of supermassive black holes in the nuclei of such discs and provides a robust determination of their mass distribution as a function of halo properties. We estimate that roughly 5\\% of discs resulting from the collapse of haloes with $M\\approx 10^7 M_{\\odot}$ should host a massive black hole with a mass $M_{\\rm BH}\\approx 10^5 M_{\\odot}$. We confirm our arguments with time-dependent calculations of the evolution of the surface density and of the accretion rate in these primordial discs. The luminosity of the outer, colder disc is expected to be very low (in the range of a few thousand $L_{\\odot}$), while the formation of the black hole is expected to produce a burst with a luminosity of a few times $10^9L_{\\odot}$. This mechanism offers an efficient way to form seed black holes at high redshift. The predicted masses for our black hole seeds enable the comfortable assembly of $10^9 M_{\\odot}$ black holes powering the luminous quasars detected by the Sloan Digital Sky Survey at $z = 6$ for a concordance cosmology. ",
        "introduction": "The local demography of black holes at the centers of galaxies suggests that black hole formation is a generic feature of galaxy formation \\citep{magorrian98,tremaine02,ferrarese2000}. Mergers and accretion processes, in fact, merger induced accretion are implicated in the assembly of black holes \\citep{kauffmann00,volonteri03,dimatteo05}. However, all models require the formation of seed black holes at high redshift. The inferred large masses ($M_{\\rm bh} \\sim 10^9 M_{\\odot}$) of the black holes powering luminous quasars at $z \\sim 6$ detected by the Sloan Digital Sky Survey (SDSS) offers a challenge to mechanisms for production of seed black holes at higher redshifts \\citep{fan04}. In a $\\Lambda$CDM cosmology, these black holes have to assemble such large masses within a Gyr. The popular picture is that the remnants of the first generation of metal-free stars that form with an initial mass function that is biased towards higher masses \\citep{abel00,bromm02} provide the initial seed black holes with masses of the order of a 10's to 100 $M_{\\odot}$ \\citep{madaurees01,volonteri03,ricotti04,mapelli06}. Growing these seeds to the requisite masses by $z = 6$ to match the abundance of observed SDSS quasars requires various kinds of fine-tuning including phases of maximal growth powered by super-Eddington accretion. \\citet{volonteri05b} consider the growth of black holes in metal-free halos with $T > 10^4$ K that cool via atomic hydrogen lines leading to the formation of a fat, gas disc. They argue that an episode of super-Eddington accretion assembles black holes with masses of $10^6\\,M_{\\odot}$ at $z \\sim 15 - 20$.  There are alternative models that predict the direct formation of more massive seeds with masses of about $10^5\\,M_{\\odot}$.  These range from scenarios based on the formation of supermassive objects formed directly out of the collapse of dense gas \\citep{haehnelt93,umemura93,loeb94,eisenstein95,bromm03,koushiappas04,begelman06}. The key limiting factor for these models is the disposal of the angular momentum. One approach taken by \\citet{eisenstein95} and \\citet{koushiappas04} is to argue that it is preferentially either low spin halos with consequently low angular momentum gas or the low spin tail of the gas distribution in halos that can cool efficiently that form these seeds. Significant transfer of angular momentum is still required to go from the central massive object to the final collapsed hole that is expected to occur via the post-Newtonian instability. Non-axisymmetric structures like bars have been proposed \\citep{shlosman90} to efficiently transfer angular momentum at these late stages. In a recent paper, \\citet{begelman06} have pursued this picture and argue that a low specific entropy `quasi-star' is produced as a result of the cascade of the bars-within-bars instability leading to the formation of a black hole of a few solar masses.  In this paper, we argue for the formation of more massive black hole seeds ($\\sim 10^5\\,M_{\\odot}$) from gas cooling in primordial halos via the growth of gravitational instabilities. We offer a picture wherein accretion processes and fragmentation criteria are addressed in a coupled fashion. The outline of our paper is as follows: in Section~2, we outline the scenario for pre-galactic disc formation in dark matter halos and discuss their accretion properties, the fragmentation criteria which determine the fate of the gas are discussed in Section~3 and the implied mass distribution of central concentrations is derived in Section~4. Detailed time-dependent models that confirm the analytic calculations in the previous sections are presented in Section~5, the resultant luminosity and potential observability of these sources is discussed in Section~6 followed by a discussion and implications of our work. ",
        "conclusions": "In this paper, we demonstrate that massive central concentrations can form naturally from pre-galactic discs that assemble in dark matter halos at high redshift. The masses of these concentrations depend on key properties of the host halo, mass, spin and gas cooling in the halo. In particular, low spin halos and massive halos are most efficient in concentrating gas in their centers. However, not all halos will be able to accrete gas into their centers. Taking into account the possibility of fragmentation furthers restricts the formation of eventual seed black holes.  Using simple stability criteria, we predict that fragmentation and subsequent star formation are the fate for gas in some fraction of haloes. We derive 3 interesting regimes that are determined by the ratio of $T_{\\rm vir}/T_{\\rm gas}$, (i) if this ratio is greater than 2.9, haloes will form discs that will fragment and form stars and not directly form black holes; (ii) if this ratio lies between 1.8 - 2.9, haloes with low spin will lead to central mass concentrations; however, the lowest spin cases might be affected by fragmentation; (iii) if this ratio is lower than 1.8 the haloes will not produce fragmenting discs but will successfully accrete gas in their centers.  Calculating the accretion rates for the fate of the accumulated gas we find super-Eddington rates leading to eventual black holes of masses up to $10^5 M_{\\odot}$.  The model proposed here has two important predecessors in the work of \\citet{eisenstein95} and in that of \\citet{koushiappas04}, but also presents significant differences with respect to both of them. Primordial gas acquires angular momentum through tidal interaction with the surrounding \\citep{peebles69}. This is generally measured through the parameter $\\lambda$, the distribution of which can be determined from numerical simulations \\citep{warren92}. Thus, the centrifugal barrier is the main obstacle to the formation of any compact object in the center of primordial galaxies. It is therefore not surprising that all models (including our own) predict that the most favourable sites for black hole formation are haloes with low spin. It is also not surprising that any black hole formation model has to deal with the problem of angular momentum removal. This issue was tackled in the work of \\citet{eisenstein95}, who calculated the viscous time-scale for primordial thin discs and assumed that black hole formation would only occur when the discs are so compact that their viscous time-scale is less than the typical time-scale for star formation. They found that only very rare haloes, with $\\lambda$ smaller by at least a factor 30 with respect to the average spin, would be able to produce black holes. This is substantially different from what we find here. Indeed, we find that even haloes with $\\lambda$ as large as 0.02-0.03 do produce large central mass concentrations. There are two main aspects with respect to which our model differs substantially from \\citet{eisenstein95}. The first one is the nature of the viscosity mechanism. \\citet{eisenstein95} explicitly neglect the contribution of gravitational instabilities and assume that viscosity is driven by ``turbulence''. In standard accretion discs, turbulence is thought to arise from MHD instabilities \\citep{balbusreview}. However, for such primordial discs, where the primordial magnetic field is very weak and the gas is predominantly atomic or molecular, it is unlikely that MHD instabilities would be dynamically important. The second, and most important, difference is that \\citet{eisenstein95} considered relatively thin discs, with $H/R\\approx 0.03$, whereas our discs are substantially thicker. Since the viscous timescale is proportional to $(H/R)^{-2}$, this means that the viscous timescale as estimated by \\citet{eisenstein95} is much larger than in our case. This in turn, is what leads to their more pessimistic estimate of the rarity of black hole forming haloes. Finally, a significant difference between our model and that of \\citet{eisenstein95} is that our model naturally leads to a robust determination of the seed black hole mass function.  \\citet{koushiappas04} propose that black hole seeds with masses of the order of $10^5M_{\\odot}$ (i.e. very similar to those obtained here) can form out of low angular momentum material in massive haloes. They assume that the lowest angular momentum material within the halo forms a compact disc which is gravitationally unstable and accretes onto the center due to the effect of an effective viscosity driven by the instability, in a way similar to what we propose here. However, they do not consider self-consistently the evolution of the surface density profile induced by the assumed viscosity mechanism. In their picture, the disc surface density simply grows linearly with time (at the same rate, independent of the radius) and at any given time it reflects the original angular momentum distribution of the gaseous component of the halo. On the contrary, what determines the surface density profile is actually the viscosity mechanism (which, being related to gravitational instabilities, in turn only depends on $Q$). As we have shown in Section \\ref{sec:time} (Figs. \\ref{fig:time1} and \\ref{fig:time2}), where we numerically follow the evolution of such discs, the final $\\Sigma$ profile (and the total accreted mass $m_{\\rm a}$) is the same whether we add mass to the disc with an exponential profile or with a $\\delta$-function, representing two extreme cases in the original angular momentum distribution within the halo. As a consequence of this, the estimates of black hole masses and their distribution given by \\citet{koushiappas04} artificially depend on the initial angular momentum distribution and also on the assumed viscosity law. On the other hand, we have also clearly shown that, as long as viscosity is driven by gravitational instabilities, the black hole mass distribution is independent of viscosity.  In the present paper, for illustrative purposes, we have taken the simplifying assumption that these primordial discs live in the potential well of a simple isothermal sphere, with a given circular velocity $V_{\\rm h}$. For a more realistic density profile, the NFW profile \\citep{nfw}, in its innermost parts, the disc mass might dominate the potential well and become radially unstable, giving rise to bar-like instabilities. This, however, will leave our results and conclusions unaffected. Firstly, the instability criterion for such global instabilities \\citep{christo95} is rather similar to our adopted criterion based on $Q$, and marginal stability is found to occur for equivalent values of $Q_{\\rm c}\\approx 2-3$, as adopted here. Secondly, since the disc mass is only a small fraction of the halo mass, only the innermost parts of the disc will be subject to bar-like instabilities, at radii much smaller than the typical disc radius $R_{\\rm d}$. The development of such instabilities might in fact enhance the accretion rate in the inner disc, but will not change the estimates of the total accreted mass obtained here.  The important and interesting consequence of our model is that black hole masses of $10^9 M_{\\odot}$ powering the luminous SDSS quasars at $z = 6$ can form comfortably within the available time of 1 Gyr in the concordance cosmology from our seed masses of $10^5 M_{\\odot}$ at $z \\sim 10$."
    },
    "0606/astro-ph0606190_arXiv.txt": {
        "abstract": "We investigate the cosmological effects of a neutrino interaction with cold dark matter.  We postulate a neutrino that interacts with a ``neutrino interacting dark matter'' (NIDM) particle with an elastic-scattering cross section that either decreases with temperature as $T^2$ or remains constant with temperature. The neutrino--dark-matter interaction results in a neutrino--dark-matter fluid with pressure, and this pressure results in diffusion-damped oscillations in the matter power spectrum, analogous to the acoustic oscillations in the baryon-photon fluid. We discuss the bounds from the Sloan Digital Sky Survey on the NIDM opacity (ratio of cross section to NIDM-particle mass) and compare with the constraint from observation of neutrinos from supernova 1987A.  If only a fraction of the dark matter interacts with neutrinos, then NIDM oscillations may affect current cosmological constraints from measurements of galaxy clustering. We discuss how detection of NIDM oscillations would suggest a particle-antiparticle asymmetry in the dark-matter sector. ",
        "introduction": "Flat cosmological models with baryons (about $5\\%$ of the total energy content of the universe), cold dark matter (CDM, $25\\%$), cosmological constant (or dark energy, $70\\%$) and an adiabatic, nearly scale-invariant spectrum of density fluctuations explain most cosmological observations. However, we still lack a satisfactory understanding of both dark matter and dark energy, a puzzle for both particle physics and cosmology.  The most favored candidates for dark matter are cold, collisionless massive particles, which are non-relativistic for most of the history of the Universe and so can cluster gravitationally during matter domination. Candidates for these dark-matter particles can be found in supersymmetric extensions of the standard electroweak model---namely neutralinos with mass on the order of 100 GeV \\cite{Jungman:1995df}---or in other theories (e.g., the axion, which may arise in the Peccei-Quinn mechanism \\cite{axion}). These cold-dark-matter models account well for cosmic microwave background (CMB) observations on the largest scales as well as measurements of the large-scale distribution of galaxies.  However, observations on galactic and sub-galactic scales may conflict with the predictions, from numerical simulations and analytic calculations, of CDM models.  Indeed, cold and collisionless dark-matter models seem to predict an excess of small-scale structures \\cite{silk1}, and numerical simulations \\cite{prada} predict far more satellite galaxies in the Milky Way halo than are observed.  Several solutions have been proposed to explain these discrepancies, for example, inflationary models with broken scale invariance \\cite{marcandrew}.  However, most other explanations invoke modifications to the properties of dark-matter particles. For example, a warm-dark-matter candidate, like a sterile neutrino, has been suggested because it suffers free streaming and suppresses the matter power spectrum on small scales~\\cite{silk2}. A dark-matter particle that results from decay of a short-lived charged particle can also suppress small-scale power \\cite{kris}. Other possibilities include a dark-matter particle that interacts with others particles such as  photons \\cite{bom1,xuelei,dipole} or neutrinos \\cite{bom2,stefano} or self-interacting dark matter \\cite{spergel}. For a review of different alternative scenarios to standard collisionless cold dark matter, see Ref.~\\cite{tesina}.  In this paper, we investigate the possibility of a neutrino interacting dark matter (NIDM) component. If dark matter and neutrinos interact, there was an epoch in the very early universe in which they were strongly coupled. Dark-matter perturbations that entered the horizon during this period would then be erased because of diffusion damping, and the suppression scale will depend on the dark-matter--neutrino interaction. Even if only a fraction of the dark matter interacts with neutrinos, a pattern of oscillations in the matter power spectrum arises, much like the oscillations in the baryon-photon fluid.  In the following, we limit our study of DM-neutrino couplings to effects on cosmological scales in the frequency range smaller than $k < 0.2\\,h$\\,Mpc$^{-1}$---i.e., on scales where linear perturbation theory is viable. We consider flat cosmological models with an adiabatic and nearly scale invariant spectrum $P(k)\\sim k^n$ of density perturbations where $n=0.97$. Unless explicitly stated, the energy content of the Universe corresponds to the standard $\\Lambda$CDM model with baryons contributing as $\\Omega_b=0.05$ and a cold-dark-matter energy density $\\Omega_{dm}=0.25$. We also choose a Hubble parameter $h=0.73$, a standard value $3.04$ for the effective number of (massless) neutrinos \\cite{Mangano:2005cc}, and dark energy in the form of a cosmological constant with equation-of-state parameter $w=-1$. The interaction between dark matter and neutrinos is given in terms of the opacity $Q=\\langle \\sigma_{dm-\\nu}|v| \\rangle /m_{dm}$, the ratio of the thermal averaged dark-matter--neutrino cross section to the mass of the dark-matter particle.  The paper is organized as follows. In the next Section, we discuss a class of models of neutrino--dark-matter interaction for both scalar and spinor dark-matter candidates and obtain an estimate for the opacity $Q$. In Section III, we outline the cosmological consequences of a NIDM component, and we compare those predictions with the latest data on galaxy clustering from the Sloan Digital Sky Survey (SDSS). In Section IV, we consider astrophysical constraints, particularly those from observation of neutrinos from supernova 1987A. Finally, in the last Section, we report our conclusions. ",
        "conclusions": "In this paper, we have studied the cosmological consequences of a possible coupling between neutrinos and light dark matter with mass in the MeV range. We considered two possible behaviors for the thermally-averaged neutrino-DM elastic-scattering cross section, either decreasing with temperature as $T^2$ or constant. We compared the NIDM scenario with the large-scale galaxy distribution and obtained upper limits on the opacity (ratio of the DM-neutrino cross section to the dark-matter mass) of $Q_2 < 10^{-42}\\, {\\rm cm}^2$~MeV$^{-1}$ and $Q_0 < 10^{-34} {\\rm cm}^2$~MeV$^{-1}$ at the $95 \\%$ C.L. These limits may be relaxed if one consider the possibility that only a fraction of the dark matter is made of NIDM. The main cosmological observable for NIDM consists in diffusion-damped oscillations in the matter power spectrum. Those NIDM oscillations may affect current cosmological constraints on neutrinos masses and dark energy from galaxy clustering. We have stressed that strongly-coupled DM particles would have a non-negligible relic abundance today only if an asymmetry between DM particle and antiparticle is produced at some early stage in the evolution of the Universe, since their density would vanish today because of effective annihilation processes into neutrinos down to temperatures much smaller than the DM mass. Detection of NIDM-induced oscillations in the LSS power spectrum would be a hint for such a non-standard scenario."
    },
    "0606/astro-ph0606645_arXiv.txt": {
        "abstract": "Hostile tidal forces may inhibit the formation of Jovian planets in binaries with semimajor axes of $\\la$$50\\au$, binaries that might be called ``close'' in this context.  As an alternative to in situ planet formation, a binary can acquire a giant planet when one of its original members is replaced in a dynamical interaction with another star that hosts a planet.  Simple scaling relations for the structure and evolution of star clusters, coupled with analytic arguments regarding binary-single and binary-binary scattering, indicate that dynamical processes can deposit Jovian planets in $<$1\\% of close binaries.  If ongoing and future exoplanet surveys measure a much larger fraction, it may be that giant planets do somehow form frequently in such systems. ",
        "introduction": "\\label{sec:intro}  Surveys using the Doppler technique have identified over 150 extrasolar planets in the last decade.  The available data reveal important clues to the formation of giant planets around {\\em single} stars \\citep[e.g.,][]{Marcy2005}.  Comparatively little is known about the population of {\\em binary} star systems that harbor planets.  Past planet searches have largely excluded known binaries with angular separations of $\\la$$1\\arcsec$, where blending of the two stellar spectra decreases the sensitivity to small velocity shifts.  Roughly 30 planets have been detected around stars in binaries \\citep{Raghavan2006}.  Most of these binaries are very wide, although several have separations of $\\la$20\\,AU, small enough to challenge standard ideas on Jovian planet formation.  Many more of these compact systems must be found before we can draw robust conclusions.  In an ongoing targeted search for planets in close, double-lined, spectroscopic binaries, \\citet{Konacki2005} discovered a ``hot Jupiter'' orbiting the outlying member of the hierarchical triple star HD 188753.  The inner binary is sufficiently compact that its influence on the third star is essentially that of a point mass.  What is intriguing about this system is that a disk around the planetary host star would be tidally truncated at a radius of only $\\simeq$1\\,AU, perhaps leaving insufficient material to produce a Jovian-mass planet \\citep[][]{Jang-Condell2005}.  Broader questions of how a binary companion impacts planet formation have been explored in the literature.  If a protoplanetary disk is tidally truncated at $\\la$10\\,AU, stirring by the tidal field may prevent the growth of icy grains and planetesimals, as well as stabilize the disk against fragmentation \\citep[][]{Nelson2000,Thebault2004,Thebault2006}.  In this case, neither the core-accretion scenario \\citep[e.g.,][]{Lissauer1993} nor gravitational instability \\citep[e.g.,][]{Boss2000} are accessible modes of giant planet formation.  However, the tidal field might also trigger fragmentation of a marginally stable disk \\citep{Boss2006}. Whether or not giant planet formation is inhibited in close binaries remains an open problem.  These uncertainties are circumvented if one member of a close binary is divorced from its original companion and acquires a new partner star with a planet in tow.  An example of such an event is an exchange interaction between a binary and a single star in a cluster environment.  \\citet{Pfahl2005} and \\citet{PortegiesZwart2005} proposed a form of this idea as a solution to the puzzle of HD~188753. Here we present a more general account of dynamical processes that deposit giant planets in binaries hostile to planet formation.  We focus exclusively on encounters that mix a binary and a single star or two binaries.  Higher order multiples are neglected here, but deserve further attention in light of HD~188753.  We define ``close binary'' in \\S~\\ref{sec:closebin}.  An overview of binary scattering dynamics is given in \\S~ \\ref{sec:scatter}.  Various aspects of star clusters are summarized in \\S~\\ref{sec:clusstat}. Ingredients from \\S\\S~\\ref{sec:scatter} and \\ref{sec:clusstat} are combined in \\S~\\ref{sec:binplanfrac} to estimate the frequency of giant planets in close binaries.  In \\S~\\ref{sec:observations}, our results are discussed in the context of current exoplanet surveys. ",
        "conclusions": ""
    },
    "0606/quant-ph0606117_arXiv.txt": {
        "abstract": "We report on a study of complementarity in a two-terminal $closed$-$loop$ Aharonov-Bohm interferometer. In this interferometer, the simple picture of two-path interference cannot be applied. We introduce a nearby quantum point contact to detect the electron in a quantum dot inserted in the interferometer. We found that charge detection reduces but does not completely suppress the interference even in the limit of perfect detection. We attribute this phenomenon to the unique nature of the closed-loop interferometer. That is, the closed-loop interferometer cannot be simply regarded as a two-path interferometer because of multiple reflections of electrons. As a result, there exist indistinguishable paths of the electron in the interferometer and the interference survives even in the limit of perfect charge detection. This implies that charge detection is not equivalent to path detection in a closed-loop interferometer. We also discuss the phase rigidity of the transmission probability for a two-terminal conductor in the presence of a detector. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/gr-qc0606091_arXiv.txt": {
        "abstract": "We examine the transition of a particle across the singularity of the compactified Milne (CM) space. Quantization of the phase space of a particle and testing the quantum stability of its dynamics are consistent to one another.  One type of transition of a quantum particle is described by a quantum state that is continuous at the singularity. It indicates the existence of a deterministic link between the propagation of a particle before and after crossing the singularity. Regularization of the CM space leads to the dynamics similar to the dynamics in the de Sitter space. The CM space  is a promising model to describe  the cosmological singularity deserving further investigation by making use of strings and membranes. ",
        "introduction": "Presently  available cosmological data indicate that known forms of energy and matter comprise only $4\\%$ of the makeup of the Universe. The remaining $96\\%$ is unknown, called `dark', but its existence is needed to explain the evolution of the Universe \\cite{Spergel:2003cb,Bahcall:1999xn}.  The dark matter, DM, contributes $22\\%$ of the mean density. It is introduced to explain the observed dynamics of galaxies and clusters of galaxies. The dark energy, DE, comprises $74\\%$ of the density and is responsible for the observed accelerating expansion. These data mean that we know almost nothing about the dominant components of the Universe!  Understanding the nature and the abundance of the DE and DM within the standard  model of cosmology has  difficulties \\cite{NGT,GDS}. These difficulties have led many physicists to seek anthropic explanations which, unfortunately, have little predictive power. An alternative model has been proposed by Steinhardt and Turok (ST) \\cite{Steinhardt:2001vw,Steinhardt:2001st,Steinhardt:2004gk}. It is based on the idea of a cyclic evolution, CE, of the Universe. The ST model has been inspired by string/M  theories \\cite{Khoury:2001bz}. In its simplest version it assumes that the spacetime can be modelled by the higher dimensional compactified Milne, CM, space. The attraction of the ST model is that it potentially provides a complete scenario of the evolution of the universe, one in which the DE and DM play a key role in both the past and the future. The ST model \\textit{requires} DE for its consistency, whereas in the standard model, DE is introduced in a totally  \\textit{ad hoc} manner. Demerits of the ST model  are extensively discussed in \\cite{Linde:2002ws}. Response to the criticisms of \\cite{Linde:2002ws} can be found in \\cite{NGT}.  The mathematical structure and self-consistency of the ST model has yet not been fully tested and understood. Such task presents a serious mathematical challenge. It is the subject of our research programme.  The CE model has in each of its cycles  a quantum phase including the cosmological singularity, CS.  The CS plays key role because it joins each two consecutive classical phases.  Understanding the nature of the CS has  primary importance for the CE model.  Each CS consists of contraction and expansion phases. A physically correct model of the CS, within the framework of string/M theory, should be able to describe  propagation of a  p-brane, i.e. an elementary object like a particle, string and membrane, from the pre-singularity to post-singularity epoch. This is the most elementary, and fundamental, criterion that should be satisfied. It presents a new criterion for testing the CE model. Hitherto, most research has focussed on the evolution of scalar perturbations through the CS.  Successful quantization of the dynamics of p-brane will mean that the CM space is a promising candidate to model the evolution of the Universe at the cosmological singularity. Thus, it could be further used in advanced numerical calculations to explain  the data of observational cosmology. Failure in quantization may mean that the CS should be modelled by a spacetime more sophisticated than the CM space.  Preliminary insight into the problem has already been achieved by studying classical and quantum dynamics of a test particle in the two-dimensional CM space \\cite{Malkiewicz:2005ii}.  The present paper is a continuation of \\cite{Malkiewicz:2005ii} and it addresses the two issues:  the Cauchy problem at the CS  and the stability problem in the propagation of a particle across the CS. Both issues concern the nature of the CS.  In Sec. II we define and make comparison of the two models of the universe: the CM space and the regularized CM space. The classical dynamics of a particle in both spaces is presented in Sec. III.  The quantization of the phase space of a particle is carried out in Sec. IV. In Sec. V we examine the stability problem of particle's dynamics both at classical and quantum levels. We summarize our results, conclude  and suggest next steps in Sec. VI. ",
        "conclusions": "The Cauchy problem at the cosmological singularity of the geodesic equations may be `resolved' by the regularization, which replaces the double conical vertex of the CM space by a space with the vertex of the big-bounce type, i.e. with non-vanishing space dimension at the singularity. We have presented a specific example of such regularization of the CM space. Both classical and quantum dynamics of a particle in the  regularized CM space  are deterministic and stable. We have examined these aspects of the dynamics at the phase space and Hamiltonian levels. The classical and quantum dynamics of a particle in the regularized CM space is similar to the dynamics in the de Sitter space \\cite{WP,Piechocki:2003hh}. We are conscious that our regularization of the singularity is rather {\\it ad hoc}. Our arguing (presented at the beginning of Sec. IIB) that taking into account the interaction of a  physical particle with the singularity may lead effectively to changing of the latter into a big-bounce type singularity should be replaced by analyzes. However, examination of this problem is beyond the scope of the present paper, but will be considered elsewhere.  The classical dynamics in the CM space is unstable (apart from the one class of geodesics). However, the quantum dynamics is well defined. The Cauchy problem of the geodesics is not an obstacle to the quantization. The examination of the quantum stability has revealed surprising result that in one case a quantum particle propagates deterministically in the sense that it can be described by a quantum state  that is continuous at the singularity. This case is very interesting as it says that there can exist deterministic link between the data of the pre-singularity and post-singularity epochs. All other states have discontinuity at the singularity of the CM space, but they can be used successfully to construct a Hilbert space. This way we have proved the stability of the dynamics of a {\\it quantum} particle.  At the quantum level the stability condition requiring the boundedness from below of the Hamiltonian operator means  the imposition of the first-class constraint onto the space of quantum states to get the space of  {\\it physical}  quantum states. The resulting equation depends on all spacetime coordinates. In the pre-singularity and post-singularity epochs the CM space is locally isometric to the Minkowski space \\cite{Malkiewicz:2005ii}. Owing to this isometry, the stability condition is in fact  the Klein-Gordon, KG, equation. The space of solutions to the KG equation in these two epochs and the corresponding Hilbert space are fortunately non-trivial ones, otherwise our quantum theory of a particle would be empty.  Quantization of the phase space carried out in Sec. IV (and in our previous paper \\cite{Malkiewicz:2005ii}), corresponds to some extent to the method of quantization in which one first solves constraints at the classical level and then quantize the resulting theory. Quantization that we call here examination of the stability at the quantum level, is effectively the method in which we impose the constraint, but at the quantum level. The results we have obtained within both methods of quantization are consistent. It means that the quantum theory of a particle in the compactified Milne space does exist. The CM space seems to model the cosmological singularity in a satisfactory way\\footnote{Our result should be further confirmed by the examination of the dynamics of a particle in a higher dimensional CM space.}.  It turns out that the  time-like geodesics of our  CM space may have interpretation in terms of cosmological solutions of some sophisticated higher dimensional field theories \\cite{Russo:2004am,Bergshoeff:2005cp,Bergshoeff:2005bt}. We have already discussed some aspects of this connection in our previous paper (see Sec. 5 of \\cite{Malkiewicz:2005ii}). Presently, we can say that in one case (see Sec. 4 of \\cite{Russo:2004am} and Sec. V.B.1 of this paper) this analogy extends to the quantum level: transition of a particle through the cosmological singularity in both models is mathematically well defined. In both cases the operator constraint is used to select quantum physical states. Elaboration of this analogy needs an extension of our results to the Misner space (that consists of the Milne and Rindler spaces) because it is the spacetime  used in \\cite{Russo:2004am}. Another subtlety is connected with the fact that we carry out analysis in the compactified space, whereas the authors of \\cite{Russo:2004am} use the covering space.   There exists another model to describe the  evolution of the universe based on string/M theory. It is called the pre-big-bang model \\cite{Gasperini:2002bn}. However, the ST model is more self-consistent and complete.  Other sophisticated model called loop quantum cosmology, LQC, is based on non-perturbative formulation of quantum gravity called the loop quantum gravity, LQG \\cite{TT,CR}. It is claimed that the CS is resolved in this approach \\cite{Bojowald:2006da}. However, this issue seems to be still open due to  the assumptions made in the process of truncating the infinite number degrees of freedom of the LQG to the finite number used in the LQC \\cite{Brunnemann:2005in,Brunnemann:2005ip}. This model has also problems in obtaining an unique semi-classical approximations \\cite{Nicolai:2005mc}, which are required to link the quantum phase with the nearby classical phase in the evolution of the Universe. For response to \\cite{Nicolai:2005mc} we recommend \\cite{Thiemann:2006cf}.  Quantization of dynamics of {\\it extended} objects in the CM space is our  next step. There exist  promising results on propagation of a string and membrane \\cite{Turok:2004gb,Niz:2006ef}. However, these results concern extended objects in the low energy states called the zero-mode states and quantum evolution is approximated by a semi-classical model. Recently, we have quantized the dynamics of a string in the  CM space rigorously \\cite{Malkiewicz:2006bw}, but our results concern only the zero-mode state of a string.  For drawing firm conclusions about the physics of the problem, one should also examine the non-zero modes. Work is in progress."
    },
    "0606/astro-ph0606703_arXiv.txt": {
        "abstract": " ",
        "introduction": "In recent years the exploration of the universe has provided detailed information about the distance-redshift relationship of many sources at very small redshifts as well as at redshifts of order unity. These observations, {\\it if} interpreted within the framework of a homogeneous Friedmann-Lema\\^itre-Robertson-Walker (FLRW) metric indicate that the universe is presently undergoing a phase of accelerated expansion \\cite{review}. Such an accelerated expansion is most commonly interpreted as evidence  for the presence of a negative pressure component (``Dark Energy'') in the mass-energy density of the universe. This has given rise to the so-called ``concordance'' flat $\\Lambda CDM$ model, where dark energy constitutes $70\\%$ of the present energy budget in the form of a cosmological constant along with $30\\%$ of matter. A second fact lending credence to the  $\\Lambda CDM$ model is the observation that CMB excludes, with high precision, the presence of spatial curvature. Applying this to the late time universe in an FLRW model leads to a mismatch between the observed matter density and  the FLRW critical density. This mismatch can also be cured by the presence of a Dark Energy component.  However, from the point of view of fundamental physics  the presence of such a tiny cosmological constant is almost absurd: first, the scale is extremely low compared to any possible fundamental scale, and second it is tuned in such a way that it shows up just at $z\\simeq 1$, where we live. This is so baffling\\footnote{There have been some ideas/efforts to address these issues. For instance, see string-gas inspired  coupled quintessence models \\cite{string} by one of the authors. However, it is fair to say that there is not yet a completely satisfactory way of solving both the ``smallness'' and the ``coincidence'' problems.} that before concluding that we are indeed facing this huge puzzle, it is worth making efforts to explore if the correct equations are being used to fit the data.  The homogeneous FLRW metric  certainly offers the simplest paradigm to interpret  cosmological observations. Indeed, the FLRW model is a very good approximation in the Early Universe as probed by the homogeneity of the CMB (density contrasts in  photons and dark matter are of the order  $10^{-5}$ and  $10^{-3}$ respectively). However at low redshifts, the density contrast in matter grows up to values of ${\\cal O}(1)$ and beyond: small scales become nonlinear first, and then more and more scales enter this regime. Today, scales of order of $60/h \\, {\\rm Mpc}$ (which is $2\\%$ of the size of the horizon) have an average density contrast of order 1. Structures observed with density contrast larger than $0.1$ (so, already in a mildly nonlinear regime) extend to few hundreds of Mpc (around $10\\%$ of the size of the horizon). It is thus not clear or obvious  that one can keep using  FLRW metric to interpret the large amount of  high precision experimental data collected in recent times in a regime where most of the mass in the Universe is clumped into structures or it is forming structures.  Three are three main physical effects that we could be missing while using the FLRW model: \\begin{description} \\item[I.\\ \\ ] The ``overall'' dynamics of a Universe with these inhomogeneities could be significantly different from the FLRW dynamics. \\item[II.\\ ] Even if the dynamics is approximatively the same,  one has to wonder about the light propagation in a clumpy Universe: since all our conclusions about the dynamics are based on observations of light, this is of crucial importance. \\item[III.] The fact that we have just one observer could influence significantly the observations: we could live in an especially underdense or overdense region, which may induce large corrections on the observations. \\end{description} In recent times there has been an on-going debate as to whether some of these effects may be large enough to give rise to an ``apparent acceleration'' leading us to believe in dark energy when there is none. The aim of this paper is to construct an {\\it exact model of non-linear structure formation} where we can systematically study these effects. We wanted a model where one can study (possibly also {\\it analytically}) the effects that inhomogeneities can cause on quantities like the {\\it redshift} and the {\\it angular distance}, thereby hopefully providing us with valuable insights to resolve the issue. The full problem of solving Einstein equations and light propagation for a realistic mass distribution is beyond present human possibilities, so some simplification is needed. Our simplification is to assume spherical symmetry and in particular our model is based on Lema\\^itre-Tolman-Bondi (LTB) metrics which are exact spherically symmetric solutions of general relativity with only dust. This approach has the obvious advantage of being exact: it does not use perturbation theory (as most of previous literature did to address {\\bf I}), it allows to study light propagation without any FLRW assumptions (as opposed to literature addressing point {\\bf II}), studying the corrections on the distance and also on the redshift of light (which has been ignored by most previous literature). The challenge however,   is to make these kind of models as realistic as possible, or at least be able to extrapolate from them the information that is relevant for our Universe.  So far such models have been used mostly to understand what happens to observations if we are living in a special position, a void for example~\\cite{celerier,pm84,kl92,Tomita00,Tomita01,wiltshire05,moffat,alnes,abtt06,m05,ps06}, and therefore can at best be thought of as describing the local universe. It is interesting to go beyond this approach. Firstly, we may be missing some important effects (such as light propagating through several structures) by looking just at the local universe. Secondly, although these models show interesting effects capable of mimicking dark energy, it is not clear how realistic are these setups~\\cite{bolejko,garfinkle} (also, some of these results have been questioned by~\\cite{Flanagan}). On the other hand, we note that a different approach (aimed at solving question {\\bf I}), based on back reaction of inhomogeneities on the overall dynamics of the Universe, suggests a much smaller correction at the first nontrivial order~\\cite{Huiseljak,rasanen,KMNR,frysiegel}, although even in this framework some authors have argued in favour of a much larger effect~\\cite{rasanen,Notari,KMR}. Thus it is imperative that we understand the origin of the large effects and in the process hopefully also bridge the gap between the perturbative back-reaction effects and the exact toy model approach.   It is clear that to address these issues we have to go beyond a ``local'' or a ``perturbative'' description. Thus, although we  work with the simplifying assumption of spherical symmetry,  we try to make our model realistic in the following sense. The universe  looks like an onion (and we frequently refer to it as the ``Onion model''): there is a homogeneous background density and on top of it density fluctuations as a function of the radial coordinate. Our model preserves large-scale homogeneity and  the shells are also distributed in a homogeneous way, making the picture a lot closer to the real world. One may wonder, nonetheless, that the centre is still a special point. In fact, we find that the centre has some special properties. For instance typically, the curvature tends to become large at the centre. In this case if we find significant corrections to, say, the redshift-luminosity distance relation near the centre, it is hard to distinguish whether the correction originates from the non-linear structure formation, or the ``excess curvature''. This therefore can jeopardise what we are trying to avoid, that of violating the cosmological principle. We deal with this problem by considering an observer who sits in a generic position and looks at  sources in the radial direction, along which we have periodic inhomogeneities. To our knowledge, the luminosity-redshift relation in LTB models has always been analyzed putting the observer at the centre of the  coordinates\\footnote{Except for an analysis at small $z$ in \\cite{humphreys}.}. For the first time we derive an exact expression for the luminosity distance of an object as seen by an off-centre observer\\footnote{In this respect we note that~\\cite{Flanagan} has shown that locally around the centre in the LTB model the deceleration parameter cannot be negative.  However, this result is only local  and so the universe may accelerate slightly far away from the observer.}.  At early times (at last scattering for example) we assume the density fluctuations to have an amplitude of the order of the fluctuations observed in the CMB. Then we are able to follow exactly the evolution of density fluctuations. At redshifts of $z \\gtrsim {\\cal O}(10)$ the density contrast grows exactly as in a perturbed FLRW. Then, when the density contrast becomes of  ${\\cal O}(1)$ nonlinear clustering appears. So at low redshift the density does what is expected: overdense regions start contracting and they become thin shells (mimicking structures), while underdense regions become larger (mimicking voids) and eventually they occupy most of the volume. One can see this completely analytically upto high density contrasts (we have checked that the analytical solution is good even with $\\de \\rho/\\rho\\sim {\\cal O}(100)$), which already makes it  an interesting model for the theory of structure formation. However, the really nice feature of the Onion model is that one can even calculate how the redshift and the various distances get corrected when we are in this non-linear regime.   In this model we can solve for the radial geodesics, therefore determining exactly what happens to the redshift of light. We place the observer in a generic location and measure the redshift as a function of the source coordinate. Then we trace a beam of light that starts from the source with some initial angle, and in particular we are able to compute the area seen by the observer. This provides us with the angular distance ($D_A$) or equivalently the luminosity distance ($D_L$) for a generic source. In this way we provide the $D_L-z$ relationship (or the $D_A-z$ relationship) exactly. Thus we are able to study the three possible effects ({\\bf I},{\\bf II}, and {\\bf III}) both numerically and  analytically. We now discuss our findings briefly.  The first issue ({\\bf I}) of whether there is an overall effect in expansion due to inhomogeneities has been investigated  only recently. Buchert~\\cite{buchert} has defined an average volume expansion and has developed a formalism to take into account of the backreaction of the inhomogeneities, although the problem of computing the size and the effect of the corrections is left unanswered. Rasanen~\\cite{rasanen} has tried to estimate the effect of inhomogeneities (extending the formalism developed in~\\cite{brand}, albeit in a different context) using perturbation theory, finding an effect of order $10^{-5}$ (in agreement with~\\cite{Huiseljak}) and speculating about nonlinear effects. A complete second order calculation has been done in~\\cite{KMNR}, and in a subsequent work one of us has shown~\\cite{Notari} that the perturbative series is likely to diverge at small $z$ due to the fact that all the perturbative corrections become of the same size (each one is $10^{-5}$), and one may hope that these corrections give rise to acceleration (which it is shown to be in principle possible~\\cite{nambu,Notari}). The fact that this may lead to acceleration has been explored in subsequent papers~\\cite{KMR} and criticized in others (\\cite{frysiegel} for example have argued the effect to be too small). Unfortunately a perturbative approach seem inadequate to answer this question as the phenomenon is intrinsically nonlinear.  On the other hand, since our model is an exact solution of  Einstein equations,  we can expect to get faithful results. Based on our analysis we do not find any significant overall effect\\footnote{In agreement with the similar findings of~\\cite{bolejko,Sugiura}}. Quantities, such as the  matter density, on an average, still behave as in the homogeneous Einstein-de Sitter (EdS) universe (\\ie a matter dominated flat Universe). We remind the reader that if the universe starts to accelerate the matter density should dilute faster than EdS models.  The second issue ({\\bf II}) has been investigated in the past already in the '60s. The main point here is that the light that travels to us in a realistic Universe is not redshifted in the same way as in a FLRW metric, nor the angles and luminosities evolve in the same manner: it is more likely that a photon travels through almost empty space rather than meeting structures which are clumped in small sizes. However, typically work in the past has focused on what happens to the angles (or equivalently on luminosities), almost ignoring the problem of the redshift. Early pioneering work on this subject was due to Zeldovich~\\cite{zeldovichL} that  estimated the effect on the angular distance ($D_A$) on light travelling through an empty cone embedded in an FLRW metric. A similar  treatment based on~\\cite{dyerroeder} has been used recently to analyze data by the Supernova Cosmology Project Collaboration (see~\\cite{Perlmutter}, fig.8). While giving a quite relevant correction (for large matter fraction $\\Omega_m$), the effect does not allow to get rid of the cosmological constant. Therefore the collaboration~\\cite{Perlmutter} estimates to have control on the inhomogeneities. However such calculations are incomplete: the dynamics is completely ignored, and more importantly the redshift experienced by a light ray is assumed to be the usual FLRW result.  Instead in the Onion model we could keep track of the corrections to both the redshift and the luminosity distance and in fact we found both of them to be significant, once the inhomogeneities become large. We found that the correction in the redshift with respect to a homogeneous model is of the same order, and even larger, than  the correction in luminosity distance. Qualitatively one can understand where such corrections come from:  in the underdense regions the expansion is faster and the photons suffer a larger redshift than it would in the corresponding EdS model, while in the overdensities the expansion is balanced by the gravitational collapse and the photons experience milder redshift (or even a blue shift!). Here comes however one of the shortcomings of the Onion model: a radial light ray in the Onion model unavoidably meets underdense and overdense structures and they tend to average out the corrections.  In the real Universe, instead, where the light hardly encounters  any structure, one does not expect such cancellations to occur. Nonetheless, based on these considerations, we can try to go beyond the Onion model: since the photon in the real world is mostly passing through voids it gets redshifted  faster as the non-linearities increase with time, thereby potentially producing the effect of apparent acceleration. In this paper we only present an estimate of such an effect leaving a more detailed study for  future work.   Next we come to the third issue ({\\bf III}) of whether the position of the observer can have any important consequences. In general, in our analysis, we found that the corrections to redshift or luminosity distance is controlled by two quantities (in fact it is the product of the two): the amplitude of density fluctuations, $\\de\\rho/\\rho$ and the ratio $L/r_{\\mt{hor}}$, where $L$ is the wavelength of density fluctuations and $r_{\\mt{hor}}$ is the horizon distance. As we mentioned earlier, the length scale at which the density fluctuations become non-linear, \\ie $\\de\\rho/\\rho\\sim {\\cal O}(1)$, the latter ratio is only around a few percent. Therefore naively one may expect these corrections to be too  small to be relevant for supernova cosmology (as argued in several literature). However the issue is more subtle: near the observer when the redshift and distances are themselves small these corrections become significant as they do not proportionately decrease. Instead of the ratio $L/r_{\\mt{hor}}$, the relative correction is now governed by the ratio $L/\\de r$, where $\\de r$ is just the distance between the source and the observer. Thus in the first few oscillations, especially the first one, when $\\de r\\sim L$ we obtain significant corrections.  In our analysis we find that, if we want to explain the Hubble diagram solely by the effect coming from the first  oscillation, then  the location of the observer is of crucial importance\\footnote{This situation can change drastically if we are also able to account for the second effect, that of light mostly propagating through voids.}. Basically, what one needs to explain is the mismatch in the measurement of the local Hubble parameter (between redshifts $0.03<z<0.07$) and the luminosity distance observed for high redshift supernovae (typically between redshifts $0.4<z<1.5$). Although we can place the observer in any generic location, we find that the best choice, to account for the larger locally measured Hubble parameter, is if we  place the observer in an underdense region. Our main result is that we found that we require $3-4$ times larger local density contrast than the average value inferred by CMB measurements. With these assumptions, however, we find that the Onion model can reasonably fit the supernova data and in this sense  our model  corroborates some of the findings of  the  models based on local voids.  We further find that the Onion model can be consistent with other observations such as local density measurements, the first peak position of the CMB (which measures the curvature of the universe) and baryon acoustic oscillations.  The Onion model also offers us the opportunity to look at other relevant observations. For instance, since we have the freedom to choose the position of the observer we can in principle calculate quantities such as anisotropies in the measurement of Hubble constant which are essential to make a realistic assessment of the viability of any  model relying on  ``us'' being in a special position.  Finally, suppose that inhomogeneities do not give a large enough effect to mimick Dark Energy, we can still estimate what is a minimal effect due to structures that typically affects distance measurements in the real Universe, \\ie an universe constructed by just propagating the CMB fluctuations  to today. We find that a typical effect gives a correction of order $0.15$ apparent magnitudes to the Hubble diagram at all redshifts. We can also explain another source of uncertainty in supernovae: the observed nearby supernovae seem to have larger intrinsic scatter~\\cite{astier} which can easily be explained by the large oscillations that we get in the $D_L(z)$ curve close to the observer.  We briefly discuss such applications in this paper but mostly focus on the problem of dark energy.  The structure of the paper is as follows. In section~\\ref{model} we introduce our LTB model, and in section~\\ref{structure} we solve it and discuss structure formation. In section~\\ref{photon} we solve for light propagation in this background, both analytically and numerically. In section~\\ref{luminosity} we discuss how to find the luminosity distance, both analytically and numerically. In section~\\ref{apparent} we  first construct a setup that could mimick Dark Energy, and then ask how realistic is such a setup, discussing observations such as local measurements of matter densities, CMB acoustic peaks and baryon oscillations. We then go on to  discuss light propagation in the real Universe as opposed to the Onion model, and we estimate how large can the missing effect be. We also estimate  minimal uncertainties that inhomogeneities can lead to in the $D_L-z$ relationship. Finally in section~\\ref{conclusions} we draw our conclusions. Appendices contain many technical details: Appendix~\\ref{Esmall} contains a discussion of the analytical approximations that we employ, Appendix~\\ref{nearcentre} contains a brief discussion of what happens near the centre of the LTB model, Appendix~\\ref{tr} and~\\ref{app-redshift} contain the analytical derivation of (respectively) the time and the redshift along the photon trajectory and Appendix \\ref{OA} contains the full calculation of $D_L$ for an off-centre observer. In appendix \\ref{sachs} we show how our results for the off-centre angular distance is consistent with the optical scalar evolution equations for the angular distance derived by Sachs in 1961~\\cite{sachs}. ",
        "conclusions": "\\label{conclusions}  In this paper we have proposed an exact model, based on LTB metric with only dust, for structure formation applicable to a Universe homogeneous and isotropic on very large scale. The structures here look like concentric shells, and we could study their evolution from the early Universe to the present epoch.  In this background we have studied exactly the propagation of light along a radial trajectory, and we have computed the distance-redshift relation as measured by an observer sitting in a generic position.  For this purpose we have derived (for the first time) an expression for the luminosity distance in an LTB metric for an off-centre observer, we have solved the system numerically, and we have found accurate analytical approximations to the dynamical equations and to the light propagation equations. We have made several conclusions on the basis of our computations.  First we notice that the corrections become larger close to the observer (far away the corrections due to underdensities tend to cancel the corrections due to overdensities).  Then, we have investigated whether there is a setup that can mimick an accelerating $\\Lambda CDM$ cosmology. We have shown that this is possible under some special conditions: the observer has to be located around a minimum of the density contrast, and has to live in a large-scale underdensity with quite high density contrast: a typical example is $\\delta\\approx 0.3$ on a scale of roughly $400/h \\, {\\rm Mpc}$, while the typical value of $\\delta$ at that scale is around $\\delta\\simeq 0.05 - 0.1$.  Without such a large density fluctuation the situation is still very interesting: one can estimate a typical effect, which affects all measurements, on the $D_L-z$ plot of about $\\Delta m\\simeq 0.15$ apparent magnitudes. By considering fluctuations of smaller length scales, $\\sim 100/h$ Mpc, we can also explain the observation that low redshift supernovae should show more scatter than high redshift ones (as seen experimentally~\\cite{astier}).  Finally, and most crucially, we notice that making the model more realistic could uncover relevant effects: the fact that light in a radial trajectory in the Onion model has to go through an equal number of structures and voids leads, in fact, to peculiar cancellations in the $D_L-z$ plot. Without these cancellations (which do not happen in the real world) we have estimated that the effect goes in the direction of apparent acceleration and that it should be large when matter fluctuations become nonlinear. We leave a more detailed analysis of this effect to a subsequent paper~\\cite{BMN2}."
    },
    "0606/astro-ph0606253_arXiv.txt": {
        "abstract": "We report our analysis of a \\Chandra\\ X-ray observation of the rich globular cluster Terzan 5, in which we detect 50 sources to a limiting 1.0-6 keV X-ray luminosity of $3\\times10^{31}$ ergs s$^{-1}$ within the half-mass radius of the cluster.  Thirty-three of these have $L_X>10^{32}$ ergs s$^{-1}$, the largest number yet seen in any globular cluster.  In addition to the quiescent low-mass X-ray binary (LMXB, identified by Wijnands et al.), another 12 relatively soft sources may be quiescent LMXBs.  We compare the X-ray colors of the harder sources in Terzan 5 to the Galactic Center sources studied by Muno and collaborators, and find the Galactic Center sources to have harder X-ray colors, indicating a possible difference in the populations. We cannot clearly identify a metallicity dependence in the production of low-luminosity X-ray binaries in Galactic globular clusters, but a metallicity dependence of the form suggested by Jord\\'{a}n et al. for extragalactic LMXBs is consistent with our data. ",
        "introduction": "\\label{s:intro}  Globular clusters are highly efficient at producing X-ray binaries through dynamical interactions \\citep{Ivanova04}. For luminous low-mass X-ray binaries (LMXBs; for the purposes of this paper, all discussion of LMXBs refers to those containing accreting neutron stars), this has been known for many years, as their production rate per unit mass is $>100$ times that of the rest of the Galaxy \\citep{Clark75}. Only in the past few years has it been possible to study the populations of faint X-ray sources in the densest globular clusters in depth, due to the high spatial resolution of the \\Chandra\\ {\\it X-ray Observatory} and optical identifications by the {\\it Hubble Space Telescope} \\citep[see][ for a review]{Verbunt04}.  These low-luminosity X-ray sources include quiescent LMXBs (qLMXBs), identified by their previous outbursts or soft blackbody-like X-ray spectra \\citep{intZand01,Rutledge02a}, cataclysmic variables (CVs) generally identified by their blue, variable optical counterparts \\citep{Cool95}, chromospherically active main-sequence binaries (ABs) identified by their main- (or binary-) sequence, variable optical counterparts \\citep{Edmonds03a}, and millisecond pulsars (MSPs) identified by their spatial coincidence with radio timing positions \\citep{Grindlay01a}.  These lower-luminosity X-ray sources are also produced through dynamical interactions, as demonstrated by the correlation between the number of X-ray sources in a cluster and its ``collision number'' \\citep{Pooley03}, a measure of the cluster's stellar interaction rate. One of the clusters with the highest collision numbers is Terzan 5, a dense cluster located 8.7 kpc away, near the Galactic center \\citep{Cohn02,Heinke03b}. This cluster hosts a transiently luminous LMXB, EXO 1745-248, first detected through X-ray bursts in 1980 \\citep{Makishima81} and irregularly active since then \\citep{Wijnands05a}.  Terzan 5 also hosts at least 30 MSPs, including the fastest known \\citep{Ransom05, Hessels06}, the largest number yet discovered in any globular cluster.  Terzan 5 also has a high metallicity of [Fe/H]=-0.21 \\citep{Origlia04}.  The incidence of bright LMXBs in globular clusters has been clearly associated with increasing metallicity, but to date the effects of metallicity on faint X-ray binaries in globular clusters have not been studied.  \\citet{Grindlay87} identified an apparent trend for LMXBs to be more common in metal-rich globular clusters in the Milky Way, confirmed for the Milky Way and M31 by \\citet{Bellazzini95}.  \\citet{Kundu02} demonstrated that metal-rich clusters are 3 times more likely than metal-poor clusters to possess LMXBs in the elliptical NGC 4472.  This result has been confirmed for various early-type galaxies by \\citet{Maccarone03, Kundu03, Sarazin03} and \\citet{Jordan04}, the last offering a scaling for the likelihood of a cluster in M87 hosting an LMXB of $(Z/Z_{\\odot})^{0.33\\pm0.1}$.  Suggested explanations for this effect are a dependence of the cluster initial mass function on the metallicity \\citep{Grindlay87}, a change in the rate of tidal captures \\citep{Bellazzini95}, a change in the strength of stellar winds \\citep{Maccarone04a}, and a change in magnetic braking rates due to differences in convective zone depths \\citep{Ivanova06}.  Terzan 5 has been observed by \\Chandra\\ in 2000 (two closely spaced observations) and 2002. The 2000 observations caught EXO 1745-248 during a bright outburst \\citep{Heinke03b}. The high resolution of \\Chandra\\ allowed the detection of nine additional low-luminosity sources within the cluster, and some useful spectral information of the transient was recovered from the readout streak.  In the 2002 observation, EXO 1745-248 was observed at a typical X-ray luminosity for a qLMXB ($L_X=2\\times10^{33}$ ergs s$^{-1}$), but with an unusually hard X-ray spectrum \\citep{Wijnands05a}. \\citet{Wijnands05a} also detected a large number of faint X-ray sources in Terzan 5, which are the focus of this paper. ",
        "conclusions": "\\label{s:concl}  Terzan 5 contains 28 X-ray sources above $L_X=10^{32}$ ergs s$^{-1}$ (0.5-2.5 keV), the richest population of X-ray sources so far observed in a globular cluster in this $L_X$ range. Twelve sources show soft X-ray colors suggesting a qLMXB nature. However, these sources are not generally well-fit by a simple hydrogen-atmosphere model, indicating that if these sources are qLMXBs, they have a substantial flux from harder nonthermal spectral components \\citep[as seen in non-cluster systems, e.g. ][]{Rutledge01a}. Several faint X-ray sources have demonstrated substantial variability, up to a factor of five, between the 2000 and 2003 \\Chandra\\  observations.  We constructed an X-ray color-color diagram for the sources in Terzan 5, for comparison with X-ray sources of similar luminosity at the Galactic center.  We find that the X-ray colors of the Galactic center sources are substantially harder than the relatively hard X-ray sources in Terzan 5, including controlling for the differences in photoelectric absorption. This suggests an intrinsic difference between the sources in Terzan 5 and the Galactic center.  Our study of the distribution of X-ray sources among globular clusters finds that likely qLMXBs, and hard X-ray sources with $L_X>10^{32}$ ergs s$^{-1}$ (which may be dominated by bright CVs) show consistency with the parametrizations by density, core radius, and metallicity of \\citet{Verbunt87} and \\citet{Jordan04}.  However, the hard X-ray sources with $10^{31}<L_X<10^{32}$ ergs s$^{-1}$ show a lesser dependence on density and metallicity."
    },
    "0606/astro-ph0606315_arXiv.txt": {
        "abstract": "We consider the generation and evolution of quantum fluctuations of a massive nonsymmetric gravitational field ($B$-field) from inflationary epoch to matter era in the simplest variant of the nonsymmetric theory of gravitation (NGT), which consists of a gauge kinetic term and a mass term. We observe that quite generically a nonsymmetric metric field with mass, $m_B\\simeq  0.03(H_I/10^{13}~{\\rm GeV})^4~{\\rm eV}$, is a good dark matter candidate, where $H_I$ denotes the inflationary scale. The most prominent feature of this dark matter is a peak in power at a comoving momentum, $k\\sim \\sqrt{m_B H_0}/(1+z_{\\rm eq})^{1/4}$, where $z_{\\rm eq}$ is the redshift at equality. This scale corresponds roughly to the Earth-Sun distance. ",
        "introduction": "Einstein's theory of relativity has passed all experimental tests on laboratory and intermediate scales. On cosmological scales however, in order to get a successful description of the Universe's dynamics one requires addition of dark matter and dark energy, both of unknown composition. An alternative is to extend Einstein's theory and hope that a modified theory of gravitation would describe gravitation on cosmological scale without a need for dark matter and/or dark energy. Examples of extended theories of gravitation which have both been used as alternatives to dark matter~\\cite{Moffat:2004bm} are the nonsymmetric theory of gravitation (NGT)~\\cite{Moffat:1978tr} and MOND~\\cite{Milgrom:1983ca,Bekenstein:2004ne}.  In Ref.~\\cite{Prokopec:2005fb} we assumed that the antiymmetric tensor field $B$ is dynamical, and considered the cosmological evolution of quantum fluctuations generated in de Sitter inflation of a massless $B$-field, which is related by a duality transformation to the Kalb-Ramond axion field. We have further shown that the evolution of a massless $B$ field mimics that of a very light field. We then presented some early results on the evolution of a massive $B$-field, and observed that the scaling changes when the field becomes nonrelativistic in radiation era. From our analysis of quantum fluctuations of the physical mode of the antisymmetric field, which corresponds to the longitudinal `magnetic' component, it follows that the field couples conformally during de Sitter inflation, such that the $B$-field correlator exhibits the spectrum of conformal vacuum fluctuations. This is contrary to what has been claimed based on studies of the Kalb-Ramond axion, which couples to gravitation just like a massless minimally coupled scalar field, and therefore exhibits a nearly scale invariant spectrum during inflation. In fact, because of the conformal coupling during inflation, there cannot be a physical observable which exhibits a scale invariant spectrum during inflation or subsequent epochs. There is amplification however, which is induced by the matching at the inflation-radiation transition, and based on which the spectral amplitude gets enhanced at superhorizon scales with respect to the conformal vacuum, but not enough to get a scale invariant observable. Finally we observed that the spectral features of the $B$ field fluctuations produced in inflation are similar to that of gravitational waves, making it thus an alternative probe of inflationary scale. There is an important difference in the spectrum between the primordial gravitons and the $B$ field fluctuations: while the primordial gravitons exhibit a scale invariant spectrum on superhubble scales, the $B$-field fluctuations are suppressed.   Building on an earlier work of Damour, Deser and McCarthy~\\cite{Damour:1992bt} and Clayton~\\cite{Clayton:1996dz}, Janssen and Prokopec~\\cite{Janssen:2006tj} have recently shown that there is no nonsymmetric geometric theory which yields a dynamical $B$ field and which is ``problem free,'' in the sense that the $B$-field dynamics is ghost free and it is not marred by instabilities. Furthermore the authors of~\\cite{Janssen:2006tj} have shown that the most general (problem free) quadratic action in the $B$ field is of the form, \\begin{eqnarray} S &=& S_{\\rm HE} + S_{\\rm B} \\,,\\qquad S_{\\rm B} = \\int d^4 x {\\cal L}_{\\rm B} \\nonumber\\\\ {\\cal L}_{\\rm B} &=& - \\sqrt{-g}\\frac{1}{12}H_{\\mu\\nu\\rho} H^{\\mu\\nu\\rho} - \\sqrt{-g}\\frac{1}{4}\\Big(m_B^2 + \\xi{\\cal R}\\Big) B_{\\mu\\nu}B^{\\mu\\nu} \\,, \\label{general lagrangian} \\end{eqnarray} where \\begin{equation} H_{\\mu\\nu\\rho} = \\partial_\\mu B_{\\nu\\rho} + \\partial_\\nu B_{\\rho\\mu} + \\partial_\\rho B_{\\mu\\nu} \\label{field strength H} \\end{equation} is the antisymmetric field strength and \\begin{equation} S_{\\rm HE} = \\int d^4x {\\cal L}_{\\rm HE} \\,, \\qquad {\\cal L}_{\\rm HE} = -\\frac{1}{16\\pi G_N}\\Big({\\cal R}+2\\Lambda\\Big) \\label{HE action} \\end{equation} is the Hilbert-Einstein action, with $\\Lambda$ being the cosmological constant. In this paper we work in the units in which, $1/(16\\pi G_N)=1$. From Eq.~(\\ref{general lagrangian}) we see that the coupling to the Ricci scalar ${\\cal R}$ is the only allowed coupling to curvature.  In order to simplify the problem further, we drop the coupling to the Ricci scalar and set $\\xi=0$ in Eq.~(\\ref{general lagrangian}). This term is expected to act as an additional mass term during inflation era and thus introduce additional breaking of conformal invariance. During radiation and matter era however, where the Ricci scalar is either zero (radiation era) or of the order of the Hubble parameter (matter era), we expect that the Ricci scalar term has no significant impact on the evolution of the $B$-field.  In a geometric theory the mass term is naturally induced by the cosmological constant, in which case~\\cite{Janssen:2006tj}, \\begin{equation} m_B^2 \\simeq 2\\Lambda (1-2\\rho + 8\\sigma) \\,, \\label{mass vs Lambda} \\end{equation} where $\\rho$ and $\\sigma$ are defined by the following decomposition of the metric tensor, \\begin{equation} \\bar g_{\\mu\\nu} =   g_{\\mu\\nu} + B_{\\mu\\nu} + \\rho B_\\mu^{\\;\\alpha} B_{\\;\\alpha\\nu} + \\sigma B^2 g_{\\mu\\nu} + {\\cal O}(B_{\\mu\\nu}^3) \\,. \\label{metric decomposition} \\end{equation} This is a general metric tensor decomposition in terms of the symmetric metric tensor ($g_{\\mu\\nu}$) and antisymmetric metric tensor ($B_{\\mu\\nu}$). Yet there is no reason to assume that the mass term is fully of geometric origin, and thus the relation~(\\ref{mass vs Lambda}) needs not to hold in general.  Equation~(\\ref{mass vs Lambda}) is also significant because it implies that, if the $B$ field is of a geometric origin, then it would be unnatural to assume that its mass vanishes. Indeed the current observations, based on the luminosity-redshift relation of distant supernovae Ia~\\cite{Perlmutter:1998np,Riess:2004nr}, suggest that the cosmological term today is of the order, $\\Lambda \\sim 10^{-84}~{\\rm GeV}^2$. To maintain generality we assume in this work that the mass $m_B$ is unspecified and study its cosmological implications. In order to do that, we canonically quantise the massive $B$ field in inflation, and follow its subsequent dynamics in radiation and matter eras. Our main finding is that the $B$ field with a mass of the order, \\begin{equation} m_B \\simeq 0.03 \\Big(\\frac{10^{13}~{\\rm GeV}}{H_I}\\Big)^4~{\\rm eV} \\end{equation} which corresponds to a lengthscale, $m_B^{-1}\\simeq 7\\times10^{-8}\\big({H_I}/{10^{13}~{\\rm GeV}}\\big)^4~{\\rm m}$, is a good dark matter candidate.  Since the $B$ field is produced in inflation and does not couple to the matter fields, its spectrum is highly nonthermal. Indeed, we find that the spectral power is peaked at a comoving momentum, $k\\simeq \\sqrt{m_B H_0}/(1+z_{\\rm eq})^{1/4}$, where $z_{\\rm eq}$ is the redshift at matter-radiation equality, and corresponds to a physical scale at structure formation ($z\\sim 10$), $k_{\\rm phys}^{-1}\\sim 2\\times 10^7~{\\rm km}$. This peak is generated as a consequence of a different nature of the vacuum states in inflation and radiation era. This is the main feature by which this dark matter can be distinguished from other dark matter candidates, which typically obey a thermal statistic.  Another important feature are Fourier space pressure oscillations, which occur after the second Hubble crossing. Although we find that the pressure of the $B$-field drops to zero before the decoupling of the cosmic photon fluid, the Fourier pressure components exhibit significant oscillations. These oscillations may have a potentially observable impact on the gravitational potentials, and they are thus the second distinct feature of our dark matter candidate. The physical significance of these spectral pressure oscillations should be further investigated. ",
        "conclusions": "We have considered the cosmology of a massive antisymmetric tensor field whose dynamics is given by the action~(\\ref{general lagrangian}). We follow the evolution of the vacuum fluctuations generated during an inflationary epoch, through inflation, radiation and matter era. We find that the antisymmetric field with a mass, $m_B=3\\times10^{-2}\\left({10^{13} \\,\\mbox{GeV}}/{H_I}\\right)^4\\, \\mbox{eV}$, results in the right energy density today to account for the dark matter of the Universe. This then implies that below the corresponding length scale, $m_B^{-1} \\simeq 0.1\\,\\,{\\rm \\mu m} \\left({H_I}/{10^{13} \\,\\mbox{GeV}}\\right)^4$, the strength of the gravitational coupling may change, as it has been argued in Refs.~\\cite{Moffat:1995dq,Moffat:2004bm}. This scale is about two orders of magnitude below the present experimental bound, which is of the order $10 \\mu\\mbox{m}$~\\cite{Will:2005va}. Note that the mass scale~(\\ref{m_B}) depends strongly on the scale of inflation, such that, if the gravitational force law below the scale of $0.1\\,\\,{\\rm \\mu m}$ remains unchanged, would imply that, either the $B$ field is nondynamical, or the inflationary scale is lower than $10^{13}~{\\rm GeV}$. Note that, in contrast to gravitons, the antisymmetric tensor field begins scaling as non-relativistic matter during radiation era, making it potentially the most sensitive probe of the inflationary scale.  The peak in the energy density power spectrum, generated as a consequence of a breakdown of conformal invariance in radiation era, corresponds to a comoving momentum scale of the order, $k\\simeq \\sqrt{m_B H_0}/(1+z_{\\rm eq})^{1/4}$, where $z_{\\rm eq}\\simeq 3230$ is the redshift at matter-radiation equality. Hence the most prominent matter density perturbations at a redshift $z=10$, when structure formation begins, occur at a scale, $k_{\\rm phys}^{-1}\\sim 2\\times 10^{7}\\,{\\rm km}$ (corresponding to a wavelength, $\\lambda_{\\rm phys} \\sim 1\\times 10^{8}~{\\rm km}$, which is (co-)incidentally the Earth-Sun distance), which may boost early structure formation on these scales. We therefore expect that, when compared to other CDM models, this type of cold dark matter may induce an earlier structure formation. Even though the mass of the antisymmetric field is quite small (of the order of the heaviest neutrino mass, $m_{\\nu} \\sim 0.1\\,\\,{\\rm eV}$), the antisymmetric tensor field dark matter is neither hot nor warm. Indeed, since the antisymmetric tensor field does not couple to matter fields, it cannot thermalise and its spectrum remains primordial, and thus highly non-thermal. Hence, in spite of its small mass the field pressure remains small, such that structure formation on small scales does not get washed out. Indeed, $P/\\rho_B \\simeq E_{kin}/(m_B c^2) \\sim 10^{-34}$ evaluated at ${k_{phys}=(k_{peak}/a)|_{today}}$, which is a tiny pressure.  Although the pressure converges to zero as the field becomes dominated by non-relativistic modes, the pressure per mode in Fourier space shows a characteristic oscillatory behaviour shown in figure~\\ref{fig:pressure}. These oscillations may influence cosmological perturbations in a manner analogous to Sakharov oscillations~\\cite{Albrecht:1992kf}, which in turn may produce an observable imprint in the cosmic background photon fluid. This question deserves further investigation."
    },
    "0606/astro-ph0606033_arXiv.txt": {
        "abstract": "s{ It is well known that in manifestly Lorentz invariant theories with nontrivial kinetic terms, perturbations around some classical backgrounds can travel faster than light. These exotic \"supersonic\" models may have interesting consequences for cosmology and astrophysics. In particular, one can show \\cite{my} that in such theories the contribution of the gravitational waves to the CMB fluctuations can be significantly larger than that in standard inflationary models. This increase of the tensor-to-scalar perturbation ratio leads to a larger B-component of the CMB polarization, thus making the prospects for future detection much more promising. Interestingly, the spectral index of scalar perturbations and mass of the scalar field considered in the model are practically indistinguishable from the standard case. Whereas the energy scale of inflation and hence the reheating temperature can be much higher compared to a simple chaotic inflation.} ",
        "introduction": "One of the main consequences of inflation is the generation of primordial cosmological perturbations \\cite{mc} and the production of long wavelength gravitational waves (tensor perturbations)~\\cite{star}. The predicted slightly red-tilted spectrum of the scalar perturbations is at present in excellent agreement with the measurements of the CMB fluctuations \\cite{cmb}. The observation of primordial gravitational waves together with the detection of a small deviation of the spectrum from flat would give us further strong confirmation of inflationary paradigm. The detection of primordial gravitational waves is not easy, but they can be seen indirectly in the B-mode of the CMB polarization (see, for example, \\cite{book}). In standard slow-roll inflationary scenarios \\cite{Chaot} the amplitude of the tensor perturbations can, in principle, be large enough to be observed. However, it is only on the border of detectability in future experiments.  There are a lot of inflationary scenarios where the tensor component produced during inflation is much less then that in the chaotic inflation. In particular, in models such as new inflation~\\cite{New} and hybrid inflation~\\cite{Hybrid}, tensor perturbations are typically small~\\cite{book}. Moreover, in the curvaton scenario \\cite{curva} and k-inflation \\cite{k-Inflation}, they can be suppressed completely.  A natural question is whether the gravitational waves can be significantly enhanced compared to standard scanarious. Recently it was argued that the contribution of tensor perturbations to the CMB anisotropy can be much greater than expected \\cite{Bartolo:2005jg,Sloth:2005yx}. However, it was found in \\cite{lms} that in the models considered in \\cite{Bartolo:2005jg,Sloth:2005yx} one cannot avoid the production of too large scalar perturbations and therefore they are in contradiction with observations.  In the paper \\cite{my} we introduced a class of inflationary models where the B-mode of polarization can exceed that predicted by simple chaotic inflation. These models resemble both k-inflation \\cite{k-Inflation} and chaotic inflation \\cite{Chaot}. Inflation occurs due to the potential term in the Lagrangian, and the kinetic term has a nontrivial structure responsible for the large sound speed of perturbations. In this talk I will review the model from \\cite{my}. ",
        "conclusions": "We have shown above that in theories where the Lagrangian is a nontrivial, nonlinear function of the kinetic term, the scale of inflation can be pushed to a very high energies without coming into conflict with observations. As a result, the amount of produced gravitational waves can be much larger than is usually expected. If such a situation were realized in nature then the prospects for the future detection of the B-mode of CMB polarization are greatly improved. Of course, the theories where this happens are very unusual. For example, the Cauchy problem is well-posed not for all initial data \\cite{rendall,adams}. Moreover, the horizons lose their universality \\cite{bmv,SS}. Therefore, future observations of the CMB fluctuations are extremely important since they will not only restrict the number of possible candidates for the inflaton but also shed light on the problem of the ``superluminal'' propagation."
    },
    "0606/astro-ph0606519_arXiv.txt": {
        "abstract": "We use semi-analytic techniques to study the formation and evolution of brightest cluster galaxies (BCGs). We show the extreme hierarchical nature of these objects and discuss the limitations of simple ways to capture their evolution. In a model where cooling flows are suppressed at late times by AGN activity, the stars of BCGs are formed very early (50 per cent at $z\\sim 5$, 80 per cent at $z\\sim3$) and in many small galaxies. The high star formation rates in these high-$z$ progenitors are fuelled by rapid cooling, not by merger-triggered starbursts. We find that model BCGs assemble surprisingly late: half their final mass is typically locked-up in a single galaxy after $z\\sim0.5$. Because most of the galaxies accreted onto BCGs have little gas content and red colours, late mergers do not change the apparent age of BCGs. It is this accumulation of a large number of old stellar populations -- driven mainly by the merging history of the dark matter halo itself -- that yields the observed homogeneity of BCG properties.  In the second part of the paper, we discuss the evolution of BCGs to high redshifts, from both observational and theoretical viewpoints. We show that our model BCGs are in qualitative agreement with high-$z$ observations. We discuss the hierarchical link between high-$z$ BCGs and their local counter-parts. We show that high-$z$ BCGs belong to the same population as the massive end of local BCG progenitors, although they are not in general the same galaxies. Similarly, high-$z$ BCGs end-up as massive galaxies in the local Universe, although only a fraction of them are actually BCGs of massive clusters. ",
        "introduction": "\\label{sec:intro}  Brightest cluster galaxies (BCGs) are the very bright galaxies that inhabit the cores of rich galaxy clusters. They are among the most luminous and most massive galaxies in the Universe at the present epoch.  At low redshift, these objects exhibit a small dispersion in their aperture luminosities (after correction for systematic effects and for dependences on galaxy structure and environment).  As a result, many early studies used BCGs as `standard candles' for classical cosmological tests \\citep{Sandage72,GunnOke75,HoesselSchneider85,PostmanLauer95}.  BCGs lie close to the peaks of the X--ray emission in concentrated X--ray bright clusters \\citep{JonesForman84,RheeLatour91} and their rest--frame velocities have small offsets from those of their host clusters \\citep*{QuintanaLawrie82,ZabludoffHuchraGeller90}.  They do not appear to be drawn from the same luminosity function as cluster ellipticals \\citep{TremaineRichstone77,Dressler78} or as bright galaxies in general \\citep{BernsteinBahvsar01}. They also exhibit different luminosity profiles than typical cluster elliptical galaxies \\citep{Oemler76,Schombert86}. All these observations suggest that the evolution of the BCGs might be appreciably distinct from galaxy evolution in general.  Early theoretical studies discussed the role of X--ray driven {\\it cooling flows} \\citep{Silk76,CowieBinney77,Fabian94} and {\\it cannibalism} due to dynamical friction \\citep{OstrikerTremaine75,White76, MalumuthRichstone84, Merritt85} in the formation of these special objects.  This early work was not however successful due to the use of a simplified cluster model.  In the now standard cold dark matter model of structure formation, we understand that local massive haloes (clusters) assembled rather late, through the merging of smaller systems. In this perspective, cooling flows are indeed the main fuel for galaxy mass-growth at high redshifts, in dense and lower mass haloes. This source is removed only at low redshifts and in group or cluster environments, possibly due to AGN feedback. In addition, galaxy-galaxy mergers are most efficient within small haloes with low velocity dispersions. They are indeed driven by dynamical friction, but it is the accretion rate of galaxies into the proto-cluster, along with cluster growth itself, that regulates and sets the conditions for galaxy merging.  Since these early studies, little theoretical work has focused on the formation and evolution of BCGs, and it was only in 1998 that the subject was revisited in a CDM hierarchical framework \\citep{Dubinski98, Aragon-SalamancaBaughKauffmann98}.  \\citet{Dubinski98} used a N--body simulation of a cluster of galaxies and showed that merging naturally produces a massive central galaxy with surface brightness and velocity dispersion that resemble those of BCGs.  \\citet*{Aragon-SalamancaBaughKauffmann98} analysed the K--band Hubble diagram of BCGs up to redshift $\\sim 1$ and interpreted the apparent lack of passive evolution as evidence for recent mass accretion by the BCGs.  Using a simple parametrisation, the authors estimated a mass growth by a factor $4-5$, a result that they found to be consistent with predictions from semi--analytic models of galaxy formation.  Later work \\citep{CollinsMann98,BurkeCollinsMann00,Brough02} questioned this agreement showing that it might reflect cluster selection.  \\citet{BurkeCollinsMann00} showed that the results of \\citet{Aragon-SalamancaBaughKauffmann98} were based on a sample of low X-ray luminosity (low--L$_{\\rm X}$) clusters.  For high--L$_{\\rm X}$ clusters, which better match the cluster mass selection used in the semi-analytic models, they found substantially lower rates of mass accretion, at most a factor of $\\sim 2$ since $z\\sim 1$.  In recent work, \\citet{Gao04a} used numerical simulations to study the formation of the inner cores of massive dark matter haloes.  Their results suggest that typical central galaxies have undergone a significant number of mergers since $z\\sim1$.  Observations do provide some examples of central galaxies in the act of merging \\citep{Nipoti03}, although not many. It is therefore still a matter of debate whether the large number of mergers predicted in a hierarchical scenario is in agreement with observational data.  In this paper we study the formation and evolution of BCGs using a combination of N--body simulation and semi-analytic techniques.  In these models, central galaxies are bound to be `special': gas that is shock heated to the virial temperature of dark matter haloes cools radiatively and condenses only onto the central galaxy.  An additional physical process, such as feedback from a central AGN, must be invoked to prevent excessive growth from such cooling flows. The haloes which BCGs inhabit originate from the gravitational collapse of the highest (and rarest) peaks of primordial density fluctuations.  In these overdense regions, haloes collapse earlier and merge more rapidly in comparison to regions of the Universe with {\\it average} density.  Taking advantage of the largest simulation of the cosmic structure carried out so far -- the {\\it Millennium Simulation} \\citep{Springel05}-- we select a large number of massive haloes at $z=0$ and study in detail how and when the stars that compose their central galaxies formed, and how the assembly of these objects relate to the assembly of the halo itself.  In the second part of the paper, we investigate if the low rate of mass accretion inferred from observations is in agreement with our model predictions.  We describe the simulation and the semi-analytic model used for this study in Secs.~\\ref{sec:sam1} and \\ref{sec:sam2}.  In Sec.~\\ref{sec:case}, we present a {\\it case--study} in order to define the terminology that we will use later on. In Sec.~\\ref{sec:stat} we present statistical results for $125$ BCGs selected from the simulation at $z=0$. In Sec.~\\ref{sec:obs}, we extend that selection to higher $z$ and we compare the luminosity evolution measured using high--L$_{\\rm X}$ clusters to that predicted by our model.  Finally, we discuss our results and give our conclusions in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this study, we have used a combination of high-resolution $N$-body simulations and semi-analytic techniques to investigate the formation and evolution of BCGs in a cosmological context. This allows us to explicitly take into account the full hierarchy of dark matter structure growth, and its impact on galaxy formation. We show that in a CDM Universe, this hierarchical build-up implies that central galaxies of massive haloes have complex merging histories that are not well described by monolithic approximations.  In the first part of our paper, we have studied in detail the formation and assembly histories of local BCGs. We find that the stars in these systems start forming very early (about $50$ per cent of the stars in these systems have already formed by redshift $\\sim 5$) although the {\\it identity} of the BCGs themselves is not defined before $z \\sim 0.7$.  A monolithic approximation clearly fails to describe these objects for most of their history.  Most of the stars are formed in separate entities in a quiescent mode and most of the mass in the local BCGs is gained through accretion of a relatively low number (typically $10$) of objects more massive than $10^{10}\\,h^{-1}{\\rm M}_\\odot$. Most of these accreted objects, with the exception of those accreted at very high redshift, have very low gas fractions and star formation rates, quite red colours and an early-type morphology.  Interestingly, the main branches of BCGs grow fast enough that very few major mergers occur along them, most of the accretion events being {\\it minor mergers}.  It is worth mentioning that our quantitative results do depend strongly on the implemented feedback models, both from AGNs and from supernovae. The AGN feedback model employed here \\citep{Croton06} is extremely efficient in switching off cooling in relatively massive haloes even at high redshifts. At the same time, the supernova feedback model we use \\citep[also from][]{Croton06} drives strong winds that blow out all the gas from galaxies on short timescales. The combination of these two processes causes the early shutdown of star formation in the progenitors of local BCGs. These two particular feedback implementations are among the most efficient proposed so far in the literature. The work presented here leaves room for less drastic alternatives.  In the second part of our work, we have studied the observable evolution of BCGs by selecting the $125$ most massive haloes at different redshifts.  From these samples of high-$z$ BCGs, we measure a mass growth of a factor of $\\sim 3$ from $z=1$ to $z=0$. This is {\\it by coincidence} very close to the growth we find following the main branch of local BCGs. These mass growth factors are higher than those inferred from observational data. However, when we use a selection that more closely matches that of high-L$_{\\rm {X}}$ clusters, we find milder evolution, consistent with the observations.  Interestingly we find that our model BCGs exhibit, at all redshifts, a weak correlation between stellar mass and stellar metallicity and an even weaker correlation between their magnitudes and their colours.  This changes slightly if a different supernova feedback model is employed. Our high-$z$ BCGs do end up as central galaxies of massive haloes at $z=0$ but, because of our cluster sample selection criterion, many of their descendants are not included in our local BCG sample.  However, we have shown that there is no substantive physical difference (1) between local BCGs and the descendents of high-$z$ BCGs, and (2) between high-$z$ BCGs and the most massive progenitors of local BCGs."
    },
    "0606/astro-ph0606205_arXiv.txt": {
        "abstract": "In this paper we explore the dynamics of ionization bounded planetary nebulae after the termination of the fast stellar wind. When the stellar wind becomes negligible, the hot, shocked bubble depressurizes and the thermal pressure of the photoionized region, at the inner edge of the swept-up shell, becomes dominant. At this stage the shell tends to fragment creating clumps with  comet-like tails and long, photoionized trails in between, while the photoionized material expands back towards the central stars as a rarefaction wave. Once that the photoionized gas fills the inner cavity, it develops a kinematical pattern of increasing velocity from the center outwards with a typical range of velocities starting from the systemic velocity to $\\sim 50 \\kms $ at the edges. The Helix nebula is a clear example of a planetary nebula at this late evolutionary stage. ",
        "introduction": "High spatial resolution imagery in recent years have revealed a number of well developed planetary nebulae that exhibit a fragmented appearance of their shells with clumps and radial spokes, as it is the case of the famous Helix Nebula, NGC 7293 (Vorontsov-Velyaminov 1968; Meaburn et al. 1992; O'Dell \\& Handron 1996). Similar examples are the Dumbbell nebula, NGC 6853 (Manchado et al.1996; Meaburn et al. 2005a) and NGC 2867 (Corradi et al. 2003). Rayleigh-Taylor instabilities have been some times invoked as the cause of this fragmentation but with limited success (Burkert \\& O'Dell 1996). Alternatively, the effects of the ionizing radiation on the nebula have been shown by Capriotti (1973) and also by Garc\\'{\\i}a-Segura \\& Franco (1996) and Garc\\'{\\i}a-Segura et al.(1999) to be much more relevant in the fragmentation process of  the nebular shell.  Recently, Meaburn et al. (2005b) have shown that the current kinematics of nebulae such as NGC 7293 cannot be explained by the classic ``two-wind model'' (Kwok 1982). Since NGC 7293 shows emission in \\Heii and \\Oiii at very low velocities close to the central star, Meaburn et al. (2005b) concluded that only after the fast wind has switched off could this global velocity structure be generated. The absence of a fast stellar wind in NGC 7293 has been confirmed from observations with the IUE satellite (Perinotto 1983; Cerruti-Sola \\& Perinotto 1985) that failed to detect P-Cygni profiles from their central stars; likewise in the cases of NGC 6853 and NGC 2867. Furthermore, the lack of extended, soft x-ray emission from NGC 7293 (Guerrero et al. 2001) discards the presence of a \"hot bubble\" and the  shocked, fast wind. In general,  Cerruti-Sola \\& Perinotto (1985) found that the hotter the central stars are the least likely to have fast stellar winds, a result indicative of the rapid fading of the stellar wind as the central star evolves.   In this letter we present  numerical calculations which show that the decay of the fast stellar winds produces self-consistently the fragmented structures and kinematics observed in evolved planetary nebulae. ",
        "conclusions": ""
    },
    "0606/astro-ph0606496_arXiv.txt": {
        "abstract": "The convective period leading up to a Type~Ia supernova (SN~Ia) explosion is characterized by very low Mach number flows, requiring hydrodynamical methods well-suited to long-time integration.  We continue the development of the low Mach number equation set for stellar scale flows by incorporating the effects of heat release due to external sources. Low Mach number hydrodynamics equations with a time-dependent background state are derived, and a numerical method based on the approximate projection formalism is presented.  We demonstrate through validation with a fully compressible hydrodynamics code that this low Mach number model accurately captures the expansion of the stellar atmosphere as well as the local dynamics due to external heat sources. This algorithm provides the basis for an efficient simulation tool for studying the ignition of SNe~Ia. ",
        "introduction": "Modeling the period of convection leading up to the ignition of Type~Ia supernovae (SNe~Ia) is critical to determining the distribution of hot spots that seed the subsequent explosion.  Multidimensional simulations of SNe~Ia explosions presently seed one or more hot spots at or near the center of a white dwarf and use a flame model to describe the subsequent evolution (see for example \\citealt{roepke2005,gamezo2005,plewa:2004}).  Variations in the size, number, and distribution of these seeds can lead to large differences in the explosion outcome \\citep{niemeyer:1996,garciasenz:2005,livne:2005}.  The convective flows leading up to ignition have Mach numbers of 0.01 or less, with temperature perturbations of only a few percent \\citep{woosley2001,Woosley:2004}---conditions that are extremely challenging for fully compressible codes.  One-dimensional statistical methods \\citep{wunschwoosley:2004} predict that off-center ignition is likely, but they cannot give information about the distribution of the hot spots.  Only recently has progress been made in multidimensional modeling of convection in white dwarfs.  The first such calculations \\citep{hoflichstein:2002} evolved a two-dimensional wedge of the star for a few hours using an implicit hydrodynamics algorithm.  Convective velocities of approximately 100~km~s$^{-1}$ developed.  They observed compression near the center of the star leading to slightly off-center ignition.  Three-dimensional anelastic calculations \\citep{kuhlen:2005}, showed a large-scale dipole flow dominating the evolution, leading to an off-center ignition.  These simulations used a spectral decomposition, with a small portion of the center of the star removed due to the coordinate singularity at $r=0$.  Neither of these calculations operated at a Reynolds number large enough to see fully developed turbulence.  Further three-dimensional studies are needed to see how robust this dipole flow is to rotation and convection at higher Reynolds and Rayleigh numbers.  The goal of the present work is the development of a new multidimensional hydrodynamics algorithm capable of evolving the full star from the convective phase, though ignition and into the early stages of flame propagation.  Long time integration is critical.  As we showed previously (\\citealt{ABRZ:I}---henceforth paper~I), the low Mach number hydrodynamics equations provide an accurate description of flows with Mach numbers less than 0.2.  By filtering out sound waves, the low Mach number approximation allows for much larger time steps ($\\sim 1/M$ larger) than corresponding compressible codes.  In contrast to the anelastic equation set, the low Mach number equations are capable of modeling flows with finite-amplitude density and temperature perturbations.  The only restriction is that the pressure perturbation be small.  Furthermore, because the compressibility effects due to both the background stratification and local heat release are included, the low Mach number equations set can self-consistently evolve the expansion of a hydrostatic atmosphere due to heat release \\citep{almgren:2000}.  In this paper, we continue the development of the low Mach number hydrodynamics algorithm to include the effects of heat release.  An energy equation is added, and an earlier assumption from paper~I is relaxed, now allowing the background state to vary in time. In \\S~\\ref{sec:lmn} we develop the low Mach number equation set.  In \\S~\\ref{sec:numer} the numerical methodology is explained. Comparisons to fully compressible calculations are provided in \\S~\\ref{sec:results} to demonstrate the accuracy and utility of our new algorithm.  We conclude in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have introduced a new algorithm for evolving low Mach number flows in the presence of local and large-scale heating.  By contrast with the previous low Mach number model, this new model allows time variation of the base state in order to account for atmospheric expansion due to large-scale heat sources.  The time evolution of the base state must be calculated at each time step in addition to the local dynamics.  Numerical comparisons of low Mach number simulations with simulations using a fully compressible code demonstrate that the low Mach number algorithm with a time-dependent base state can accurately capture the hydrostatic adjustment of an atmosphere as well as local dynamics in response to large- and small-scale heat release.  Our long-term goal is to develop the capability for full star simulation using the new low Mach number approach.  The fundamental low Mach number approach has been validated with a number of simplified test cases, but further development is necessary to begin to perform detailed physics investigations of ignition and other problems of interest.  This development will include extension to three dimensions with adaptive griding, radial representation of gravity and the base state within the three-dimensional setting, non-constant gravity, and the calculation of internal heating due to reaction networks.  The tests presented above are quite demanding, and provided a challenge to both the low Mach number and the compressible solvers. The energy generation and resulting temperature/density contrasts during the convective phase of an SN Ia are much smaller.  Based on the agreement demonstrated on these difficult tests, we are confident that the low Mach number hydrodynamics method will be a useful and efficient tool in exploring the problem of SNe~Ia ignition. In addition, this algorithm is applicable to a wide range of problems outside of our target application (SNe Ia), including Type I X-ray bursts, classical novae, and convection in stars."
    },
    "0606/astro-ph0606175_arXiv.txt": {
        "abstract": "The James Webb Space Telescope (JWST) is a large (6.6m), cold ($<$50K), infrared-optimized space observatory that will be launched early in the next decade into orbit around the second Earth-Sun Lagrange point. The observatory will have four instruments: a near-infrared camera, a near-infrared multi-object spectrograph, and a tunable filter imager will cover the wavelength range, 0.6 $<$ $\\lambda$ $<$ 5.0 $\\mu$m, while the mid-infrared instrument will do both imaging and spectroscopy from 5.0 $<$ $\\lambda$ $<$ 29 $\\mu$m.  The JWST science goals are divided into four themes. The key objective of The End of the Dark Ages: First Light and Reionization theme is to identify the first luminous sources to form and to determine the ionization history of the early universe. The key objective of The Assembly of Galaxies theme is to determine how galaxies and the dark matter, gas, stars, metals, morphological structures, and active nuclei within them evolved from the epoch of reionization to the present day. The key objective of The Birth of Stars and Protoplanetary Systems theme is to unravel the birth and early evolution of stars, from infall on to dust-enshrouded protostars to the genesis of planetary systems. The key objective of the Planetary Systems and the Origins of Life theme is to determine the physical and chemical properties of planetary systems including our own, and investigate the potential for the origins of life in those systems. Within these themes and objectives we have derived representative astronomical observations.  \\noindent{\\bf Abstract, cont.} To enable these observations, JWST consists of a telescope, an instrument package, a spacecraft and a sunshield. The telescope consists of 18 beryllium segments, some of which are deployed. The segments will be brought into optical alignment on-orbit through a process of periodic wavefront sensing and control. The instrument package contains the four science instruments and a fine guidance sensor. The spacecraft provides pointing, orbit maintenance and communications. The sunshield provides passive thermal control. The JWST operations plan is based on that used for previous space observatories, and the majority of JWST observing time will be allocated to the international astronomical community through annual peer-reviewed proposal opportunities. ",
        "introduction": "The James Webb Space Telescope (JWST; table~\\ref{tabacro}) will be a large, cold, infrared-optimized space telescope designed to enable fundamental breakthroughs in our understanding of the formation and evolution of galaxies, stars, and planetary systems. It is a project led by the United States National Aeronautics and Space Administration (NASA), with major contributions from the European and Canadian Space Agencies (ESA and CSA). It will have an approximately 6.6 m diameter aperture, will be passively cooled to below 50 K, and will carry 4 scientific instruments: a Near Infrared Camera (NIRCam), a Near Infrared Spectrograph (NIRSpec), a near-infrared Tunable Filter Imager (TFI), and a Mid Infrared Instrument (MIRI). It is planned for launch early in the next decade on an Ariane 5 rocket to a deep space orbit around the Sun-Earth Lagrange point L2, about 1.5 $\\times$ 10$^{6}$ km from Earth. The spacecraft will carry enough fuel for a 10-year mission.  In this paper, we describe the scientific capabilities and planned implementation of JWST. The scientific planning is based on current theoretical understanding, interpretation of Hubble Space Telescope (HST) and Spitzer Space Telescope observations (e.g., Werner et al. 2004 and additional papers in that volume), and results from studies using ground-based and other space-based facilities. Many classes of targets for JWST have never been observed before, so the scientific plans are necessarily derived from theoretical predictions and extrapolations from known objects. In these cases, we outline the basis of the predictions and describe the observatory capabilities needed to verify them.  Additional results from HST and Spitzer, and other advances in theory and observation will continue to refine the observational plans for JWST. The scientific investigations we describe here define the measurement capabilities of the telescope, but they do not imply that those particular observations will be made. In this paper we do not list all potential applications of JWST. Instead, the scientific programs we discuss here are used to determine the key parameters of the mission such as collecting area, spatial and spectral resolution, wavelength coverage, etc. A mission which provides these capabilities will support a wide variety of astrophysical investigations. JWST is a facility-class mission, so most of the observing time will be allocated to investigators from the international astronomical community through competitively-selected proposals.  The plans for JWST reflect scientific and engineering discussions, studies, and development over the last 16 years. At a workshop held in 1989, the astronomical community began discussions of a scientific successor to HST (B\\'ely, Burrows \\& Illingworth, 1989). In the mid-1990s, the ``HST and Beyond'' committee recommended that NASA build an infrared-optimized telescope to extend HST discoveries to higher redshift and longer wavelength (Dressler 1996). Initial studies of the mission, then called the Next Generation Space Telescope, were reported by Stockman et al. (1997). Its scientific program was given top priority by the National Academy of Sciences survey ``Astronomy and Astrophysics in the New Millennium'' (McKee \\& Taylor 2001). In 2002, the Next Generation Space Telescope was renamed the James Webb Space Telescope in honor of the Administrator of NASA during the Apollo era (Lambright 1995). This paper provides an update to these previous studies.  The capabilities and performance specifications for JWST given here are preliminary. The mission is currently in its detailed design phase and has not yet been given official permission to proceed beyond that stage. NASA missions receive authority, and the accompanying budget, to proceed to launch only after they pass a nonadvocate review and transition from detailed design into development phases. Hence, the specifications given in this paper are subject to change based upon technical progress of individual elements of the observatory and available budgets. JWST is expected to receive final approval and make the transition to the development phase in 2008.  The scientific objectives of JWST fall into four broad themes: The End of the Dark Ages: First Light and Reionization; The Assembly of Galaxies; The Birth of Stars and Protoplanetary Systems; and Planetary Systems and the Origins of Life.  We organize this paper as follows. The remainder of this section provides a scientific introduction to these themes. In sections 2, 3, 4 and 5, we expand on the themes, and describe the scientific capabilities JWST will use to address them. In section 6 we describe the planned implementation of JWST.  \\subsection{The End of the Dark Ages: First Light and Reionization}  Theory and observation have given us a simple picture of the early universe. The Big Bang produced (in decreasing order of present mass-energy density): dark energy (the cosmic acceleration force), dark matter, hydrogen, helium, cosmic microwave and neutrino background radiation, and trace quantities of lithium, beryllium, and boron. As the universe expanded and cooled, hydrogen molecules formed, enabling the formation of the first individual stars, at about 180 million years after the Big Bang (Barkana \\& Loeb, 2001). According to theory, the first stars were 30 to 300 times as massive as the Sun and millions of times as bright, burning for only a few million years before meeting a violent end (Bromm \\& Larson 2004). Each one produced either a core-collapse supernova (type II) or a black hole. The supernovae enriched the surrounding gas with the chemical elements produced in their interiors, and future generations of stars contained these heavier elements. The black holes started to swallow gas and other stars to become mini-quasars, which grew and merged to become the huge black holes now found at the centers of nearly all massive galaxies (Magorrian et al. 1998). The supernovae and the mini-quasars could be individually observable by JWST.  Some time after the appearance of the first sources of light, hydrogen in the intergalactic medium was reionized. Results from the Wilkinson Microwave Anisotropy Probe (WMAP; Kogut et al. 2003; Page et al. 2006; Spergel et al. 2006) combined with data on quasars at z \\ensuremath{\\sim} 6 from the Sloan Digital Sky Survey (SDSS; Fan et al. 2002) show that this reionization had a complex history (Cen 2003b). Although there are indications that galaxies produced the majority of the ultraviolet radiation which caused the reionization, the contribution of quasars could be significant.  JWST will address several key questions in this theme: What are the first galaxies? When and how did reionization occur? What sources caused reionization? JWST will conduct ultra-deep near-infrared surveys with spectroscopic and mid-infrared follow-up to find and identify the first galaxies to form in the early universe. It will determine the processes that caused reionization through spectroscopy of high-redshift quasars or galaxies, studies of the properties of galaxies during that epoch.  \\subsection{The Assembly of Galaxies.}  Theory predicts that galaxies are assembled through a process of the hierarchical merging of dark matter concentrations (e.g., White \\& Frenk 1991; Cole et al. 1994). Small objects formed first, and were drawn together to form larger ones. This dynamical build-up of massive systems is accompanied by chemical evolution, as the gas (and dust) within the galaxies are processed through successive generations of stars. The interaction of these luminous components with the invisible dark matter produces the diverse properties of present-day galaxies, organized into the Hubble Sequence of galaxies. This galaxy assembly process is still occurring today, as the Magellanic Clouds fall into the Milky Way, and as the Andromeda Nebula heads toward the Milky Way for a possible future collision. To date, galaxies have been observed back to times about one billion years after the Big Bang. While most of these early galaxies are smaller and more irregular than present-day galaxies, some early galaxies are very similar to those seen nearby today.  Despite all the work done to date, many questions are still open. We do not really know how galaxies are formed, what controls their shapes, and what makes them form stars. We do not know how the chemical elements are generated and redistributed through the galaxies, and whether the central black holes exert great influence over the galaxies. We do not know the global effects of violent events as small and large parts join together in collisions.  JWST will address several key questions in this theme: When and how did the Hubble Sequence form? How did the heavy elements form? What physical processes determine galaxy properties? What are the roles of starbursts and black holes in galaxy evolution? To answer these questions, JWST will observe galaxies back to their earliest precursors, so that we can understand their growth and evolution. JWST will conduct deep-wide imaging and spectroscopic surveys of thousands of galaxies to study morphology, composition and the effects of environment. It will conduct detailed studies of individual ultra-luminous infrared galaxies (ULIRGs) and active galactic nuclei (AGN) to investigate what powers these energetic sources.  \\subsection{The Birth of Stars and Protoplanetary Systems.}  While stars have been the main topic of astronomy for thousands of years, only in recent times have we begun to understand them with detailed observations and computer simulations. A hundred years ago, we did not know that stars are powered by nuclear fusion, and 50 years ago we did not know that stars are continually being formed. We still do not know the details of how stars are formed from clouds of gas and dust, why most stars form in groups, or how planetary systems form. Young stars within a star-forming region interact with each other chemically, dynamically and radiatively in complex ways. The details of how they evolve and liberate the ``metals'' back into space for recycling into new generations of stars and planets remains to be determined through a combination of observation and theory.  We also know that a substantial fraction of stars, solar-type and later, have gas-giant planets, although the discovery of large numbers of these planets in very close orbits around their stars was a surprise. The development of a full theory of planet formation requires substantially more observational input, including observations of young circumstellar disks and older debris disks in which the presence of planets can be traced.  JWST will address several key questions in this theme: How do protostellar clouds collapse? How does environment affect star formation and vice versa? What is the initial mass function of stars at sub-stellar masses? How do proto-planetary systems form? How do gas and dust coalesce to form planetary systems? JWST will observe stars at all phases of their evolution, from infall onto dust-enshrouded protostars, through the formation of planetary systems, penetrating the dust to determine the physical processes that produce stars, planets and debris disks.  \\subsection{Planetary Systems and the Origins of Life.}  Understanding the origin of the Earth and its ability to support life is an important objective for astronomy. Key parts of the story include understanding the formation of planetesimals, and how they combine to form larger objects. We do not know how planets reach their present orbits, and how the large planets affect the smaller ones in solar systems like our own. We want to learn about the chemical and physical history of the small and large objects that formed the Earth and delivered the necessary chemical precursors for life. The cool objects and dust in the outer Solar System are evidence of conditions in the early Solar System, and are directly comparable to cool objects and dust observed around other stars.  JWST will address several key questions in this theme: How do planets form? How are circumstellar disks like our Solar System? How are habitable zones established? JWST will determine the physical and chemical properties of planetary systems including our own, and investigate the potential for the origins of life in those systems. JWST will use coronagraphy to investigate exosolar planets and debris disks, and will compare these observations to objects within our own Solar System. ",
        "conclusions": "The JWST science requirements are divided into four themes. The key objective of The End of the Dark Ages: First Light and Reionization theme is to identify the first luminous sources to form and to determine the ionization history of the early universe. The key objective of The Assembly of Galaxies theme is to determine how galaxies and the dark matter, gas, stars, metals, morphological structures, and active nuclei within them evolved from the epoch of reionization to the present day. The key objective of The Birth of Stars and Protoplanetary Systems theme is to unravel the birth and early evolution of stars, from infall on to dust-enshrouded protostars to the genesis of planetary systems. The key objective of the Planetary Systems and the Origins of Life theme is to determine the physical and chemical properties of planetary systems including our own, and investigate the potential for the origins of life in those systems. Within these themes and objectives, we have derived representative astronomical observations.  To enable these observations, JWST consists of a telescope, an instrument package, a spacecraft and a sunshield. The telescope consists of 18 beryllium segments, some of which are deployed after launch. The segments will be brought into optical alignment on-orbit through a process of periodic wavefront sensing and control. The instrument package contains the four science instruments and a fine guidance sensor. The spacecraft provides pointing, orbit maintenance and communications. The sunshield provides passive thermal control. The JWST operations plan is based on that used for previous space observatories, and the majority of JWST observing time will be allocated to the international astronomical community through annual peer-reviewed proposal opportunities.  In this paper we have described the astronomical observations that JWST is designed to make and the implementation that will enable those observations. We have provided scientific justifications for these observations within 4 science themes. The themes are chosen to span the range of science that we expect JWST will do, and the observations are chosen to be representative of that science. We wish to emphasize, however, that these observations will not necessarily be done. The majority of observing time on JWST will be allocated to the worldwide astronomical community through competitive selection of peer-reviewed proposals. Opportunities for proposals will begin one or two years before launch, and will continue annually for the lifetime of the mission. JWST will be a highly capable general-purpose observatory able to address a very wide range of scientific investigations. It represents a major contribution to scientific progress by the governments of the United States of America, of the European nations and of Canada. Regular, competitive peer-reviewed proposal selection will ensure that this international resource will address the most relevant and strongly justified scientific questions, and will leave a legacy of knowledge and discovery for future generations.  \\clearpage"
    },
    "0606/astro-ph0606343_arXiv.txt": {
        "abstract": "The current measurements of the cosmic ray energy spectrum at ultra-high energies ($\\text{E}>10^{19}$\\, eV) are characterized by large systematic errors and poor statistics.  In addition, the experimental results of the two experiments with the largest published data sets, AGASA and HiRes, appear to be inconsistent with each other, with AGASA seeing an unabated continuation of the energy spectrum even at energies beyond the GZK cutoff energy at $10^{19.6}$\\,eV.  Given the importance of the related astrophysical questions regarding the unknown origin of these highly energetic particles, it is crucial that the extent to which these measurements disagree be well understood.  Here we evaluate the consistency of the two measurements for the first time with a model-independent method that accounts for the large statistical and systematic errors of current measurements.  We further compare the AGASA and HiRes spectra with the recently presented Auger spectrum.  The method directly compares two measurements, bypassing the introduction of theoretical models for the shape of the energy spectrum.  The inconsistency between the observations is expressed in terms of a Bayes Factor, a standard statistic defined as the ratio of a separate parent source hypothesis to a single parent source hypothesis. Application to the data shows that the two-parent hypothesis is disfavored.  We expand the method to allow comparisons between an experimental flux and that predicted by any model. ",
        "introduction": "By measuring the energy spectrum of cosmic rays above $10^{19}$\\,eV we hope to shed some light on the yet unknown sites and acceleration mechanisms that can produce a particle of such energies.  These measurements are intrinsically difficult, as the cosmic ray flux at ultra-high energies is low and we rely on earthbound detectors which cannot directly detect the primary particle.  Rather, its properties are reconstructed by measuring the secondary particles of the extensive air shower produced when the primary enters Earth's atmosphere. Low statistics and large systematic errors in the energy determination are typical for these measurements, which should consequently be treated with care.  The existing measurements of the energy spectrum above $10^{19}$\\,eV have received much attention.  A longstanding hypothesis predicts that the flux of the highest energy cosmic rays should be suppressed above $10^{19.6}$\\,eV as cosmic rays from distant sources will interact with the cosmic microwave background via photopion production until their energy drops below this threshold energy \\cite{greisen,zatsepin}.  This so-called GZK suppression is in itself not a controversial prediction, but it has not been detected unambiguously, and several cosmic rays with substantially higher energies have been observed with various detectors over the years.  The most recent results disfavoring the GZK hypothesis come from the Akeno Giant Air Shower Array (AGASA) collaboration, whose published energy spectrum shows no indication of a high-energy suppression \\cite{sasaki,agasa}.  In contrast, the monocular-mode energy spectra measured by the High Resolution Fly's Eye detectors (HiRes 1 and HiRes 2) support the existence of a GZK feature \\cite{hiresI,hiresII}.  The current world data set is still small.  The AGASA claim that the GZK supression is not observed is based on only 11 events above $10^{20}$\\,eV.  In the near future, the Pierre Auger Observatory, currently under construction in Malargue, Argentina, will dramatically increase the world data set. The Auger collaboration has published a first energy spectrum based on 1.5 years of data taken during construction \\cite{auger1}.  A second problem is that although typically not shown in plots of the energy spectrum, the errors on the energy determination are also large: 30\\,\\% at $3\\times 10^{19}$\\,eV and 25\\,\\% above $10^{20}$\\,eV in AGASA \\cite{agasa}, 30\\,\\% in HiRes \\cite{hiresI,hiresII} and 50\\,\\% in Auger \\cite{auger1}. Furthermore, a correct evaluation of the systematic errors is difficult. It should for example be noted that the systematic uncertainty for the AGASA energy scale includes a 10\\,\\% systematic uncertainty due to the hadronic model.  This uncertainty is calculated by comparing different hadronic event generators and defining the systematic uncertainty by their difference. This may not be a reliable estimate of this uncertainty, which could potentially be much larger and is, in any case, unknown at this point. Since the AGASA and HiRes experiments use radically different techniques to determine the primary energy -- AGASA is a ground array and HiRes an air fluorescence detector -- the errors have fundamentally different sources. \\footnote{It should be noted that the published Auger spectrum has an energy scale that is calibrated with its fluorescence detector.  If simulations are used to determine the shower energies from surface detector data alone, the energies are systematically higher by at least 25% \\cite{auger1}.}  The large statistical and systematic errors quoted by the experiments raise the question how much significance should be attached to the discrepancy between the HiRes and AGASA spectra.  Several authors have addressed this question. In \\cite{demarco1,demarco2}, the experimental results are compared individually to model predictions of the shape of the energy spectrum at GZK energies.  Such an analysis requires assumptions on the nature of the cosmic ray sources, as the shape of the GZK suppression will of course depend on the source distribution.  Before one asks if any of the current measurements can be used as evidence for or against a specific model for the origin of cosmic rays, one needs to answer the question of whether or not the two experiments actually disagree at all, given their uncertainties.  This requires a method to compare the two spectra that is both model-independent and considers the uncertainties quoted by the experiments by accounting for the probability of any systematic shift in energy scale applied to the data.  This paper presents a general spectrum comparison technique that for the first time addresses these points. The technique naturally incorporates the Poisson errors in the observed fluxes, and can probabilistically treat systematic errors in the absolute energy scale.  The technique is developed in Section 2 and applied to the energy spectra of AGASA \\cite{agasa}, HiRes 1 and HiRes 2 in monocular mode \\cite{hiresI,hiresII} and Auger \\cite{auger2} in Section 3.  Section 4 summarizes the results and describes how the method can be expanded for testing all experimental measurements against a theoretical prediction.  \\begin{figure}[t] \\includegraphics[width=0.7\\textwidth]{Plot_0.eps} \\caption{\\label{fig:plot_0}  The AGASA \\cite{agasa}, Auger \\cite{auger2} and HiRes monocular spectra \\cite{hiresI,hiresII}.  } \\end{figure} ",
        "conclusions": "\\label{discussion} This result is similar to those obtained for the AGASA and HiRes spectra in \\cite{demarco1,demarco2} by a rather different method.  Our method directly compares two experimental results rather than comparing each to a theoretical prediction.  It also treats systematic errors correctly by weighting possible systematic shifts in the energy distribution by the probability of that shift.  This result does not say anything about what the specific (in)consistencies would be of either experiment given a theory. However, our results indicate that there is room for a theory that agrees with each pair of spectra. As the statistical power of the world data improves, the methods discussed here will provide a convenient framework to compare experimental results with one another and with theoretical models.  The Auger experiment is currently the only operating ultrahigh energy cosmic ray detector.  The Auger data set will increase dramatically over the next years. Figure\\,\\ref{fig:projected} shows the projected value of the $BF$ between Auger and the other experiments as a function of the fraction of the current Auger exposure. The $BF$ is calculated assuming the current spectra, the existing 30\\,\\% energy uncertainties in AGASA and HiRes and the projected 15\\,\\% energy uncertainties in the Auger spectrum. In the limit where the exposure for Auger becomes large, the $BF$ decreases (increasing support for the one-source hypothesis). That is, the statistical and energy scale uncertainties in HiRes and AGASA are so large that, barring a significant change in the Auger spectrum, one will never be able to tell whether or not the Auger and AGASA (HiRes) spectra are derived from different parent distributions provided that we limit our analysis to events above $10^{19.6}$\\,eV. \\begin{figure}[t] \\includegraphics[width=0.7\\textwidth]{projected.eps} \\caption{\\label{fig:projected}  The projected $BF$ between the current Auger spectrum and the AGASA and HiRes spectra as a function of the fraction of the current Auger exposure. The calculations assume 30\\,\\% energy uncertainties in AGASA and HiRes and the projected 15\\,\\% energy uncertainties in the Auger spectrum. } \\end{figure} Note, however, that if one fits a model to one of the experimental spectra, our result does not imply consistency of the model with the other experiment. \\footnote{In order to compare a model with both experiments, one must compare the model to both data sets simultaneously. The above method can be modified to provide a statistic analogous to a $\\chi ^2$ probability that could measure the consistency of theory given all experimental results. In this case, Bayes' theorem is used to obtain \\begin{equation} P(\\hat{t}|A,H)=\\frac{P(A,H|\\hat{t})q(\\hat{t})}{\\sum_{All~theories} P(A,H|t)q(t)}. \\end{equation} where, because of the definition of $f$ ($f^\\prime$), each $\\hat{t}_i$ is the total number of events expected in the $i^{th}$ bin for both experiments, and, likewise, each $t_i$ is the number of events expected from any physically meaningful theory. Using the experimental Poisson and Gaussian errors, \\begin{align} P(\\hat{t}|A,H)= \\frac{\\int \\int \\left[ \\prod_{i=1}^{N} \\frac{f_i^{A_i}(1-f_i)^{H_i}\\hat{t}_i^{A_i+H_i}e^{-\\hat{t}_i}}{A_i!H_i!} \\right] N(s_A;0,\\sigma_A)N(s_H;0,\\sigma_H) ds_A ds_H} {\\int \\int \\left[ \\prod_{i=1}^{N} \\frac{1}{A_i+H_i+1} \\right] N(s_A;0,\\sigma_A)N(s_H;0,\\sigma_H) ds_A ds_H} \\notag \\\\ \\end{align} where all theories are effectively marginalized by integrating $t_i$ within the product. To obtain a $\\chi ^2$-like probability, one would calculate $P(\\hat{t}|A,H)$, and then calculate many $P(t^\\prime|A,H)$'s where each $t^\\prime_i$ is the Poisson fluctuated $\\hat{t}_i$. The fraction of events where $P(t^\\prime_i|A,H)<P(\\hat{t}_i|A,H)$ would be the degree of confidence in the theory. }  This project is supported by the National Science Foundation under contract NSF-PHY-0500492."
    },
    "0606/astro-ph0606669_arXiv.txt": {
        "abstract": "We have carried out a sensitive search for gas emission lines at infrared and millimeter wavelengths for a sample of 15 young sun--like stars selected from our dust disk survey with the {\\it Spitzer Space Telescope}. We have used mid--infrared lines to trace the warm (300--100\\,K) gas in the inner disk and millimeter transitions of $^{12}$CO to probe the cold ($\\sim$20\\,K)  outer disk. We report no gas line detections from our sample. Line flux upper limits are first converted to warm and cold gas mass limits using simple approximations allowing a direct comparison with values from the literature. We also present results from more sophisticated models following Gorti and Hollenbach (2004) which confirm and extend our simple analysis. These models show that the \\si{} line at 25.23\\,\\micron{} can set constraining limits on the gas surface density at the disk inner radius and traces disk regions up to a few AU. We find that none of the 15 systems have more than 0.04\\,M$_{\\rm J}$ of gas within a few AU from the disk inner radius for disk radii from 1\\,AU up to $\\sim$40\\,AU. These gas mass upper limits even in  the 8 systems younger than $\\sim$30\\,Myr suggest that most of the gas is dispersed early. The  gas mass upper limits in the 10--40\\,AU region, that is mainly traced by our CO data, are $<$\\,2\\,M$_\\oplus$. If these systems are analogs of the Solar System,  either they have already formed Uranus-- and Neptune--like planets or they will not form them beyond 100\\,Myr. Finally, the gas surface density upper limits at 1\\,AU are smaller than 0.01\\% of the minimum mass solar nebula for most of the sources.  If terrestrial planets form frequently and their orbits are circularized by gas,  then circularization occurs early. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606339_arXiv.txt": {
        "abstract": "{Stochastic acceleration is thought to be a key mechanism in the energization of solar flare electrons.}% {We study whether stochastic acceleration can reproduce the observed soft-hard-soft evolution of the spectral features of the hard X-ray emitted by suprathermal electron. We pay special attention to the effects of particle trapping and escape.}% {The Fokker-Planck equation for the electron distribution is integrated numerically using the coefficients derived by Miller et al. for transit-time damping acceleration. The electron spectra are then converted to photon spectra for comparison with RHESSI observation of looptop sources.}% {The presence of particle escape softens the model spectra computed in the stochastic acceleration framework. The ratio between the efficiency of trapping and acceleration controls the spectral evolution which follows a soft-hard-soft pattern. Furthermore, a \\emph{pivot point} (that is, a common crossing point of the accelerated particle spectra at different times) is found at around 10~keV. It can be brought into agreement with the observed value of 20~keV by enhanced trapping through an electric potential.}% {The model proposed here accounts for the key features observed in the spectral evolution of hard X-ray emission from looptop sources.}% ",
        "introduction": "Solar flares emit large amounts of continuum hard X-ray bremsstrahlung from energetic electrons up to several hundreds of keV. The thermal energy of the ambient electrons in the corona is substantially lower, with electron temperatures observed up to 6~MK in nonflaring active regions (Brosius et al. \\cite{brosius96}, Benz \\& Grigis \\cite{benz02}). Clearly, a mechanism is needed to accelerate the electrons out of the thermal distribution up to relativistic energies.  More than $10^{31}$~erg (Lin et al. \\cite{lin03}, Emslie et al. \\cite{emslie04}, Kane et al.  \\cite{kane05}) of energy can be transferred into non-thermal electrons during a major solar flare over a period of a few hundred seconds.   These numbers set stringent requirements to acceleration models, and it is by no means a simple task to identify the physical processes involved. Most of the proposed models fall into three broad classes: electric DC field acceleration, stochastic acceleration and shock acceleration (see e.g. the review by Aschwanden \\cite{aschwanden02}).   Miller et al. (\\cite{miller96}, hereafter MLM) proposed a stochastic acceleration mechanism where electrons are energized by small amplitude turbulent fast-mode waves, the transit-time damping model. MLM showed that their model could successfully account for the observed number and energy of electrons accelerated above 20 keV in subsecond spikes or energy release fragments in impulsive solar flares. However, they made no attempt to explain the observed hard X-ray spectra (which are softer than predicted by the transit-time damping model) and did not consider spectral evolution. The MLM approach does not account for escape and transport processes.   What is the effect of escape on the electron spectrum in the accelerator? In stochastic acceleration, each electron describes a random walk in energy space. The effect of escape is to shorten the dwell time of the particles in the accelerator, and therefore it limits the average energy of the electrons. This results in a softer spectrum. Therefore, we need to take in account escape to compare quantitatively the model predictions with hard X-ray observations. Furthermore, transport effects account for the modification of the electron spectrum from the time they leave the accelerator until they reach the hard X-ray emitting regions. The usual interpretation is that the particles are accelerated near the top of magnetic loops, but most of the X-rays are emitted in the loop footpoints, where the density is much larger.  Nevertheless, looptop sources are also observed (Masuda et al.  \\cite{Masuda94}).   Observations from the Reuven Ramaty High Energy Imaging Spectrometer (RHESSI, Lin et al. \\cite{lin02}) deliver hard X-ray spectra of both footpoint and looptop sources (Emslie et al. \\cite{emslie03}, Liu et al. \\cite{liu04}). Looptop sources are better suited for comparison with accelerator models than footpoint sources, since one does not need to take into account the spatial transport from the accelerator. The emission from the modeled electron distribution can be directly compared with the observed spectra of the looptop source.  \\begin{table} \\caption{Observational constraints for looptop sources (taken from Battaglia \\& Benz \\cite{battaglia06}).} \\label{tab:obsres} \\centering \\begin{tabular}{l l l} \\hline\\hline Parameter Description & Average & Range\\\\ \\hline Pivot-point energy $\\phe_*$ & 20 keV & 16--24 keV\\\\ Pivot-point flux$^\\mathrm{a}$ $J_*$ &  2 & 1--4  \\\\ Mean temperature & 22 & 18--25 MK\\\\ Nonthermal fitting range: \\\\ \\hskip 0.5cm Lower energy $\\phe_\\mathrm{min}$ & 25 keV & 20--30 keV\\\\ \\hskip 0.5cm Upper energy $\\phe_\\mathrm{max}$ & 60 keV & 40--80 keV\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$]In units of  photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$. \\end{list} \\end{table}    Battaglia \\& Benz (\\cite{battaglia06}) present a careful and comprehensive study of the spectral evolution of looptop and footpoint sources for 5 well observed RHESSI flares.  They find that looptop sources systematically show the soft-hard-soft (SHS) behavior observed spectroscopically for the total flare emission (Grigis \\& Benz \\cite{grigis04} and references therein). On the other hand, some footpoint sources do not follow the SHS pattern. These observations suggest that the SHS behaviour is a property of the acceleration process.   The observed SHS behavior manifests itself as a tight correlation between the curves describing the time evolution of the negative spectral index $\\gamma$ and the normalization $J_{E_0}$ of the non-thermal power-law component of the photon spectra, measured at a fixed reference energy $E_0$. RHESSI observations show a linear relation between $\\gamma$ and $\\log J_{E_0}$ (Grigis \\& Benz \\cite{grigis04}). This is equivalent (Grigis \\& Benz \\cite{grigis05}) to the presence of a \\emph{pivot point}, that is, a common point where the power-law component of the photon spectra observed at different times intersect. Therefore, the position of the pivot point can be found observationally by fitting the regression line to $\\gamma$ and $\\log J_{E_0}$. Since the observations show that during the impulsive phase of most flares the scatter around such a model is relatively small, such fittings provide us with two new meaningful observational parameters which can be used for comparison with theories. The best-fit pivot-point energy can be measured by RHESSI with an accuracy of 10--20\\% (Battaglia \\& Benz \\cite{battaglia06}).   The measured energies of the pivot points lie in the range of 18--24 keV for looptop sources and 13--15 keV for footpoint sources. In Table \\ref{tab:obsres} we summarize the observational results from Battaglia \\& Benz (\\cite{battaglia06}) which we will use to constrain the transit-time damping acceleration model.   In this paper we study how well the transit-time damping model endowed with an escape mechanism can account for the hard X-ray spectra of the looptop sources seen by RHESSI and, in particular, for their spectral evolution. We investigate whether the spectrum of the photons emitted by the modeled electrons in the accelerator shows SHS behavior and how well this quantitatively agrees with the observations.   We proceed as follows: a modified version of the MLM model with a better characterization of electron trapping is presented in Sect. \\ref{model}.  The evolution of the electron energy distribution function is given by a diffusion equation which is integrated numerically and transformed into photon spectra. In Sect. \\ref{results} we present the resulting photon spectra and compare their behaviour with the observations. The results are discussed in Sect. \\ref{discussion} and conclusions are drawn in Sect. \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}   The model studied assumes that turbulent fast-mode waves are present in the accelerator region. The equilibrium spectrum of electrons subjected to Coulomb collisions and transit-time damping interactions with the waves are computed allowing for trapping/escape and replenishment of ``cold'' particles. As such the model does not distinguish between the accelerator and the radiator and considers the photon spectra emitted by the electrons in the accelerator itself. The relatively dense looptop sources observed by RHESSI fit this scenario as accelerators/emitters, and the observed high densities support a scenario in which trapping plays an important role in acceleration.  Therefore, spectral observations of these looptop sources may well deliver a snapshot of the electrons in the accelerator.  To match the low-energy part of the observed spectrum, we have to add an additional thermal component at the observed temperature. This component fills a much larger volume (factor of 1000). The heating may have occurred by earlier acceleration or may be the result of waves escaping from the accelerator region with energy density too low to yield significant energy gain per particle.   The goals of this work were to test whether a stochastic acceleration mechanism can account for the observed soft-hard-soft behavior and, in particular, to find a minimal set of modifications to the transit-time damping mechanism yielding a pivot-point like behavior of the electron spectrum in the deka-keV range.  Therefore we have not fully investigated the physics of escape, and in particular we leave open the question of the connection between the acceleration process and the trapping mechanism.  The results shown in Section \\ref{results} confirm that the hardness of the spectrum of accelerated particles in a stochastic acceleration model depends strongly on particle trapping. This can be understood by recalling that stochastic acceleration can be thought as a transport of particles in energy space due to a random walk process. While acceleration pushes particles toward higher energies, the escape acts against this flow, because fast particles are more likely to be lost, and replaced by particles at lower energies. Therefore, when the losses are stronger, the average time a particle spends random-walking in energy space is shorter, thus gaining less energy. This results in softer spectra.  We have used one specific kind of stochastic acceleration model and a simple escape scenario, but the physical mechanisms explored by the specific models are general enough that other stochastic models and escape terms are expected to follow the same general behavior leading to a soft-hard-soft effect, manifesting itself as a correlation of the time evolution of the spectral index and the non-thermal X-ray flux.  A simple escape model cannot account for the observed value of the photon pivot-point energy, which is about a factor 2 higher. The modification in the escape term featuring no escape below a threshold energy $E_\\mathrm{T}\\simeq 40$~keV can increase the energy of the pivot point to the observed values, but at the price of introducing a spectral hardening below 30~keV, which is not observed for large $\\gamma$.  One possible interpretation of the threshold energy used in the alternative escape model is the presence of an electric field driving the return current. It was treated as a free parameter, but physically the escape process and the return current are linked by the electric conductivity of the loop legs.  In this sense the model is not self consistent. The issues of transport of the escaped particles to the footpoint and return currents are important (see for example Zharkova \\& Gordovsky \\cite{zharkova05}) and need to be addressed in future work to be able to derive footpoint spectra.   Thus the model is not able to account quantitatively for all the observed features of the intricate spectral behavior of the hard X-ray emission from looptop sources of solar flares. Nevertheless, the model qualitatively accounts for the key features of the spectral evolution, going one step further toward the solution of the acceleration problem. This allows to interpret the impressive RHESSI observations showing spectral variation in looptop sources.  The simplifications introduced by the model may well be responsible for its shortcoming in reproducing the observed spectra for large $\\gamma$. Future improvements may include an isotropization process and magnetic trapping. This will require the extension of the model including at least one spatial dimension and the pitch angle distribution of the particles."
    },
    "0606/astro-ph0606613_arXiv.txt": {
        "abstract": "We present the results of numerical simulations of stationary, spherically outflowing, $e^\\pm$ pair winds, with total luminosities in the range $10^{34}- 10^{42}$ ergs~s$^{-1}$. In the concrete example described here, the wind injection source is a hot, bare, strange star, predicted to be a powerful source of $e^\\pm$ pairs created by the Coulomb barrier at the quark surface. We find that photons dominate in the emerging emission, and the emerging photon spectrum is rather hard and differs substantially from the thermal spectrum expected from a neutron star with the same luminosity. This might help distinguish the putative bare strange stars from neutron stars. ",
        "introduction": "\\label{intro} There is now compelling evidence that electron-positron ($e^\\pm$) pairs form and flow away in the vicinity of many compact astronomical objects (radio pulsars, accretion disk coronae of Galactic X-ray binaries, soft $\\gamma$-ray repeaters, active galactic nuclei, cosmological $\\gamma$-ray bursters, etc.). The estimated luminosity in $e^\\pm$ pairs varies greatly depending on the object and the specific conditions: from $\\sim 10^{31}-10^{36}$ ergs~s$^{-1}$ for radio pulsars up to $\\sim 10^{50}-10^{52}$ ergs~s$^{-1}$ in cosmological $\\gamma$-ray bursters.  For a wind out-flowing spherically from a surface of radius $R$ there is a maximum (isotropic, unbeamed) pair luminosity beyond which the pairs annihilate significantly before they escape. This is given by \\begin{equation} L_\\pm^{\\rm max} = {4\\pi m_ec^3R\\Gamma^2 \\over\\sigma_{\\rm T}}\\simeq 10^{36}(R/10^6{\\rm cm})\\Gamma^2\\,{\\rm ergs~s}^{-1}, \\label{Lpmmax} \\end{equation} where $\\Gamma$ is the pair bulk Lorentz factor, and $\\sigma_{\\rm T}$ the Thomson cross section. When the injected pair luminosity, $\\tilde L_\\pm$, greatly exceeds this value the emerging pair luminosity , $L_\\pm$ , cannot significantly exceed $L_\\pm^{\\rm max}$; in this case photons strongly dominate in the emerging emission: $L_\\pm < L_\\pm^{\\rm max}\\ll \\tilde L_\\pm\\simeq L_\\gamma$. Injected pair luminosities typical of cosmological $\\gamma$-ray bursts (e.g., Piran 2000), $\\tilde L_\\pm\\sim 10^{50}-10^{52}$ ergs~s$^{-1}$, greatly exceed $L_\\pm^{\\rm max}$. For such a powerful wind the pair density near the source is very high, and the out-flowing pairs and photons are nearly in thermal equilibrium almost up to the wind photosphere (Paczy\\'nski 1990). The outflow process of such a wind may be described fairly well by relativistic hydrodynamics (e.g., Grimsrud \\& Wasserman 1998; Iwamoto \\& Takahara 2002). \\par In contrast if  $\\tilde L_\\pm \\ll L_\\pm^{\\rm max}$ annihilation of the outflowing pairs is negligible. It is now commonly accepted that the magnetospheres of radio pulsars contain such a very rarefied ultra-relativistic ($\\Gamma_\\pm \\sim 10-10^2$) pair plasma that is practically collisionless (e.g., Melrose 1995).  Recently we developed a numerical code for solving the relativistic kinetic Boltzmann equations for pairs and photons. Using this we considered a spherically out-flowing, non-relativistic ($\\Gamma\\sim 1$) pair winds with the total luminosity in the range $10^{34}-10^{42}$ ergs~s$^{-1}$, that is $\\sim (10^{-2}-10^{6})L_\\pm^{\\rm max}$ (Aksenov et al. 2004, 2005). (A brief account of the emerging emission from such a pair wind has been given by Aksenov et al. 2003.) While our numerical code can be more generally employed, the results presented in this paper are for a hot, bare, strange star as the wind injection source. Such stars are thought to be powerful sources of pairs created by the Coulomb barrier at the quark surface (Usov 1998, 2001).  \\begin{figure} \\centering \\includegraphics[width=0.47\\textwidth]{Aksenovf1.eps} \\caption{Luminosities of a hot, bare, strange star in $e^+e^-$ pairs (dotted line), in thermal equilibrium photons (dashed line), and the total (solid line) as functions of the surface temperature $T_{_{\\rm S}}$. The theoretical upper limit on the luminosity in non-equilibrium photons, $10^{-6}L_{\\rm BB}$ (Cheng \\& Harko 2003), is shown by the dot-dashed line, $L_{\\rm BB}$ being the blackbody luminosity.} \\label{Aksenovf2.eps}       % \\end{figure} ",
        "conclusions": ""
    },
    "0606/astro-ph0606049_arXiv.txt": {
        "abstract": "{We report on a densely time sampled polarimetric flux density monitoring of the BL~Lac object  \\object{S5\\,0716+71} at 86\\,GHz and 229\\,GHz. The source was observed with the IRAM 30\\,m telescope at Pico Veleta within a coordinated multi-frequency observing campaign, which was centred around a 500\\,ks INTEGRAL observation during November 10 to 16, 2003.  The aim of this campaign was to search for signatures of inverse-Compton catastrophes through the observation of the broad-band variability of the source. At 86\\,GHz,  S5\\,0716+71 showed no intra-day variability, but showed remarkable inter-day variability with a flux density increase of $34$\\,\\% during the first four observing days, which can not be explained by source extrinsic causes. At this frequency, making use of a new calibration strategy, we reach a relative rms accuracy of the flux density measurements of $1.2$\\,\\%. Although the flux density variability at 229\\,GHz was consistent with that at 86\\,GHz, the larger measurement errors at 229\\,GHz do not allow us to detect, with high confidence, inter-day variations at this frequency. At 86\\,GHz, the linear polarization fraction of  S5\\,0716+71 was unusually large $\\rm{(}15.0\\pm1.8\\rm{)}\\,\\%$. Inter-day variability in linear polarization at 86\\,GHz, with significance level $\\apgt  95\\,\\%$; $\\sigma_{P}/<P>=15$\\,\\% and $\\sigma_{\\chi}=6^{\\circ}$, was observed during the first four observing days. From the  total flux density variations at the synchrotron turnover frequency ($\\sim 86\\,\\rm{GHz}$) we compute an apparent brightness temperature $T_{\\rm{B}}^{\\rm{app}} > 1.4 \\times {10}^{14}$\\,K at a redshift of 0.3, which exceeds  by two orders of magnitude the inverse-Compton limit. A relativistic correction for $T_{\\rm{B}}^{\\rm{ app}}$ with a Doppler factor $\\delta > 7.8$ brings  the observed brightness temperature down to the inverse Compton limit. A more accurate lower limit of $\\delta > \\rm{ 14.0}$, consistent with previous  estimates from VLBI observations, is obtained from the comparison of  the 86\\,GHz synchrotron  flux density and  the upper limits for the synchrotron self-Compton flux density obtained from the INTEGRAL observations. The relativistic beaming of the emission by this high Doppler factor explains the non-detection of ``catastrophic\" inverse-Compton avalanches by INTEGRAL. ",
        "introduction": "\\label{int}  The effect of intra-day variability (IDV) of radio loud active galactic nuclei (AGN) in the radio bands was discovered in 1986 (Witzel et al.~\\cite{Wit86}; Heeschen et al.~\\cite{Hee87}) and since then led to controversial discussions about its physical origin. Most of the objects presenting radio IDV appear compact -when are observed in the radio and optical bands-, their relativistic jets are highly core-dominated on VLBI scales and exhibit high brightness temperatures (e.g. Wagner \\& Witzel~\\cite{Wag95}). IDV is a common phenomenon in extragalactic flat spectrum radio sources and is observed, at centimetre wavelengths, in about 10\\% to 25\\% of all objects of this class (Quirrenbach et al.~\\cite{Qui92}; Kedziora-Chudczer et al.~\\cite{Ked01}; Lovell et al.~\\cite{Lov03}). The ``classical IDV\" (of type II, as defined in Heeschen et al. \\cite{Hee87}) is characterised by variability amplitudes $\\aplt 20$\\,\\% and variability time-scales between 0.5\\,days and 2\\,days. In parallel to the variability of the total flux density, similar or even faster variations of the linear polarization (Quirrenbach et al.~\\cite{Qui89}; Kraus et al. \\cite{Kra99a},~\\cite{Kra99b},~\\cite{Kra03}) and of the circular polarization (Macquart et al.~\\cite{Mac00}) are observed. Intensity and polarization variations can be correlated or anti-correlated (e.g. Wagner et al.~\\cite{Wag96}; Macquart et al.~\\cite{Mac00}; Qian et al.~\\cite{Qia02}).  Observations of IDV sources in the radio and millimetre bands are particularly important, since they reveal the highest apparent brightness temperatures ($T_{\\rm{B}} \\propto \\nu^{-2}$; e.g. Readhead~\\cite{Rea94}). However, the observed radio emission might also be severely affected by interstellar scintillation (ISS), which is unavoidable because of the small intrinsic sizes of IDV sources (Rickett \\cite{Ric90}). It is known that, for high Galactic latitude blazars, the amplitude of IDV due to the effect of interstellar scintillation rapidly decreases with frequency for $\\nu \\apgt 5$\\,GHz (e.g. Rickett~\\cite{Ric90}; Rickett et al.~\\cite{Ric95}; Beckert et al.~\\cite{Bec02}). In such case and for a source whose size is constant with frequency, the ISS-induced variability amplitude scales as $\\propto \\nu^{-2}$ (e.g. Beckert et al.~\\cite{Bec02}). Hence, any rapid and large amplitude source variability seen at millimetre (or shorter) wavelengths should not be strongly affected by ISS. Therefore, monitoring of IDV sources at mm-bands is particularly important to determine whether the observed IDV is mainly source extrinsic (scintillation), due to source intrinsic processes, or a mixture of both.  \\object{S5\\,0716+71} (hereafter \\object{0716+714}) is an extremely active BL~Lac object which varies on time scales from less than one hour to months, from radio to X-rays (e.g. Wagner et al.~\\cite{Wag96}, Kraus et al.~\\cite{Kra03}, Raiteri et al.~\\cite{Rai03}). Although the optical spectrum of \\object{0716+714} appears featureless, even with 4\\,m class telescopes, the absence of any signature of a host galaxy (in deep images) sets a lower limit to its redshift of $z>0.3$ (Wagner et al.~\\cite{Wag96}). Therefore, only lower limits to the brightness temperature of the source can be derived. The strongest constraint on the brightness temperature of this source comes from the observed short variability time-scales at radio wavelengths, which sets, via the causality argument, an upper limit to the source size of a few tens of micro-arcseconds. During the last decade the source has been observed repeatedly in simultaneous centimetre wavelength campaigns (Wagner et al. \\cite{Wag90},~\\cite{Wag96}; Otterbein et al.~\\cite{Ott99}; Kraus et al.~\\cite{Kra03}; Raiteri et al.~\\cite{Rai03}), during which it typically showed IDV. \\object{0716+714} is the only IDV source which has shown simultaneous variability-time-scale transitions in the radio and optical bands, which have been interpreted as evidence for source intrinsic variability (Quirrenbach et al.~\\cite{Qui91}; Wagner \\& Witzel~\\cite{Wag95}; Qian et al.~\\cite{Qia96}). The correlation between the radio spectral index and the optical flux suggested that the radiation was produced by the same particle population (Qian et al.~\\cite{Qia96}). Additional arguments for the relevance of the intrinsic origin of the IDV in \\object{0716+714} come from the variability amplitude, which increases from radio to optical, and the recently detected IDV at a wavelength of 9\\,mm (Kraus et al.~\\cite{Kra03}). Both effects contradict expectation for the frequency dependence of ISS.  If IDV is produced by the intrinsic properties of AGN, very small angular sizes of the emitting regions (of a few micro-arcseconds) are inferred from their short variability time-scales. This usually implies apparent brightness temperatures several orders of magnitude larger than the inverse-Compton (IC) limit ($\\aplt 10^{12}$\\,K, Kellermann \\& Pauliny-Toth~\\cite{Kel69}; Readhead~\\cite{Rea94}; Kellermann~\\cite{Kel03}). For incoherent synchrotron sources of radiation, this limit could be violated only during short time ranges, because the high energy photon densities at the IC-limit should lead to the rapid cooling of the emitting region through the IC scattering of the synchrotron radiation. This process is known as the \\emph{inverse-Compton catastrophe}. The relativistic beaming of the radiation coming from the emitting region (Rees~\\cite{Ree66}) has been proposed as a likely cause to explain the systematic violation of the IC-limit (e.g. Wagner \\& Witzel~\\cite{Wag95}). Other scenarios, as coherent synchrotron radiation, which is an unconventional but not impossible process in AGN, have been proposed (e.g. Benford~\\cite{Ben92}). However, the inverse-Compton catastrophe scenario can not still be ruled out (Wagner \\& Witzel~\\cite{Wag95}).  A coordinated broad-band observing campaign -centred around a 500\\,ksec INTEGRAL\\footnote{INTErnational Gamma-Ray Astrophysics Laboratory} observation during November 10 to 16, 2003- was performed to search for signatures of inverse-Compton catastrophes on a flaring state of \\object{0716+714} through the correlation of the broad-band variability of the source. The ground-based observations were performed with the VLBA, the Effelsberg 100\\,m, the Mets\\\"ahovi 13.7\\,m, the IRAM 30\\,m, the JCMT, the HHT, and the Kitt Peak 12\\,m telescopes and the optical-IR telescopes of the WEBT\\footnote{Whole Earth Blazar Telescope} collaboration. The first analysis of the broad-band (from radio to soft $\\gamma$-rays) observations (Ostorero et al.~\\cite{Ost06}) do not show obvious correlation between the intra-day optical variability and the inter-day radio variability displayed by the source. Although the apparent brightness temperatures of \\object{0716+714} largely exceeded the IC-limit at radio wavelengths, no evidence of IC avalanches was seen in the INTEGRAL data.  In this paper we report on the results and implications of the first polarimetric millimetre wavelength IDV observations of \\object{0716+714}, which were performed (within the above-mentioned campaign) with the IRAM 30\\,m radio telescope during the period November 10 to 18, 2003. A combined discussion of the radio, millimetre and sub-millimetre single dish data set will be presented by Fuhrmann et al.~(\\cite{Fuh06}). Forthcoming papers will describe more detailed analysis of the optical observations and more sophisticated theoretical modelling of the broad-band data set. ",
        "conclusions": "We have presented the results from millimetre observations of \\object{0716+714}, which were performed on 2003 November 10 to 18 with the IRAM 30\\,m telescope. Our observation strategy, based on the rapid time sampling of both the target source and the calibrators, enabled us to reach a relative calibration { accuracy} of $1.2$\\,\\% at 86\\,GHz, which demonstrates the good performance of this telescope and its ability for future accurate IDV studies in the millimetre range.  { During our first four observing days, the source displayed large amplitude ($Y=34$\\,\\%) and monotonous inter-day variability at 86\\,GHz. As such large amplitude variability is not expected to be produced by standard ISS at mm wavelengths (Rickett et al. \\cite{Ric95}), this variation should be considered as intrinsic to the source and not due to the influence of the interstellar medium. The similar 32\\,GHz and 37\\,GHz behaviour reported by Ostorero et al.~(\\cite{Ost06}) during the same observing time range could hence be explained by intrinsic causes also.  Ostorero et al.~(\\cite{Ost06}) also report clear evidence of IDV in the optical range, which is not matched either in the mm or in the cm bands (see also Fuhrmann et al.~\\cite{Fuh06}). This indicates that the radio-mm emitting region was located in a different region than the optical one or that their radiation behaviours were driven by different physical processes during our observations.}  We have reported { an unusually large linear polarization degree $<p>=\\rm{(}15.0\\pm1.8\\rm{)}\\,\\%$ of \\object{0716+714} at 86\\,GHz, which suggests a large level of magnetic field alignment. { At such frequency, linear polarization inter-day variability, with significance level $\\apgt 95\\,\\%$; $\\sigma_{P}/<P>=15$\\,\\% and $\\sigma_{\\chi}=6^{\\circ}$, was observed during the first four observing days.} We have also shown} $2 \\sigma$ evidence of { simultaneous} polarization flux density and polarization angle IDV within a time $\\aplt 24$\\,h. If such rapid polarization variations that are uncorrelated with the total flux density variability are confirmed by future observations, then equally rapid changes of the magnetic field configuration of the source or changes of opacity around $\\nu_{m}$ would be required to explain the phenomenon. In both cases, inhomogeneous models would probably have to be invoked.  The synchrotron spectrum of \\object{0716+714} peaked at $\\nu_{m} \\approx 86$\\,GHz during 2003 November 10 to 15 { (Ostorero et al.~\\cite{Ost06})}, and { showed} an optically thin spectral index $\\bar{\\alpha}_{86-229}=-0.23 \\pm 0.10$ between 86\\,GHz and 229\\,GHz. The apparent brightness temperature derived from our 86\\,GHz light curve, $T_{\\rm{B}}^{\\rm{app}}>1.4 \\times {10}^{14}$\\,K for a redshift $z=0.3$, exceeds, at least by two orders of magnitude, the IC-limit of $3\\times {10}^{11}$\\,K (Kellemann \\& Pauliny-Toth~\\cite{Kel69}; Readhead~\\cite{Rea94}; Kellermann~\\cite{Kel03}). This mismatch can be explained by the relativistic motion or expansion of the source with a minimum Doppler factor $\\delta > 7.8$. The upper limits from the soft $\\gamma$-ray simultaneous INTEGRAL observations in the 3\\,keV to 200\\,keV energy range enabled us to compute { an } independent and more { robust} limit for the source Doppler factor, $\\delta_{\\rm{IC}} \\apgt \\rm{{ 14}}$. This limit is consistent with previous estimates { of} the Doppler factor of \\object{0716+714} measured from the kinematics of VLBI-scale jet features (Bach et al.~\\cite{Bac05}), which ranged from $\\delta_{\\rm{VLBI}}=20$ to $\\delta_{\\rm{VLBI}}=30$. Such high Doppler factors have also been measured with 43\\,GHz-VLBI by Jorstad et al.~(\\cite{Jor05}), who detected $\\delta_{\\rm{VLBI}}>20$ in 9 of the 13 monitored blazars.  { As} no soft-$\\gamma$-ray IC-avalanches were detected by INTEGRAL during our observations and the reported large amplitude 86\\,GHz variability can not be ascribed to the ISS, the relativistic beaming of the radiation coming from the emitting region in \\object{0716+714} { offers} a robust explanation of the { apparent} violation of the IC-limit { of the brightness temperature in the mm range}. Note, however, that (total or partial) coherent synchrotron-emission scenarios can not be ruled out.  Finally, we should stress that we have proven that \\object{0716+714}, in particular, and blazars, in general, can display apparent brightness temperatures two orders of magnitude larger than the theoretical limits { in the millimetre range} and that the influence of the interstellar medium is not always necessary to explain the mismatch between observations and theory. Following the above arguments, we are rather confident that, apart from inaccuracies in our assumptions, the estimates and limits for the physical parameters of \\object{0716+714} reported in the previous sections correspond to those governing the observed source behaviour."
    },
    "0606/astro-ph0606080_arXiv.txt": {
        "abstract": "We have studied the face-on, barred spiral NGC\\,7424 (site of the rare Type IIb SN\\,2001ig) with {\\it Chandra}, {\\it Gemini} and the Australia Telescope Compact Array. After giving revised X-ray colours and luminosity of the supernova, here we focus on some other interesting sources in the galaxy: in particular, our serendipitous discovery of two ultraluminous X-ray sources (ULXs). The brighter one ($\\sim 10^{40}$ erg s$^{-1}$) has a power-law-like spectrum with photon index $\\Gamma \\approx 1.8$. The other ULX shows a spectral state transition or outburst between the two {\\it Chandra} observations, 20 days apart. Optical data show that this ULX is located in a young (age $\\approx 7$--$10$ Myr), bright complex rich with OB stars and clusters. An exceptionally bright, unresolved radio source ($0.14$ mJy at $4.79$ GHz, implying a radio luminosity twice as high as Cas A) is found slightly offset from the ULX ($\\approx 80$ pc). Its radio spectral index $\\alpha \\approx -0.7$ suggests optically-thin synchrotron emission, either from a young supernova remnant or from a radio lobe powered by a ULX jet. An even brighter, unresolved radio source ($0.22$ mJy at $4.79$ GHz) is found in another young, massive stellar complex, not associated with any X-ray sources: based on its flatter radio spectral index ($\\alpha \\approx -0.3$), we suggest that it is a young pulsar wind nebula, a factor of $10$ more radio luminous than the Crab. ",
        "introduction": "The physical interpretation of ultra-luminous X-ray sources (ULXs) remains a fervently debated and unsolved problem (for recent reviews, see King 2006; Fabbiano \\& White 2006; Colbert \\& Miller 2004). They could be scaled-up equivalents of Galactic black hole (BH) X-ray binaries, with a higher-mass accretor (up to $\\sim$ a few $10^2 M_{\\odot}$ for the brightest sources, if the accretion is isotropic and Eddington-limited). Alternatively, they could be beamed sources---either mild geometrical beaming or relativistic Doppler boosting have been suggested (King et al.~2001; K\\\"{o}rding, Falcke \\& Markoff~2002; Fabrika \\& Mescheryakov 2001). A direct comparison with Galactic BH X-ray binaries is hampered by at least two problems. Firstly, the mass and evolutionary state of the donor star is unknown in almost all cases: some ULXs appear to be associated with OB associations or young star clusters; others, however, do not have optical counterparts down to typical detection limits $M_V \\sim -5$ mag. Secondly, most ULXs have been observed in the X-rays only once or few times over the last decade; thus, it is still unclear whether they follow the same patterns of spectral evolution and state transitions as daily-monitored Galactic BH X-ray binaries, and over what time-scale.  In this paper, we report the serendipitous discovery of two ULXs in the face-on SABcd galaxy NGC\\,7424. Henceforth, we adopt a distance $d = 11.5$ Mpc (Tully 1988); a slightly lower distance of $10.9$ Mpc was estimated by B\\\"oker et al. (2002). The galaxy was observed twice with {\\it Chandra}, on 2002 May 21--22 and 2002 June 11, to study the X-ray supernova SN\\,2001ig (Evans, White \\& Bembrick 2001; Schlegel \\& Ryder 2002). For the same reason, it was monitored on a regular basis for three years with the Australia Telescope Compact Array (ATCA) (Ryder et al.~2004). We studied the X-ray spectral state of the two ULXs found near SN\\,2001ig, and their variability between the two observations. We then searched for their radio and optical counterparts, using the ATCA dataset, together with archival {\\it HST}/WFPC2 and new {\\it Gemini} images (Ryder, Murrowood \\& Stathakis~2006)\\footnote{High-resolution colour images of NGC\\,7424 from the {\\it Gemini} observations are available at http://www.gemini.edu/2001igpr}. Finally, we determined the X-ray luminosity and colours of the other bright X-ray sources detected in the field of this galaxy. ",
        "conclusions": "We studied the brightest X-ray sources and their optical and radio counterparts and environments in the face-on spiral galaxy NGC\\,7424. The galaxy was originally monitored in various energy bands to study the evolution of SN\\,2001ig. In particular, we studied it with {\\it Gemini} in the $u'$, $g'$ and $r'$ optical bands, in 2004; with the ATCA in four radio bands, over many observations between 2001 and 2004; and with {\\it Chandra}, twice, on 2002 May 21--22 and June 11. Optical and radio results on SN\\,2001ig are reported elsewhere (Ryder et al.~2004, 2006); here we only revise its X-ray colours and luminosity, $\\approx 6 \\times 10^{37}$ erg s$^{-1}$ on 2002 May 21--22, declining to $\\approx 3 \\times 10^{37}$ erg s$^{-1}$ on 2002 June 11, for a $0.5$ keV thermal plasma.  Our monitoring of NGC\\,7424 has led us to the serendipitous discovery of two ULXs with interesting multi-band properties. ULX\\,1 shows a $75\\%$ increase in flux between the two observations, reaching an emitted luminosity $\\approx 9 \\times 10^{39}$ erg s$^{-1}$ in the second one. It may have been even brighter when observed by {\\it ROSAT}, in 1990. Its X-ray spectrum is well fitted with a power-law of photon index $\\Gamma = 1.8\\pm0.1$. This is significantly flatter (harder) than the typical values found in Galactic BHs when they are bright ($L_{\\rm X} \\sim L_{\\rm Edd}$) {\\it and} power-law dominated. Various other bright ULXs have shown this feature, which is yet to be properly understood: we have briefly discussed some possible explanations. ULX\\,1 is in a relatively empty interarm region, far from any bright clusters or star-forming complexes; the brightest optical source in the {\\it Chandra} error box has an absolute magnitude $M_V \\sim -4.5$ mag, consistent with a main-sequence B star.  ULX\\,2 is more peculiar for at least three reasons. Firstly, it showed some kind of outburst or state transition between the two {\\it Chandra} observations, increasing its brightness by an order of magnitude. Thermal plasma emission is also detected in its high state. Secondly, unlike ULX\\,1, it is located in an exceptionally bright, young stellar complex (age $\\approx 7$--$10$ Myr depending on the assumed metal abundance), which also includes a few compact clusters. It is still unclear why some ULXs are in low-density, probably older environments while others are associated with OB associations and young star clusters. A similar dichotomy is sometimes found within the same galaxy. For example, the late-type spiral galaxy NGC\\,4559 also contains two bright ULXs (Cropper et al.~2004), one of them located in an interarm region relatively devoid of bright stars, the other in a bright star-forming complex a few hundred pc in size (Soria et al.~2005).  In the case of ULX\\,2 in NGC\\,7424, the biggest star cluster in its environment has a mass $\\sim$ a few $10^4 M_{\\odot}$. In general, ULXs seem to be associated with young medium-size clusters much more often than with super star-clusters ($M \\sim 10^6 M_{\\odot}$). This may give us a clue to the most likely formation process for the accreting BHs. In fact, the observed X-ray luminosities only require BH masses up to $\\sim 150 M_{\\odot}$. The necessary (but not sufficient) condition to end up with such a BH is to have a massive stellar progenitor, perhaps up to $\\sim 300 M_{\\odot}$, preferebly with low metal abundance. We speculate that the most efficient process to create such a massive progenitor could be an externally-triggered, dynamical gas collapse in a medium-size protocluster, together with a few Class-0 protostellar mergers, rather than runaway mergers of main-sequence O stars in a super star-cluster.  The third reason why ULX\\,2 is remarkable is the presence of an exceptionally bright (twice as luminous as Cas A), unresolved radio source slightly offset from the ULX ($\\approx 80$ pc). Its radio spectral index ($\\alpha \\approx -0.7$) is consistent with optically-thin synchrotron emission. Two alternative interpretations for the radio source are either lobe emission, powered by a ULX jet, or a young, bright SNR (or perhaps hypernova remnant) associated with the young stellar complex or maybe even the progenitor of the ULX itself. Compact radio sources of similar brightness and spectral index have been found associated with a handful of other ULXs, in some cases implying energies up to $\\sim 10^{52}$ erg. In almost all cases, it is impossible to rule out either of those two scenarios altogether. It may even be possible that jet lobes and an underlying SNR coexist, like in the Galactic microquasar SS\\,433.  Another bright (in fact, even brighter), unresolved radio source was found in the southern spiral arm, also associated with a young, massive stellar complex. In that case, however, there are no X-ray sources nearby. Moreover, its radio spectral index is flatter ($\\alpha \\approx -0.25$). We speculate that it is most likely to be a very young pulsar-wind nebula, $\\sim 10$ times more luminous than the Crab Nebula. However, further radio and X-ray observations of this peculiar source will be necessary to test this hypothesis and rule out other scenarios, for example beamed radio core emission from a microblazar.  This year, {\\it Chandra} is carrying out a complete snapshot survey of ULXs in about 150 galaxies within 15 Mpc, to pinpoint their locations (PI: D. Swartz). A systematic optical survey of the same galaxies is also being carried out or planned, to classify and study the ULX optical counterparts. However, much more work remains to be done to plan and implement a correspondingly deep, complete radio survey, which would determine what fraction of ULXs are associated with optically-thin synchrotron sources, whether those sources are statistically different from normal SNRs, and perhaps also help distinguishing between SNR and radio lobe scenarios. The long-term goal of combined radio/X-ray studies is to determine the balance of power in accreting BHs, between radiative (X-rays) and mechanical luminosity (jet power) in different spectral states."
    },
    "0606/astro-ph0606741_arXiv.txt": {
        "abstract": "We present the results of a numerical magnetohydrodynamic simulation that demonstrates a mechanism by which magnetic fields tap rotational energy of a stellar core and expel the envelope.  Our numerical setup, designed to focus on the basic physics of the outflow mechanism, consists of a solid, gravitating sphere, which may represent the compact core of a star, surrounded by an initially hydrostatic envelope of ionized gas.  The core is threaded by a dipolar magnetic field that also permeates the envelope.  At the start of the simulation, the core begins to rotate at 10\\% of the escape speed. The magnetic field is sufficiently strong to drive a magneto-rotational explosion, whereby the entire envelope is expelled, confirming the expectation of analytical models.  Furthermore, the dipolar nature of the field results in an explosion that is enhanced simultaneously along the rotation axis (a jet) and along the magnetic equator.  While the initial condition is simplified, the simulation approximates circumstances that may arise in astrophysical objects such as Type II supernovae, gamma ray bursts, and proto-planetary nebulae. ",
        "introduction": "}    Over the last two decades magnetic fields have been identified as the principle, universal agent for creating collimated astrophysical outflows. When a magnetic field is anchored in a rapidly rotating object near the bottom of a gravitational potential, the field can act as a drive-belt, tapping rotational kinetic energy and launching plasma back up the potential well.  This magneto-rotational (MR) scenario for jet launching has been explored by numerous authors, both analytically \\citep{blandfordpayne82, pelletierpudritz92} and numerically \\citep*{ouyed3ea97, krasnopolsky3ea99}.  These authors showed this mechanism can produce steady flows of matter, in a sense that the outflow engine still operates while the flow is observable---as in jets from young stellar objects \\citep{reipurthbally01} and active galactic nuclei \\citep[][]{begelman3ea84}.  On the other hand, this mechanism can also operate in a transient event, linked to a rapid evolution of the source and driving an explosion.  In this case, the engine loses a significant fraction of its power by the time the outflow (or its interaction with the environment) is detected.  The proposed mechanism for gamma ray bursts \\citep[GRBs;][]{piran05} and supernovae (SNe), lie in the explosive regime.  There is also growing evidence \\citep{bujarrabalea01} that collimated outflows from planetary nebulae (PNs) are explosive.  Transient MR explosions have not been explored in as much detail as the steady-state models.  In a collapsing star, differential rotation near the core may amplify the field linearly \\citep[when turbulence is unimportant;][]{kluzniakruderman98, wheeler3ea02} or exponentially \\citep[when turbulence is important;][]{akiyamaea03, blackman3ea06}. When differential rotation twists a poloidal magnetic field, generating toroidal field $B_\\phi$, enough rotational energy may be tapped to power a supernova explosion.  The role of toroidal field pressure in then driving an outflow has been highlighted by \\citet{lyndenbell96}, who explored magneto-static ``magnetic towers'' to understand jet properties \\citep[see also][]{liea01}, and \\citet{uzdenskymacfadyen06} has extended this work to exploding stars. Numerical simulations by \\citet{leblancwilson70} and more recently by \\citet*{ardeljan3ea05} support the idea that MR explosions can be important for driving SNe and GRBs.  However, the inherent time-dependence and general complexity of the system presents a significant challenge to our understanding of the basic mechanism and identification of key parameters.  It is clear that a MR explosion is ultimately driven by the rotational kinetic energy extracted from the material that is left behind (the stellar remnant).  But, notably, there is still some uncertainty as to whether an accretion disk must form inside the star \\citep[as in][]{uzdenskymacfadyen06} or whether the rotation of the stellar core \\citep[e.g., a protopulsar;][]{wheeler3ea02} alone will drive a MR explosion.  In this Letter, we present a magnetohydrodynamic (MHD) simulation of a MR explosion.  While our setup is simple, similar to the analytical ``protopulsar jet'' model described by \\citet{wheeler3ea02}, our simulation captures the nonlinear dynamics, as the magnetic field is twisted at the shear layer between the core and envelope.  This heuristic approach enables a better understanding of the basic MR physics and complements the more phenomenological approach of previous, more complex, numerical studies. ",
        "conclusions": "}    Our numerical setup is simple, which allows us to understand the basic physical principles at work.  The requirements for the mechanism to operate are a) a shear layer in which the inner region is rotating at a fraction of breakup speed and b) a magnetic field that results in an Alfv\\'en speed comparable to the escape speed.  Because we focus on the MR mechanism, and ignore other physics (e.g., neutrino heating), we can scale our results to various astrophysical systems.  This is instructive, as it determines what conditions would be necessary for a MR explosion alone to drive off the envelope.  Here, we briefly apply our results to the progenitor stars of SNe and PNs.  A simple model of a core-collapse SN suggests that core material conserves magnetic flux and angular momentum during collapse and naturally forms a rapidly-rotating, highly-magnetized proto-neutron star.  We adopt a neutron star (``core'') mass of 1.4 $M_\\odot$ and a core radius of $R_{\\rm c} = 10^6$ cm.  If the overlying envelope contains 1 $M_\\odot$ between $r = R_{\\rm c}$ and $1 R_\\odot$, our simulation corresponds to a dipole field strength at the core surface equal to $B = 3 \\times 10^{15}$ G.  Note that the envelope mass is comparable to the core mass, but the self-gravity of the envelope near the core should not be important, and so the simulation still applies. In our simulation, we find that the core loses rotational energy at a constant rate of $B^2 R_{\\rm c}^3 / 3 = 3 \\times 10^{48}$ ergs per rotation, which should extract most of the rotational energy of the core ($\\sim 10^{51}$ ergs) in 1 second.  These numbers are consistent with previous work and should be sufficient to eject the envelope \\citep[e.g.,][]{wheeler3ea02, ardeljan3ea05, blackman3ea06}.  The evolutionary phase preceding PN formation is marked by high mass loss rates, leading to an expansion of the stellar envelope, which overlies a proto-white dwarf.  A shear layer at the core-envelope boundary is expected \\citep{blackmanea01}, and as the envelope density decreases, the Alfv\\'en speed should rise.  There may be a threshold density, below which the remaining envelope may be expelled via a MR explosion.  We adopt a white dwarf (``core'') mass of 0.5 $M_\\odot$, and $R_{\\rm c} = 10^9$ cm.  At this time the envelope is essentially an extension of the massive outflow, and if it contains 0.05 $M_\\odot$ between $r = R_{\\rm c}$ and $10^4$ AU, our simulation corresponds to $B = 9 \\times 10^6$ G.  Most of the rotational energy of the core ($\\sim 10^{48}$ ergs) should be transferred to the envelope in a few hundred years.  The linear momentum of the swept up shell will be $\\sim 10^{39}$ g cm s$^{-1}$, which is consistent with the observations of proto-PNs by \\citet{bujarrabalea01}.  Finally, we note that the large-scale field in our simulations is strong enough to suppress turbulence due to shear instabilities near the core, and the field therefore is neither subject to decay nor the exponential growth that is important in some other models.  We have instead focused on the physics of how the field actually drives the outflow."
    },
    "0606/astro-ph0606577_arXiv.txt": {
        "abstract": "In the hypothesis that the 5.4m binary \\rxj\\, consists of a low mass helium white dwarf (donor) transferring mass towards its more massive white dwarf companion (primary), we consider as possible donors white dwarfs which are the result of common envelope evolution occurring when the helium core mass of the progenitor giant was still very small ($\\simlt 0.2$\\msun), so that they are surrounded by a quite massive hydrogen envelope ($\\simeq$1/100\\msun or larger), and live for a very long time supported by proton--proton burning. Mass transfer from such low mass white dwarfs very probably starts during the hydrogen burning stage, and the donor structure will remain dominated by the burning shell until it loses all the hydrogen envelope and begins transferring helium.  We model mass transfer from these low mass white dwarfs, and show that the radius of the donor decreases while they shed the hydrogen envelope. This radius behavior, which is due to the fact that the white dwarf is not fully degenerate, has two important consequences on the evolution of the binary: 1) the orbital period decreases, with a timescale consistent with the period decrease of the binary \\rxj ; 2) the mass transfer rate is a factor of about 10 smaller than from a fully degenerate white dwarf, easing the problem connected with the small X--ray luminosity of this object. The possibility that such evolution describes the system \\rxj\\, is also consistent with the possible presence of hydrogen in the optical spectrum of the star, whose confirmation would become a test of the model. ",
        "introduction": "\\label{sec:intro} The X ray source RX~J0806.3+1527, discovered by ROSAT in 1990 \\citep{beuermann1999}, is variable with a period of 321.5s \\citep{israel1999}, which resulted to be also the only variability period in the optical and infrared light curves \\citep{ramsay2002,israel2002}. This promoted the interpretation of the 321.5s as an orbital period \\citep{burwitz2001}, in a system similar to V407~Vul (RX~J1914.4+2456, P=570s), for which \\cite{cropper1998} had proposed a ``double degenerate polar\" model. In this interpretation, the two systems RXJ0806.3+1527 and V407~Vul would be the shortest period double degenerate white dwarf (DDWD) systems, progenitors of the class dubbed AM CVn, having helium dominated spectra and \\Porb\\ from 10 to 65m. The DDWD systems (both interacting and detached) might be a dominant source of low frequency gravitational waves in the Galaxy \\citep{hils1990, nelemans2001gr} and a fraction of them could even be progenitors of Type I supernovae.  The nature of the soft X-ray emission detected from RX J0806.3+1527 and RX J1914.4+2456 is still debated. Several models have been proposed (see Cropper et al. 2003 and reference therein). In addition to the polar--like model, in which the accreting WD is magnetic, a ``direct--impact\" accretion model on a non magnetic WD has been proposed \\citep{marsh-steeghs2002}.  One problem with the mass transfer interpretation of both RXJ0806.3+1527 and RX J1914.4+2456 is their period derivative, which is negative in both cases, while it is to be expected that stable mass transfer between WDs will produce increasing orbital periods. \\cite{han-webbink1999} notice that indeed finite- temperature white dwarfs depart significantly from zero temperature white dwarfs only in their partially or non degenerate outer envelopes: as soon as these layers are stripped away by mass loss, the interiors behave practically indistinguishably from fully degenerate white dwarfs, and their adiabatic mass radius exponent is then negative. If the radius increases when the mass decreases due to mass loss, the orbital period must indeed increase, unless the mass transfer is unstable. The problem of the decreasing period of this system has been a motivation to fully develop alternative models, such as the unipolar inductor model \\citep{wu2002,dall'osso2006} or the intermediate polar model \\citep[e.g.][]{norton2004}, in which the 5.4m period is suggested to be the WD spin period. Notice that direct accretion --no disk-- models (which probably apply to this case) worsen the problem of the period derivative, and indeed act to destabilize the mass transfer. The unipolar inductor model has also been criticized by \\cite{marsh-nelemans2005}, who also propose that the negative period derivative can be explained if the mass transfer rate can be pushed away from its equilibrium value.  There are two other problems for the interpretation of RXJ0806.3+1527 as a DDWD:\\par {\\it i)} the X-ray luminosity of the source is quite low: in the range 0.5--2.5keV  it is only $\\simeq 2 \\times 10^{33} (\\rm d/1kpc)^2$ erg s$^{-1}$, that is $\\sim 5 \\times 10^{32}$erg s$^{-1}$ for a distance of 500pc \\citep{israel2003}, while the value predicted in the case of mass transfer driven by gravitational radiation (GR) is $\\sim 10^{35}$erg s$^{-1}$ \\citep{israel2002}. This value can be reduced to $\\sim 2 \\times 10^{33}$erg s$^{-1}$ in the case of highly non conservative mass transfer and very low mass ($\\simlt 0.35$\\msun) accreting WD primary  \\citep{willems-kalogera2005}. Another possible caveat to this problem is that the primary WD may be affected by compressional heating, which could bring it at an absolute M$_v \\sim 4.7$ and $T_{\\rm eff} \\sim 140000$K according to  \\cite{bildsten2006}. In this case the distance would be much larger, $\\sim$20kps, 8kpc above the galactic plane, and the X-ray luminosity would be consistent with that predicted by a GR driven mass transfer rate.\\par {\\it ii)} There is some evidence that the optical spectrum of \\rxj\\ shows the presence of hydrogen \\citep{norton2004}, with a non negligible abundance \\citep{steiper2005}: also this feature, if confirmed, {\\it apparently} argues against the DDWD scenario.\\par In this paper we study the evolution of DDWDs by following the mass transfer phases with a complete stellar evolution code, and assuming a quasi--evolutionary structure for the donor WD. One or two common envelope phases must have occurred in the binary story before the present phase of mass transfer, so it is not possible to model consistently the whole evolution of the binary with hydrostatic codes. Nevertheless, we provide insight about which are possible starting conditions for the mass loss from a degenerate dwarf. We propose that RXJ0806.3+1527, and possibly also V407~Vul, are indeed DDWDs, but that we see them during those  phases of mass transfer, during which the structure of the external layers of the donor WD is dominated by the p-p hydrogen shell burning, so that the stellar radius contracts in response to mass loss. Thus we obtain both a period derivative correct in sign and order of magnitude, and the solution of the conundrum of the optical spectrum and of the low X--ray emission. ",
        "conclusions": "The optical spectrum  of of the V $\\simeq$ 21~mag counterpart of \\rxj, obtained with FORS1 at the ESO VLT, shows a blue continuum with faint emission lines of HeI and HeII which are taken as strong evidence for a hydrogen depleted binary (Israel et al. 2002). \\cite{norton2004}, examining this spectrum, notice that the fluxes of the emission lines corresponding to the odd terms of the He II Pickering series are at least a factor of 1.5-3 less than the fluxes of the neighboring even term transitions, indicating that the even terms may be blended by emission lines from the H Balmer series. First results of a detailed modeling of the spectrum \\citep{steiper2005} yield a He/H abundance number ratio 0.1 $<$(He/H) $<$ 0.3. \\cite{reinsch2004} suggest that such a ratio is not consistent with a helium WD donor but rather that we see emission from a hot hydrogen-rich plasma and that the dominance of He II emission is just a consequence of the high plasma temperature.  Figure 4 shows the variation of helium abundance Y (mass fraction) at the surface of Seq.~2 as a function of the stellar mass. As in this run we include helium diffusion, the helium abundance is initially zero, but it increases as soon as the layers in which helium depletion is not complete are exposed. We see that Y$\\sim$0.35 when P=5.4m, corresponding to a number ratio He/H$\\sim$0.12, consistent with \\cite{steiper2005} analysis. Of course, the spectral evidence is not so pregnant, and modeling of the physical conditions of this system is difficult. Nevertheless, we urge new observations, as our new models {\\it require} the presence of hydrogen in the spectrum! The CNO abundances in the spectrum might also become a constraint of the evolutionary status of this intriguing binary (D'Antona et al., in preparation).  The novelty of the present models for the shortest periods DDWDs is that we have assumed that the donor white dwarf has a very small initial mass, so that it is a helium white dwarf with a massive hydrogen envelope which is not subject to diffusion induced hydrogen shell flashes. Due to the long phase of p-p burning, prolonged by helium diffusion, the donor may be still in this burning phase when it begins mass transfer to the primary WD. Until the whole hydrogen envelope is lost, the donor WD {\\it contracts} in response to mass loss, the orbital period decreases, and the mass transfer rate is smaller by a factor up to $\\sim 10$ than in the case of mass transfer from a fully degenerate helium white dwarf. This model is able to explain the decreasing orbital period of \\rxj, its low X ray luminosity, and the possible presence of hydrogen in the spectrum. This latter feature becomes a requirement of the model, so that it is necessary to confirm it by new spectroscopic observations and its careful model analysis. The lifetime of a system like RXJ0806.3+1527 in this phase is not more than a factor two shorter than the lifetime at the same orbital period, but when the period is increasing. This model suggests that a fraction of the double degenerate systems could be formed from common envelope evolution, ending up in the formation of a quite low mass WD with a massive hydrogen envelope. There is an important observational evidence for the existence of very low mass and long-lived hydrogen burning WDs, namely the optical companions of a few millisecond pulsars in Globular Clusters (e.g. Ferraro 2006) and the system PSR J1012+5307, in which the spin down age of the MSP is compatible with the companion WD cooling age only if this is stably burning hydrogen. Lack of a population of luminous low mass WD remnants of common envelope evolution does not necessary mean that this evolution is rare, but it may be telling us that the duration of the non interacting phase, which depends on the distribution of orbital periods following the common envelope, is short. Further exploration of the modalities of formation of these systems is necessary, if we wish to understand the consequences of this model for the background of gravitational waves emission by compact objects in the Galaxy."
    },
    "0606/astro-ph0606094_arXiv.txt": {
        "abstract": "{We examine the predictions of the core accretion - gas capture model concerning the efficiency of planet formation around stars with various masses. First, we follow the evolution of gas and solids from the moment when all solids are in the form of small grains to the stage when most of them are in the form of planetesimals. We show that the surface density of the planetesimal swarm tends to be higher around less massive stars. Then, we derive the minimum surface density of the planetesimal swarm required for the formation of a giant planet both in a numerical and in an approximate analytical approach. We combine these results by calculating a set of representative disk models characterized by different masses, sizes, and metallicities, and by estimating their capability of forming giant planets. Our results show that the set of protoplanetary disks capable of giant planet formation is larger for less massive stars. Provided that the distribution of initial disk parameters does not depend too strongly on the mass of the central star, we predict that the percentage of stars with giant planets should increase with decreasing stellar mass. Furthermore, we identify the radial redistribution of solids during the formation of planetesimal swarms as the key element in explaining these effects.} ",
        "introduction": "Radial velocity surveys led to the discovery of over 150 extrasolar planets around main sequence stars. Published descriptions of most of them can be found in the references given by \\citet{marcy05} and \\citet{mayor04}. Those surveys have been the most successful in the case of G dwarf stars, because such stars have well-defined spectroscopic features and show only a little photospheric activity. Consequently, most of the known extrasolar planets orbit stars similar to our Sun. Due to the constant progress in the detection techniques, the observational programs recently started to also include stars with lower masses on a larger scale, namely M dwarfs. Moreover, some of these surveys are now particularly dedicated to lower-mass stars \\citep[e.q.][]{Endl03,Bonfils04}. So far, these efforts have led to the discovery of three planets around two M dwarf stars: Gliese 876b,c \\citep{marcy98,marcy01} and GJ436b \\citep{butler04}.  From the theoretical point of view, the problem of giant planet formation around M dwarfs was studied recently by \\citet{laughlin04}. They addressed it within the core accretion - gas capture model (CAGCM) that provides the most widely accepted scenario explaining the formation of giant planets in both the Solar System and extrasolar planetary systems. This model predicts that first a solid planetary core is formed by collisional accumulation of planetesimals. When the core reaches a mass of a few Earth masses, it starts to accrete gas, and an extended hydrostatic envelope is built around it.  As the accretion rate of gas is greater than the accretion rate of solids at this time, the envelope eventually becomes more massive than the core. When this happens, a runaway accretion of gas ensues, which is terminated either by tidal interactions of the planet with the protoplanetary disk or by the dissipation of the disk. CAGCM has found supporting evidence in the discovery that stars with planets have higher metallicities than field stars \\citep{santos00,debra03}. This is because the formation time of giant planets decreases with increasing surface density of the planetesimal swarm \\citep{pollack96}, which in turn increases with the primordial metallicity of the protoplanetary disk. Thus, giant planets are expected to form more easily in disks with higher metallicity.  \\citet{laughlin04} conclude that M dwarfs have a limited ability to form Jupiter-mass planets. This is a direct consequence of their assumption that the surface density of the planetesimal swarm out of which planetary cores are formed scales linearly with the mass of the central star. However, the solid component of the protoplanetary disk evolves in a different way than the gaseous component \\citep{weiden93}. Due to the gas drag, a significant redistribution of solids takes place, and in the inner disk their surface density can be substantially enhanced compared to the initial one \\citep{weiden03,SV97}. In general, the efficiency of the processes responsible for the redistribution of dust depends on the mass of the central star. An obvious conclusion is that analysis of the formation of giant planets around stars with various masses should include the global evolution of solids in protoplanetary disks. A simple model of the evolution of solids was proposed by \\citet{kac2,kac05}. Applying it to solar-like central stars, these authors reproduced the observed correlation between stellar metallicity and the probability of a planet occurring \\citep[a similar result was independently obtained by][]{ida04b}.  The rapid progress in observational techniques opens up the possibility of testing the correctness and predictive power of the model proposed by \\citet{kac05}. To that end, we extend their analysis and calculate probabilities of planet occurrence around stars with different masses (both smaller and larger than $1 M_\\odot$), which may be compared to future observational data. In Sect. \\ref{s:methods} we explain our approach to the evolution of protoplanetary disks and planet formation. The results of our calculations are presented in Sect. \\ref{s:res} and discussed in Sect. \\ref{s:conc}. ",
        "conclusions": "\\label{s:conc} Based on a simple approach to the evolution of solids in protoplanetary disks, we investigated the influence of the mass of the central star on the formation of giant planets. We showed that due to the more efficient redistribution of solids the planetesimal swarms around less massive stars tend to have higher surface densities. Next, we derived the minimum surface density of the planetesimal swarm needed to enable formation of a giant planet within the lifetime of the protoplanetary disk, and we found that at distances from the star smaller than $\\sim\\!10$~AU it increases with the stellar mass. Farther away from the star the minimum density becomes anticorrelated with the mass of the star. However this effect is offset by the anticorrelation mentioned already between the mass of the star and the surface density of the planetesimals.  These two effects determine the set of initial parameters characterising protoplanetary disks that are capable of giant planet formation within the core accretion - gas capture scenario. We showed that this set is larger for less massive stars. This means that the percentage of stars with massive planets should increase with decreasing stellar mass (at least in the range 0.5 $M_\\odot$ -- 4 $M_\\odot$). However, as discussed below, in the currently accessible range of orbital radii ($<$ 5 AU), the situation is not all that clear.  Based on the sets obtained for different metallicities, we determined the occurrence rate of planets with orbits smaller than $5$~AU as a function of the mass and metallicity of the star. We took into account the fact that  the outer region of the disk is gravitationally unstable in some models. Such regions are located farther than $5$~AU from the central star, and planets formed there by disk fragmentation are not included in our occurrence rate. However, their presence reduces the amount of solid material available for the formation of planetesimals. For less massive stars this effect is so strong that it overcomes factors promoting planet formation, so that for metal-poor disks the rate of planet occurrence decreases with the mass of the central star. As a result, the minimum metallicity at which giant planets can form at orbits smaller than $5\\ \\mathrm{AU}$ decreases from $\\sim 0.6$ for stars with masses of $0.5 M_\\odot$ to $\\sim 0.2$ for $4 M_\\odot$.  In the metal-rich regime the percentage of entirely stable disks in which formation of giant planets is possible is larger, and stable regions of partly unstable disks contain enough solids to produce planetesimal swarms capable of giant planet formation. Consequently, both factors promoting planet formation around less massive stars are in play, and a clear anticorrelation between the stellar mass and planet occurrence rate is observed. At the same time, our model does not account for the presence of giant planets around metal-poor stars. This may be due to the fact that we do not include planets that have formed beyond 5 AU and later migrated inward. Such an assumption is valid as long as the number of these planets is small compared to the number of planets that have formed within 5 AU. However, as we move to lower metallicities, the percentage of giant planets with silicate cores decreases (the silicates simply become too scarce), while the percentage of planets forming from ice grains increases. Thus, in metal-poor systems the number of planets with ice cores that migrated from large orbits can become a large fraction of planets at orbits smaller than 5 AU.  Obviously, our description of the evolution of solids is very simplified. The basic underlying assumptions like the single-size distribution of solid grains or the neglect of planet migration already have been discussed by \\citet{kac4} and \\citet{kac05}. The main additional assumption introduced in the present paper is the independence of the initial parameters of protoplanetary disks on the mass of the central star. While admittedly {\\it ad hoc}, it seems to be better than the one adopted by \\citet{laughlin04}, who scaled their initial surface density of planetesimals linearly with the mass of the star. They did not take into account the antecedent evolution of solids leading to the formation of planetesimal swarms, and concluded that the formation of giant planets around low-mass stars is difficult. Recent observations suggest that masses of protoplanetary disks do not strongly depend on masses of the central stars \\citep{guilloteau}. Nevertheless, to investigate the influence of our assumption, we performed additional set of calculations with a mass of the central star of $0.5 M_\\odot$ and with the initial masses of disks scaled by factor of $0.5$. The results are shown in Fig. \\ref{f:rout_sc_t}. In this case the probability of finding a planet does not seem to depend strongly on the mass of the central star,  which is true for the whole range of metallicities we have considered. Still, our models show that the evolution of solids leading to the formation of planetesimal swarms is a vital factor facilitating the formation of giant planets, whose role should be particularly clear for low-mass stars.  Our models of gaseous disks do not reproduce recent observations by \\citep{muzerolle}, which show that the accretion rate in protoplanetary disks increases with the mass of the central star. However, in the mass range considered here this dependence is very weak, and for a given value of stellar mass the spread in accretion rates reaches two orders of magnitude. In our opinion these data do not invalidate our basic assumption that initial disk parameters do not depend on the mass of the star. We also assumed that heating by stellar radiation is negligible, whereas at least in some cases it can be a dominant source of energy in the outer regions of the disk (more efficient than the turbulent dissipation). As such, it may substantially change the structure of the disk and the radial velocities of solids. Currently we are working on models that will take these effects into account.     \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{P_p_scale.ps}} \\caption{The rate of planet occurrence as a function of the primordial metallicity of protoplanetary disks. Disks around  stars with masses of $1 M_\\odot$ are represented by a solid line. Disks around stars with masses of $0.5 M_\\odot$  are represented by dotted and dashed lines. In the first case the range of initial masses of the disks is the same as in the case of solar type stars, while in the second it was scaled according to the mass of the central star.} \\label{f:rout_sc_t} \\end{figure}"
    },
    "0606/astro-ph0606607_arXiv.txt": {
        "abstract": "We perform one-zone simulations of the infall epoch of a pre-supernova stellar core in the presence of neutrino flavor changing scattering interactions.   Our calculations give a self-consistent assessment of the relationship between flavor changing rates and the reduction in electron fraction and re-distribution of initial electron lepton number among the neutrino flavors.  We discuss and include in our calculations sub-nuclear density medium corrections for flavor changing scattering coherence factors. We find that flavor changing couplings $\\epsilon > 3\\times10^{-4}$ in either the $\\nu_e\\leftrightarrow\\nu_{\\mu}$ or $\\nu_e\\leftrightarrow\\nu_{\\tau}$ channels result in a dynamically significant reduction in core electron fraction relatively soon after neutrino trapping and well before the core reaches nuclear matter density. ",
        "introduction": "Core collapse supernovae are exquisitely sensitive to lepton number violating processes.  This is because the infall (collapse) epoch of the pre-supernova core is characterized by low entropy\\cite{bbal} and large lepton (electron and electron neutrino) degeneracy.  Nearly all of the pressure support stems from these degenerate leptons. The effects of including neutrino flavor changing neutral current (FCNC) interactions in the infall stage of a core collapse supernova have recently been investigated in Ref. \\cite{afg}.  It was noted there that neutrinos in the core of a collapsing star could undergo large numbers of scatterings due to the coherent amplification of the neutrino-quark flavor changing neutral current cross section for elastic scattering on heavy nuclei.   Such interactions could cause significant numbers of electron neutrinos in the core to be converted to mu and tau neutrinos.   In turn, this would open phase space for further electron capture and thereby significantly impact the pressure, homologous core mass, and the initial shock energy.  The explosion of core collapse (Type II, Ib, and Ic) supernovae is believed to be the result of gravitational collapse, subsequent hydrodynamic bounce of the star's core, and release of gravitational binding energy into neutrinos which ultimately provide the energy to revive and sustain the shock \\cite{c&w,wilson,b&w,bethe,b&y,m&b,jkr}. One important feature of the model is that the entropy of the core is low ($s/k\\sim 1$) and nucleons remain bound in nuclei during most of the collapse. The number of electrons in the core (hence, the pressure and homologous core mass) is governed by the electron capture reaction $e^- + p \\leftrightarrow \\nu_e + n$.  When the neutrino mean free path becomes smaller than the size of the core (because of scattering on heavy nuclei) the neutrinos become trapped.  They thermalize quickly and comprise a degenerate Fermi-Dirac sea.  When the $\\nu_e$ Fermi level becomes high enough, electron capture is blocked and \\emph{net} reduction in $Y_e$ (where $Y_f\\equiv (n_f-n_{\\bar{f}})/n_b$) no longer occurs on dynamical time scales.  However, re-distribution of electron lepton number between $\\nu_e$'s and electrons will still occur as the density rises and the nuclear composition changes.  Any further changes in the core's electron fraction during the collapse could result in a change in the collapse dynamics and explosion mechanism\\cite{hix}. Including neutrino FCNC interactions in the collapse model causes greater  reduction in $Y_e$ during infall.  This is because when electron neutrinos change flavor by scattering, holes open in the $\\nu_e$ sea and the electron capture reaction can procede.  Neutrino-quark FCNCs of the form \\begin{equation} \\mathcal{L}=\\frac{G_F}{\\sqrt 2}\\bar{\\nu}^j\\gamma^\\mu\\nu^i\\bar{q}\\gamma_\\mu(\\epsilon^q_{V_{ij}}+\\epsilon^q_{A_{ij}}\\gamma^5)q \\label{fcnclag} \\end{equation} were considered in Ref. \\cite{afg}. Here, the parameters $\\epsilon^q_{V_{ij}}$ and $\\epsilon^q_{A_{ij}}$ quantify the strength of the FCNC relative to the Fermi constant $G_F$.  Current experimental constraints\\cite{davidson} on the FCNC couplings are $\\epsilon^q_{V_{e\\mu}} < 10^{-3}$ for the channel $\\nu_e\\leftrightarrow\\nu_\\mu$ and $\\epsilon^q_{V_{e\\tau}} < 5\\times 10^{-1}$ for the channel $\\nu_e\\leftrightarrow\\nu_\\tau$. (Similar, and in some cases better, constraints on these interactions may be possible from solar and atmospheric neutrinos\\cite{f&l}.)  The cross section for neutrino flavor changing elastic scattering on heavy nuclei, mediated by the FCNCs of Eq. (\\ref{fcnclag}), was calculated in Ref. \\cite{afg} and a coherent amplification was found. Using this cross section, and employing values of the coupling constant up to and beyond current experimental constraints, Ref. \\cite{afg} gave estimates for the number of neutrino flavor changing scattering events which could occur in the core.  The resulting reduction in $Y_e$ and implications for the stellar collapse model were then discussed in a qualitative sense.  In this paper we present results of a one-zone calculation of the infall epoch of a pre-supernova star with neutrino-quark FCNCs included. Our code gives a more accurate accounting of scattering rates and the change in $Y_e$ than do the estimates of Ref. \\cite{afg} and we are able to account for some of the  feedback in the system.  We model neutrino scattering with nuclei in the core medium and account for sub-nuclear matter density structure effects.  By contrast, Ref. \\cite{afg} employed neutrino-nucleus vacuum cross sections with no accounting for medium effects.  Reference \\cite{afg} estimated which values of $\\epsilon$ would give a fast enough FCNC scattering rate such that reduction in $Y_e$ would be possible. Here we actually compute what the reduction in $Y_e$ is for various values of $\\epsilon$, including values below the best experimental bounds. We have discovered that maximal reduction in $Y_e$ is possible for values of $\\epsilon$ smaller than the best experimental bound in the $\\nu_e\\leftrightarrow\\nu_{\\tau}$ channel, and that dynamically significant reduction in $Y_e$ is possible for values of $\\epsilon$ smaller than the best experimental bound in the $\\nu_e\\leftrightarrow\\nu_{\\mu}$ channel. In section II we describe our code and method of computing the change in electron fraction.  In section III we discuss our results and their meaning for the stellar collapse model.  In section IV we list the key approximations in our calculation and give an assessment of the possible impact of the potential uncertainties introduced by these. In section V we give conclusions. ",
        "conclusions": "We have used a one-zone core collapse simulation to investigate some effects of including neutrino flavor changing interactions in the supernova model.  We have calculated the reduction in $Y_e$ as a function of the coupling constant $\\epsilon$ for collapse simulations that run up to density $\\rho= 3.8\\times 10^{13} {\\rm g} / {\\rm cm}^3$.  For values of the interaction coupling constant $\\epsilon \\gtrsim 5\\times 10^{-4}$ in either the $\\nu_e \\leftrightarrow \\nu_\\mu$ or $\\nu_e \\leftrightarrow \\nu_\\tau$ channel we have found that maximal reduction in the core's electron fraction can occur.  (See Fig. \\ref{epcon}.)  This work gives a more accurate and quantitative calculation of the effects of FCNCs than do the qualitative estimates given in Ref. \\cite{afg}. Here, we are able to account for the FCNC rate's dependence on density, and the feedback on the rates as the core becomes more neutron-rich as a result of increased net electron capture.   However, a more accurate treatment of these interactions is possible and is warranted.  Some of our approximations were made for ease of calculation, while others were made to handle physics that is not yet well understood.  Even with our conservative treatment (e.g., not following FCNCs beyond a density $\\rho = 3.8\\times 10^{13} {\\rm g} / {\\rm cm}^3$), we were able to demonstrate how strong the effect of neutrino flavor changing interactions can be on the  infall epoch physics. At very high densities near core bounce, we expect FCNCs still to be appreciable and to continue to cause net reduction in $Y_e$.  If a full simulation was preformed which included FCNCs and ran all the way to bounce density, properly accounting for neutrino scattering with nuclear matter, our results lead us to believe that significant reduction in $Y_e$ would occur for values of $\\epsilon$ even smaller than we have found here.  The EOS and neutrino scattering cross sections in nuclear matter in the core are open areas of research.  Some current simulations have more accurate treatments of these issues than we have used here.  However, obtaining reliable cross sections for neutrino scattering with nuclear matter via Standard Model interactions remains problematic, in part because of the difficulty inherent in modeling nuclear matter. The standard neutrino interactions are treated with approximations, just as we have treated non-standard interactions with approximations.   Our approximations are not a result of mysterious properties of FCNCs, but rather stem from uncertainties in matter at high density.  The biggest uncertainty in our calculation does not come from the computational approximations in our model, but rather from this lack of knowledge. We point out these issues to differentiate {\\em physical} approximations from {\\em computational} approximations. When accurate and reliable EOS and compositions in nuclear matter in the core are available, standard and non-standard types of neutrino scattering can be correctly accounted for.  A full supernova simulation, with neutrino transport and hydrodynamics, is needed to properly show all the effects of neutrino FCNCs. There are many pieces of known physics that are being tested for relevance in explaining supernova explosions \\cite{open}. The supernova model cannot be used as a means of discovering or constraining new physics until known physics has been included and tested in simulations. Such simulations can treat neutrino trapping more realistically than we have done. By keeping track of neutrino distributions, such a simulation could handle the issue, discussed in Sec. \\ref{dyefcnc}, of mu and tau neutrinos changing flavor and filling holes in the $\\nu_e$ sea before electron captures can occur.  There is also a neutrino FCNC interaction with electrons \\cite{afg}. This additional opacity source for neutrino flavor changing could be modeled easily in a full simulation. For all of these reasons, a better result for the reduction in $Y_e$ could be obtained.  More sophisticated simulations also may reveal the fate of the shock, as well as changes to the thermodynamic profile of the core. We used a constant collapse rate in our simulation, but in fact the pressure changes resulting from a continually decreasing $Y_e$ would cause a non-uniform collapse rate.  A full simulation would be able to follow the actual rate of collapse, and any consequences of a non-uniform collapse rate. Finally, a full simulation would provide neutrino specta which could reveal some signature of FCNCs in a supernova signal.  The work presented in this paper will serve as a guide to preparing such a full simulation.  Our results cannot be construed as either favoring or eliminating the existence of FCNCs.  However, they do show that including FCNCs in the current supernova model could cause major changes to the model and its predictions.  It is possible that data from a supernova signal could be used to constrain new physics such as FCNCs. On the other hand, new physics, such as what may be discovered at the LHC, might be required for successful explanation of supernovae."
    },
    "0606/hep-th0606006_arXiv.txt": {
        "abstract": "{ We look for spherically symmetric star or black hole solutions on a Randall-Sundrum brane from the perspective of the bulk. We take a known bulk solution, and analyse possible braneworld trajectories within it that correspond, from the braneworld point of view, to solutions of the brane Tolman-Oppenheimer-Volkoff equations. Our solutions are therefore embedded consistently into a full bulk solution. We find the full set of static gravitating matter sources on a brane in a range of bulk spacetimes, analyzing which can correspond to physically sensible sources. Finally, we look at time-dependent trajectories in a Schwarzschild--anti de Sitter spacetime as possible descriptions of time-dependent braneworld black holes, highlighting some of the general features one might expect, as well as some of the difficulties involved in getting a full solution to the question. } ",
        "introduction": "The idea that spacetime may not be simply four-dimensional, but have extra dimensions as yet undetected by experiment, has become essentially accepted as fact over the past two decades, largely as a consequence of string or M-theory, but also as a result of earlier work on supergravity. Over the past few years, a new alternative has emerged in our understanding of how extra dimensions can be compactified -- braneworlds and warped compactifications. Rather than the old Kaluza-Klein (KK) idea of wrapping up extra dimensions so small we only see them through extra massless fields, the braneworld idea allows us to have relatively large extra dimensions, possibly even up to the microscale, with standard particle physics confined to the ``brane'' and thereby unable to detect the extra dimensions at ordinary energy scales \\cite{EBW,ADD,RS}. Gravity however can sample these extra dimensions, and one of the most alluring aspects of warped compactifications is the possibility of unusual gravitational phenomenology not only at small scales (akin to the KK picture) but at large scales, too \\cite{oxmill,GRS,DGP}.  The braneworld paradigm views our universe as a slice of some higher dimensional spacetime, in which we have standard four-dimensional physics confined to the brane, and only gravity (plus possibly a small number of other fields) propagating in the bulk. Confinement to the brane, while at first sounding counter-intuitive, is in fact a common occurrence. The first braneworld scenarios \\cite{EBW} used topological defects to model the braneworld, and zero-modes on the defects to produce confinement, and in string theory, D-branes have `confined' gauge theories on their worldvolumes.  The new phenomenology of braneworld scenarios is then primarily located in the gravitational sector, with a particularly nice possible resolution of the hierarchy problem being its primary motivation. Clearly however, the scenario has far outgrown these initial particle phenomenology motivations, and has proved a fertile testbed for new possibilities in cosmology, astrophysics, and quantum gravity. One of the most popular models to explore, and the one which we will be using, has been the Randall-Sundrum scenario, \\cite{RS}, which consists of a domain wall universe living in five-dimensional anti-de Sitter (adS) spacetime. This model can be loosely motivated by the Horava-Witten compactification of M-theory \\cite{HW}, and many of the ideas tested and developed in this simple, calculationally explicit model underly more recent string theory motivated compactifications \\cite{KKLT}.  The Randall-Sundrum model has one (or two) domain walls situated as minimal submanifolds in adS spacetime. In its usual form, the metric of the braneworld is \\be ds^2 = e^{-2k|z|} \\left [ -dt^2 + d{\\bf x}^2 \\right ] + dz^2\\,. \\label{rsmet} \\ee Here, the spacetime is constructed so that there are four-dimensional flat slices stacked along the fifth $z$-dimension, which have a $z$-dependent conformal pre-factor known as the warp factor. Since this warp factor has a cusp at $z=0$, this indicates the presence of a domain wall -- the braneworld -- which represents an exactly flat Minkowski universe. The reason for choosing this particular slicing of adS spacetime was to have a flat Minkowski metric on the brane -- i.e.\\ to choose the ``standard vacuum''.  Randall and Sundrum showed that although gravity was inherently five-dimen\\-sio\\-nal, and the spacetime was strongly warped, as far as a four-dimensional brane-world observer was concerned, the gravitational potential of a particle on the brane was the Newtonian $1/r$ potential to leading order. A complete analysis shows that the graviton propagator has the correct tensor structure, and that the effect of the KK modes is to introduce a $1/r^3$ correction to the gravitational potential \\cite{GT}.  In astrophysics and cosmology we are not so much interested in the perturbative graviton propagator as in issues such as cosmological models or black holes, which are questions of non-perturbative, or strong gravity. The braneworld generalization of the FRW universe has been well explored and understood \\cite{GCOS}; the high degree of symmetry present renders the full five-dimensional problem fully integrable \\cite{BCG}, and the general cosmological braneworld is fully understood in terms of a slice of a five-dimensional adS black hole \\cite{BCOS}. The mass of this bulk black hole then generates a radiation-style source for the Friedman equation. Interestingly, this understanding feeds in to the second question: {\\it What is the metric of a braneworld black hole?}  At first sight it might seem that this question should be very similar to answer; as both reduce to a two-dimensional problem, however, the symmetry groups of the two spacetimes are crucially different. For cosmology, the metric splits into two parts -- the two dimensions on which it depends, and the spatial part of the universe, which has constant curvature. This problem is equivalent to a two-dimensional field theory which turns out to be totally integrable. For the black hole however, the metric splits into three parts -- the two dimensions on which it depends, the time coordinate and the remaining spatial part in which the horizon resides. Thus there are two fields in the two-dimensional theory, and there is no longer a simple solution \\cite{CG}.  The first attempt, \\cite{CHR}, to find a black hole solution replaced the Minkowski metric in (\\ref{rsmet}) by the Schwarzschild metric, thus creating a black string sticking out of the brane.  Unfortunately, as suspected by the authors, this string is unstable to classical linear perturbations \\cite{BSINS}. Chamblin et.\\ al.\\ realised that the true static localised black hole would actually be a slice of a five-dimensional accelerating black hole metric (known as the C-metric, \\cite{CMET}, in four dimensions), however no such metric has as yet been found. A lower dimensional version of a black hole living on a $2+1$-dimensional braneworld was however presented by Emparan, Horowitz and Myers \\cite{EHM}, using this four-dimensional C-metric. Since then, several authors have attempted to find the full metric using numerical techniques \\cite{BHNUM}, although the main drawback seems to be that it is a very sensitive numerical system. Nonetheless, the results of \\cite{KU} for small black holes are encouraging.  Analytically, progress has mostly (though see \\cite{bulkbh}) been limited to considering the brane metric equations of motion, with the only bulk input coming from the projection of the Weyl tensor, the ``{\\it Weyl term}'' \\cite{SMS}, onto the brane. Since this system contains an unknown bulk dependent term, assumptions have to be made either in the form of the metric or the Weyl term \\cite{BBH,GWBD}. There is no clear consensus on what the brane black hole metric is, however, some interesting features which do occur are wormholes and singular horizons \\cite{IFeat,GWBD}.  One of the reasons this braneworld black hole metric is so interesting is that it is believed to correspond to a quantum corrected black hole. In string theory, it has been realized for some time that there is a correspondence between string theory on adS space, and a CFT on the boundary of that adS space \\cite{MAL}. In other words, all of the information contained in the five-dimensional gravitational spacetime is encoded in a pure quantum field theory living on a four-dimensional spacetime. In the braneworld picture, the brane is not at the adS boundary, but at a finite distance, and the theory on this brane now contains gravity, as well as a conformal energy-momentum tensor -- the Weyl term. The effect of the brane on the adS/CFT correspondence therefore is that the bulk theory of gravity in five dimensions corresponds to the four dimensional brane theory of a CFT, with a UV cutoff, interacting with gravity \\cite{DL}.  Since the brane theory is a quantum theory, the holographic correspondence suggests that the classical bulk solution projects to a quantum corrected solution on the brane \\cite{EFK}. Indeed, it was this type of argument that led Tanaka \\cite{TAN} to argue that the braneworld black hole metric would be time-dependent, corresponding to the back-reaction of Hawking radiation on the black hole metric.  For cosmological solutions, this holographic interpretation works very well; the presence of a black hole horizon in the bulk (which is the only allowed class of bulk solution \\cite{BCG}) induces a corresponding source in the Friedman equation which has the form of a radiation source. This source can be interpreted as a CFT in a thermal state corresponding to the Hawking temperature of the bulk black hole. The brane cosmological metric has a constant curvature spatial part, and its symmetries demand that only a radiation energy Weyl term is allowed. From the bulk perspective, this means that every point on the brane is at the same distance from the bulk black hole. Thus a flat universe corresponds to a `flat' bulk black hole, a closed universe to a conventional spherical bulk black hole.  Transporting this intuition over to the brane black hole situation, one can imagine that the black hole becomes displaced from the ``center of gravity'' of the spherical wall, causing an anisotropy in the brane Weyl term, ${\\cal E}_{\\mu\\nu}$. As the black hole gets closer to the brane, this anisotropy increases, possibly becoming more important than the radiation term. This reasoning argues for a near-horizon equation of state for ${\\cal E}_{\\mu\\nu}$ which leads to a singular `event horizon' \\cite{GWBD}, which possibly corresponds to a Boulware choice of vacuum in the quantum corrected black hole \\cite{EFK}.  However, there is another holographic interpretation possible, and one which is far more intriguing and experimentally relevant -- one which incorporates black hole evaporation. Instead of imagining a quasi-static transport of the bulk black hole towards the brane, consider a bulk geodesic. From the point of view of the brane, these trajectories have constant acceleration away from the brane, so a particle in the bulk moving along a geodesic initially moves towards the brane, can touch the brane, but then moves back into the bulk accelerating away (see section \\ref{sec:tdep}). Thus a black hole would move towards the brane, hit the brane, then recoil away back off to infinity. From the brane point of view, this would correspond to collapse of conformal matter, localized around the lightcone, formation of an horizon, and subsequent evaporation of the black hole, again localized around the lightcone. This picture was indeed obtained in perturbation theory in \\cite{grs}, where the metric of a particle leaving the brane was obtained to leading order.  Such a picture should be a reasonable approximation for small mass black holes, which, coincidentally, are precisely the type of black holes that are believed to be important in LHC and cosmic ray phenomenology (for some early works, see \\cite{BHC,CRS}; for a more complete list of references, see \\cite{Kanti}). Such small black holes are thought to be produced after brane-localized particles scatter at high energies and undergo gravitational collapse. A horizon is then formed engulfing the two particles, which can never escape their mutual gravitational attraction. Due to their small mass, these black holes quickly evaporate through the emission of Hawking radiation \\cite{Kanti}: for a black hole with mass $M_{BH} \\simeq 5$ TeV, and fundamental gravity scale $M_*=1$ TeV, their lifetime is only $\\tau_{BH} \\simeq 10^{-26}$ sec, and therefore exist on our brane only momentarily. Even for a higher mass, that would result in a longer lifetime, the corresponding black hole may still `disappear' from our brane due to the so-called recoil effect \\cite{Frolov-recoil, RECOIL, Stojkovic}. Due to the absence of an analytic solution describing a black hole localized on a brane with a non-vanishing self-energy, all studies of the evaporation of brane-world black holes have been restricted to the case where the black hole mass is assumed to be significantly larger than the brane self-energy. In addition, by assuming that the black hole horizon is much smaller than the inverse adS radius, the bulk warping has also been ignored. As a result, all studies up to now have failed to consider the complete bulk-brane-black-hole gravitational system.  In the present work, we study the aforementioned gravitational system in full. Our analysis will be complementary to work on probe branes \\cite{Flachi}, and develops the work on specific brane trajectories in black hole backgrounds \\cite{Seahra,Galfard}. We restrict our study to the case of a 5-dimensional spacetime in which a 3-brane with a non-negligible energy-momentum tensor is embedded. By using the Israel's junction conditions \\cite{Israel}, we derive a set of equations corresponding to a spherically symmetric brane with additional matter content corresponding to a homogeneous and isotropic fluid, in other words the brane equivalent of the Tolman-Oppenheimer-Volkoff (TOV) equations. The main difference between this work and the brane based work of \\cite{BBH} is that we have not only a complete brane solution to the TOV equations, but also the full bulk solution. In other words a genuine brane star. Clearly the general brane-bulk system has infinitely many degrees of freedom, so our approach here is to restrict to a spherically symmetric bulk solution, and a variety of bulk backgrounds are considered with the final objective being the consistent embedding of a 3-brane into a Schwarzschild--anti de Sitter spacetime.  The outline of the paper is as follows: In the next section we derive the brane equations of motion for a brane with a general isotropic fluid source living in a (general) spherically symmetric bulk. We then consider the static system in section \\ref{sec:genstat}, and show how the static brane is completely integrable. In section \\ref{sec:star} we specialize to the physically relevant case of a Schwarzschild--anti de Sitter bulk, exploring possible black hole and stellar solutions. We then briefly consider time dependent solutions in section \\ref{sec:tdep} before concluding. ",
        "conclusions": "In this work, we have analyzed spherically symmetric brane solutions in a known bulk spacetime with the aim of finding a consistent black hole solution for the brane. We found that the problem of a static braneworld slicing a known spherically symmetric bulk was completely integrable, with the solution being given in terms of an implicit function of the bulk radial variable. Thus, we have found all possible complete brane and bulk solutions for a brane with a perfect fluid matter source living on it -- in other words complete brane TOV solutions. These solutions have the interpretation of braneworld stars, and correspond to static slicings of a Sch-adS bulk spacetime, with the bulk solution corresponding to the part of the Sch-adS spacetime {\\it not} containing the event horizon of the black hole. Thus our solutions are completely nonsingular. We have also found solutions in which the event horizon of a bulk black hole impinges upon the brane, but these typically have divergent pressure on the brane, reminiscent of the singularity in the TOV system when we try to solve for too big or compact a star.  All of our solutions contain excess pressure at large radii, this seems to be a feature of the slicing of the pure adS bulk, and it is related to the fact that the Randall-Sundrum solution, a pure Minkowski brane, is in fact not a static slicing of adS in global coordinates. The only possibility for having a well-behaved asymptopia is to have a subcritical Karch-Randall brane. These however, cannot be extended into positive mass sources. We have therefore been unable to find a solution which has all the features we would desire in a braneworld star, however, we have made crucial progress by first demonstrating how to find exact and complete solutions to the brane TOV problem, as well as classifying these according to their energy-momentum.  Probe brane calculations of the interaction of a wall with a black hole indicate the possibility of brane excision, that is, that as the black hole leaves the brane, the brane is distorted so much that it self-intersects and part of the brane is excised, falling into the black hole, with the remainder moving away towards infinity. Among our various solutions are closed bubbles as well as open branes, and it is tempting to try to model this process using a quasi-static approach -- taking a sequence of the static solutions we have found as approximate solutions (such as in \\cite{CritVe}). Unfortunately however, positivity of energy requires that the interior of the bubbles be kept, and for the branes extending to infinity, the bulk does not contain a black hole, therefore these solutions are not suitable for such an approximation. If we wished to keep the exterior of the bubble and the black hole inside the bulk spacetime, we would need negative energy branes. It would appear that time-dependence is key to finding consistent solutions, as in those cases, the black hole is actually retained in the bulk.  We have also explored the time dependent brane trajectories, using the RS brane as a starting point, to try to model the process of gravitational collapse, and to explore the issues involved in black hole evaporation. We found that the effect of a bulk black hole on the RS brane was to induce a negative energy source, however, by bending the RS brane by a small amount, we could restore the brane-DEC, although these solutions had anisotropic pressure. Provided we allow matter on the brane, we can form trajectories which now have black holes in the bulk, and which can intersect with the brane, although in this case the DEC is violated.  Finally, an important point to note is that in all our results, we have made the simplifying assumption of $\\mathbb{Z}_2$-symmetry around the brane. The RS model is $\\mathbb{Z}_2$-symmetric, and many of the investigations into gravitational braneworld solutions are also $\\mathbb{Z}_2$-symmetric. However, it is obviously important to check and explore if any of our conclusions change significantly if we drop this restriction. In particular, for black hole recoil off the brane, we would expect the black hole to recoil on one side only of the brane, and hence for $\\mathbb{Z}_2$-symmetry to be broken. It may be that many of the restrictions we have found with our solutions can be evaded if we remove $\\mathbb{Z}_2$-symmetry. This is currently under investigation."
    },
    "0606/astro-ph0606161_arXiv.txt": {
        "abstract": "In 2005 and 2006, the MAGIC telescope has observed very high energy gamma-ray emission from the distant BL~Lac object PG 1553+113. The overall significance of the signal is $8.8~ \\sigma$ for 18.8~h observation time. The light curve shows no significant flux variations on a daily time-scale, the flux level during 2005 was, however, significantly higher compared to 2006. The differential energy spectrum between $\\sim 90$~GeV and 500~GeV is well described by a power law with photon index $\\Gamma = 4.2 \\pm 0.3$. The combined 2005 and 2006 energy spectrum provides an upper limit of $z=0.74$ on the redshift of the object. ",
        "introduction": "\\subsection{The BL~Lac object PG 1553+113} The Active Galactic Nucleus (AGN) PG 1553+113 was first reported in the Palomar-Green catalogue of UV bright objects \\citep{green}. It was the only new BL~Lac object found in the survey and the first BL Lac object found in an optical survey.  Its spectrum is, typical for BL~Lac objects, featureless \\citep{miller83} and the optical variability strong ($m_p = 13.2-15.0$; \\cite{miller88}). The spectral characteristics are close to those of X-ray selected BL~Lacs \\citep{falomo90} and it is classified in the literature as intermediate BL~Lac \\citep{laurent99, nieppola06} or high-frequency peaked BL~Lac \\citep{giommi95}, as its synchrotron peak frequency lies on the borderline of these two groups.  Despite several attempts, no emission or absorption lines have been found in the spectrum of PG 1553+113 \\citep{falomo90}. Thus only indirect methods can be used to determine the redshift $z$ (e.g. \\cite{sbarufatti05, sbarufatti06}). The host galaxy was not resolved in Hubble Space Telescope (HST) images \\citep{urry00}, it is therefore safe to assume $z>0.25$. The observation of very high energy (VHE, defined here as $E \\gtrsim 100$ GeV) $\\gamma$-ray emission, on the other hand, may permit to set an upper limit on $z$. The $\\gamma$-ray absorption in the Extragalactic Background Light (EBL) by means of $e^+\\ e^-$ pair production \\citep{stecker92, hessebl} can significantly affect the shape of the observed energy spectrum depending on the source redshift. Based on present-day EBL models and the observed $\\gamma$-ray spectrum, one can derive the intrinsic spectrum as a function of $z$. Physical constraints on e.g. the slope of the intrinsic spectrum may then permit to set upper limits on the possible redshift \\citep{hess1553}.  PG 1553+113 belongs to a catalog of X-ray bright objects \\citep{donato05} and, based on its Spectral Energy Distribution (SED) properties, was one of the most promising candidates from a list of VHE $\\gamma$-ray emitting candidates proposed by \\cite{costamante02}. So far, upper limits on the $\\gamma$-ray emission have been reported by the Whipple collaboration (19\\% Crab flux above 390 GeV, \\cite{deperez03}) and Milagro \\citep{williams04}. Recently the H.E.S.S. collaboration has presented evidence for a $\\gamma$-ray signal at the $4\\sigma$ level (up to $5.3\\ \\sigma$ using a low energy threshold analysis) above 200 GeV corresponding to about 2\\% of the Crab flux \\citep{hess1553}. The energy spectrum was found to have a steep slope with $\\Gamma = 4.0\\pm0.6_{stat}$ and an upper limit on the redshift of $z<0.74$ was derived.  \\subsection{The MAGIC telescope}  The MAGIC telescope is located on the Canary Island of La Palma ($28.75^{\\circ}$~N, $17.86^{\\circ}$~W, at 2225~m asl.). The telescope comprises a 17~m diameter tessellated, parabolic mirror with a total surface of 234~m$^2$, a light-weight space-frame made from carbon fiber-epoxy tubes, and a camera with 576 hemispherical photo-multiplier tubes (PMT) with enhanced quantum efficiency ~\\citep{lacquer}. The field of view of the camera is 3.5$^\\circ$ while the trigger area covers about 2.0$^\\circ$ in diameter. The fast PMT analog signals are routed via optical fibers to the DAQ-system electronics in the counting house 80~m away. The signals are digitized by dual range 300~MHz FADCs. MAGIC can explore $\\gamma$-rays at energies down to 50~GeV (trigger threshold, depending on the zenith angle), critical for the observation of medium redshift VHE sources with steeply falling spectra like PG 1553+113. The MAGIC telescope parameters and performance are described in more detail in \\cite{magiccomm} and \\cite{magictech}.\\\\ Simultaneous with MAGIC, optical observations were performed with the KVA telescope on La Palma, operated remotely from Tuorla Observatory. The main instrument is a 60~cm (f/15) Cassegrain telescope equipped with a CCD capable of polarimetric measurements. A 35~cm auxiliary telescope (f/11) is mounted on the same RA axis. This telescope is used for BVRI CCD photometry. ",
        "conclusions": "Discussion}  The BL~Lac object PG 1553+113 has been detected at $8.8~ \\sigma$ with the MAGIC telescope in 18.8 hours of observation during 2005 and 2006. This confirms the tentative signal seen by H.E.S.S. at a higher energy threshold with data taken at about the same time as MAGIC in the 2005 period \\citep{hess1553}. The source, therefore, can now be considered as firmly detected.\\\\ The agreement between the measured H.E.S.S. and MAGIC energy spectra of PG 1553+113 in 2005 is reasonably good. While the spectral slope is consistent within errors, the absolute flux above 200~GeV in 2005 is by a factor 4 larger compared to H.E.S.S. This difference may in part be explained by the systematic errors of both measurements but also by variations in the flux level of the source (the observations with H.E.S.S. were commenced after MAGIC). The observed energy spectrum is steeper than that of any other known BL~Lac object. This may be an indication of a large redshift ($z \\gtrsim 0.3$), but can as well be attributed to intrinsic absorption at the AGN or, more naturally, to an inverse Compton peak position at lower energies. The spectrum can, however, be used to derive an upper limit on the source redshift from physical constrains on the intrinsic photon index ($\\Gamma_{int} > 1.5$) as discussed in \\cite{hess1553}. Using the lower limit on the evolving EBL density from \\cite{kneiske04} we derived a $2 \\sigma$ upper limit on the redshift of $z < 0.74$. The same value was reported by \\cite{hess1553} where a slightly different EBL model was used.  The Broad band SED of PG 1553+113 together with the results from a model calculation are shown in Fig.~\\ref{fig:spectrum_2}.The VHE data points correspond to the intrinsic spectrum of PG 1553+113 as derived for a redshift of $z = 0.3$. The black points at low energies denote the average optical and X-ray flux taken at the same time as the MAGIC observations. The gray hatched radio, optical and X-ray non-simultaneous data were taken from \\cite{giommi02}. The solid line shows the result of a model fit to the simultaneously recorded data (black points) using a homogeneous, one-zone Synchrotron Self-Compton (SSC) model as provided by \\cite{krawczynski04}. As can be seen from Fig.~\\ref{fig:spectrum_2}, the $\\gamma$-ray, X-ray and optical data are well described by the model. This is not the case for the radio data where intrinsic absorption requires a much larger emitting volume compared to X-rays and $\\gamma$-rays. Except for a somewhat smaller radius of the emitting region, identical model parameters as in \\cite{costamante02} have been used: Doppler factor $D=21$, magnetic field strength $B=0.7\\ \\mathrm{G}$, radius of the emitting region $R=1.16^{+0.62}_{-0.21} \\cdot 10^{16}\\ \\mathrm{cm}$, electron energy density $\\rho_e = 0.11^{+0.18}_{-0.06} \\ \\mathrm{erg / cm^3}$ slope of the electron distribution $\\alpha_e=-2.6$ for $8.2 < \\log \\left( E / \\mathrm{eV} \\right) < 9.8^{+0.2}_{-0.05}$ and $\\alpha_e=-3.6$ for $9.8^{+0.2}_{-0.05} < \\log \\left( E / \\mathrm{eV} \\right) <10.6^{+1.6}_{-0.0}$. The limits on some of these parameters indicate the change of the SED model parameters when varying the assumed redshift from z=0.2 up to z=0.7 (parameters without limits were kept constant for all fits). In the case of $z \\ge 0.56$ the SED model can not accurately describe the data and, based on the obtained $\\chi^2$ value, a redshift of 0.56 is excluded on the $4.5\\sigma$ level. For a comparison of the model parameters with those from other BL Lacs we refer to \\cite{costamante02}.\\\\ PG 1553+113 was in a high state in the optical in both years showing a strong flare at the end of March 2006. The high linear polarization of the optical emission ($8.3 \\pm 0.2 \\%$) indicates that a sizeable fraction of the optical flux is indeed synchrotron radiation. In $\\gamma$-rays only a significant change in the flux level from 2005 to 2006 is found while there is no evidence for variability in X-rays. As a result, a possible correlation between the different energy bands can not be established. A possible connection between the $\\gamma$-ray detection and the optical high state can, however, not be excluded. The optical flare without X-ray or $\\gamma$-ray counterpart may still be explained by external-inverse-Compton (EIC) models which predict a time lag of the X-rays and $\\gamma$-rays with respect to the optical emission."
    },
    "0606/astro-ph0606482_arXiv.txt": {
        "abstract": "We derive the evolution of the energy deposition in the intergalactic medium (IGM) by dark matter (DM) decays/annihilations for both sterile neutrinos and light dark matter (LDM) particles.  At $z > 200$ sterile neutrinos transfer a fraction $f_{\\rm abs}\\sim{} 0.5$ of their rest mass energy into the IGM; at lower redshifts this fraction becomes $\\lesssim{} 0.3$ depending on the particle mass. The LDM particles can decay or annihilate. In both cases $f_{\\rm abs}\\sim{} 0.4-0.9$ at high ($> 300$) redshift, dropping to $\\approx 0.1$ below $z=100$. These results indicate that the impact of DM decays/annihilations on the IGM thermal and ionization history is less important than previously thought.  We find that sterile neutrinos (LDM) decays are able to increase the IGM temperature by $z=5$ at most up to $4$~K ($100$~K), about 50-200 times less than predicted by estimates based on the assumption of complete energy transfer to the gas. ",
        "introduction": "According to 3-yr WMAP ({\\it Wilkinson Microwave Anisotropy Probe}) results (Spergel \\etal 2006), the dark matter (DM) constitutes about 20\\% of the cosmic energy density. However, the nature of such elusive component remains unclear.  As proposed DM candidates might induce drastically different evolutionary effects depending on their properties (\\eg velocity dispersion), one hopes to be able to select suitable DM candidates from the comparison between their predicted effects and observations.  Potentially important cosmological effects might be induced if DM particles either decay or annihilate, as predicted by fundamental physics theories.  For example, sufficiently light DM particles (mass $\\simlt{}100$ MeV; Boehm \\etal 2004; Ascasibar \\etal 2006) can annihilate, or decay into lighter particles (Hooper \\& Wang 2004; Picciotto \\& Pospelov 2005; Ascasibar \\etal 2006) remaining good DM candidates. The products of DM decays/annihilations can be photons, neutrinos, electron-positron pairs, and/or more massive particles, depending on the mass of the progenitor.  The decay/annihilation of DM particles into $e^+-e^-$ pairs has been recently invoked to explain the observation, by the SPI spectrometer aboard ESA's INTEGRAL satellite, of an excess in the 511-keV line emission from the Galactic bulge (Kn\\\"odlseder \\etal 2005). Although exotic, this idea has triggered many theoretical studies (Boehm \\etal 2004; Hooper \\& Wang 2004; Picciotto \\& Pospelov 2005; Kawasaki \\& Yanagida 2005; Kasuya \\& Takahashi 2005; Cass\\'e \\& Fayet 2006; Ascasibar \\etal 2006), aimed at constraining DM properties through SPI/INTEGRAL observations.  The products of decays and/or annihilations are expected to interact with the intergalactic medium (IGM), transferring part of their energy. If so, DM decays and annihilations might change the IGM thermal/ionization history in a sensible and detectable way. Various flavors of this mechanism have been investigated in a considerable number of studies (Hansen \\& Haiman 2004; Chen \\& Kamionkowski 2004; Pierpaoli 2004; Padmanabhan \\& Finkbeiner 2005; Mapelli \\& Ferrara 2005; Biermann \\& Kusenko 2006; Mapelli, Ferrara \\& Pierpaoli 2006, hereafter MFP06; Zhang \\etal 2006).  The above studies, though, made the simplifying assumptions that either (i) the energy injected by DM decays/annihilation is {\\it entirely} absorbed by the IGM (Hansen \\& Haiman 2004; Pierpaoli 2004; Biermann \\& Kusenko 2006; MFP06), or (ii) leave the absorption efficiency as a free parameter (Padmanabhan \\& Finkbeiner 2005; Zhang \\etal 2006), or (iii) make a partial treatment of the physical processes responsible for the energy redistribution (Chen \\& Kamionkowski 2004; Mapelli \\& Ferrara 2005).  In this paper we model in detail, for the first time, the physical processes governing the interaction between the IGM and the decay/annihilation products, and we derive the fraction of energy actually absorbed (Section \\ref{section method}). We restrict our analysis to the case in which the decay/annihilation products are photons, electron-positron pairs, or neutrinos (which are assumed to have negligible interactions with matter), because of the uncertainties in modeling the cascade associated with more massive product particles. For photons (Section \\ref{photon subsection}) we include the effects of Compton scattering and photo-ionization; for pairs, the relevant processes are inverse Compton scattering, collisional ionizations, and positron annihilations (Section \\ref{pair subsection}).  Our model exhaustively describes the behavior of relatively light DM candidates (i.e. warm and cold DM particles with mass lower than $\\sim{}100$ MeV), whose only decay/annihilation products are photons, pairs and neutrinos. As an application, we explore the effects of sterile neutrinos (mass 2-50 keV) and of viable light dark matter particles (LDM; mass 1-10 MeV) on the IGM thermal and ionization history (Section \\ref{section applications}).  We adopt the best-fit cosmological parameters after the 3-yr WMAP results (Spergel \\etal 2006), i.e. $\\Omega{}_{\\rm b}=0.042$, $\\Omega{}_{\\rm M}=0.24$, $\\Omega{}_{\\rm DM}\\equiv{}\\Omega{}_{\\rm M}-\\Omega{}_{\\rm b}=0.198$, $\\Omega{}_\\Lambda{}=0.76$, $h=0.73$, $H_0=100\\,{}h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have modelled the absorption rate of the energy released in the IGM by DM decays/annihilations producing photons, electron-positron pairs and neutrinos. Our model suitably describes the energy deposition of a wide class of DM candidates, such as sterile neutrinos and LDM.  Useful fits to the absorbed energy fraction as a function of redshift for the various particles are given in Appendix B.   In the case of radiatively decaying sterile neutrinos (mass 2-50 keV), at $z > 200$ a fraction $f_{\\rm abs}\\simeq{}0.5$ of the particle energy is transferred to the IGM, predominantly via ionizations; at lower redshifts $f_{\\rm abs}$ decreases rapidly to values of 0.0005-0.3 depending on the neutrino mass.  LDM particles can decay or annihilate. In both cases $f_{\\rm abs}\\approx 1$ at high ($> 300$) redshift, due to positron annihilation and inverse Compton scattering, and it drops to values around 0.1 below $z=100$.  Our determination of $f_{\\rm abs}$ has a dramatic impact on the results of previous studies (which adopted naive assumptions for this parameter) concerned with the IGM heating by DM.  To illustrate this point, we have re-calculated the IGM thermal and ionization history induced by either sterile neutrinos or LDM particles, using the previous findings. We find that sterile neutrino (LDM) decays are able to increase the IGM temperature by $z=5$ at most up to $4$~K ($100$~K). Both these values are 50-200 times lower than the estimates based on the assumption of complete energy transfer to the gas.  In addition, significant departures from the adiabatic temperature evolution induced by the Hubble expansion occur only below $z \\approx 30$, at an epoch when heating and ionization by conventional sources (stars or accretion-powered objects) are likely to swamp the DM signal.  LDM annihilations instead produce a very extended ($5 < z < 800$) electron fraction plateau, at a level of 5-10 times the relic one. The main effect of these extra electrons is to extend the cosmic time interval during which the IGM kinetic temperature is coupled to the CMB one down to $z\\approx 100$. Although the electron scattering optical depth in this case is large (0.08-0.10), its effects on the CMB temperature/polarization spectra are hardly appreciable.    The detailed computation of the $f_{\\rm abs}$ presented in this paper and summarized by the fits given in the Appendix B, might be useful for a large number of future applications in which the cosmological role of the DM is investigated.  Among these are the effects of DM decays/annihilations on the 21 cm emission (Shchekinov \\& Vasiliev 2006) and on the structure formation history (Shchekinov \\& Vasiliev 2004; Biermann \\& Kusenko 2006; Ripamonti, Mapelli \\& Ferrara 2006)."
    },
    "0606/astro-ph0606211_arXiv.txt": {
        "abstract": "We present the first data  release of the Radial Velocity Experiment (RAVE), an ambitious  spectroscopic survey to measure radial  velocities and stellar atmosphere parameters  (temperature, metallicity, surface gravity)  of up to one million  stars using the 6dF  multi-object spectrograph on  the 1.2-m UK Schmidt  Telescope  of the  Anglo-Australian  Observatory  (AAO).  The  RAVE program  started  in  2003,  obtaining  medium  resolution  spectra  (median R=7,500) in  the Ca-triplet  region ($\\lambda\\lambda$ 8,410--8,795  \\AA) for southern hemisphere  stars drawn from the Tycho-2  and SuperCOSMOS catalogs, in the magnitude range  $9\\!\\!<\\!\\!I\\!\\!<\\!\\!12$.  The first data release is described in this paper and contains radial velocities for 24,748 individual stars (25,274 measurements when including re-observations).  Those data were obtained on 67 nights between 11 April 2003 to 03 April 2004.  The total sky coverage  within  this data  release  is  $\\sim$4,760  square degrees.   The average signal to  noise ratio of the observed spectra is  29.5, and 80\\% of the radial  velocities have  uncertainties better than  3.4~km/s.  Combining internal  errors and  zero-point errors,  the mode  is found  to  be 2~km/s. Repeat observations are used to  assess the stability of our radial velocity solution,  resulting in  a variance  of 2.8~km/s.   We demonstrate  that the radial velocities derived for the first  data set do not show any systematic trend  with  color or  signal  to noise.   The  RAVE  radial velocities  are complemented  in the  data release  with  proper motions  from Starnet  2.0, Tycho-2  and SuperCOSMOS,  in addition  to photometric  data from  the major optical and infrared catalogs (Tycho-2,  USNO-B, DENIS and 2MASS).  The data release can be accessed via the RAVE webpage: http://www.rave-survey.org. ",
        "introduction": "\\label{s:introduction}  Within the past decade it is  being increasingly recognised that many of the clues to the  fundamental problem of galaxy formation  in the early Universe are contained in the motions and chemical composition of long-lived stars in our Milky Way galaxy (see  e.g. Freeman \\& Bland-Hawthorn 2002).  The recent discovery of several instances of  tidal debris in our Galaxy challenges the view laid down in the seminal  paper by \\citet{els}, who envision the Galaxy to be formed in one major monolithic collapse at an early epoch, followed by a period of  relative quiescence, lasting many Gyr.   These examples include the discovery of the tidally distorted/disrupted Sagittarius dwarf galaxy \\citep{sgr}, the photometrically identified low-latitude Monoceros structure in  the Sloan  Digital Sky  Survey  \\citet{monoceros} and  the multitude  of features in higher latitude fields \\citep{belokurov}.  \\clearpage  \\begin{figure*}[hbtp] \\centering \\includegraphics[width=8.7cm,angle=270]{f1.eps} \\caption{The RAVE spectrum of a typical  field star, HD~154837 (K0~III), illustrating the properties  of   the  chosen  wavelength   interval  around  the   Ca~II  IR triplet. The strongest other absorption lines are identified.  } \\label{f:nicespectrum} \\end{figure*}   \\clearpage   Furthermore,  within   the  context   of  the  concordance   LCDM  scenario, sophisticated computer simulations of structure growth within a CDM universe have now  begun to shed light on  how the galaxy formation  process may have taken  place in  a hierarchical  framework (see  e.g.  Steinmetz  \\& Navarro 2002,  Abadi  et  al.~2003,   Brook  et  al.~2005,  Governato  et  al.~2004, Sommer-Larsen et  al.~2003).  These analyses  lead to a  reinterpretation of structures such  as the Eggen  moving groups \\citep{navarro04} or  the Omega Cen globular cluster  \\citep{meza05} in terms of merger  remnants.  In fact, in the  extreme case, structures, and  old stars, in even  the thin Galactic disk  are   attributed  to  accretion   events  \\citep{abadi03,meza05}.   In addition,  strong  evidence  for  accretion and  assimilation  of  satellite galaxies  can  also be  seen  for  other galaxies  in  the  Local Group,  in particular the  `great stream'  in M31 \\citep{sgrstream}.   However, whether the Galaxy can indeed be formed by a sequence of merging events as predicted by  current cosmological  models of  galaxy  formation, or  whether the  few well-established accretion and merging remnants  -- which account for only a small fraction of the stellar mass of the Galaxy -- are all there is, is still a largely  unanswered question.  Large  kinematic surveys  are needed,  as are large  surveys  that  derive  chemical  abundances,  since  both  kinematics signatures and elemental abundance signatures persist longer than do spatial over-densities.  It  is still  unclear (Wyse \\&  Gilmore 2006),  whether the observed chemical properties of the stars in the thick disk (Gilmore, Wyse \\& Jones 1995)  can be brought into agreement with a  scenario that sees the thick disk primarily  as the result of accretion  events \\citep{abadi03}.  A similar  question arises owing  to the  distinct age  distribution (Unavane, Wyse \\& Gilmore 1996, see however  Abadi et al.~2003) and chemical elemental abundance distributions of  stars in the stellar halo and  in low-mass dwarf galaxies of the Local Group  (Tolstoy et al.~2003; see however, Robertson et al.~2005, Bullock \\& Johnston 2005).  Stellar clusters, spiral arms, and the  Galactic bar leave an imprint in the chemical   and  stellar   velocity  distribution  in   the  solar neighborhood (Dehnen 2000; Quillen \\&  Minchev 2005; de Simone et al.~2004) as well. Multidimensional databases are required to  investigate  and differentiate  between these  processes  and structure  caused by  satellite accretion.  The growing awareness of the importance  of the `fossil record' in the Milky Way  Galaxy in  constraining galaxy  formation  theory is  reflected by  the increasing number of  missions designed to unravel the  formation history of the Galaxy.  Stellar spectroscopy plays a crucial role in these studies, not only providing  radial velocities  as a key  component of  the 6-dimensional phase  space  of  stellar  positions  and  velocities,  but  also  providing much-needed  information   on  the  gravity  and   chemical  composition  of individual stars.  An example of  the power of such multidimensional stellar datasets  have  recently  been  shown  by  \\citet{helmi06}  who,  by  using  a combination of proper motions and distances from the Hipparcos catalog \\citep{hip} and  spectra from  the Geneva-Copenhagen Survey  (Nordstr\\\"om et al.~ 2005,  hereafter  GCS)  were  able to  identify  several  accretion candidates within the immediate neighborhood of the Sun.  However, despite  the importance of  stellar spectroscopy, the  past decades have seen  only limited progress. Soon after  \\citet{vogel1873} measured the radial  velocities  of  Sirus  and  Procyon,  \\citet{oldgeorge}\\footnote{The great-great-grandfather of G. Seabroke,  co-author of this paper.} performed one of the first surveys,  measuring 68 radial velocities for  29 stars, followed by 699 observations of  40 stars \\citep{oldgeorge2} and 866  observations of 49 stars  \\citep{oldgeorge3}.  Since  then,  over the  next  125 years,  radial velocities  for  some 50,000  stars  have  become  available in  the  public databases of the Centre  de Donn\\'ees astronomiques de Strasbourg (hereafter CDS).   This is  surprisingly few,  compared to  the more  than  one million galaxy redshifts  measured within the  past decade.  This sample  of stellar radial  velocities has  recently been  increased substantially  by  the GCS, containing radial velocities for 16,682 nearby dwarf stars, and by \\citet{famay} publishing 6,691 radial velocities for apparently bright giant stars.  Both catalogs were part of the Hipparcos follow-up campaign.  With the advent of wide  field multi-object spectroscopy (MOS) fiber systems in the  1990's, pioneered  particularly at the  AAO with FOCAP,  AUTOFIB and most  recently  with  the 2dF  and  6dF  instruments  on  the AAT  and  UKST respectively  (e.g.   Lewis  et  al.~  2002 and  Watson  et  al.~2000),  the possibility  of  undertaking  wide-area  surveys with  hemispheric  coverage became  feasible.   Initially,  such   projects  were  more  concerned  with large-scale galaxy  and quasar redshift  surveys (e.g.  Colless  et al.~2001 for 2dF and  Jones et al.~2004 for 6dF).  Apart from  the samples of several hundred  to a  few thousand  stars obtained  prior to  the  commissioning of AAOmega at the AAT (see for example Kuijken \\& Gilmore 1989a, b, c; Gilmore, Wyse \\& Jones  1995; Wyse \\& Gilmore 1995; Gilmore, Wyse  \\& Norris 2002) no large-scale, wide area stellar  spectroscopy projects had been undertaken in our own  galaxy.  This  has now changed  with new surveys  like SDSSII/SEGUE already under way, with a planned delivery of 240,000 spectra by mid 2008 \\citep{segue}, and the capabilities of  AAOmega on the AAT.  Over a slightly longer time frame,  the RAdial Velocity Experiment (RAVE), the survey we describe in this paper,  is expected to provide spectra for  up to 1 Million stars by 2011.  This trend for large stellar surveys will culminate with the ESA cornerstone mission Gaia, which, in addition to astrometric information, will provide  multi-epoch radial  velocities for up  to one  hundred million stars, by  2018. Each of these surveys  has its own unique  aspect, and they are largely complementary in capabilities and target sample.  With  a radial  velocity error  of about  2~km/s and  80\\%  of measurements better than 3.4~km/s, the RAVE velocities are accurate enough for almost any Galactic kinematical study.   Radial velocities are however just  one of the necessary  stellar  parameters:   proper  motions,  distances  and  chemical abundances  are  also  needed.   Proper  motions  of  varying  accuracy  are available  for most  of  the RAVE  stars,  via Starnet2.0,  Tycho-2 or  SSS. Already, some  of the  cooler dwarfs ($J-K>0.5$)  with more  accurate proper motions can be identified as  dwarfs from their reduced proper motions.  For these  stars, it is  possible to  estimate all  six phase  space coordinates using their photometric parallaxes.  For most of these stars,  there is no previous spectroscopic information, so the RAVE sample provides many scientific opportunities.  Some of the science programs in progress by RAVE team members include: \\begin{itemize} \\item Discovery of extreme velocity stars and estimates of the local escape velocity and total mass of the Galaxy \\item The 3D velocity distribution function of the local Galactic disk \\item Kinematics of the main stellar components of the Galaxy \\item Characterization of the local Galactic disk  potential and the structure of the disk components \\item Substructure in the disk and halo of the Galaxy, including the Arcturus, Sagittarius and other star streams \\item Elemental abundances of high velocity stars \\item Calibration of stellar atmospheric parameters and correspondence with the MK scheme through the HR diagram \\item Searches for spectroscopic binaries and cataclysmic variables \\item The $\\lambda8,620$\\AA~diffuse interstellar band as an estimator of interstellar reddening. \\end{itemize}  In this  paper we describe the first  data release of the  RAVE survey which contains radial velocities obtained  from RAVE spectra (the spectra, stellar parameters and additional  information will be part of  the further releases as  the  first  year  spectra  are  contaminated  by  second  order  light). Photometric and proper motion data  from other surveys are also provided for ease of  use. The  structure of  the paper is  as follows:  In Section  2 we describe the survey layout,  technical equipment and input catalog.  Section 3 is devoted to the actual observations, followed by a section detailing the data reduction. Section 5 discusses  the data quality and compares RAVE data with  independent data  taken  with other  telescopes.   Finally, Section  6 provides  a detailed  description  of the  data  product of  the first  data release (henceforth DR1) and concludes with longer term perspectives. ",
        "conclusions": "This  first data release  presents radial  velocities for  24,748 individual stars in the range $9\\simlt  I_{DENIS}\\simlt 12.5$, obtained from spectra in the infrared  Calcium-triplet region at  a median resolution of  7,500.  The total  sky  coverage of  this  catalog  is  $\\sim$4,760 square  degrees.  We demonstrated that  the radial  velocities are not  affected by any  trend in color  or SNR using  both external  data and  RAVE repeat  observations. The resulting variance  for each set of  validation data is  consistent with our estimated errors.  DR1  does  not  include  information  about chemical  abundances  and  other atmospheric  parameters, for reasons  described above.   The quality  of the currently acquired spectra  is good enough for derivation  of $T_{\\rm eff}$, gravity,  [M/H], [$\\alpha$/Fe],  $V_{\\rm rot}\\sin  i$,  and micro-turbulence \\citep{fiorucci}, and we expect to  include chemical and atmospheric data in subsequent data releases.  RAVE  is planned  to observe  until  2010 and  will acquire  up to 1,000,000 spectra.   Incremental  releases,   containing  radial  velocities,  stellar parameters as  well as spectra, are  now planned on  an approximately yearly basis, providing an  unprecedented sample of stellar kinematics and chemical abundances in the range of  magnitudes probing scales between  the  very local  (Hipparcos  based) radial  velocity surveys (GCS and Famaye et al.~ 2005) and the more distant SDSSII/SEGUE and surveys with AAOmega, therefore completing our picture of the Milky Way."
    },
    "0606/astro-ph0606027_arXiv.txt": {
        "abstract": "We identify ``red clump stars'' -- core helium-burning giants -- among 2MASS stars and use them to measure the run of reddening with distance in the direction of each of the Galactic Anomalous X-ray Pulsars (AXP).  We combine this with extinction estimates from X-ray spectroscopy to infer distances and find that the locations of all AXP are consistent with being in Galactic spiral arms.  We also find that the 2--10\\,keV luminosities implied by our distances are remarkably similar for all AXP, being all around $\\sim\\!1.3\\times10^{35}{\\rm\\,erg\\,s^{-1}}$.  Furthermore, using our distances to estimate effective black-body emitting radii, we find that the radii are tightly anti-correlated with pulsed fraction, and somewhat less tightly anti-correlated with black-body temperature.  We find no obvious relationship of any property with the dipole magnetic field strength inferred from the spin-down rate. ",
        "introduction": "The Anomalous X-ray Pulsars (AXPs) are young, energetic, X-ray bright isolated neutron stars, with spin periods of the order 10\\,s.  They are called {\\em anomalous} since their luminosity far exceeds the energy available from spin-down, and no binary companions are seen. AXPs (along with the related Soft Gamma-ray Repeaters or SGRs) are now believed to be {\\em magnetars} (Thompson \\& Duncan, 1996). Magnetars have huge external magnetic fields ($\\sim 10^{14}$G) and even larger internal fields. It is the decay of the magnetic flux which provides the luminosity seen, and is responsible for a whole array of observational effects such as bursting and giant flares. See Woods \\& Thompson (2004) for a summary of recent observational data on magnetars, and how they are modeled.  Since they are young remnants of massive, short-lived progenitors, all of the AXPs are found in the Galactic plane (except for CXOU~J010043.1$-$721134 which is in the Small Magallanic Cloud). This causes a major obstacle to observations: high interstellar extinction, manifested as photo-electric edges from different elements in the soft X-ray band, and as continuum extinction from dust in the optical and near-infrared.  Since the amount of extinction has been difficult to estimate accurately, the spectral energy distributions of AXP have been subject to large uncertainties (Hulleman et al.\\ 2004; Durant \\& van Kerkwijk 2005, 2006).  Furthermore, even if the interstellar absorption is well-characterized, distances and therefore absolute fluxes are difficult to determine.  The simplest distance estimate is made by requiring that the black-body component typically inferred from the X-ray spectrum arises from a neutron-star sized area.  We do not, however, expect the surfaces of AXPs to be homogeneous, both on observational grounds (they pulsate) and from theoretical considerations (the magnetic field, which affects the heat conduction, will vary across the surface).  For AXPs that are associated with supernova remnants or other interstellar structure, a more direct distance estimate can be made using 21\\,cm \\ion{H}{1} measurements and the Galactic rotation curve. Convincing cases for associations with supernova remnants have been made for two AXPs: 1E~2259+589 with CTB~109 (Gregory \\& Fahlman, 1980), and 1E~1841$-$045 with Kes~73 (Sanbonmatsu \\& Helfand 1992).  Furthermore, Gaensler et al. (2005) found an \\ion{H}{1} bubble coincident with the direction of AXP 1E~1048.1$-$5937, which they suggest was created by the winds of the massive progenitor of the AXP (see Muno et al. 2006 for a discussion of possible massive progenitors to AXPs). Even for these systems, however, distance estimates can be rather uncertain, with different authors finding inconsistent results (we discuss this further in \\S4 and~5).  For sources in the Galactic plane, a different clue to the distance is available if one has a good measure of the interstellar extinction, and the run of extinction with distance can be determined independently. This works because the extinction increases with distance, more so towards star-forming regions and molecular clouds.  The so-called red clump method provides a means for deriving the function of reddening versus distance in any given line of sight, based on field stars over a relatively small area.  L\\'opez-Corredoira et al.\\ (2002) noted that in an infrared color-magnitude diagramme of a stellar cluster ($J-K$ versus $K$ for example), core helium-burning giants, or red clump stars, form a well-defined and easily-identified concentration of stars redward of the main sequence.  Stars spend up to 10\\% of their lifetime in this phase, and are much more luminous than typical main sequence stars.  Because their helium cores all have roughly the same mass, their luminosities are largely independent of the total stellar mass.  Furthermore, their infrared colours are insensitive to metallicity.  As a result, they are good infrared standard candles and if the red clump can be identified at each of a range of distances, then the reddening at each distance can be calculated.  The method has been used not only by Lopez-Corredoira et al.\\ (2002) to measure the density distribution of stars in the Galaxy, but also by, e.g., Drimmel et al.\\ (2003) to map the distribution of dust.  Here, we apply the red-clump method to measure distances to the Anomalous X-ray Pulsars.  In section \\S2, we first describe how we applied the red-clump method, focussing on the AXP for which the method was most tricky, and in section \\S3 we use this field to test the reliability of the method, and to estimate uncertainties in the derived values.  In \\S4, we present the results of the red clump method applied to each of the Galactic AXPs, and we discuss the implications in terms of distance. In \\S5, we compare our new distances with those in the literature, in particular for the controversial case of 1E~2259+589 and the associated supernova remnant CTB~109.  In \\S6, we use our results to infer luminosities and emitting radii, and discuss how these depend on other AXP properties. We draw conclusions in \\S7. ",
        "conclusions": "We have applied the ``red-clump'' method to 2MASS data to construct the run of reddening versus distance in the directions of each of the Galactic AXPs.  Combined with estimates of the reddenings to the AXPs, half of which are from our recent, model-independent determinations from photo-electric absorption edges in high-resolution X-ray spectra, we inferred distances.  We found that two of these estimated distances are inconsistent with ones in the literature, but found likely reasons for the discrepancy in both cases and concluded that our results were reliable.  From the reddening versus distance diagrammes, we find that the AXPs tend to fall in regions of rapidly rising reddening that are associated with spiral arms.  This is not surprising, since they are the young remnants of short-lived massive stars.  In particular, 4U~0142+61 has a position consistent with the Perseus Arm, 1E~1048.1$-$5937 lies on the far side of the Carina Arm, and 1XTE~J1810$-$197 and 1RXS~J170849.0$-$400910 fall along the Crux-Scutum arm.  1E~1841$-$045 could either lie in the Scutum arm or in the Molecular Ring, which dominates the gas and dust density at galactocentric distances of about 4\\,kpc (Dame et al.\\ 2001). 1E~2258+586 falls near the end of the Outer Arm, known to exist in this direction beyond the Perseus arm (e.g., Kimeswenger \\& Weinberger 1989).  From our distances, we infer 2--10\\,keV luminosities and we find that these cluster tightly around $1.3\\times10^{35}{\\rm\\,erg\\,s^{-1}}$, consistent with the prediction in the context of the magnetar model that a saturating luminosity must exist above which rapid internal neutrino cooling is effective.  Furthermore, we calculated effective emitting radii for the thermal components in the X-ray spectra, and find that these are inversely correlated with temperature, while the corresponding areas are inversely proportional to the pulsed fraction. This suggests the internal heating is released predominantly through one or more hot polar caps, with sizes that differ between the different AXPs.  The red-clump method can be applied to any line of sight in the Galactic plane, and is particularly useful combined with reddening estimates from X-ray spectroscopy.  More accurate results may be obtained using deeper infrared imaging of selected fields, although, unfortunately, it may not be possible to extend the method to regions with very high reddening, since the red clump stars may well become confused with highly reddened main-sequence stars.  As a result of the latter limitation, the method may not be generally useful for the other class of magnetars, the soft gamma-ray repeaters, which generally suffer very high extinction.  It should be useful, however, for point sources such as the Compact Central Objects.  Generally, for further analysis of distances and structure within the Milky Way, it would be useful to cross-calibrate results from the red-clump method with those from X-ray absorption studies, X-ray dust scattering haloes, and \\ion{H}{1} and CO measurements.  \\medskip\\noindent{\\bf Acknowledgments:} We would like to acknowledge the thoroughness and many useful suggestions and comments of the anonymous referee. We thank David Kaplan for pointing out the paper describing the red-clump method to us, and Martin L\\'opez-Corredoira for information on the applicability of the red-clump method.  We also thank Bryan Gaensler for help with interpreting CO and \\ion{H}{1} data and for discussions about 1E~1048.1$-$5937, Manami Sasaki and Terrance Gaetz for discussing the case of CTB~109, and Tom Dame and Peter Martin for general discussions of Galactic gas and dust. This publication makes use of the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by NASA and NSF.  We acknowledge financial support from NSERC."
    },
    "0606/astro-ph0606556_arXiv.txt": {
        "abstract": "We address the question of to what accuracy remote sensing images of the surface of planets can be matched, so that the possible displacement of features on the surface can be accurately measured. This is relevant in the context of the libration experiment aboard the European Space Agency's BepiColombo mission to Mercury.  We focus here only on the algorithmic aspects of the problem, and disregard all other sources of error (spacecraft position, calibration uncertainties, \\textit{etc.}) that would have to be taken into account. We conclude that for a wide range of illumination conditions, translations between images can be recovered to about one tenth of a pixel \\textit{r.m.s.} ",
        "introduction": "\\label{sec:introduction}  One of the goals of the European Space Agency's BepiColombo mission to Mercury is the measurement of the amplitude of the libration of Mercury. In order to do this images of the same surface areas will be taken at different times during the libration cycle and compared. When all other effects---spacecraft position, Mercury's rotation, spacecraft attitude, \\textit{etc.}---are taken into account, any remaining discrepancy between the positions of features on the surface must be due to the libration of the crust of the planet.  Here we address the question of to what accuracy images can be matched, and we focus only on the algorithmic aspects of the problem, disregarding all other sources of error that would have to be taken into account to solve the scientific problem. We shall show that by using a shape-based matching algorithm images taken under a wide range of illumination conditions can be matched to one tenth of a pixel root-mean-square. Based on this we conclude that the accuracy of the pattern matching algorithm is not the limiting factor in the ultimate accuracy that can be achieved by the libration experiment on BepiColombo. ",
        "conclusions": "\\label{sec:conclusions}  We have performed a study of the accuracy with which a shape-based pattern matching algorithm can identify translations between remote sensing images of the same planetary features. We have applied the algorithms in a Monte Carlo fashion to digital elevation models (both real and synthetic) in order to investigate the statistical performance of the procedure.  We find that for a broad range of illumination conditions translations between images can be recovered with an accuracy of 0.1\\,pixel \\textit{r.m.s}.  The algorithm performs best for translations along the projected direction to the Sun on the image plane. This study shows that translations along both image axes at the same time can be recovered with the same accuracy of 0.1\\,pixel as long as the projected direction to the Sun lies more than $\\approx 20^\\circ$ away from the same image axes.  Finally, this study demonstrates that the images to be compared need not be taken under the very same illumination conditions in order to be effectively matched. For a given Sun azimuth, any pair of images taken with Sun elevation angles larger than $10^\\circ$ can be used; images taken when the Sun is at the zenith must also be avoided. The range of useful illumination conditions is further broadened because this study concludes that differences in Sun azimuth of at least $30^\\circ$ do not affect the accuracy of the matching algorithm.  The error contributed by the matching algorithm is but one of the several error contributions to be taken into account during the analysis of the data pertaining to the measurement of the possible libration of the surface of Mercury. This study shows that the accuracy of the pattern matching algorithm is not a limiting factor in the ultimate accuracy of the libration experiment aboard the BepiColombo mission to Mercury."
    },
    "0606/astro-ph0606283_arXiv.txt": {
        "abstract": "We present a combined analysis of the X-ray emission of the Capella corona obtained with \\textit{XMM-Newton} RGS, \\textit{Chandra} HETGS, and LETGS. An improved atomic line database and a new differential emission measure (DEM) deconvolution method are developed for this purpose. Our new atomic database is based on the Astrophysical Plasma Emission Database and incorporates improved calculations of ionization equilibrium and line emissivities for L-shell ions of abundant elements using the Flexible Atomic Code. The new DEM deconvolution method uses a Markov Chain Monte-Carlo (MCMC) technique which differs from existing MCMC or $\\chi^2$-fitting based methods. We analyze the advantages and disadvantages of each individual instrument in determining the DEM and elemental abundances. We conclude that results from either RGS or HETGS data alone are not robust enough due to their failure to constrain DEM in some temperature region or the lack of significant continuum emission in the wavelength band of the spectrometers, and that the combination of HETGS and RGS gives more stringent constraints on the DEM and abundance determinations. Using the LETGS data, we show that the recently discovered inconsistencies between the EUV and X-ray lines of Fe XVIII and XIX also exist in more highly charged iron ions, up to Fe XXIII, and that enhanced interstellar absorption due to partially ionized plasma along the Capella line of sight may explain some, but not all, of these discrepancies. ",
        "introduction": "The Capella system is one of the strongest X-ray emitting coronal sources, and has been observed numerous times with the past and current X-ray observatories, including both High and Low Energy Transmission Grating Spectrometers (HETGS, LETGS) on board \\chandra and the Reflection Grating Spectrometer (RGS) of \\xmm. Capella is a close spectroscopic binary with an orbital period of 104 days and a distance of 12.9~pc \\citep{hummel94}.  The general features of the temperature distribution of the Capella coronal plasma were reasonably well determined, even with the previous generation X-ray and EUV observatories, equipped with limited spectral resolutions. Observations with the \\euve spacecraft have shown a continuous differential emission measure (DEM) distribution over a temperature range of $10^5$ to $10^7$~K \\citep{dupree93}. \\citet{brickhouse00} analyzed the simultaneous observations of \\euve and \\asca, and concluded that the DEM is sharply peaked near $10^{6.8}$~K. The abundances of Mg, Si, S, and Fe were found to be consistent with solar photospheric values, and Ne was found to be underabundant by a factor of $\\sim 3$ to 4 in that analysis.  Since the launch of \\chandra and \\xmm, high resolution X-ray spectra have become available, and numerous analyses of the Capella coronal X-ray emission have appeared using all three grating instruments. The first light observations were presented by \\citet{canizares00} for HETGS, \\citet{brinkman00} for LETGS, and \\citet{audard01} for RGS, which gave an overview of the spectral data using relatively simple analysis methods. \\citet{behar01} investigated in detail the Fe L-shell line emission using the HETGS data and theoretical calculations of the Hebrew University Lawrence Livermore Atomic Code (HULLAC). \\citet{phillips01} made detailed comparisons between the HETGS X-ray spectra and extreme-ultraviolet emission from \\euve. \\citet{mewe01} presented temperature, density and abundance diagnostics using the LETGS observation and a line ratio based analysis method. \\citet{argiroffi03} studied the structure and variability of the X-ray corona using multiple LETGS observations and a Markov Chain Monte-Carlo (MCMC) method for the DEM deconvolution developed by \\citet{kashyap98}. The x-ray emission are found to be constant to within a few percent on both short and long time scales. \\citet{audard03} studied the coronal abundances and the first ionization potential (FIP) effects of several RS CVn binaries, including Capella, using the \\xmm observations.  All previously mentioned analyses have relied on a single grating instrument in the X-ray band. \\citet{desai05} combined multiple HETGS and LETGS observations and investigated various line ratios of Fe XVIII and XIX. Large discrepancies of a factor of two were found between the observed and theoretical ratios involving $3\\to 2$ X-ray transitions and $2\\to 2$ EUV resonance lines. It was assumed that such discrepancies reflect the uncertainties of theoretical atomic data. However, this combined analysis was not aimed at deriving coronal properties of Capella.  Because of the different spectral coverage of the three grating instruments, it is conceivable that a joint analysis of all instruments may yield more stringent constraints on the DEM distribution and elemental abundances of the capella corona than that of individual instruments. Although observations with three gratings are generally not simultaneous, a joint analysis is possible due to the lack of variability of X-ray emission as indicated by multiple observations with individual instruments. In this paper, we present such a combined analysis of observations with HETGS (ObsID 5040, 28 ks), LETGS (ObsID 1248, 84 ks), and RGS (ObsID 0121920101, 51 ks). Many more similar observations exist for all instruments. The three are chosen more or less randomly. The primary goal of the present work is to introduce our analysis method and discuss the complementary nature of the three instruments. The statistical quality of these individual observations are sufficient for this purpose. We leave a systematic investigation of all available data for future work.  In the course of this work, we have developed an improved atomic line database and a new DEM deconvolution method. The new atomic line database is based on the line list of astrophysical plasma emission code (APEC) of \\citet{smith01} with the L-shell emission lines from Ne, Mg, Al, Si, S, Ar, Ca, Fe, and Ni replaced with calculations using the Flexible Atomic Code (FAC) developed by \\citet{gu03a}. Our new MCMC DEM deconvolution method is an adaptation of the spatial-spectral analysis procedure of \\citet{peterson06} to the analysis of pure spectral data, and is different from the MCMC method of \\citet{kashyap98} in technical details.  In \\S\\ref{sec:atomic}, we describe our improvements to the atomic line database of APEC; in \\S\\ref{sec:dem}, the details of the MCMC method for DEM deconvolution are discussed; \\S\\ref{sec:rgs} and \\S\\ref{sec:hetg} present the analysis of the \\xmm RGS and \\chandra HETGS observations respectively; \\S\\ref{sec:rhetg} presents the joint analysis of RGS and HETGS data; In \\S\\ref{sec:letg}, the additional constraints on the physical properties of the interstellar medium along the Capella sight line imposed by LETGS data are analyzed; we conclude with brief discussions of the results in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} During this reanalysis of Capella data, we developed a new atomic line list based on APED/APEC. The improvements include better ionization balance calculations for all L-shell ions, more level population mechanisms for Fe L-shell ions such as resonance excitation, recombination, and ionization processes, better wavelengths based on either accurate many-body perturbation theory or laboratory measurements for L-shell lines of S, Ar, Fe, and Ni. We have tested the new line database on many other coronal sources, and found it to always perform better than APED/APEC. However, there are still some serious deficiencies in the atomic data. For example, the L-shell lines of elements other than Fe only include direct excitation as the level population mechanism; the line emissivities are calculated in the low density limit; and the wavelengths of high-$n$ transitions of Fe L-shell ions have not been well determined as those of $3\\to 2$ transitions. Some of these problems are being addressed with the ongoing laboratory and theoretical work. We expect to further improve the database in the near future.  We also developed a new MCMC based DEM deconvolution method. It is different from existing methods in several aspects, including the parameterization of DEM and the utilization of the raw spectra instead of line fluxes. We demonstrated that the HETGS data alone often cannot constrain the total emission measure and the absolute abundances independently due to the lack of significant continuum emission. The DEM derived from RGS alone often has a high temperature tail which may bias the abundance measurements. The combination of HETGS and RGS data were shown to give the most robust results. The abundances derived from the joint analysis of HETGS and RGS indicate the presence of solar-like FIP effect.  We investigated the problem of intensity ratios of $2\\to 2$ EUV lines to $3\\to 2$ X-ray lines of Fe L-shell ions using the LETGS data. The overestimation of $2\\to 2$ lines are shown to exist not only for Fe XVIII and XIX, but also for higher charge ions up to Fe XXIII. The discrepancy appears to grow larger for lines at longer wavelengths. We proposed a partial explanation by assuming that the interstellar medium along the Capella line sight is partially ionized, and therefore has a larger effective absorption column density than if the medium is neutral. However, constraints on the ionization fraction of H II dictate that it can only account for about half of the discrepancy. It is plausible that the remaining discrepancy is due to either the LETGS calibration uncertainties or systematic errors in theoretical line ratios."
    },
    "0606/astro-ph0606410_arXiv.txt": {
        "abstract": "We present matched-resolution maps of HI and CO emission in the Virgo Cluster spiral NGC~4647. The galaxy shows a mild kinematic disturbance in which one side of the rotation curve flattens but the other side continues to rise. This kinematic asymmetry is coupled with a dramatic asymmetry in the molecular gas distribution but not in the atomic gas. An analysis of the gas column densities and the interstellar pressure suggests that the \\htoo/HI surface density ratio on the east side of the galaxy is three times higher than expected from the hydrostatic pressure contributed by the mass of the stellar disk. We discuss the probable effects of ram pressure, gravitational interactions, and asymmetric potentials on the interstellar medium and suggest it is likely that a $m=1$ perturbation in the gravitational potential could be responsible for all of the galaxy's features. Kinematic disturbances of the type seen here are common, but the curious thing about NGC 4647 is that the molecular distribution appears more disturbed than the HI distribution. Thus it is the combination of the two gas phases that provides such interesting insight into the galaxy's history and into models of the interstellar medium. ",
        "introduction": "Galaxies that are rich in cold gas tend to have roughly equal amounts of gas in the atomic and in the molecular phase. But, to the best of our knowledge, star formation only occurs in the molecular phase.  Therefore, the global balance between atomic and molecular gas is crucial to the long-term evolution of galaxies and their stellar populations.  What determines the balance between atomic and molecular gas in a galaxy, and how does that balance evolve as the environment of the galaxy changes?  On theoretical grounds we would expect that the density and temperature of the gas, the strength of the dissociating UV field, and the metallicity all affect the relative amounts of molecular and atomic gas \\citep{hidaka2002,elmegreen93}. The gas should also be in hydrostatic equilibrium with a pressure that is mostly determined by the stellar distribution, since stars contribute the bulk of the mass in the inner parts of galaxies where the \\htoo---HI transition occurs.  If a galaxy is disturbed by gravitational interactions or by falling into an intracluster medium, the balance between atomic and molecular gas may be altered. Indeed, \\citet{miller03a} and \\citet{miller03b} observe increased star formation activity in the members of a galaxy group or subcluster that recently merged into a larger cluster. In a bottom-up (cold dark matter) cosmology such disturbances should be common, and they may be important both to the star formation histories and the morphological evolution of galaxies. In turn, the disturbances and asymmetries in galaxies can inform our understanding of cosmology as they allow us to place constraints on the merger/interaction rate.  The semi-analytic simulations of galaxy formation (e.g.\\ \\citet{somerville99,kauffmann93} and successors) are attempts to understand the observed properties of galaxies by predicting their star formation histories in a cosmological context. To this end the semi-analytic simulations typically assume that the star formation rate can be calculated from empirically motivated recipes involving the gas content and the recent merger history.  However, these empirical recipes may not capture important phases of galaxy evolution. Thus it is desirable to have a more physically motivated understanding of the interstellar medium (ISM) and star formation that can be applied to both semi-empirical studies and future detailed N-body/hydrodynamic simulations.  \\citet{krumholz05} have recently put together a detailed model for the star formation rate in galaxies. They begin with the molecular cloud microphysics, in which fractal turbulence drives a small fraction of the gas to high enough density that it collapses to form stars. They extend the physics to galaxy-wide scales by then considering the molecular clouds as gravitationally bound entities formed in a disk that is marginally stable against gravitational collapse using the Toomre criterion \\citep{Toomre64,BT}. We note that this model still depends critically on knowing how much of the gas disk is in molecular form and how much is atomic.  In the absence of observations, that information must be obtained from a hypothesis such as that of \\citet{br04}, which predicts atomic and molecular column density ratios as a function of the midplane hydrostatic pressure in a disk. Future work on simulations of galaxy evolution could be made more realistic by incorporation of such physically-based models of a multiple-phase ISM and its effects on the star formation rate.  Our current contribution to this effort is an observational study that begins to test such physical models on a gently disturbed galaxy. We present an analysis of the kinematics of the neutral interstellar medium in the Virgo Cluster spiral NGC 4647, along with an analysis of the relationships between atomic gas, molecular gas, and star formation in the galaxy. We investigate whether the observed molecular/atomic balance and the disturbed kinematics can be understood in the context of the galaxy's environment. The form and magnitude of the kinematic disturbance in NGC 4647 are common, so the processes affecting this galaxy might be applicable to many spirals. However, few mildly disturbed spiral galaxies have been mapped at matching resolutions in both atomic and molecular gas, so opportunities for this type of analysis are rare. ",
        "conclusions": "In a $\\Lambda$CDM universe galaxy-galaxy mergers and encounters are recognized as major processes that have determined the spectrum of galaxy properties.  From an observational perspective, galaxy asymmetries can serve as indicators of the merger/interaction rate.  One may measure the incidence of asymmetries and then make the leap to the merger/interaction rate via several important assumptions such as the timescale over which asymmetries will remain detectable. That timescale depends crucially on an understanding of the processes responsible for the asymmetries and how they affect the gas, stars, and star formation activity.  In this context our observations of the mildly disturbed galaxy NGC 4647 raise several questions. NGC 4647 shows two kinds of lopsidedness, in its kinematics and in its gas distribution.  More specifically, the gas distribution is asymmetric in two ways:  the total gas surface density is significantly higher on the east side than on the west, and the gas disk is also more highly molecular on the east side than on the west.  How common are these types of asymmetries, and what causes them?  We address the problems in two stages. In section \\ref{lopsided}, a comparison to other mildly lopsided galaxies suggests processes that may create the kinematic asymmetry and enhanced gas surface densities on the east side of the galaxy (or the depressed surface densities on the west side). Section \\ref{press-asym} discusses local pressure enhancements that may drive the ISM into a more highly molecular state on the east side of the galaxy.  In section \\ref{leftoverQ}  we list some remaining questions and propose tests that would give greater insight into NGC 4647 in particular and spirals in general.  \\subsection{Lopsided Galaxies}\\label{lopsided}  It has been known for some time that many spirals are asymmetric or lopsided. For example, \\citet{rz95} and \\citet{zr97} find that some 30\\% of field spirals have lopsided stellar distributions. \\citet{baldwin80} show asymmetries in the HI distributions of spirals, and \\citet{rs94} find that 50\\% of field spirals show asymmetric global HI profiles that are probably attributable to lopsided gas distributions. \\citet{rubin99} studied \\halpha\\ rotation curves for Virgo spirals, and they find that half show significant kinematic disturbances.  In many of these, the disturbance takes the form of an asymmetry in which the velocity field is different on the two sides of the galaxy. From studies of HI position-velocity diagrams and velocity fields \\citet{swaters99} estimate that the incidence of kinematic lopsidedness is 15\\%--30\\% among non-interacting field spirals.  The relatively high incidence of asymmetries (even among field spirals) is significant, since it implies that the responsible mechanisms must either be frequent or long-lasting.  Typical dynamical timescales in spiral disks are only on the order of $10^8$ yr. In this context, four general processes have been proposed as sources of kinematic and/or morphological lopsidedness: ram pressure in a cluster environment, gravitational interactions between galaxies, asymmetric accretion of gas from the cosmic web or from satellites, and asymmetric or off-center gravitational potentials.  NGC 4647 is undoubtedly experiencing some ram pressure in its passage through the Virgo intracluster medium.  Section \\ref{press-asym} makes quantitative estimates of the ram pressure.  At a purely qualitative level, the sharp HI edge on the northeast side and the shallow dropoff in the HI column density on the southwest (Section \\ref{HIdist}) are suggestive of gentle stripping. The kinematic asymmetry in NGC 4647 may also be evidence of the non-circular velocities induced by significant ram pressure \\citep{hidaka2002}. Finally, the form of the CO asymmetry in NGC 4647 is very like that in NGC~4419, which was attributed to ram pressure by \\citet{kenney1990}.  However, it is not at all obvious that ram pressure by itself is a convincing explanation for the ensemble of properties displayed in NGC 4647. The east side of the galaxy appears strongly disturbed, with its enhanced CO emission and rising rotation curve.  At 40\\asec\\ radius the total gas density is twice as high on the east side as on the west side. In contrast, the west side appears very gently (if at all) disturbed.  Its rotation curve is well behaved out to at least 70\\asec. Unless the galaxy is proceeding exactly edge-on through the intracluster medium, so that the west side is shielded by the east side, we might expect that the effects of ram pressure would be more noticeable on the low density side of the galaxy than on the high density side. In this respect the ram-pressure stripped galaxy NGC~4419 is actually quite different from NGC~4647; both have asymmetric molecular distributions, but in NGC~4419 the atomic gas is very severely stripped as well \\citep{kenney04,chung06}.  Asymmetric gravitational potentials can also produce kinematic features like those observed in NGC 4647 and other spirals. For example, \\citet{schoenmakers97} add a perturbation of the form $\\cos(m\\theta)$ to the potential of an axisymmetric galaxy and calculate the effects on its velocity field. Based on that work \\citet{swaters99} have shown that a $m=1$ perturbation in the potential, where the axis of the perturbation is aligned with the kinematic major axis of the galaxy, can produce a velocity field that flattens on one side but continues to rise on the other.  A 5\\%--10\\% perturbation in the potential can induce 10\\%--20\\% amplitude differences in the rotation curves of the two sides of the galaxy. This kind of a lopsided potential is the explanation favored by \\citet{swaters99} to explain the HI velocity fields of DDO~9 and NGC~4395, which are similar in character to NGC 4647. \\citet{noordemeer01} have also shown that the lopsided kinematics of NGC~4395 (and, by analogy, NGC~4647) could be explained by setting the galaxy disk off-center with respect to the dark matter halo.  \\citet{battaglia06} also propose that the halo of NGC~5055 is offset from the center of its disk.  A lopsided gravitational potential would affect the stellar distribution in NGC~4647 in addition to the gas distribution.  There are subtle hints of a $m=1$ asymmetry in the stellar distribution of the galaxy; for example, in the broadband red image of Figure \\ref{tilefig} the outermost contours are compressed on the southern side of the galaxy and extended on the northwest.  The galaxy nucleus is therefore not quite centered in the middle of the outer contours.  NGC 4649 has been modeled and subtracted, but more careful analysis should be done to quantify any degree of asymmetry in the stellar distribution of NGC 4647.  NGC~4647, of course, is also an obvious candidate for a gravitational interaction with NGC~4649. \\citet{bournaud05} have modelled distant, high velocity encounters of similar-mass galaxies, and they find that for impact parameters 130--450 kpc and velocities 160--450 \\kms, the retrograde in-plane encounters can make $m=1$ asymmetries in which the ratio of the perturbed to the unperturbed gravitational force is greater than 0.1 for several Gyr. These amplitudes are similar to those required by \\citet{swaters99}.  A statistical analysis of the distribution of galaxy lopsidednesses leads \\citet{bournaud05} to propose that asymmetric accretion of gas from the cosmic web must be a significant contributor to lopsidedness in general. The mechanism simply requires that the incoming gas have a nonzero impact parameter. They find that if accretion rates are fairly high --- several \\solmass~\\peryr, or enough to double the galaxy masses over a Hubble time --- a galaxy may develop an $m = 1$ asymmetry with an amplitude of 10\\% to 20\\%.  The asymmetries in the gas and stellar distributions can persist even for several Gyr after accretion stops.  As NGC~4647 is embedded in the Virgo cluster's hot ICM it is unclear how this process might apply to NGC~4647, what the phase of the accreted gas would be, or whether the effect on galaxy kinematics would match the observations.  We suspect that the process that is responsible for the kinematic asymmetry also enhances the gas densities on the eastern side of the galaxy.  But the opposite perspective is also interesting: the pile-up in gas surface density may contribute to the kinematic asymmetry. We estimate the dynamical mass of the galaxy taking a rotation speed of 115 \\kms\\ at a radius of 70\\asec\\ on the approaching (flatter, probably less disturbed) side of the rotation curve; these values give 1.8\\e{10} \\solmass\\ within 5.8 kpc. Some 7.7\\% of this dynamical mass is attributable to molecular gas plus helium.  Since most of the molecular gas is found on the east side of the galaxy, that gas distribution may well contribute to a few percent perturbation in the gravitational potential.   The gas density asymmetry is also a $m=1$ perturbation whose axis is roughly aligned with the kinematic major axis, so it has the correct form to explain the kinematic asymmetry according to the models of \\citet{swaters99}. Therefore, even if the original cause produced both the kinematic and the gas density asymmetries together, the resulting density asymmetry may help to maintain the kinematic asymmetry over timescales longer than a dynamical time.  \\subsection{Producing the Pressure Asymmetry}\\label{press-asym}  If pressure regulates the chemical state of the ISM, as suggested by \\citet{br06}, then the asymmetry in molecular gas fraction is produced by an asymmetry in ISM pressure in the disk.  Pressure, by itself, cannot explain the enhanced gas column density on the east side of the galaxy. However, once the gas is piled up on one side, pressure could explain the amount of the ISM that is molecular.  The analysis of section \\ref{pressure-sec} finds that the observed asymmetry in the molecular gas fraction would imply an excess pressure of $P/k=2.4\\times 10^{4}~\\mbox{K cm}^{-3}$ beyond that required for the hydrostatic support of the disk of NGC 4647. Thus, we consider again the effects of the intracluster medium, the proximity of NGC 4649, and asymmetric gravitational potentials as possible sources of extra pressure on the east side of NGC 4647.  At the $\\sim 1$ Mpc separation between M87 and NGC 4647, the intercluster medium (ICM) of Virgo has a density of $n\\sim 10^{-4}~\\mbox{ cm}^{-3}$ and a temperature of $\\sim 2$ keV \\citep{kenney89}. The isotropic pressure in the ICM is significant: $P_{ICM}/k\\sim 2\\times 10^{3}\\mbox{ cm}^{-3}\\mbox{K}$, on the order of 10\\% of the midplane hydrostatic pressure in the outer parts of NGC 4647's disk. If NGC 4647 is moving with respect to the ICM, it suffers additional ram pressure \\begin{equation} P_{ram} = \\frac{1}{2}\\mu m_{\\mathrm{H}} n v^2 = 1.3\\times 10^{4}~ \\mbox{cm}^{-3}\\mbox{ K} \\left(\\frac{v}{800 \\mbox{ km s}^{-1}}\\right)^2. \\end{equation} The equation has been normalized to the one-dimensional velocity dispersion of the Virgo Cluster \\citep{huc85}.  NGC 4647 has a line-of-sight motion relative to M87 of only 100 km s$^{-1}$, but if its tangential velocity carries it eastward on the sky at 1000 \\kms\\ or so then the ram pressure could be strong enough to effect the inferred pressure difference. It remains to be seen whether the extra pressure could be applied delicately enough to avoid major disruptions in the kinematics of the downstream side of the galaxy.  Provided that NGC 4647 is physically close to NGC 4649, the hot gas in NGC 4649 can increase the pressure on the eastern side of NGC 4647 through two mechanisms. The first is ram pressure between NGC 4647 and the halo of NGC 4649. Indeed, the diffuse X-ray emission from NGC 4649 extends out to the projected distance of NGC 4647 with a suggestion of a bow shock between the two \\citep{randall04}. If the galaxies are approaching each other, a lower limit on the relative velocities of the gas systems is $V_{4647}+V_{rot}\\cos i-V_{4649}$=400 km s$^{-1}$. An enhancement in the plasma density of 0.002 particles cm$^{-3}$ on the eastern side of NGC 4647 would produce the necessary pressure to account for the observed asymmetry.  Using the density profile of the X-ray halo of NGC 4649 \\citep{bri97}, we estimate that this particle density occurs at roughly 12 kpc from the center of NGC 4649. The ISM of NGC 4647 may also be compressed by the isotropic pressure of the hot halo gas; again, the thermal pressure of NGC 4649 is $2.4\\times 10^4\\:\\mbox{cm}^{-3}$ at a distance of $\\sim 12$ kpc from the center of the galaxy \\citep{bri97}. The two galaxies would have to be no farther apart than their current projected separation for these mechanisms to drive the molecular asymmetry.  The difficulty with appealing to a close approach of NGC 4647 and NGC 4649 to explain the kinematic and pressure asymmetries is, again, it seems unlikely that the disk of NGC 4647 could survive the tidal forces and remain unscathed. We take the mass of NGC 4649 to be at least $10^{12}$ \\solmass\\ based on the total $V$ magnitude and a $V$-band mass-to-light ratio of 9.5--16 \\citep{debruyne01}.  Strictly speaking, the dynamical analysis of \\citet{debruyne01} measures the mass-to-light ratio within the effective radius so that this dynamical mass of $10^{12}$ \\solmass\\ is probably an underestimate.  For galaxy separations as small as 30 kpc, the tidal acceleration at the edge of NGC 4647's HI disk is two-thirds as large as the spiral's own centripetal acceleration. We therefore consider it extremely unlikely that NGC 4647 and NGC 4649 are presently receding from a past encounter at a few tens of kpc.  It may be possible that they are still approaching each other.  Alternatively, if something (NGC 4649, or an off-center dark matter halo) generates a mild $m=1$ mode in the potential of NGC 4647 as described in section \\ref{lopsided}, the gas must speed up and slow down as it executes its nearly circular orbit.  The ram pressure of gas flowing down into the trough in the potential may be sufficient to account for the expected asymmetry in pressure.  This possibility is especially appealing since it naturally connects the asymmetries in the velocities, the gas surface density, and the molecular fraction. The ram pressure of such a flow is \\begin{equation} P_{ram} = \\frac{1}{2}\\mu m_{\\mathrm{H}} n v^2 = 3\\times 10^{4}~ \\mbox{cm}^{-3}\\mbox{ K} \\left(\\frac{n_{\\mathrm{H2}}}{\\mbox{ cm}^{-3}}\\right) \\left(\\frac{v_{flow}}{10 \\mbox{ km s}^{-1}}\\right)^2 \\end{equation} where $v_{flow}$ is the relative velocity at which gas approaches the potential minimum. For flow velocities comparable to the observed velocity differences on the two sides of the galaxy, the combination of the ram pressure with the ICM and the pressure of the gas accumulating in the $m=1$ potential could explain the observed gas distributions.  Finally, since the implied hydrostatic pressure scales linearly with the gas velocity dispersion, we note that the molecular asymmetry could be explained without recourse to an additional source of pressure if the gas velocity dispersion is higher on the east side of the galaxy than on the west. However, to account for the molecular asymmetry entirely, the velocity dispersion would have to be $>20$ \\kms\\ on the east side. Such high dispersions are ruled out by our observations (section \\ref{veldisp}).  \\subsection{Remaining Questions}\\label{leftoverQ}  How long the asymmetries in the kinematics and the gas distribution of NGC 4647 will last depends critically on their causes.  If the disturbance is temporary, such as a fly-by of NGC 4649, then the spiral can be expected to settle back to a more symmetric state eventually.  Alternatively, an offset dark matter halo might produce a long-lasting asymmetric state. The discussions above do not provide a definitive answer to the source of the asymmetries in NGC 4647, but a variety of observational and simulation-based tests can help confirm or rule out various processes.  Part of the difficulty with understanding NGC 4647 is that the separation between NGC 4647 and NGC 4649 is still unknown (as is the true relative velocity of the two).  However, if the asymmetries of NGC 4647 are due to a galaxy-galaxy interaction then we would expect the stellar kinematics of NGC 4649 to reflect that interaction.  N-body simulations could be used to predict the character and the magnitude of those interaction signatures.  We have also discussed whether hydrodynamic processes such as ram pressure stripping or off-center gas accretion could be responsible for the kinematic and gas distribution features of NGC 4647, but we do not know the history of NGC 4647's interaction with the ICM.  The importance of hydrodynamic processes could be assessed by comparing the gas kinematics to the kinematics of the old stellar populations.  If the stellar kinematics are symmetric, then the asymmetries in the gas can probably be attributed to temporary hydrodynamic effects. In addition, hydrodynamic simulations could be used to test whether the ensemble of properties displayed by NGC~4647 implies the presence of both ram pressure and gravitational disturbances (as \\citet{vollmer03} has suggested for NGC 4654).  A comparison of the $m=1$ asymmetry in the gas surface density to the models of \\citet{swaters99} suggested that the lopsided gas distribution might contribute to the observed kinematic asymmetry in the gas.  That idea could be tested with appropriate hydrodynamic simulations. In addition, simulations that can follow heating, cooling, and formation and dissociation of molecules could investigate whether the conversion from HI into \\htoo\\ follows naturally from the differences in the gas velocity from one side to the other.  H$\\alpha$ images of NGC 4647 are noticeably asymmetric (Figure \\ref{tilefig}), and we have shown that the H$\\alpha$ surface brightness roughly traces the CO surface brightness.  If the \\htooco\\ conversion factor is roughly constant throughout the galaxy, the star formation efficiency is constant as well. Yet the epicyclic frequency, an important component of a local gravitational stability analysis, may be significantly different on the two sides of the galaxy. More detailed analyses of the stability of this asymmetric disk could give valuable insight into large-scale star formation processes.  The pressure model of \\citet{br04} is consistent with our observations of NGC 4647, to the extent that there are several viable candidate sources for the implied extra pressure on the east side of the galaxy. Further progress in testing the model could be stimulated by a better knowledge of how many spirals show relatively symmetric atomic gas but strongly asymmetric molecular gas.  We note that NGC 4647 is actually classified as kinematically regular (non-disturbed) by \\citet{rubin99}.  Its kinematic asymmetries are indeed milder than those of many Virgo spirals, as one can see by perusing the rotation curves presented by those authors.  But the asymmetries of NGC 4647 are more obvious in two-dimensional velocity field data than in one-dimensional longslit data. The properties of NGC 4647 therefore suggest that the true incidence of disturbances among Virgo spirals is higher than the 50\\% measured by \\citet{rubin99}.  Furthermore, NGC 4647 would probably not have been identified as asymmetric on the basis of its global HI profile alone (Figure \\ref{spectra}).  The galaxy therefore suggests that the incidence of asymmetric gas distributions may be higher than the 50\\% rate inferred by \\citet{rs94}.  It is only the combination of the HI and the CO data that revealed the true peculiarities in its gas distribution. Likewise, it is the combination of these two gas phases that will give the tightest constraints on the galaxy's history. Additional matched resolution HI and CO maps of Virgo spirals would be very helpful for understanding the evolution of the ISM in mildly disturbed galaxies."
    },
    "0606/astro-ph0606626_arXiv.txt": {
        "abstract": "Recent detections of the Integrated Sachs-Wolfe effect through the correlation of the cosmic microwave background temperature anisotropy with traces of large scale structure provided independent evidence for the expansion of the universe being dominated by something other than matter. Even with perfect data, statistical errors will limit the accuracy of such measurements to worse than $10$\\%. On the other hand, the extraordinary sensitivity of the ISW effect to the details of structure formation should help to make up for the lack of precision. In these conference proceedings I discuss the extent to which future ISW measurements can help in testing the physics responsible for the observed cosmic acceleration. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606189_arXiv.txt": {
        "abstract": "Past years have brought an increasingly wider recognition of the ubiquity of relativistic outflows (jets) in galactic nuclei, which has turned jets into an effective tool for investigating the physics of nuclear regions in galaxies. A brief summary is given here of recent results from studies of jets and nuclear regions in several active galaxies with prominent outflows. ",
        "introduction": "\\label{lobanov2:sec1}  Substantial progress achieved during the past decade in studies of active galactic nuclei (see~\\cite{lobanov2006} for a review of recent results) has brought an increasingly wider recognition of the ubiquity of relativistic outflows (jets) in galactic nuclei~\\cite{falcke2001,zensus1997} turning them into an effective probe of nuclear regions in galaxies~\\cite{lobanov2005}. Emission properties, dynamics, and evolution of an extragalactic jet are intimately connected to the characteristics of the supermassive black hole, accretion disk and broad-line region in the nucleus of the host galaxy~\\cite{lobanov2006}.  The jet continuum emission is dominated by non-thermal synchrotron and inverse-Compton radiation~\\cite{unwin1997}. The synchrotron mechanism plays a more prominent role in the radio domain, and the properties of the emitting material can be assessed using the turnover point in the synchrotron spectrum~\\cite{lobanov1998b}, synchrotron self-absorption~\\cite{lobanov1998a}, and free-free absorption in the plasma~\\cite{kadler2004,walker2000}.  High-resolution radio observations access directly the regions where the jets are formed~\\cite{junor1999}, and trace their evolution and interaction with the nuclear environment~\\cite{mundell2003}. Evolution of compact radio emission from several hundreds of extragalactic relativistic jets is now systematically studied with dedicated monitoring programs and large surveys using very long baseline interferometry (such as the 15\\,GHz VLBA\\footnote{Very Long Baseline Array of National Radio Astronomy Observatory, USA} survey~\\cite{kellermann2004} and MOJAVE~\\cite{lister2005}).  These studies, combined with optical and X-ray studies, yield arguably the most detailed picture of the galactic nuclei. Presented below is a brief summary of recent results from studies outlining the relation between jets, supermassive black holes, accretion disks and broad-line regions in prominent active galactic nuclei (AGN).   \\begin{table}[t] \\caption{Characteristic scales in the nuclear regions in active galaxies} \\label{lobanov2:tb1} \\begin{center} \\begin{tabular}{rccccc}\\hline\\hline & $l$ & $l_8$ & $\\theta_\\mathrm{Gpc}$ & $\\tau_c$ & $\\tau_\\mathrm{orb}$ \\\\ & [$R_\\mathrm{g}$] & [pc] & [mas]& [yr] & [yr] \\\\ \\hline Event horizon:           &1--2          &$10^{-5}$           &$5\\times 10^{-6}$    &0.0001     & 0.001 \\\\ Ergosphere:              &1--2          &$10^{-5}$           &$5\\times 10^{-6}$    &0.0001     & 0.001 \\\\ Accretion disk:          &10$^1$--10$^3$&$10^{-4}$--$10^{-2}$&$0.005$              &0.001--0.1 & 0.2--15 \\\\ Corona:                  &10$^2$--10$^3$&$10^{-3}$--$10^{-2}$&$5\\times 10^{-3}$    &0.01--0.1& 0.5--15 \\\\ Broad line region:       &10$^2$--10$^5$&$10^{-3}$--1        &$0.05$               &0.01--10   & 0.5--15000 \\\\ Molecular torus:         &$>$10$^5$     &$>$1                &$>$$0.5$             &$>$10      & $>$15000 \\\\ Narrow line region:      &$>$10$^6$     &$>$10               &$>$5                 &$>$100     & $>$500000 \\\\ Jet formation:           &$>$10$^2$     &$>$$10^{-3}$        &$>$$5\\times 10^{-4}$ &$>$0.01    & $>$0.5 \\\\ Jet visible in the radio:&$>$10$^3$     &$>$$10^{-2}$        &$>$$0.005$           &$>$0.1     & $>$15 \\\\ \\hline \\end{tabular} \\end{center} {\\bf Column designation:} $l$ -- dimensionless scale in units of the gravitational radius, $G\\,M/c^2$; $l_8$ -- corresponding linear scale, for a black hole with a mass of $5\\times 10^8\\,$M$_{\\odot}$; $\\theta_\\mathrm{Gpc}$ -- corresponding largest angular scale at 1\\,Gpc distance; $\\tau_c$ -- rest frame light crossing time; $\\tau_\\mathrm{orb}$ -- rest frame orbital period, for a circular Keplerian orbit. Table is reproduced from~\\cite{lobanov2006} \\end{table} ",
        "conclusions": "Extragalactic jets are an excellent laboratory for studying physics of relativistic outflows and probing conditions in the central regions of active galaxies. Recent studies of extragalactic jets show that they are formed in the immediate vicinity of central black holes in galaxies and carry away a substantial fraction of the angular momentum and energy stored in the accretion flow and rotation of the black hole. The jets are most likely collimated and accelerated by electromagnetic fields. Relativistic shocks are present in the flows on small scales, but dissipate on scales of $\\lesssim 100$\\,pc. Plasma instabilities dominate the flows on larger scales.  Convincing observational evidence exists connecting ejections of material into the flow with instabilities in the inner accretion disks. In radio-loud objects, continuum emission from the jets may also drive broad emission lines generated in sub-relativistic outflows surrounding the jets.  Magnetically confined outflows may preserve information about the dynamics state of the central region, allowing detailed investigations of jet precession and binary black hole evolution to be made. This makes studies of extragalactic jets a powerful tool for addressing the general questions of physics and evolution of nuclear activity in galaxies."
    },
    "0606/astro-ph0606140_arXiv.txt": {
        "abstract": "We study influence of a galactic central supermassive black hole (SMBH) binary on gas dynamics and star formation activity in a nuclear gas disk by making three-dimensional Tree+SPH simulations. Due to orbital motions of SMBHs, there are various resonances between gas motion and the SMBH binary motion. We have shown that these resonances create some characteristic structures of gas in the nuclear gas disk, for examples, gas elongated or filament structures, formation of gaseous spiral arms, and small gas disks around SMBHs. In these gaseous dense regions, active star formations are induced. As the result, many star burst regions are formed in the nuclear region. ",
        "introduction": "In the recent high-resolution observations (e.g., $Chandra$ $X-ray$ $Observatory$), the evidence of supermassive black hole (here after we use SMBH for them) binaries are shown in several galaxies, e.g., in NGC 6240 \\citep{kom03}, Arp 220 \\citep{cle02}, M83 \\citep{sak03, mas05}, and 3C 66B \\citep{sud03}. Particularly, NGC 6240 has been well observed in the wide range of wave length \\citep{tac99,kom03}. In the high-resolution X-ray observation by $Chandra$ $X-ray$ $Observatory$, strong two peaks of hard X-ray are detected in the galactic center and this is the strong evidence of a SMBH binary \\citep{kom03}. By the radio continuum, the near-infrared, and the soft X-ray observations of this galaxy, nuclear star burst has been indicated \\citep{lir02, pas04}. Nuclear gas rich disks have been observed around the SMBH binary in NGC 6240 \\citep{tac99} and Arp 220 \\citep{sco97, sak99}.  The SMBH binaries in galactic central regions are expected to be formed in merging process of galaxies each of which has a SMBH in its galactic center. After galaxies merge, SMBHs sink into the center of the merging galaxy by dynamical friction between SMBHs and field stars \\citep{ebi01, esc04, esc05}. These SMBHs will form a SMBH binary and finally merge due to emission of gravitational wave \\citep{mat04, eno04, esc04, esc05}. In the process, \\citet{esc04, esc05} have shown an important role of dynamical interaction between SMBHs and gas, especially in very gas rich regions.  \\citet{kaz05} have made simulations of the merging process of two disk galaxies with SMBHs at each galactic center. Their results indicate that much gas flows into the center of merging galaxy and star burst is triggered. They also show that in the process, a SMBH binary is formed in the center of merging galaxy and a nuclear gas disk with radius of $1-2$ kpc is formed at the galactic center. In the gas disk, effects of the SMBH binary on a nuclear gas disk are not studied yet and it may be very important for star formation in the galactic center.  Since the gravitational potential of a SMBH binary has non-axisymmetric component, we expect that gas motion in a nuclear gas disk is strongly influenced by a SMBH binary. Change of gravitational potential due to orbital motions of SMBHs may induce resonance phenomena in the disk as in barred galaxies. In barred galaxies, gas motion is influenced by resonances \\citep{ath92}. It was indicated that these resonances trigger nuclear star burst in the barred galaxies \\citep{fuk91,wad92,elm94,wad95,fuk98}. In a SMBH binary case, similar resonances may trigger active star formation in the nuclear gas disk. Moreover, it is interesting that the SMBH binary may yield some peculiar gaseous features in the nuclear disk which can be used as an evidence of a SMBH binary.  In this paper, we study the influence of a galactic central SMBH binary on gas motion in a nuclear gas disk by hydrodynamic simulations by using 3-dimensional Tree+SPH code. We show that the resonances due to a SMBH binary trigger formation of gas concentrations in the disk and as the results star formation rate increases.  In section 2, we present our simulation model. In section 3, we present the results of our simulations. In section 4, we summarize our results and give some discussions. ",
        "conclusions": "We study influence of the galactic central SMBH binary on gas dynamics in the nuclear gas disk by numerical simulations. We calculate various cases for initial SMBH orbits which are the circular orbit case and the elliptical orbit cases. We have shown that in the all cases, the SMBH binary has large influence on gas motion in the gas disk. The SMBH binary induces some resonances on gas motion in a nuclear gas disk. Due to these resonances, various dense gas structures are formed in the nuclear gas disk and gaseous spiral arms are formed near the vicinity of Outer Lindblad Resonance. In these dense gas regions, star formation becomes very active. In the high eccentric orbit cases of the SMBH binary, gaseous narrow filaments and dense clump structures are developed well, and active star formation occurs in these regions. These features can be strong evidence of existence of a SMBH binary. It is very interesting to compare these features with very high-resolution observations of galaxies which are proposed to have a SMBH binary. Dense gas structures and distribution of star formation cites will inform us dynamical state of the SMBH binary. It should be noticed that these features are not appeared in the case of which the SMBH mass is smaller than about $1\\times 10^8$ $M_{\\odot}$ in our model. We note that when we compare whole SFR of ultra luminous infrared galaxies, the SFRs induced by the SMBH binary are not high. However, the SFR is as high as nuclear star bursts in nearby star burst galaxies.  In our simulations, gaseous ridges are formed by shocks. In the elliptical cases of SMBHs, there are the collisions between gas clumps in the galactic center. Shocks are exited by these collisions of gas clumps. H$_2$ emission line is expected to be excited in the shocks. \\citet{van93} and \\citet{sug97} have observed bright H$_2$ emission line in the galactic central region of NGC 6240 and they conclude that the H$_2$ emission is excited by shock.  From our numerical simulations, small gas disks are formed around SMBHs in the elliptical orbit cases. Such gas disks around SMBHs have been observed in Arp 220 which has the double nuclei in the galactic center \\citep{sak99}. In our simulations, star formation is very active in the small gas disks around SMBHs. The active star formation in the gas disks around SMBHs may correspond to radio continuum sources observed around SMBHs in NGC6240 \\citep{tac99}. $Chandra$ $X-ray$ $Observatory$ observed hard X-ray from double nuclei in NGC 6240. It is possible to excite AGN activity by gas accretion onto SMBHs in the small gas disks. If AGN activities are highly excited and AGN feedback becomes very strong, active star formation will be quenched \\citep{mat05}.  Since active star formation occurs in very compact regions in the highly eccentric elliptical orbit cases, we expect that these newly formed stars concentrate in compact massive star clusters. If these star clusters interact gravitationally with the SMBH binary, the interaction may induce losing of orbital angular momentum of SMBHs due to the unstableness of three body problem in which SMBHs and the star cluster interact with each other and ejection of the star cluster occurs. If these star clusters are massive enough, this process may be very effective and the binary can evolve to a more tightly binding state. This process may have an important role in merging process of SMBHs.  \\citet{esc05} studied the effect of hydrodynamic drag force by dense gas on evolution of a SMBH binary by numerical simulations. They have shown that after SMBHs gradually fall into the galactic central dense gas region by the dynamical friction, effect of hydrodynamical drag becomes very effective in the central region. They suggested that finally SMBHs can be close enough to merge by the hydrodynamic effect. In their simulations, they don't consider effects of radiative cooling and star formations on gas. By the effect of radiative cooling, many dense clump structures will be formed and those distribution is more complicated. The dense clumps may interact with SMBHs and play an important role in the coalesce of SMBHs. Moreover, active star formation will occur in the clumps and gas mass will decrease in the galactic central region. After the active star formation, it is not clear that in the galactic center gas remains enough for hydrodynamic interaction with SMBHs or not. It is needed to make simulations of evolution of a SMBH binary in more realistic model.  In our simulations, we didn't consider the dynamical friction between field stars and SMBHs. The dynamical friction induces decay of orbital radius of SMBHs. If timescale of dynamical friction is larger than timescale of rotation motion of SMBHs, the resonances between SMBH motions and gas motion will be effective. In this case, similar process appeared in our simulations will occur. On the other hand, if the dynamical friction timescale is smaller than the rotation timescale, orbits of SMBHs shrink very rapidly and the resonance phenomena are not important.  We have assumed that initially gas disk is a circularly rotating. However, since a galaxy with a SMBH binary is expected to be formed due to merging of galaxies with SMBHs, gas motion is more complex in merging galaxies. To simulate more realistic evolution of a galaxy with a SMBH binary, we will study merging process of galaxies with SMBHs. In this process, radiation drag \\citep{kaw05} and the influence of AGN feedback \\citep{spr05} should be considered."
    },
    "0606/astro-ph0606695_arXiv.txt": {
        "abstract": "{A new sample of hard X-ray sources in the Galactic Plane is being revealed by the regular observations performed by the {\\it INTEGRAL} satellite. The full characterization of these sources is mandatory to understand the hard X-ray sky. Here we report new multifrequency radio, infrared and optical observations of the source \\object{IGR~J18027$-$1455}, as well as a multi-wavelength study from radio to hard X-rays. The radio counterpart of \\object{IGR~J18027$-$1455} is not resolved at any observing frequency. The radio flux density is well fitted by a simple power law with a spectral index $\\alpha=-0.75\\pm0.02$. This value is typical of optically thin non-thermal synchrotron emission originated in a jet. The NIR and optical spectra show redshifted emission lines with $z=0.034$, and a broad H$\\alpha$ line profile with FWHM $\\sim$3400~km~s$^{-1}$. This suggests an Active Galactic Nucleus (AGN) of type~1 as the optical counterpart of \\object{IGR~J18027$-$1455}. We confirm the Seyfert~1 nature of the source, which is intrinsically bright at high energies both in absolute terms and when scaled to a normalized 6~cm luminosity. Finally, comparing its X-ray luminosity with isotropic indicators, we find that the source is Compton thin and AGN dominated. This indicates that {\\it INTEGRAL} might have just seen the tip of the iceberg, and several tens of such sources should be unveiled during the course of its lifetime. ",
        "introduction": "\\label{introduction}  Unidentified high energy sources have been a subject of interest from the early days of the {\\it COS-B} era. In the 1990s, with the advent of X-ray/$\\gamma$-ray satellites like {\\it ASCA} and {\\it CGRO} the number of sources with unidentified counterparts at other frequencies increased considerably. During the first year of observations, the IBIS/ISGRI instrument on board the {\\it INTEGRAL} satellite \\citep{winkler03} detected 123 hard X-ray/$\\gamma$-ray point sources, 28 of which had no clear identification with known objects in other ranges of the electromagnetic spectrum \\citep{bird04}. These X-ray/$\\gamma$-ray emitters could be high or low mass X-ray binaries, radio quiet pulsars, clusters of galaxies, or a significant fraction of any class of AGNs heavily obscured, at few keV, by the absorbing material of the galactic plane. The possibility that several unidentified IBIS sources were of extragalactic nature was early suggested by some authors (\\citealt{ribo04}; \\citealt{combi04}; \\citealt{masetti04a}; \\citealt{masetti04b}; \\citealt{bassani04}; \\citealt{combi05}).  The source \\object{IGR~J18027$-$1455} is one of such sources. It was discovered in the energy range from 20 to 100~keV during 769~ks of observations. Looking for possible counterparts \\cite{combi04} found two weak point-like radio sources from the 20~cm NRAO VLA Sky Survey (NVSS, \\citealt{condon98}) inside its 2 arcmin-radius position error circle (see Fig.~\\ref{fig:nvss}). One of them, \\object{NVSS~J180247$-$145451}, lies inside and near the edge of the 2$\\sigma$ position error circle of the faint {\\it ROSAT} X-ray source \\object{1RXS~J180245.5$-$145432} \\citep{voges00}, which is the only soft X-ray source well within the IBIS/ISGRI error circle. In addition, inside the 2$\\sigma$ position error ellipse of this radio source, it is located an extended near infrared (NIR) source, \\object{2MASS~J18024737$-$1454547} \\citep{cutri03,skrutskie06}, with standard aperture magnitudes $J=12.78\\pm0.01$, $H=11.52\\pm0.01$, and $K_s=10.72\\pm0.01$. Its optical counterpart has average magnitudes $B=19.3\\pm1.0$, $R=14.9\\pm0.8$ and $I=13.8\\pm0.5$ in the USNO-B1.0 catalog \\citep{monet03}. The photometry of the NIR/optical counterpart is not consistent with a stellar spectrum \\citep{combi04}. On the basis of spectroscopic optical observations \\cite{masetti04b} have tentatively classified this source as a Seyfert~1 galaxy at redshift $z=0.035\\pm0.001$.  \\begin{figure}[t!] % \\center \\resizebox{1.0\\hsize}{!}{\\includegraphics[angle=0]{f1.eps}} \\caption{Image of the NVSS data obtained with the VLA at 20~cm on 1997 October 13 around \\object{IGR~J18027$-$1455}. The image size is 7\\farcm2$\\times$5\\farcm5. Contours represent $-$3, 3, 5, 8, 11, 15, 18, and 22 times the rms noise level of 0.5~mJy~beam$^{-1}$. The circle in the bottom left corner represents the 45 arcsec of Full Width at Half Maximum (FWHM) of the convolving beam. Two NVSS sources fall inside the 90\\% error circle in position of \\object{IGR~J18027$-$1455}, and one of them is within the 2$\\sigma$ uncertainty error circle of a {\\it ROSAT} source.} \\label{fig:nvss} \\end{figure}  An important characteristic of AGNs is that they radiate over a wide range of frequencies, from radio to gamma-rays. For this reason, multi-wavelength observations are an important tool to discriminate between objects of different classes. Here we report multi-wavelength observations of \\object{IGR~J18027$-$1455} and discuss the obtained results. The structure of the paper is as follows. In Sect.~\\ref{observations} we describe our radio continuum, NIR and optical observations and present the results. In Sect.~\\ref{discussion} we discuss the nature of all the detected multi-wavelength emissions and we summarize our conclusions in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{discussion}  The position agreement between the \\object{IGR~J18027$-$1455}, \\object{1RXS~J180245.5$-$145432} and \\object{NVSS~J180247$-$145451} indicates that these are the multi-wavelength manifestations of the same source radiating at different bands of the electromagnetic spectrum. Moreover, the recently published second IBIS/ISGRI catalog \\citep{bird06} lists a position and error circle for \\object{IGR~J18027$-$1455} that clearly exclude the other radio source, namely \\object{NVSS~J180239$-$145453}, as a possible counterpart (even at the 90\\% confidence level). The NIR source \\object{2MASS~J18024737$-$1454547} has a position in agreement within errors with our new precise radio position of \\object{NVSS~J180247$-$145451}. Its optical counterpart is \\object{USNO-B1.0~0750-0506536}. Both the NIR and optical objects are clearly extended (this work and \\citealt{masetti04b}, respectively). In addition, the redshifted emission lines seen in the NIR and optical spectra reveal an extragalactic object at a redshift of $z=0.034\\pm0.001$, with a broad H$\\alpha$ emission line with FWHM~$\\simeq3400\\pm300$~km~s$^{-1}$ (these values are compatible with those of \\citealt{masetti04b}, although they reported a slightly narrower emission line with FWHM~$\\sim$2700~km~s$^{-1}$). We are therefore seeing the broad line region (BLR), and the object is classified as a type~1 AGN.  \\begin{figure}[t!] % \\center \\resizebox{1.0\\hsize}{!}{\\includegraphics[angle=0]{f5.eps}} \\caption{Averaged optical spectrum of the optical counterpart of \\object{IGR~J18027$-$1455} acquired with CAFOS on the 2.2m telescope at CAHA on 2004 July 19. The spectrum has been smoothed with a Gaussian filter. The identified emission lines are indicated. A broad H$\\alpha$ line dominates the spectrum. A redshift of $z=0.034\\pm0.001$ is obtained.} \\label{fig:caha} \\end{figure}  The steep radio spectral index ($\\alpha=-075\\pm0.02$) strongly supports a non-thermal emission mechanism of synchrotron nature. This is clearly compatible with optically thin extended jet emission from an extragalactic source. On the other hand, the NVSS flux density of the source at 1.4~GHz is $10.5\\pm0.6$~mJy, to be compared with our measurement of $7.5\\pm0.3$~mJy. Both values are only marginally consistent at the 3$\\sigma$ level, suggesting that the source is variable at radio wavelengths. We note that 0.61~GHz (49~cm wavelength) observations conducted 3 months later provided a detection at a level of $5.0\\pm0.35$~mJy \\citep{pandey06}, either supporting the variability of the source or indicating that there is a low frequency turnover.  The NIR spectrum is very similar to other well studied Seyfert~1 galaxies such as \\object{NGC~863} or \\object{Mrk~335} \\citep{rodriguez02}, although the poor signal-to-noise ratio of our NIR observations is not enough to discriminate weak lines as in these cases. It is interesting to note that the permitted \\ion{O}{i} $\\lambda$11287 line is a feature completely associated with the BLR of Seyfert galaxies. In our case this line is marginally detected, as in \\object{NGC~863} \\citep{rodriguez02}. The non-detection of the Br$\\gamma$ line seems to suggest that thermal UV heating is not important, as it also happens in the case of \\object{NGC~1097} \\citep{reunanen02}.  The multiwavelength properties of \\object{IGR~J18027$-$1455} strongly support an AGN nature and more specifically a type~1 Seyfert galaxy. In order to compare the broadband emission of the object with that one of the mean for type~1 Seyfert galaxies, we have determined the nuclear spectral energy distribution (SED), from the radio to the gamma-ray band. The observations used to build the SED have been discussed in Sect.~\\ref{observations}. The observed magnitudes in the NIR and optical bands were corrected for reddening from Galactic extinction based on the estimated hydrogen column density $N_{\\rm H}=(5.0\\pm1.0) \\times 10^{21}$ cm$^{-2}$ \\citep{dickey90} and the \\cite{predehl95} relationship $A_V = (5.59\\pm0.10) \\times 10^{-22}~N_{\\rm H}$, which gives $A_V=2.8\\pm0.6$ magnitudes. The transformation between the absorption in the optical $A_V$ and that at other wavelengths was computed according to the \\cite{rieke85} interstellar extinction law.  At soft X-ray energies, between 0.1--2.4~keV, the flux was obtained using the {\\it ROSAT}/PSPC count rate of $(3.25\\pm1.39)\\times10^{-2}$~count~s$^{-1}$ \\citep{voges00} and a photon index of $\\Gamma=+1.9\\pm01$, typical of Seyfert~1 galaxies \\citep{malizia03}. We used the web based tool PIMMS~v3.7a\\footnote{\\tt http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html} and the $N_{\\rm H}$ value given above to obtain an unabsorbed flux of $2.4^{+2.0}_{-1.3}\\times10^{-12}$~erg~cm$^2$~s$^{-1}$ (propagating all possible uncertainties). In addition, extrapolation of the {\\it ROSAT}/PSPC count rate with the same photon index and absorption, considering all possible uncertainties, provides a flux of $1.5^{+1.4}_{-0.9}\\times10^{-12}$ erg~cm$^{-2}$~s$^{-1}$ in the 2--10~keV energy range.  The average fluxes detected by {\\it INTEGRAL} in the 20--40 and 40--100~keV energy ranges are $3.0\\pm0.2$ and $3.3\\pm0.3$~mCrab, respectively \\citep{bird06}. These can be converted to the cgs fluxes (assuming a Crab-like spectrum and the values in \\citealt{bird06}) $(2.3\\pm0.2)\\times10^{-11}$~erg~cm$^2$~s$^{-1}$ and $(3.1\\pm0.3)\\times10^{-11}$~erg~cm$^2$~s$^{-1}$, respectively, which provide a total flux in the 20--100~keV range of $(5.4\\pm0.4)\\times10^{-11}$~erg~cm$^2$~s$^{-1}$. We note that an analysis with more {\\it INTEGRAL} data reveals the following values in the 20--100~keV range \\citep{bassani06}: $2.6\\pm0.1$~mCrab and $(4.4\\pm0.2)\\times10^{-11}$~erg~cm$^2$~s$^{-1}$. As can be seen there are hints of variability in the hard X-ray/gamma-ray flux of \\object{IGR~J18027$-$1455}, and the average of $(4.9\\pm0.5)\\times10^{-11}$~erg~cm$^2$~s$^{-1}$ will be used when plotting the SED.  To compute the monochromatic luminosities we have adopted the cosmological parameters from \\cite{spergel03}: $H_{0}=71$~km~s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73$ and $\\Omega_{\\rm m}=0.27$. Using our measured redshift of $z=0.034\\pm0.001$ we obtain\\footnote{\\tt http://www.astro.ucla.edu/$\\sim$wright/CosmoCalc.html} a luminosity distance of $147\\pm5$~Mpc for \\object{IGR~J18027$-$1455}, leading to a hard X-ray luminosity of $(1.3\\pm0.1)\\times10^{44}$~erg~s$^{-1}$. The source is one of brightest Seyfert~1 galaxies detected so far: it is brighter than any of the Seyfert~1 galaxies detected with {\\it BeppoSAX} \\citep{panessa04t}, and the 4th brightest one among the 14 detected with {\\it INTEGRAL} \\citep{bassani06}.  We show in Fig.~\\ref{fig:sed} the overall nuclear SED of \\object{IGR~J18027$-$1455} in a $\\log(\\nu)$--$\\log(\\nu L_{\\nu})$ representation, but normalized at 6~cm for comparison with the mean SEDs of Seyfert galaxies \\citep{panessa04t,panessa04}. The real luminosities of \\object{IGR~J18027$-$1455} are 1.21 dex lower than those shown. Although in the NIR domain we have used nuclear magnitudes, the optical magnitudes ($I$ and $B$ from USNO-B1.0 and $R=16.55\\pm0.01$ from \\citealt{masetti04b}) could be strongly contaminated by the host galaxy due to a limited angular resolution of $\\sim$1\\arcsec~pixel~$^{-1}$. We have not plotted the extrapolated 2--10~keV flux. Although the archival {\\it IRAS} data does not allow us to detect the far infrared bump, the NIR/{\\it ROSAT} data clearly show a SED typical of Seyfert~1 galaxies. However, the {\\it INTEGRAL} data show that the source is clearly high-energy bright. Therefore, \\object{IGR~J18027$-$1455} is not only among the brightest type~1 Seyfert galaxies at high energies in absolute terms, but also when normalized to the 6~cm luminosity.  \\begin{figure}[t!] % \\center \\resizebox{1.0\\hsize}{!}{\\includegraphics[angle=0]{f6.eps}} \\caption{Overall nuclear spectral energy distribution of \\object{IGR~J18027$-$1455} (filled symbols) from the radio to the hard X-ray/gamma-ray band. Optical stands for the $I$ and $B$ USNO-B1.0 and $R$ magnitude from \\cite{masetti04b}. The SED has been normalized at 6~cm (assuming $\\log(\\nu L_{\\nu}[{\\rm erg~s}^{-1}])=39.75$) in order to compare between type 1 and type~2 Seyfert galaxies (open symbols; adapted from \\citealt{panessa04t}; vertical bars are errors of the averages, not the standard deviation of the samples). The real luminosities of our target source are 1.21 dex lower than those shown. The SED of \\object{IGR~J18027$-$1455} resembles the average one for Seyfert~1 galaxies, although it is brighter in the hard X-ray domain. Optical data is probably contaminated by the host galaxy.} \\label{fig:sed} \\end{figure}  In the basic scheme for unification of AGNs, Seyfert galaxies are divided in two class. Those that have narrow forbidden lines and a BLR in their optical spectrum (Seyfert~1) and those that only have narrow lines (Seyfert~2). While the broad lines originate near the central massive black hole located at $\\leq0.1$~pc, narrow lines arise far from the nuclear engine at a distance $\\leq100$~pc. Specifically they are the same type of object but, according to the standard model, in Seyfert~2 galaxies the BLR is obscured by a molecular torus \\citep{antonucci93}. For this reason, the majority of these objects are Compton thick, that is, the medium is thick to Compton scattering so that the transmitted component is dramatically suppressed below 10~keV down to the NIR domain.  We can thus further check the Seyfert~1 nature of \\object{IGR~J18027$-$1455} and its agreement with unification schemes by comparing its soft X-ray luminosity with isotropic indicators. This allows us to discriminate if starburst or AGN is the dominant component, and at the same time to assess if the source is Compton thin or Compton thick. If the presence of a molecular torus around the central region is important, the X-ray emission coming from the central engine will be negligible and it should be coming from a more extended zone like the NLR or a starburst region. In this case, the column density could be obtained indirectly from the flux ratio between the X-ray flux and isotropic emission measurements like the [\\ion{O}{iii}]$\\lambda$5007 and far-infrared fluxes. Both are good isotropic indicators, while [\\ion{O}{iii}]$\\lambda$5007 emission is produced by photons originated in the central nucleus, infrared emission is mainly associated to star-forming activity, and therefore produced in a larger region than that of the molecular torus.  To compute the [\\ion{O}{iii}]$\\lambda$5007 flux we have used our normalized optical spectrum and the average optical spectrum of \\cite{masetti04b}. The line is clearly detected at $\\lambda$=5175~\\AA. Smoothing our spectrum with a Gaussian function, its equivalent width is $10\\pm2$~\\AA\\ and the [\\ion{O}{iii}]$\\lambda$5007 flux is 7.1 $\\times$ 10$^{-16}$ erg~cm$^{-2}$~s$^{-1}$. This flux has been corrected for extinction using the equation given by \\cite{bassani99}. Using a H$\\alpha \\cong$ 8.5 $\\times$ 10$^{-16}$ erg~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$ and a H$\\beta \\leq$ 0.5 $\\times$ 10$^{-16}$ erg~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$, the observed flux ratio H$\\alpha$/H$\\beta$ $\\geq$ 17 and $F_{\\rm [OIII],~cor} \\geq$ 1.2 $\\times$ 10$^{-13}$ erg~cm$^{-2}$~s$^{-1}$. Since the $F_{\\rm X}$, between 2--10~keV, is 0.6--2.9 $\\times$ 10$^{-12}$ erg~cm$^{-2}$~s$^{-1}$ the $F_{\\rm X}/F_{\\rm [OIII]}$ ratio is between 5 and 24. According to the flux diagnostic diagrams introduced by \\cite{panessa04t} for type 1 and 2 Seyferts, the source is Compton thin. To calculate the far-infrared flux we have adopted the equation of \\cite{mulchaey94}. As a result the infrared flux is $F_{\\rm IR} \\leq$ 8 $\\times$ 10$^{-11}$ erg~cm$^{-2}$~s$^{-1}$. Therefore, the flux ratio $F_{\\rm [OIII]}/F_{\\rm IR} \\geq$ 1.5 $\\times$ 10$^{-3}$. According to \\cite{panessa04t}, this shows that AGN contribution, not starburst, is the dominant component."
    },
    "0606/gr-qc0606041_arXiv.txt": {
        "abstract": "The braneworld model of Dvali-Gabadadze-Porrati (DGP) provides an interesting alternative to a positive cosmological constant by modifying gravity at large distances. We investigate the asymptotic behavior of homogeneous and anisotropic cosmologies on the DGP brane. It is shown that all Bianchi models except type IX isotropize, as in general relativity, if the so called $E_{\\mu\\nu}$ term satisfies some energy condition. Isotropization can proceed slower in DGP gravity than in general relativity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/gr-qc0606088_arXiv.txt": {
        "abstract": "We develop a general formalism for the parameter-space metric of the multi-detector $\\F$-statistic, which is a matched-filtering detection statistic for continuous gravitational waves. We find that there exists a whole \\emph{family} of $\\F$-statistic metrics, parametrized by the (unknown) amplitude parameters of the gravitational wave. The multi-detector metric is shown to be expressible in terms of noise-weighted \\emph{averages} of single-detector contributions, which implies that the number of templates required to cover the parameter space does \\emph{not} scale with the number of detectors. Contrary to using a longer observation time, combining detectors of similar sensitivity is therefore the computationally cheapest way to improve the sensitivity of coherent wide-parameter searches for continuous gravitational waves.  We explicitly compute the $\\F$-statistic metric family for signals from isolated spinning neutron stars, and we numerically evaluate the quality of different metric approximations in a Monte-Carlo study. The metric predictions are tested against the measured mismatches and we identify regimes in which the local metric is no longer a good description of the parameter-space structure. ",
        "introduction": "Continuous gravitational waves (GWs), which would be emitted, for example, by spinning non-axisymmetric neutron stars, or by solar-mass binary systems, are generally expected to be so weak that they will be buried several orders of magnitude below the noise of even the most sensitive detectors. The detection of such signals therefore requires the exact knowledge of their waveform, in order to be able to coherently correlate the data with the expected signal by \\emph{matched filtering}.  In a wide-parameter search for unknown sources, however, we typically only know the family of possible waveforms (or an approximation thereof), parametrized by unknown signal parameters (such as the frequency or sky position of the source). The corresponding parameter space needs to be covered by a finite number of ``templates'' for which a search will be performed. These templates must be placed densely enough, so that for any possible signal, no more than a certain fraction of the signal-to-noise ratio (SNR) is lost at the closest template. On the other hand, coherently correlating the data with every templates is computationally expensive and increases the number of statistical false-alarm candidates. Therefore an \\emph{optimal} covering is desirable, which minimizes the number of templates but guarantees the required ``minimal match''.  In order to solve this \\emph{covering problem}, it is essential to understand the underlying parameter-space structure. Most studies on the construction of optimal template banks have been performed in the context of binary-inspiral searches. It was realized early on that a geometric approach to this problem is the most natural, in particular the introduction of a \\emph{metric} on the parameter space by \\citet{bala96:_gravit_binaries_metric} and \\citet{owen96:_search_templates}, building on the earlier concept of the  ``fitting factor'' introduced by \\citet{apostolatos95:_search_templates}. Note, however, that this definition of the metric differs subtly from the ``canonical'' definition used in the present work (and also in \\cite{krolak04:_optim_lisa}), which is derived directly from the detection statistic (see appendix~\\ref{sec:altern-proj-onto} for more details). The canonical definition of the metric assigns the relative loss of SNR due to an offset in signal parameters as an invariant ``distance'' measure, which can locally be expressed  as a metric tensor. The metric is closely related to the well-known concept of the ``Fisher information matrix'',  which quantifies the statistical errors in the parameter estimation of signals:  the (canonical) metric is identical to the \\emph{normalized} Fisher matrix, even though it describes conceptually rather different aspects of the detection statistic.  A somewhat related question is the \\emph{global} parameter-space structure, which was studied in \\cite{prix05:_circles_sky} for the case of isolated neutron-star signals. This study found that the global structure (the ``circles in the sky'') deviates significantly from the local metric picture. The global structure is relevant, for example, for deciding whether different detection candidates are consistent with the same signal, i.e. whether they are ``coincident candidates''. Obviously, the metric description is the local approximation to this global parameter-space structure.  In this paper we consider gravitational-wave signals that are nearly monochromatic and sinusoidal in the frame of the GW source, and which are of long duration (i.e. typically longer than the observation time $T$). This class of signals is usually referred to as ``continuous waves'', and the prime examples are GWs from non-axisymmetric spinning neutron stars (e.g. see \\cite{prix06:_cw_review} for a review) and stellar-mass binary systems in the LISA frequency band (e.g. see \\cite{krolak04:_optim_lisa}).  The phase of the signal received at the detector is Doppler modulated by the rotation and orbital motion of the detector. The observed phase therefore depends not only on the intrinsic frequency evolution of the signal, but also on its sky position. In addition to the phase modulation, there is a time-varying amplitude modulation, due to the rotating antenna pattern of the detector. This amplitude modulation depends on the polarization angle $\\psi$ and the polarization amplitudes $A_+$ and $A_\\times$ of the GW. However, as shown by \\citet{jks98:_data}, these unknown parameters (together with the initial phase $\\phi_0$), can be eliminated by analytically maximizing the detection statistic. The resulting reduced parameter space includes only the parameters affecting the time evolution of the signal phase, which we will refer to as the ``Doppler parameters''. This amplitude-maximized detection statistic is generally known as the ``$\\F$-statistic'', which has been used in several searches for continuous GWs from spinning neutron stars (e.g. \\cite{lsc04:_psr_j1939,lsc06:_coher_scorp_x,2005CQGra..22S1243A}). After two earlier (partly successful) attempts to generalize the $\\F$-statistic to a coherent network of detectors \\cite{jks98:_data,krolak04:_optim_lisa}, this problem was fully solved more recently by \\citet{cutler05:_gen_fstat}.  Somewhat surprisingly, however, there has not been much work on the metric of the $\\F$-statistic, neither in the single- nor the multi-detector case: the single-detector $\\F$-statistic metric was derived on a formal level by \\citet{krolak04:_optim_lisa}, but was not evaluated explicitly or studied further. A single-detector $\\F$-statistic metric was used (without giving any details) in \\cite{cornish05:_detec_lisa} to numerically estimate the number of templates in galactic-binary searches with LISA. An earlier study by \\citet{brady98:_search_ligo_periodic} of the metric for isolated neutron-star signals introduced a metric approximation based only on the phase modulation of the signal and neglecting the amplitude modulation. We will refer to this approximation as the ``phase metric''. This metric has a simpler structure than the full $\\F$-metric, and it can be computed analytically \\cite{jones05:_ptole_metric} if one assumes a circular orbital motion. This is the only type of continuous GW metric that is currently implemented in LAL/LALApps~\\cite{lalapps}.  As we will see in this study, the amplitude modulation cannot always be neglected, but ``on average'' the phase metric seems to be a good approximation, and its quality improves with longer observation times and with the number of detectors. With the recent multi-detector generalization of the $\\F$-statistic formalism~\\cite{cutler05:_gen_fstat} and its subsequent implementation into LAL/LALApps, the need to understand the \\emph{multi-detector} $\\F$-statistic metric has become more urgent. The most important question in this context is whether the metric resolution increases with the number of detectors, i.e. whether a denser covering of the parameter space is required, which would increase the computational cost.  The main result of this work is to show that the metric resolution does \\emph{not} scale with the number of detectors. Therefore, sensitivity can be gained at the cost of only a linear increase in the required computing power (as the signal has to be correlated with the data stream from each detector). This has to be contrasted with the (at least) $\\O(T^5)$ scaling \\eqref{eq:113} of the number of templates with observation time $T$, in the case isolated neutron-star signals with one spindown.  In order to improve the sensitivity of a coherent search for continuous gravitational waves, increasing the number of similar-sensitivity detectors is therefore computationally much cheaper than to increase the observation time.  The plan of this paper is as follows: in Sect.~\\ref{sec:multi-detector-f} we introduce the formalism and notation of the multi-detector $\\F$-statistic, following \\cite{cutler05:_gen_fstat} and \\cite{krolak04:_optim_lisa}. In Sect.~\\ref{sec:f-metric} we derive the $\\F$-statistic metric family for high-frequency signals (relevant for ground-based detector networks). We compute the extremal range of this metric family, its average metric, and the long-duration limit, in which the $\\F$-metric family reduces to a simple ``orbital metric''. In Sect.~\\ref{sec:appl-isol-puls} we apply this framework to the special case of GWs from isolated spinning neutron stars, and we evaluate the quality of the metric predictions (and different approximations) by comparing them against measured mismatches in a Monte-Carlo study. The main results are summarized in Sect.~\\ref{sec:conclusions}. Appendix~\\ref{sec:altern-proj-onto} presents an alternative, more elegant derivation of the $\\F$-statistic metric, and Appendix~\\ref{sec:keep-ampl-funct} gives the general expressions for the $\\F$-metric, which would be valid also for low-frequency signals relevant for LISA. ",
        "conclusions": "\\label{sec:conclusions}  We have derived a formalism for the general parameter-space metric of the multi-detector $\\F$-statistic, and we have explicitly computed the metric for signals from isolated spinning neutron stars. We have shown that there exists a family of $\\F$-metrics, parametrized by the two (unknown) amplitude parameters $\\psi$ and $\\cos\\iota$. We explicitly derived the extremal ``mismatch bounds'' (i.e. the maximum and minimum possible mismatches) of the $\\F$-metric family, and we introduced an average $\\F$-metric, which is independent of the unknown amplitude parameters. We have shown that the multi-detector $\\F$-metric does not scale with the number of detectors. Combining detectors coherently therefore does not increase the required number of templates. In the long-duration limit ($T\\gtrsim 1$~month), we found that the $\\F$-metric family converges towards a simple orbital metric $g^\\orb_{ij}$, which neglects both the amplitude modulation and the phase modulation caused by the diurnal rotation of the Earth.  Both the orbital and the closely related phase metric provide relatively good mismatch approximations in practice, and the quality of these approximations improves with longer observation times and with the number of coherently-combined detectors.  The orbital metric, however, while virtually identical to the phase metric for almost all directions in parameter space, was found to be substantially more degenerate for observation times shorter than a month. This has important consequences for the covering problem of the parameter space and requires further study. Finally, we have identified two regimes in which the local metric approximation itself is not very reliable: namely, for observation times $T\\lesssim 1$~day, and for large angular offsets $\\dOm \\gtrsim 5$ (in natural units).  \\appendix"
    },
    "0606/astro-ph0606006_arXiv.txt": {
        "abstract": "{} {We investigate how ellipticity, asymmetries and substructures separately affect the ability of galaxy clusters to produce strong lensing events, i.e. gravitational arcs, and how they influence the arc morphologies and fluxes. This is important for those studies aiming, for example, at constraining cosmological parameters from statistical lensing, or at determining the inner structure of galaxy clusters through gravitational arcs.} {We do so by creating two-dimensional smoothed, differently elliptical and asymmetric versions of some numerical models. By subtracting these smoothed mass distributions from the corresponding numerical maps and by gradually smoothing the residuals before re-adding them to the clusters, we are able to see how the lensing properties of the clusters react to even small modification of the cluster morphology. We study in particular by how much ellipticity, asymmetries and substructures contribute to the strong lensing cross sections of clusters. We also investigate how cluster substructures affect the morphological properties of gravitational arcs, their positions and fluxes.} {On average, we find that the contributions of ellipticity, asymmetries and substructures amount to $\\sim 40\\%$, $\\sim 10\\%$ and $\\sim 30\\%$ of the total strong lensing cross section, respectively. However, our analysis shows that substructures play a more important role in less elliptical and asymmetric clusters, even if located at large distances from the cluster centres ($\\sim 1h^{-1}$Mpc). Conversely, their effect is less important in highly asymmetric lenses. The morphology, position and flux of individual arcs are strongly affected by the presence of substructures in the clusters. Removing substructures on spatial scales $\\lesssim 50h^{-1}$kpc, roughly corresponding to mass scales $\\lesssim 5 \\times 10^{10}h^{-1}\\,M_\\odot$, alters the image multiplicity of $\\sim 35\\%$ of the sources used in the simulations and causes position shifts larger than $5''$ for $\\sim 40\\%$ of the arcs longer than $5''$.} {We conclude that any model for cluster lens cannot neglect the effects of ellipticity, asymmetries and substructures. On the other hand, the high sensitivity of gravitational arcs to deviations from regular, smooth and symmetric mass distributions suggests that strong gravitational lensing is potentially a powerful tool to measure the level of substructures and asymmetries in clusters.} ",
        "introduction": "Thanks to the improvements in the quality and in the depth of astronomical observations, in particular from space, an increasing number of gravitational arcs has recently been discovered near the centres of many galaxy clusters \\citep[see e.g.][]{BR05.1}. Since the appearance of these images reflects the shape of the gravitational potential which is responsible for their large distortions, strong lensing is, in principle, a very powerful tool for investigating how the matter, in particular the dark component, is distributed in the inner regions of cluster lenses.  Determining the inner structure of galaxy clusters is one of the major goals in cosmology, because it should allow us to set important constraints on the growth of the cosmic structures in the Universe. Moreover, constraining the mass distribution in the centre of dark matter halos has become increasingly important in the recent years, since observations of the dynamics of stars in galaxy-sized systems revealed the presence of a potential problem within the Cold-Dark-Matter (CDM) scenario. While numerical simulations in this cosmological framework predict that dark matter halos in a large range of masses should develop density profiles characterised by an inner cusp, observations of the rotation curves of dwarf and low-surface-brightness galaxies suggest that these objects rather have flat density profiles \\citep{FL94.1,MO94.1,BU95.1,BU97.1,MG98.1,DA00.1,FI01.1}.  While the centres of galaxies are dominated by stars, which renders it extremely complicated to derive constraints on the distribution of their dark matter, galaxy clusters are an alternative and, in many respects, preferable class of objects for testing the predictions of the CDM model. In fact, several authors already tried to investigate the inner structure of these large systems, using and often combining several kinds of observations. Apart from lensing, the gravitational potential of galaxy clusters can be traced with several other methods, for example through the emission in the X-ray band by the hot intra-cluster gas. However, while gravitational lensing directly probes the matter content of these objects, the other techniques usually rely on some strong assumptions about their dynamical state and the interaction between their baryonic and dark matter. For example, it must be often assumed that the gas is in hydrostatic equilibrium within the dark matter potential well and that the system is spherically symmetric.  Some ambiguous results were found when comparing the constraints on the inner structure of clusters as obtained from X-ray and lensing observations. First, masses estimated from strong lensing are usually larger by a factor of 2-3 than the masses obtained from X-ray observations \\citep{CH03.1,OT04.1}. Deviations from axial symmetry and substructures are known to be important factors in strong lensing mass estimates \\citep[see e.g.][]{BA95.2,BA96.2,ME03.1,OG05.1,GA05.1}. Second, the constraints on the inner slope of the density profiles seem to be compatible with a wide range of inner slopes \\citep{ET02.1,LE03.1,AR02.1,SA03.1,BA04.1,GA05.1}.  Apart from the above-mentioned uncertainties affecting the X-ray measurements, strong lensing observations also have several potential weaknesses. First of all, arcs are relatively rare events. Frequently, all the constraints which can be set on the inner structure of clusters via strong lensing depend on a single or on a small number of arcs and arclets observed near the cluster core. Second, arcs are the result of highly non-linear effects. This implies that their occurrence and their morphological properties are very sensitive to ellipticity, asymmetries and substructures of the cluster matter distribution.  Reversing the problem, this means that, in order to reliably describe the strong lensing properties of galaxy clusters, all of these effects must be taken into account. Fitting the positions and the morphology of gravitational arcs to derive the underlying mass distributions of the lensing clusters, usually requires to build models with multiple mass components, each of which is characterised by its ellipticity and orientation \\citep[see e.g.][]{KN93.1,CO05.1,BR05.1}. Even describing the cluster lens population in a statistical way requires to use realistic cluster models \\citep{ME00.1,ME03.1,ME03.2,OG02.1,OG03.1,DA04.1,HE05.1}.  Despite the fact that the importance of ellipticity, asymmetries and substructures for strong lensing appears clearly in many previous studies, many questions still remain. For example, what is the typical scale of substructures which contribute significantly to the strong lensing ability of a cluster?  Where are they located within the clusters? What is the relative importance of asymmetries compared to ellipticity? Moreover, how do substructures influence the appearance of giant arcs? All of these open problems are important for those studies aiming at constraining cosmological parameters from statistical lensing, or at determining the inner structure of galaxy clusters through gravitational arcs.  This paper aims at answering to these questions. To do so, we quantify the impact of ellipticity, asymmetries and substructures by creating differently smoothed models of the projected mass distributions of some numerical clusters. We gradually move from one smoothed model to another through a sequence of intermediate steps.  The plan of the paper is as follows. In Sect.~\\ref{sect:nummod}, we discuss the characteristics of the numerically simulated clusters that we use in this study; in Sect.~\\ref{sect:raytr}, we explain how ray-tracing simulations are carried out; Sect.~\\ref{sect:smooth} illustrates how we obtain smoothed versions of the numerical clusters; in Sect.~\\ref{sect:power}, we suggest a method to quantify the amount of substructures, asymmetry and ellipticity of the cluster lenses, based on multipole expansions of their surface density fields; Sect.~\\ref{sect:resu} is dedicated to the discussion of the results of our analysis. Finally, we summarise our conclusions in Sect.\\ref{sect:conclu}. ",
        "conclusions": "\\label{sect:conclu}  In this paper we have quantified the impact of several properties of realistic cluster lenses on their strong lensing ability. In particular, our goal was to separate the effects of substructures, asymmetries and ellipticity. For doing that, we analysed the lensing properties of one numerical cluster simulated with very high mass resolution. In addition, we studied four other clusters obtained from N-body simulation with a lower mass resolution.  Each cluster was projected along three independent directions.  For each projection, we constructed three completely smoothed versions. Each of them conserves the mean surface density profile of the mass distribution of the cluster. However, the first reproduces the variations of the ellipticity and of the position angle of the isodensity contours as functions of the distance from the centre; the second has elliptical isodensity contours with fixed ellipticity and orientation; the third is an axially symmetric model.  The lensing properties of the numerical clusters, of their smoothed analogues and of several intermediate versions were investigated using standard ray-tracing techniques.  Our main results can be summarised as follows: \\begin{itemize} \\item Substructures, asymmetries and ellipticity contribute to increase the ability of clusters to produce strong lensing events. Substructured, asymmetric and highly elliptical clusters produce more extended high magnification regions in the lens plane where long and thin arcs can form. Indeed, substructures, asymmetries and ellipticity determine the location and the shape of the lens caustics around which sources must be located in order to be strongly lensed by the clusters. \\item The impact of substructures, asymmetries and ellipticity on the lensing cross section for producing giant arcs is different for different lenses. The lensing properties of the most symmetric clusters appear to be particularly influenced by the substructures. On the contrary, substructures are less important in asymmetric lenses. \\item On average, we quantify that substructures account for $\\sim 30\\%$ of the total cluster cross section, asymmetries for $\\sim 10\\%$ and ellipticity for $\\sim 40\\%$. \\item The substructures that typically contribute to lensing are on scales $\\lesssim 150 - 200 \\,h^{-1}$kpc. Assuming a virial overdensity of $\\sim 123$ for $z=0.3$, this corresponds to mass scales of the order of $\\sim 10^{12} h^{-1}\\,M_\\odot$. Substructures on larger scales are not as frequent in our cluster sample, but, if present, they can boost significantly the lensing cross section \\citep[see e.g.][]{TO04.1,ME05.1}. \\item Substructures play a more important role when they are located close to the cluster centre. However, the lensing cross section for giant arcs is sensitive to substructures within a wide region around the cluster core. In particular, our simulations show that the sensitivity to substructures far from the centre is particularly high in those clusters whose inner regions are unperturbed. In these cases, the loss of strong lensing efficiency due to removing the substructures from the clusters is correlated with substructures within a region of $\\sim 1 \\, h^{-1}$Mpc in radius; on the contrary, clusters containing substructures in the inner regions are ``screened'' against external perturbers. \\item Even small substructures ($l\\lesssim 50$kpc, $M\\lesssim 5\\times 10^{10}h^{-1}\\,M_\\odot$) influence the appearance and the location of gravitational arcs. The perturbations to the projected gravitational potential of the cluster induced by the substructures alter the multiplicity of the images of individual sources. Moreover, they change the morphology and the flux of the images themselves. Finally, they can shift the position of arcs with significant length to width ratios by several arcseconds on the sky. \\end{itemize}  These results highlight several important aspects of strong lensing by clusters. First, any model for cluster lenses cannot neglect the effects of asymmetries, ellipticity and substructures. Clusters which may appear as relaxed and symmetric, for example in the X-rays, are potentially those which are most sensitive to the smallest substructures, located even at large distances from the inner cluster regions, critical for strong lensing. Even subhalos on the scales of galaxies can influence the strong lensing properties of their hosts and alter the shape and the fluxes of gravitational arcs. Therefore, if the lens modelling is not carried out at a very high level of detail, it may result in being totally incorrect.  Second, the high sensitivity of gravitational arcs to deviations from regular, smooth and symmetric mass distributions suggests that strong gravitational lensing is potentially a powerful tool to measure the level of substructures and asymmetries in clusters. Since, as we said, the sensitivity to substructures is higher in the case of more symmetric lenses, we conclude that dynamically active clusters, like those undergoing major merger events, should be quite insensitive to ``corrugations'' in the projected mass distribution but highly sensitive to asymmetries. Arcs could then be used to diagnose mergers in clusters. Conversely, substructures should become increasingly important for the arc morphology as clusters relax. Then the level of substructures in clusters should be quantified by measuring their effect on the arc morphology. This is particularly intriguing since measuring the fine structures of gravitational arcs has become feasible thanks to the high spatial resolution reached in observations from space.  Third, the strong impact of asymmetries and substructures on the lensing properties of clusters and the wide region in the cluster where these last can be located in order to produce a significant effect further support the picture that mergers might have a huge impact on the cluster optical depth for strong lensing, as suggested in several previous studies \\citep{TO04.1,ME04.1,FE05.1}."
    },
    "0606/astro-ph0606538_arXiv.txt": {
        "abstract": "Using a set of multifrequency cross-spectra computed from the three year \\WMAP\\ sky maps, we fit for the unresolved point source contribution.  For a white noise power spectrum, we find a Q-band amplitude of $A = 0.011 \\pm 0.001$ $\\mu$K$^2$ sr (antenna temperature), significantly smaller than the value of $0.017 \\pm 0.002$ $\\mu$K$^2$ sr used to correct the spectra in the \\WMAP\\ release.  Modifying the point source correction in this way largely resolves the discrepancy \\citet{eriksen:2006} found between the \\WMAP\\ V- and W-band power spectra.  Correcting the co-added \\WMAP\\ spectrum for both the low-$\\ell$ power excess due to a sub-optimal likelihood approximation---also reported by \\citet{eriksen:2006}---and the high-$\\ell$ power deficit due to over-subtracted point sources---presented in this letter---we find that the net effect in terms of cosmological parameters is a $\\sim0.7\\sigma$ shift in $n_{\\textrm{s}}$ to larger values: For the combination of \\emph{WMAP}, BOOMERanG and Acbar data, we find $n_{\\textrm{s}} = 0.969 \\pm 0.016$, lowering the significance of $n_{\\textrm{s}} \\ne 1$ from $\\sim2.7\\sigma$ to $\\sim 2.0\\sigma$. ",
        "introduction": "The results of \\emph{Wilkinson Microwave Anisotropy Probe} have made an inestimable impact on the science of cosmology, highlighted by the very recent release of the three year data: maps, power spectra, and consequent cosmological analysis \\citep{\\jarosik,\\page,\\hinshaw,\\spergel}.  Precisely because these results play so prominent a role, it is important to check and recheck their consistency.   Recently \\cite{eriksen:2006} reanalyzed the \\WMAP\\ three year temperature sky maps, and noted two discrepancies in the \\WMAP\\ power spectrum analysis.  On large angular scales there is a small power excess in the \\emph{WMAP} spectrum (5--10\\% at $\\ell \\lesssim 50$), primarily due to a problem with the % likelihood approximation used by the \\emph{WMAP} team. On small angular scales, an unexplained systematic difference between the V- and W-band spectra (few percent at $\\ell \\gtrsim 300$) was found.  In this Letter, we suggest this second discrepancy is at least partially due to an excessive point source correction in the \\WMAP\\ power spectrum. ",
        "conclusions": "\\label{sec:conclusions}  Using a combination of cross-spectra of maps from the Q-, V-, and W-bands of \\WMAP\\ three year data, we fit for the amplitude of the power spectrum of unresolved point sources in Q-band, finding $A = 0.011 \\pm 0.001$ $\\mu$K$^2$ sr.  This fit has significantly less power than the fit used to correct the \\WMAP\\ final co-added power spectrum used for cosmological analysis.  We compute and apply the proper point source correction, noting the corrected V- and W-bands are more consistent than before.  The improper point source correction conspires with a low-$\\ell$ estimator bias to impart a spurious tilt to the \\WMAP\\ temperature power spectrum.  With the revised corrections, we find evidence for spectral index $n_\\textrm{s} \\neq 1$ at only $\\sim 2\\sigma$, while other parameters remain largely unchanged."
    },
    "0606/astro-ph0606012_arXiv.txt": {
        "abstract": "We have undertaken the simulation of hydrodynamic flows with bulk Lorentz factors in the range 10$^2$--10$^6$. We discuss the application of an existing relativistic, hydrodynamic primitive-variable recovery algorithm to a study of pulsar winds, and, in particular, the refinement made to admit such ultra-relativistic flows. We show that an iterative quartic root finder breaks down for Lorentz factors above 10$^2$ and employ an analytic root finder as a solution. We find that the former, which is known to be robust for Lorentz factors up to at least 50, offers a 24\\% speed advantage. We demonstrate the existence of a simple diagnostic allowing for a hybrid primitives recovery algorithm that includes an automatic, real-time toggle between the iterative and analytical methods. We further determine the accuracy of the iterative and hybrid algorithms for a comprehensive selection of input parameters and demonstrate the latter's capability to elucidate the internal structure of ultra-relativistic plasmas. In particular, we discuss simulations showing that the interaction of a light, ultra-relativistic pulsar wind with a slow, dense ambient medium can give rise to asymmetry reminiscent of the Guitar nebula leading to the formation of a relativistic backflow harboring a series of internal shockwaves. The shockwaves provide thermalized energy that is available for the continued inflation of the PWN bubble. In turn, the bubble enhances the asymmetry, thereby providing positive feedback to the backflow. ",
        "introduction": "\\label{solver} Using the method above we created a SRHD primitive algorithm called ``REST\\_ FRAME''. Given the speed advantage of the iterative root finder (see $\\S$\\ref{timing}), it a desirable choice over the analytical method within its regime of applicability, i.e. for low Lorentz factors. As Fig.~\\ref{nr_acc} shows, the iterative root finder is accurate to order 10$^{-4}$ (see $\\S$\\ref{solacc}) for a sizable region of parameter space including all $R/E$ above the diagonal line between the points (0, $-7$) \\& (9, 0) in the $\\log(R/E)$ vs. $-\\log(1-M/E)$ plane (i.e. for $\\log(R/E) \\ge -(7/9)\\times\\log(1-M/E)-7$). Therefore, for a given $M/E$ and $R/E$, we check if this inequality is true; if (not) so, we call the (analytical) iterative root finder (see $\\S$\\ref{pcode}).  \\subsection{Pseudo-code}\\label{pcode} REST\\_FRAME calculates the primitive variables given the conservative variables and the adiabatic index as represented in the following pseudo-code (note this is a 2D example):  \\hspace*{8mm}PROCEDURE REST\\_FRAME\\\\ \\hspace*{8mm}RECEIVED FROM PARENT PROGRAM: $Y$, $Z$\\\\ \\hspace*{8mm}RETURNED TO PARENT PROGRAM: $\\gamma$, $v$, $C$\\\\ Comment: recall $Y\\equiv M/E$ and $Z\\equiv R/E$\\\\ Comment: $C$ is returned $< 0$ for code failures\\\\ \\hspace*{8mm}GLOBAL VARIABLE: $\\Gamma$\\\\ \\hspace*{8mm}SET VALUE OF $m_{underflow}$\\\\ \\hspace*{8mm}SET VALUE OF $v_{tol}$\\\\ Comment: determines iterative method velocity accuracy\\\\ Comment: we set $v_{tol}$ = 10$^{-8}$, 10$^{-10}$, 10$^{-12}$, 10$^{-14}$\\\\ Comment: for $-\\log(1-Y) < 8.3$, $< 10.3$, $< 12.3$, otherwise, respectively\\\\ \\hspace*{8mm}SET $M$ = $\\sqrt{M_x^2+M_y^2}$\\\\ \\hspace*{8mm}IF $M < m_{underflow}$ THEN\\\\ \\hspace*{8mm}\\hspace*{8mm}$v$ = 0, $\\gamma$ = 1\\\\ Comment: avoids code failure if $v$ is numerically zero\\\\ \\hspace*{8mm}ELSE\\\\ \\hspace*{8mm}\\hspace*{8mm}TEST FOR UNPHYSICAL PARAMETERS\\\\ \\hspace*{8mm}\\hspace*{8mm}IF PASSED, SET $C$ NEGATIVE AND RETURN\\\\ \\hspace*{8mm}\\hspace*{8mm}IF $\\log(Z) \\ge -(7/9)\\times\\log(1-Y)-7$, THEN\\\\ Comment: check to see if input parameters are within the acceptable\\\\ Comment: accuracy region of the iterative routine\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}CALL ITERATIVE\\_QUARTIC($Y,Z,v_{tol},v$,$C$)\\\\ Comment: updates $v_{n-1}$ to $v_{n}$ using $n$ cycles of Newton-Raphson iteration\\\\ Comment: returns $v$ = $v_{n}$ when $|v_{n}-v_{n-1}|$ $\\le$ $v_{tol}$\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}IF $C < 0$, THEN\\\\ Comment: this means the iteration failed to converge\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}RETURN\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}ELSE \\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}$\\gamma$ = $\\sqrt{\\frac{1}{1-v^2}}$\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}END IF\\\\ \\hspace*{8mm}\\hspace*{8mm}ELSE\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}CALL ANALYTICAL\\_QUARTIC($Y,Z,\\gamma$)\\\\ Comment: calculates $\\gamma$ using analytical solution -- see below\\\\ \\hspace*{8mm}\\hspace*{8mm}\\hspace*{8mm}$v$ = $\\sqrt{1-\\frac{1}{\\gamma^2}}$\\\\ \\hspace*{8mm}\\hspace*{8mm}END IF\\\\ \\hspace*{8mm}END IF\\\\ \\hspace*{8mm}END PROCEDURE REST\\_FRAME\\\\  \\hspace*{8mm}PROCEDURE ANALYTICAL\\_QUARTIC\\\\ Comment: see $\\S$\\ref{solveqe} for equations\\\\ \\hspace*{8mm}RECEIVED FROM PARENT PROGRAM: $Y,Z$\\\\ \\hspace*{8mm}RETURNED TO PARENT PROGRAM: $\\gamma$\\\\ \\hspace*{8mm}GLOBAL VARIABLE: $\\Gamma$\\\\ \\hspace*{8mm}$\\tilde{a}_3$ = $2\\Gamma(\\Gamma-1)Z(Y^{-2}+1)$\\\\ \\hspace*{8mm}$\\tilde{a}_2$ = $(\\Gamma^2-2\\Gamma(\\Gamma-1)Y^2-(\\Gamma-1)^2Z^2)(Y^{-2}+1)$\\\\ \\hspace*{8mm}$\\tilde{a}_1$ = $-a_3$\\\\ \\hspace*{8mm}$\\tilde{a}_0$ = $(\\Gamma-1)^2(Y^2+Z^2)(Y^{-2}+1)$\\\\ \\hspace*{8mm}$\\tilde{a}_4$ = $1+Y^2-a_0-a_2$\\\\ Comment: coefficients recast to counter subtractive cancellation -- see $\\S$\\ref{solacc}\\\\ \\hspace*{8mm}NORMALIZE COEFFICIENTS TO $a_4$\\\\ Comment: e.g., $a_{3N}$ = $a_3/a_4$\\\\ \\hspace*{8mm}CALCULATE CUBIC RESOLVENT COEFFICIENTS \\\\ \\hspace*{8mm}CALCULATE DISCRIMINANT, $D$\\\\ \\hspace*{8mm}IF $D\\le$0 THEN\\\\ \\hspace*{8mm}\\hspace*{8mm}WRITE ERROR MESSAGE AND STOP\\\\ Comment: exploration suggests $D\\le 0$ is unphysical but formal proof is elusive\\\\ Comment: thus, we leave $D\\le 0$ uncoded with a error flag just in case\\\\ \\hspace*{8mm}ELSE\\\\ Comment: $D>0$ $\\Rightarrow$ $Q(\\gamma)$ has 2 real roots (see Tab.~\\ref{charqsolns} \\& \\ref{charcsolns})\\\\ \\hspace*{8mm}\\hspace*{8mm} CALCULATE ROOTS OF CUBIC RESOLVENT\\\\ Comment: the cubic has one real root and a pair of complex conjugate roots\\\\ \\hspace*{8mm}\\hspace*{8mm} IF REAL ROOT $< 0$, SET REAL ROOT = 0\\\\ Comment: the real root cannot be less than zero analytically\\\\ Comment: numerically, however, it can have a very small negative value\\\\ \\hspace*{8mm}\\hspace*{8mm} CALCULATE THE TWO REAL ROOTS OF THE QUARTIC\\\\ \\hspace*{8mm}\\hspace*{8mm} TEST FOR TWO OR NO PHYSICAL ROOTS\\\\ \\hspace*{8mm}\\hspace*{8mm} IF PASSED, WRITE ERROR MESSAGE, AND RETURN\\\\ \\hspace*{8mm}\\hspace*{8mm} IF FAILED, SET $\\gamma$ = PHYSICAL ROOT\\\\ \\hspace*{8mm}END IF\\\\ \\hspace*{8mm}END PROCEDURE ANALYTICAL\\_QUARTIC\\\\  \\subsection{Code timing}\\label{timing} Using the Intel Fortran library function CPU\\_TIME, we calculated the CPU time required to execute 5$\\times$10$^7$ calls to REST\\_FRAME for $Y$ = 0.9975 \\& $Z$ = 1$\\times 10^{-4}$ ($\\gamma \\sim$10) using the Newton-Raphson iterative method with $Q(v)$ and 8-byte arithmetic, and the analytical method with $Q(\\gamma)$ and both 8-byte \\& 16-byte arithmetic (we investigated the use of 16-byte arithmetic due to an issue with subtractive cancellation -- see $\\S$\\ref{solacc}). The CPU time for each of these scenarios was 29.5, 36.5 (averaged over ten runs and rounded to the nearest half second), and $\\sim$11650 seconds (one run only), respectively. This indicates that while using the 8-byte analytical method is satisfactory, it is advantageous to use the iterative method when Lorentz factors are sufficiently low, and that the use of 16-byte arithmetic is a nonviable option. This result is not surprising as the accuracy of Newton-Raphson iteration improves by approximately one decimal place per iterative step \\citep{dun94} and the relative inefficiency of 16-byte arithmetic is a known issue \\citep[e.g.][]{per06}.  \\subsection{Solver Accuracy}\\label{solacc} The input parameters for our primitives algorithm are the ratios of the laboratory-frame momentum and mass densities to the laboratory-frame energy density (recall $Y \\equiv M/E$ and $Z \\equiv R/E$) both of which must be less than unity in order for solutions of Eqn.~\\ref{lt} to exist. In addition, the condition $Y^2 + Z^2 < 1$ must be met. Along with the fact that $Y$ and $Z$ must also be positive, this defines the comprehensive and physical input parameter space to be $0 < Y,Z < 1$ such that $Y^2 + Z^2 < 1$ (we identify a particular region of parameter space applicable to pulsar winds in the next section). We tested the accuracy of our iterative and hybrid primitives algorithms within this space as follows.  First, as we are most interested in light, highly relativistic flows (i.e. $Z$ small and $Y$ close to unity), to define the accuracy-search space we elected to use the quantities $-\\log(1-Y)$, which for values greater than unity gives $0.9 < Y < 1$, and  $\\log(Z)$, which for values less than negative unity gives $Z \\ll 1$. We selected $0 < -\\log(1-Y) < 13$ and $-13 < \\log(Z) < 0$ corresponding to Lorentz factors ($\\gamma$) between 1 and $2\\times10^6$. We chose a range with a maximal $\\gamma$ slightly above $1\\times10^6$ in order to completely bound the pulsar wind nebula parameter space defined in the next section.  Choosing a relativistic equation of state $\\Gamma$= 4/3 and using 1300 points for both $-\\log(1-Y)$ and $\\log(Z)$, we tested the accuracy of REST\\_FRAME by passing it $Y$ and $Z$, choosing $E = 1$, and using the returned primitive quantities to derive the calculated energy density $Ec$, and calculating the difference $|1-Ec/E| \\equiv \\delta E/E$. We chose this estimate of the error because $\\delta E/E \\sim \\delta\\gamma/\\gamma$ and $\\delta\\gamma/\\gamma$ is tied to the accuracy of the numerical, hydrodynamic technique (see the final paragraph in this section).  \\begin{figure}[hb] \\caption{The accuracy (estimated as $\\delta E/E$) of the Newton-Raphson (N-R) iterative primitives algorithm where white, light grey, medium grey, dark grey, and hatched regions correspond, respectively, to an accuracy of order at least 10$^{-4}$, at least 10$^{-3}$, worse than 10$^{-3}$, failure, and unphysical input ($R^2/E^2 \\ge 1-M^2/E^2$). Note that the Lorentz factor varies from order 1 at the far left to order 10$^6$ at the far right. There is a sizable white region representing $M/E < 0.999999$ ($\\gamma < 500$) and $R/E > 5\\times 10^{-8}$ within which accuracy is generally significantly better than 10$^{-4}$. N-R iteration is unreliable due to sporadic failures for all $M/E$ and $R/E$ such that $R/E < 5\\times 10^{-8}$ and for an ever increasing fraction of $R/E > 5\\times 10^{-8}$ as $M/E$ increases until accuracy becomes unacceptable or the code fails outright for $M/E$ and $R/E$ such that $M/E > 0.999999$. Failures are due to divide by zero (see \\S\\ref{sfix}) or nonconvergence within a reasonable number of iterations.} \\centerline{\\includegraphics[width=4in]{fig2.ps}} \\label{nr_acc} \\end{figure} \\begin{figure}[hb] \\caption{The accuracy (estimated as $\\delta E/E$) of the hybrid primitives algorithm where white, light grey, and hatched regions correspond, respectively, to an accuracy of order at least 10$^{-4}$, at least 10$^{-3}$, and nonphysical input ($R^2/E^2 \\ge 1-M^2/E^2$). Note that the Lorentz factor varies from order 1 at the far left to order 10$^6$ at the far right. The space between the parallel lines represents PWNe input parameter space. The accuracy degradation at the extreme right is due to subtractive cancellation in the 4$^{th}$-order coefficient of the Lorentz-factor quartic as $M/E\\rightarrow$1.} \\centerline{\\includegraphics[width=4in]{fig3.ps}} \\label{hy_acc} \\end{figure}  Our results for the Newton-Raphson (N-R) and hybrid methods are given in Figs.~\\ref{nr_acc} \\& \\ref{hy_acc} which show where the accuracy is of order at least 10$^{-4}$, at least 10$^{-3}$, worse than 10$^{-3}$, failure, and unphysical input ($Z^2 \\ge 1-Y^2$), respectively. We chose an accuracy of order 10$^{-4}$ as the upper cutoff because N-R iteration returns accuracies on this order for $\\gamma < 50$ and relativistic, hydrodynamic simulations of galactic jets by \\cite{dun94} produced robust results for Lorentz factors of at least 50 using N-R iteration. An additional result of interest is that the ultra-relativistic approximation for $v$ (i.e. taking $R = 0$ thereby reducing $Q(v) = 0$ to a quadratic equation) manages an accuracy of at least 10$^{-4}$ for a large portion of the physical $Y-Z$ plane (see Fig.~\\ref{ur_acc}).  \\begin{figure}[hb] \\caption{The accuracy (estimated as $\\delta E/E$) of the ultra-relativistic approximation of the flow velocity where white, light grey, medium grey, and hatched regions correspond to an accuracy of order at least 10$^{-4}$, at least 10$^{-3}$, worse than 10$^{-3}$, and unphysical input ($R^2/E^2 \\ge 1-M^2/E^2$), respectively. Note that the Lorentz factor varies from order 1 at the far left to order 10$^6$ at the far right. The accuracy degradation at the extreme right is due to the fact that the fractional error in the Lorentz factor is proportional to the fractional error in the velocity divided by $1-v^2$ which diverges as $v\\rightarrow$1.} \\centerline{\\includegraphics[width=4in]{fig4.ps}} \\label{ur_acc} \\end{figure}  Fig.~\\ref{nr_acc} shows the accuracy of the N-R iterative method. There are several noteworthy features. First is the presence of a sizable region corresponding to $\\gamma < 500$ within which accuracy is generally significantly better than 10$^{-4}$. Second is that N-R iteration is unreliable due to sporadic failures for increasing Lorentz factors until accuracy becomes unacceptable or the code fails outright due to divide by zero (see \\S\\ref{sfix}) or non-convergence within a reasonable number of iterations. In addition, though N-R iteration has been widely established as the primitives recovery method of choice for flows with Lorentz factors less than order 10$^2$, we found that for a subset of parameters, corresponding to $\\gamma < 2$, our N-R algorithm suffered an unacceptable degradation in accuracy. The key to this problem lies in the how the flow velocity ($v$) is initially estimated for the first iterative cycle as follows: \\begin{enumerate} \\item The established approach \\citep{dun94,sch93} is to bracket $v$ with  \\ber v_{max} &=& min(1,Y+\\delta)\\rm,\\nonumber\\\\ v_{min} &=& \\frac{\\Gamma - \\sqrt{\\Gamma^2-4(\\Gamma-1)Y^2}}{2Y(\\Gamma-1)}\\rm, \\eer where $\\delta \\sim 10^{-6}$ and $v_{min}$ is derived by taking the ultra-relativistic limit (i.e. $R = 0$) \\item The initial velocity is then $v_i = (v_{min}+v_{max})/2+\\eta$, where $\\eta = (1-Z)(v_{min}-v_{max})$ for $v_{max} > \\epsilon$ and $\\eta = 0$ otherwise ($\\epsilon$ order 10$^{-9}$) \\item This method fails due to selection of the incorrect root when the roots converge. \\item Thus, we make a simpler initial estimate of $v_i = v_{max}$, which guarantees that $v_i$ is ``uphill'' from $v$ for all physical $Y-Z$ space and that N-R iteration converges on $v$. \\end{enumerate}  Fig.~\\ref{hy_acc} shows that our hybrid algorithm REST\\_FRAME is accurate to at least 10$^{-4}$ for all but a smattering of the highest Lorentz factors. In fact, it is significantly more accurate over the majority of the physical portion of the $Y-Z$ plane. The space between the parallel lines represents the PWN input parameters discussed in the next section. We find that multiplying $Q(\\gamma)$ by $(Y^2-Y^{-2})$ and rewriting the new $a_4$ ($\\tilde{a}_4$) in terms of the new $a_2$ ($\\tilde{a}_2$) and new $a_0$ ($\\tilde{a}_0$), e.g. $\\tilde{a}_4 = 1+Y^2-\\tilde{a}_0-\\tilde{a}_2$, improves the accuracy somewhat, but does not entirely mitigate the problem. The issue of accuracy loss at large Lorentz factors in 8-byte primitives algorithms is a known issue \\citep[e.g.][]{nob03} for which we know of no complete 8-byte solution. Employing 16-byte arithmetic provides spectacular accuracy, but introduces an unacceptable increase in run time (see $\\S$\\ref{timing}).  The issue of what constitutes an acceptable error in the calculated Lorentz factor is decided by the fact that a fractional error in $\\gamma$ translates to the same fractional error in $p$ and $n$ which are needed to calculate the wave speeds that form the basis of the numerical, hydrodynamic technique, a Godunov scheme \\citep{god59} which approximates the solution to the local Riemann problem by employing an estimate of the wave speeds. We do not know a priori how accurate this estimate needs to be, and so procede with 8-byte simulations of pulsar winds with the expectation of using shock-tube tests \\citep{tho86} to validate the accuracy of the computation of well-defined flow structures as we approach the highest Lorentz factors. It is also noteworthy that while $\\gamma$ = 10$^6$ is the canonical bulk Lorentz factor for pulsar winds, $\\gamma$ = 10$^4$ and 10$^5$ are still in the ultra-relativistic regime, and it may very well prove to be that these Lorentz factors are high enough to elucidate the general ultra-relativistic, hydrodynamic features of such a system. The hybrid algorithm achieves accuracies of at least 10$^{-6}$ for $\\gamma \\sim 10^5$, which is safely in the acceptable accuracy regime. ",
        "conclusions": "\\label{sec:conc} We discussed the application of an existing special relativistic, hydrodynamic (SRHD) primitive-variable recovery algorithm to ultra-relativistic flows (Lorentz factor, $\\gamma$, of 10$^2$--10$^6$) and the refinement necessary for the numerical velocity root finder to work in this domain. We found that the velocity quartic, $Q(v)$, exhibits dual roots in the physical velocity range that move progressively closer together for larger $\\gamma$ leading to a divide by zero and the failure of the Newton-Raphson iteration method employed by the existing primitives algorithm. Our solution was to recast the quartic to be a function, $Q(\\gamma)$, of $\\gamma$. We demonstrated that $Q(\\gamma)$ exhibits only one physical root. However, Newton-Raphson iteration also failed in this case at high $\\gamma$, due to the extreme slope of the quartic near the root, necessitating the use there of an analytical numerical root finder.  Our timing analysis indicated that using $Q(\\gamma)$ with the 8-byte analytical root finder increased run time by only 24\\% compared to using $Q(v)$ with the 8-byte iterative root finder (based on 10 trial runs), while using $Q(\\gamma)$ with the 16-byte analytical root finder ballooned run time by a factor of approximately 400. The iterative root finder is accurate to order 10$^{-4}$ for a sizable region of parameter space corresponding to Lorentz factors on the order of 10$^2$ and smaller. Therefore, we implemented a computational switch that checks the values of $M/E$ and $R/E$ and calls the iterative or analytical root finder accordingly thereby creating a hybrid primitives recovery algorithm called REST\\_FRAME.  In addition, our exploration of parameter space suggests that the discriminant of the cubic resolvent (as defined by Eqn.~\\ref{disc} in $\\S$\\ref{solveqe}) will always be positive for physical flows. Therefore, we did not include code for negative discriminants in our routine. Formal proof remains elusive, however, leaving potential for future work.  We have shown that REST\\_FRAME is capable of calculating the primitive variables from the conserved variables to an accuracy of at least $O(10^{-4})$ for Lorentz factors up to 10$^6$ with significantly better accuracy for Lorentz factors $\\leq 10^5$, and slightly worse (order 10$^{-3}$) for a small portion of the space corresponding to the highest Lorentz factors. We traced the degradation in accuracy for larger Lorentz factors to the effect of subtractive cancellation. Past studies have shown that an accuracy of order 10$^{-4}$ is capable of robustly capturing hydrodynamic structures. We have applied the refined solver to an ultra-relativistic problem and have shown that it is capable of reproducing observed structures and is well-suited to our study of the internal structure of diffuse pulsar wind nebulae.  Our main conclusions are as follows: \\begin{itemize} \\item Relativistic, hydrodynamic simulations have shown that the relatively slow, dense ISM flow resulting from the space motion of a pulsar can set up an interaction with the \\textit{extremely} light, ultra-relativistic pulsar wind leading to an asymmetric nebula with a morphology reminiscent of the Guitar nebula. \\item Simulations have validated the interpretation that a relativistic backflow behind PSR1929+10 is responsible for the X-ray morphology. Results further show that the backflow can harbor a series of internal shockwaves that inflates a nebular bubble, and that the bubble provides positive feedback to the backflow, explaining how the Guitar bubble persists. \\item The evolution of the bubble/backflow structure is sensitive to the choice of input parameters justifying a future series of simulation runs that will determine what pulsar velocities and wind/ISM density ratios are required for the bubble/backflow feedback loop to arise. \\end{itemize}"
    },
    "0606/astro-ph0606154_arXiv.txt": {
        "abstract": "The complexity and accuracy of current and future ``precision cosmology'' observational campaigns has made it essential to develop an efficient technique for directly combining simulation and observational datasets to determine cosmological and model parameters; a procedure we term {\\em calibration}. Once a satisfactory calibration of the underlying cosmological model is achieved, independent predictions for new observations become possible. For this procedure to be effective, robust characterization of the uncertainty in the calibration process is highly desirable. In this {\\em Letter}, we describe a statistical methodology which can achieve both of these goals. An application example based around dark matter structure formation simulations and a synthetic mass power spectrum dataset is used to demonstrate the approach. ",
        "introduction": "It is widely recognized that, beginning in the last decade, a transition to an era of ``precision cosmology'' is well underway. Ongoing and upcoming surveys such as the Wilkinson Microwave Anisotropy Probe (WMAP, Sper\\-gel et al. 2006), the Sloan Digital Sky Survey (SDSS, Adelman-McCarthy et al. 2006), Planck, the Dark Energy Survey (DES), the Joint Dark Energy Mission (JDEM), the Large Synoptic Survey Telescope (LSST), and Pan-STARRS constitute superb sources of cosmological statistics. These sources include (galaxy, cluster, and mass) power spectra and cluster mass functions, from which roughly 25 cosmological parameters have to be constrained (see, e.g., Spergel et al. 2006, Tegmark et al. 2003, Abazajian et al. 2005). The promised accuracy from future observations is remarkable, as some parameters can be measured at the 1\\% level or better, posing a major challenge to cosmological theory. Predictions and analysis methods must at least match -- and preferably substantially exceed -- the observational accuracy. For many observables, this can only be achieved by simulations incorporating physical effects beyond the reach of analytic modeling.  Cosmological simulations already play a key role in the design and interpretation of observations. Controlling systematics is a necessary first step, followed by combining simulations with observations to extract cosmological and model parameters. This cannot be accomplished by brute force. For example, if every parameter is sampled only ten times in a twenty-dimensional parameter space, it would require $10^{20}$ large-scale simulations, which is currently -- and in the near-term -- quite impossible. Even the variation of only a subset of the parameters over a sufficient range is infeasible. The need to develop and employ reliable statistical methods to determine and constrain parameters robustly is therefore manifest.  In this {\\em Letter} we describe a statistical framework to determine cosmological and model parameters and associated uncertainties from simulations and observational data (for an overview of the basic ideas see, e.g., Kennedy \\& O'Hagan 2001 and Goldstein \\& Rougier 2004). The framework integrates a set of interlocking procedures: (i) simulation design -- the determination of the parameter settings at which to carry out the simulations; (ii) emulation -- given simulation output at the input parameter settings, how to estimate the output at new, untried settings; (iii) uncertainty and sensitivity analysis -- determining the variations in simulation output due to uncertainty or changes in the input parameters; (iv) calibration -- combining observations (with known errors) and simulations to estimate parameter values consistent with the observations, including the associated uncertainty; (v) prediction -- using the calibrated simulator to predict new cosmological results with a set of uncertainty bounds.  For concreteness, we discuss the framework methodology in terms of a simple example application: Estimation of five parameters from dark matter structure formation simulations and a synthetic set of ``WMAP + SDSS'' measurements of the matter power spectrum. A detailed description will be provided elsewhere~(S.~Habib et al. in preparation). ",
        "conclusions": "We have introduced a new, very powerful method for determining cosmological and model parameters from simulations and observations. The key idea is to extract maximum utility from a necessarily finite set of expensive simulations. The implementation of this idea includes several valuable features: (i) a design to optimally sample the simulation parameter space; (ii) an accurate emulator capable of generating the required outputs in between the sampled simulation points; (iii) an uncertainty and sensitivity analysis; (iv) the parameter constraints themselves, with associated uncertainty bounds.  In order to demonstrate the basic approach, we used a set of 128 dark matter structure formation simulations and a homogeneous synthetic ``observational'' dataset to determine five cosmological parameters. The next step is to use the framework for analyses of real data, especially of combined datasets such as the CMB and large scale structure observations.  There are many ways to enhance the method and improve its performance. One is the melding of information from codes with different degrees of resolution and input physics, such as in the extraction of information about the mass distribution from the Lyman-$\\alpha$ forest. Here, complex hydrodynamics simulations are certainly desirable, but much faster approximate methods such as hydro-particle mesh (HPM) are available. Thus, a first analysis based on HPM can be performed, narrowing the parameter range of interest sufficiently to make hydro runs feasible. Interesting offshoots of the methodology include the exploitation of certain intermediate results. For instance, a large set of N-body simulations can be performed with several input parameters such as the equation of state for dark energy. An emulator can then be constructed from these and publicly released. This emulator can then be conveniently used instead of real simulations for planning observations and data analysis."
    },
    "0606/hep-th0606033_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\hspace*{15pt}\\par The role of rolling homogeneous scalar field has been widely discussed in the various epoch for a variety of purposes[1]. Recently, with the surprising discovery of an accelerating expansive and spatially flat universe, the scalar field has gained another newly discussion as a candidate for dark energy. It can drive current accelerating expansion while its energy density can fill in the universe as \"missing matter density\". The most popular models with scalar field may be the linear scalar field model( a canonical scalar field described by the lagrangian $L=\\frac{1}{2}\\dot{\\phi}^2-V(\\phi)$)[2-6], the K-essence model( a scalar field with a non-canonical kinetic energy terms)[7-22] and the \"phantom\" model(a scalar field with the negative kinetic energy terms)[23-52]. The potentials in these models are chosen non-negative to avoid negative potential energy density. It is shown that the expanding universe with non-negative potentials have a common property that they will expand for ever, though the evolutional behavior of future universe has significant differences corresponding to different potentials. However, research shows that negative potentials can also lead to a viable cosmology[53-56]. Moreover, the universes with negative potentials are entirely different with the universes with non-negative potentials. They can trigger our flat universe from expansion($H>0$) to contraction ($H<0$), which will never occur in standard FRW model( previous oscillatory model only appears in a close universe in standard FRW model). Hence negative potentials are used to propose the \"cyclic universe\" model. In this cyclic scenario, when the scalar field rolls to a minimum of its effective potential with $V(\\phi)<0$, the universe will stop expanding and contract to a singularity eventually. Additionally, negative potentials also appeared in supergravity theory and in brane cosmology. It is theoretically important to continue investigating the cosmological features in other models where the effective potential $V(\\phi)$ may become negative for some values of the field $\\phi$. \\par Nonlinear Born-Infeld scalar field theory is firstly proposed by W.Heisenberg in order to describe the process of meson multiple production connected with strong field regime[57-59] and then is discussed in cosmology[60-64]. It shows that the lagrangian density of this NLBI scalar field posses some interesting characteristics[65-66]. In Ref[65], the author showed that a singular horizon exists for a large class of solution in which the scalar field is finite. Naked singularities with everywhere well-behaved scalar field in another class of solution have also been found in Ref[65]. Lately the quantum cosmology with the NLBI scalar field has been considered[67]. In the extreme limits of small and large cosmological scale factor the wave function of the universe was found by applying the methods developed by Vilenkin, Hartle and Hawking. The result has suggested a non-zero positive cosmological constant with largest probability, which is consistent with current observational data. The classical wormhole solution and wormhole wavefunction with the NLBI scalar filed has been obtained in Ref[68]. The phantom cosmology based on NLBI scalar field with a special potential had been considered in Ref[69-70]. The results show that the universe will evolve to a de-sitter like attractor regime in the future and the phantom NLBI scalar field can survive till today without interfering with the nucleosynthesis of the standard model. Very recently, with the analysis to Gold supernova data, we show that maybe the NLBI scalar field model is superior to conventional quintessence model[71]. Furthermore it is showed that in another analogous NLBI theories with the lagrangian $p(\\phi,X)=\\alpha^2(\\sqrt{(1+\\frac{2X}{\\alpha^2}}-1)-\\frac{1}{2}m^2\\phi^2$(where $X=\\frac{1}{2}\\dot\\phi^2$) the contribution of the gravitational waves to the CMB fluctuations can be substantially larger than that naively expected in simple inflationary models, which make the prospects for future detection much more promising[72-73]. It is also showed that with the same lagrangian, one can send information from inside a black hole[74]. In Ref[75-76], authors consider a non-Abelian Einstein-Born-Infeld-dilaton theory, where they concern a non-abelian vector field which couples to the dilaton and then describe a dark energy mechanism in a cosmological framework. \\par The key idea of NLBI scalar field theory is that the conventional quintessence scalar field can not describe the reality correctly in the case of strong field. The lagrangian of conventional quintessence model(here we also call it linear scalar field):\\begin{equation}L=\\frac{1}{2}{\\dot\\phi}^2-V(\\phi)\\end{equation} should be substituted by the lagrangian of NLBI scalar field \\begin{equation}L_{NLBI}=\\frac{1}{\\eta}[1-\\sqrt{1-\\eta{\\dot\\phi}^2}]-V(\\phi)\\end{equation} which can recover to conventional case when $\\dot\\phi\\rightarrow 0$. In fact, the lagrangian of NLBI scalar field(Eq.2) implies that there exists a maximum constant value $\\frac{1}{\\sqrt{\\eta}}$ for field velocity $\\dot\\phi$, which is very analogous to the universal constant velocity $c$. It means that $\\dot\\phi$ never reaches infinity while in linear scalar field model there are no such constraint. \\par In this paper, we combine the two ideas(negative potentials and NLBI scalar field theory) and consider the cosmology based on the NLBI scalar field with negative potentials. We think it may be very interesting and meaningful to see what will happen in this case. The paper is organized as follows: In section 2 we will describe theoretical model in NLBI scalar field theory and consider several basic regimes which are possible to happen in NLBI scalar field: the potential energy dominated regime, the kinetic energy dominated regime and the transient regime that the universe switches from expansion to contraction. In section 3, we investigate the different cosmological evolution in different cases and plot corresponding evolutive behaviors in detail. For the potential $V(\\phi)=\\frac{1}{2}m^2\\phi^2+V_0(V_0<0)$ We consider the universe evolution with different slope $m$ and different potential well $V_0$. the cases that $V_0>0$ and $V_0=0$ are also presented to compare with the case $V_0<0$, moreover we compare the different evolution between NLBI scalar field and linear scalar field. in section 4 we mention the cyclic model and consensus model. Conclusion and summary is also presented in section 4. ",
        "conclusions": "The main goal of this paper is to perform a general investigation of the NLBI scalar field cosmology with negative potentials. The cosmological solutions in different regime have obtained through some approximate approach. The results obtained in NLBI scalar field theory are quite different with that obtained in linear scalar field theory. A notable characteristic is that NLBI scalar field behaves as ordinary matter nearly the singularity while the linear scalar field behaviors as \"stiff\" matter. We also find that, due to the nonlinear effect, the oscillatory motion of $\\phi$ in the vicinity when the universe evolve to contraction from expansion is different to linear scalar field. Moreover, the value of Hubble parameter $H_i$ at time $i$ in NLBI scalar field theory is large than the one in linear scalar field theory. With the investigation of evolution with different value of $m$ and $V_0$, we find that in order to accommodate an accelerating expansive universe in which the large scale structure had formed, the value of $m$ and $|V_0|$ must have a {\\it upper bound}. Finally we review the negative potentials and the new cyclic model."
    },
    "0606/astro-ph0606632_arXiv.txt": {
        "abstract": "Hydrodynamical simulations played an important role in understanding the dynamics and shaping of planetary nebulae in the past century. However, hydrodynamical simulations were just a first order approach. The new millennium arrived with the generalized understanding that the effects of magnetic fields were necessary to study the dynamics of planetary nebulae. Thus, B-fields introduced a whole new number of physical possibilities for the modeling. In this paper, we review observational works done in the last 5 years and several works on magnetohydrodynamics about proto-planetary nebulae, since all the effort has been focused on that stage, and discuss different scenarios for the origin of magnetized winds, and the relation binary-bipolararity. ",
        "introduction": "The origin and evolution of proto-planetary nebulae (PPNs)  and planetary nebula (PNs) represent one of the key questions in our understanding of stellar physics. Modeling the fascinating features displayed by these objects requires not only a better knowledge of stellar structure at the AGB stage (and beyond) but also a proper consideration of the driving mechanisms for mass ejection. The transition from AGB to Post-AGB to PN central stars involves drastically different conditions at every stage. Whereas radiation pressure on dust grains is the most likely mechanism at the AGB phase, as are line-driven winds in the case of PN central stars, for Post-AGB stars the details of the driving force has been relatively unexplored.   To begin this review, it is interesting to mention the increasing number of articles related to magnetic fields in PNs during the last years. Before the IAU Symp. 180 (Groningen 1996), the average number of paper per year was 0.41 (base line of 39 years). For the next 5 years until the IAU Symp. 209 (Canberra 2001), this number increased up to 6.2. In the last 5 years, this numbers growths up to 11.8 at the present IAU Symp. 234 (Hawaii 2006). What is even more interesting is that, in the last five years, 75 \\% of all the papers were observations. Finally, {\\bf measurements of magnetic field intensities and their orientations have been observed !!}. This fact will be important in the next five years, since theoretical models have something robust to start with.  All the theoretical work done in PNs up to 2001 were focus on the line-driven wind theory including weak or moderate magnetic fields frozen in the winds (see review by Garc\\'{\\i}a-Segura 2003). However, since the paper by Bujarrabal et al. (2001), in which it is stated that PPNs could not be explained by radiation forces on the winds, most of the work done in the last five years is foccus on PPNs and the winds from post-AGBs stars.  In this paper, we first make a short review of the observations done in the last five years, and then proceed with the MHD work. ",
        "conclusions": ""
    },
    "0606/astro-ph0606318_arXiv.txt": {
        "abstract": "We develop a series expansion of the plasma screening length away from the classical limit in powers of $\\hbar^{2}$. It is shown that the leading order quantum correction increases the screening length in solar conditions by approximately $2\\%$ while it decreases the fusion rate by approximately $ 0.34\\%$. We also calculate the next higher order quantum correction which turns out to be approximately $0.05\\%$. ",
        "introduction": "Salpeter\\cite{salpeter54} wrote a seminal paper  more than half a century ago, concerning screening effects on thermonuclear reaction rates. He made the basic point that screening effects are small at the center of the sun.  There has been renewed interest more recently in utilizing the sun as a source of neutrinos to test the standard model of unification of electro-weak forces.  The measured neutrino flux deviates from predictions of the standard model by a factor of two\\cite{bethe90}. The measurement uncertainty is\\cite{bethe90,gruzinov98} $\\sim 1\\%$. Therefore it would be meaningful to quantify the theoretical estimate with equal precision. The structure and dynamics of the sun are complex\\cite{bahcall94}. Various phenomena need to be identified and estimated correctly.  Screening of Coulomb repulsion between nuclei at extremely short distances is one of them. Many calculations of screening have been made since Salpeter's original paper, attempting to refine the degree of screening\\cite{brown97,brown06,fiorentini04}, dynamic effects\\cite{carraro88}, quantum fluctuations\\cite{gruzinov98,gervino05}, etc\\cite{bahcall02}.  Here we shall focus on quantum corrections to screening.  The most sophisticated calculation of this effect is that of Gruzinov and Bahcall\\cite{gruzinov98}. In this paper, the electronic density matrix was evaluated accurately using Feynman's formulation in terms of a Schroedinger equation with the inverse temperature playing the role of imaginary time. Fermion statistics are ignored due to the high solar temperature. They sustain Salpeter's original conclusion that quantum corrections are minor.  These calculations are essentially correct, but cannot estimate in a systematic fashion the next higher order quantum correction.  We shall correct that deficiency in this paper.  This paper was written for the sake of completeness, since the super-Kamiokande experiment has been successful in obtaining evidence for neutrino mass (see Hosaka et al\\cite{kam1} for recent results). Nevertheless, the results of this paper may yet prove useful for more precise quantitative interpretation of stellar experimental data\\cite{pinsonn06}. ",
        "conclusions": "Systematic quantum corrections to screening in thermonuclear fusion were derived  in powers of $\\hbar^{2}$, and estimated for solar conditions. Leading order corrections were shown to be less than $0.34\\%$ under solar conditions, while the next leading order term is $\\sim 0.05\\%$. Our corrections are consistent with those previously obtained by Gruzinov and Bahcall\\cite{gruzinov98}. They complement the results of Brown et al\\cite{brown06} who show that classical, non-linear effects on screening are small."
    },
    "0606/astro-ph0606404_arXiv.txt": {
        "abstract": "{}{We aim to study the geometry and kinematics of the disk around the Be star $\\alpha$ Arae as a function of wavelength, especially across the Br$\\gamma$ emission line. The main purpose of this paper is to answer the question about the nature of the disk rotation around Be stars.}{We use the VLTI/AMBER instrument operating in the K band which provides a gain by a factor 5 in spatial resolution compared to previous VLTI/MIDI observations. Moreover, it is possible to combine the high angular resolution provided with the (medium) spectral resolution of AMBER to study the kinematics of the inner part of the disk and to infer its rotation law. }{We obtain for the first time the direct evidence that the disk is in keplerian rotation, answering a question that occurs since the discovery of the first Be star $\\gamma$ Cas by father Secchi in 1866. We also present the global geometry of the disk showing that it is compatible with a thin disk + polar enhanced winds modeled with the SIMECA code. We found that the disk around $\\alpha$ Arae is compatible with a dense equatorial  matter confined in the central region whereas a polar wind is contributing along the rotational axis of the central star. Between these two regions the density must be low enough to reproduce the large visibility modulus (small extension) obtained for two of the four VLTI baselines. Moreover, we obtain that $\\alpha$ Arae is rotating very close to its critical rotation. This scenario is also compatible with the previous MIDI measurements.}{} ",
        "introduction": "The star $\\alpha$ Arae (HD\\,158\\,427, HR\\,6510, B3\\,Ve), one of the closest (d=74 pc, Hipparcos, Perryman et al. \\cite{perryman}) Be stars, was observed with the VLTI/MIDI instrument at 10 $\\mu$m in June 2003 and its circumstellar environment was unresolved even with the 102m baseline (Chesneau et al. \\cite{chesneau}, hereafter paper I). $\\alpha$ Arae was a natural choice as first target due to its proximity but also its large mid-IR flux and its high infrared excess among other Be stars, e.g. E(V-L)$\\sim$1.8 and E(V-12$\\mu$m)$\\sim$2.23. These first IR interferometric measurements indicated that the size of the circumstellar environment was smaller than predicted by Stee \\cite{Stee4} for the K band. The fact that $\\alpha$~Arae remain unresolved, but at the same time had strong Balmer emission, have put very strong constraints on the parameters of its circumstellar disk. Independently of the model, they have obtained an upper limit of the envelope size in the N band of $\\phi_{max}$= 4 mas, i.e. 14 R$_{\\star}$ if the star is at 74 pc according to Hipparcos parallax or 20 R$_{\\star}$ if the star is at 105 pc as suggested by the model presented in paper I. \\\\  They finally propose a scenario where the circumstellar environment remains unresolved due to an outer truncation of the disc by an unseen companion. Nevertheless, this companion would be too small and too far away to have any influence on the Be phenomenon itself.\\\\  In order to study the inner part of this circumstellar truncated disk we have taken advantage of the higher spatial resolution by observing at 2 $\\mu$m with the VLTI/AMBER instrument in February 2005. It  provides a gain by a factor 5 in spatial resolution compared to VLTI/MIDI observations. We present in this paper these measurements showing, for the first time, a fully resolved circumstellar envelope in the Br$\\gamma$ emission line and a clear signature of a Keplerian rotating disk around $\\alpha$ Arae. We also discuss the challenging question on the nature of the geometry of the Be disks and particularly their opening angle since it is still an active debate.\\\\  Following the Wind Compressed Disk model (WCD) by Bjorkman \\& Cassinelli \\cite{Bjorkman93}, most authors have considered geometrically thin disks (half opening angle of 2-5 degrees) even if Owocki et al. (1996) have found that the equatorial wind compression effects are suppressed in any radiatively driven wind models for which the driving forces include a significant part from optically thick lines. Moreover, they found that gravity darkening effects can lead to a reduced mass loss, and thus a lower density in the equatorial regions. A wind compression effect is, however, not required to produce small opening angle of the disk. The investigation of accretion disks has shown that discs in hydrodynamical equilibrium and Keplerian rotation will not have much larger opening angles, since their scale height is governed by the vertical gas pressure only. For a disk to be thicker, either additional mechanisms have to be assumed, or it might not be in equilibrium at the radii in question (Bjorkman \\& Carciofi 2004).\\\\  On the other side, Stee et al. (1995; 1998) claimed that Be disks must be more ellipsoidal in order to reproduce the strong IR excess observed and interpret the possibility for a Be star to change its spectral type from Be to B, and more rarely Be to Be-Shell type, i.e. where the disk is dense enough to produce a strong \"shell\" absorption.\\\\   In the following we adopt, as a starting point, the same parameters for the modeling of $\\alpha$ Arae, e.g. the central star and its circumstellar envelope used in paper I and summarized in Table 1.  Following the polarization measurements Pl$\\approx$0.6\\% and Position Angle (PA) of 172$^\\circ$ by McLean\\&Clarke \\citealp{mclean} and Yudin \\citealp{yudin2} the disk major-axis orientation is expected to be at about PA$\\approx82^\\circ$ (Wood et al.  \\citealp{wood96a}, \\citealp{wood96b}, Quirrenbach et al. \\citealp{quirrenbach2}). Assuming a stellar radius of 4.8~R$_{\\sun}$ and an effective temperature $T_{\\rm eff}=18\\,000$\\,K, the photospheric angular diameter is estimated to be 0.7~mas (Cohen et al. \\citealp{cohen}, Chauville et al. \\citealp{chauville}). For the distance of 74~pc and a baseline of 60m at 2~$\\mu$m, Stee \\cite{Stee4} predicts the visibility of $\\alpha$~Arae to be lower than 0.2, i.e. fully resolved.\\\\  The paper is organized as follows. In Section~\\ref{secobs} and \\ref{datared} we present the interferometric AMBER observations and the data reduction. In Section~\\ref{toymodels} we try to obtain a first estimate of $\\alpha$ Arae's envelope geometry using very simple \"toy\" models. Section~\\ref{secsimeca} describes briefly the SIMECA code. In Section \\ref{bestmodel} we present the best model we obtain with SIMECA that fits both the Br$\\gamma$ line and the visibility modulus and phase as a function of wavelength which allows us to infer the disk kinematics and its rotational velocity. Finally, Section \\ref{conclusion} draws the conclusions from these first spectrally resolved interferometric measurements of a Be star at 2 $\\mu$m.  \\vspace{0.3cm} { \\begin{table} {\\centering \\begin{tabular}{cc} \\hline parameter/result    & value \\\\ \\hline Spectral type& B3Ve\\\\ $T_{\\rm eff}$& 18\\,000\\,K\\\\ Mass& 9.6 M\\( _{\\sun } \\)\\\\ Radius& 4.8 R\\( _{\\sun } \\)\\\\ Luminosity& 5.8 10\\( ^{3} \\)L\\( _{\\sun } \\)\\\\ Inclination angle i & 45$\\degr$\\\\ Photospheric density ($\\rho_{phot}$)&1.2 10\\( ^{-12} \\)g cm\\( ^{-3} \\)\\\\ Photospheric expansion velocity& 0.07 km s\\( ^{-1} \\) \\\\ Equatorial rotation velocity & 300 km s\\( ^{-1} \\)  \\\\ Equatorial terminal velocity & 170 km s\\( ^{-1} \\) \\\\ Polar terminal velocity & 2000 km s\\( ^{-1} \\) \\\\ Polar mass flux & 1.7 10\\( ^{-9} \\)M\\( _{\\sun } \\) year\\( ^{-1} \\) sr\\( ^{-1} \\) \\\\ m1 & 0.3 \\\\ m2 & 0.45 \\\\ C1 & 30\\\\ Mass of the disk & 2.3 10\\( ^{-10} \\)M\\( _{\\sun } \\) \\\\ Mass loss & 6.0 10\\( ^{-7} \\)M\\( _{\\sun } \\) year\\( ^{-1} \\)\\\\ \\hline \\end{tabular}\\par}  \\caption{Model parameters for the $\\alpha$ Arae central star and its circumstellar environment obtained from paper I} \\label{midi_model} \\end{table} \\par} \\vspace{0.3cm} ",
        "conclusions": "\\label{conclusion}  \\begin{enumerate} \\item Thanks to these first spectrally resolved interferometric measurements of a Be star at 2 $\\mu$m we are able to propose a possible scenario for the Be star $\\alpha$ Arae's circumstellar environment which consist in a thin disk + polar enhanced winds that is successfully modeled with the SIMECA code. \\item We found that the disk around $\\alpha$ Arae is compatible with a dense equatorial  matter confined in the central region whereas a polar wind is contributing along the rotational axis of the central star. Between these two regions the density must be low enough to reproduce the large visibility modulus (small extension) obtained for two of the four VLTI baselines. This new scenario is also compatible with the previous MIDI measurements and the fact that the outer part of the disk may be truncated by a unseen companion at 32~R$_\\star$. \\item We obtain for the first time the clear evidence that the disk is in Keplerian rotation, closing a debate that occurs since the discovery of the first Be star $\\gamma$ Cas by father Secchi in 1866. \\item We found that that $\\alpha$ Arae must be rotating very close to its critical velocity. \\item These observations were done using the medium (1500) spectral resolution of the VLTI/AMBER instrument and are very promising for the forthcoming AMBER high spectral resolution observational mode (10000) and the coupling of the Auxiliary Telescopes (ATs) on the VLTI array. \\end{enumerate}"
    },
    "0606/astro-ph0606297_arXiv.txt": {
        "abstract": "We present deep radio observations of the most distant complete quasar sample drawn from the Sloan Digital Sky Survey. Combining our new data with those from literature we obtain a sample which is $\\sim 100$ per cent complete down to $\\rm S_{1.4GHz} = 60\\mu Jy$ over the redshift range $3.8 \\le z \\le 5$.  The fraction of radio detections is relatively high ($\\sim 43$ per cent), similar to what observed locally in bright optical surveys. Even though the combined radio and optical properties of quasars remain overall unchanged from $z \\sim 5$ to the local Universe, there is some evidence for a slight over-abundance of radio-loud sources at the highest redshifts when compared with the lower-z regime.  Exploiting the deep radio VLA observations we present the first attempt to directly derive the radio luminosity function of bright quasars at $z \\simgt 4$. The unique depth -- both in radio and optical -- allows us to thoroughly explore the population of optically bright FR~II quasars up to $z \\sim 5$ and opens a window on the behaviour of the brightest FR~I sources. A close investigation of the space density of radio loud quasars also suggests a differential evolution, with the more luminous sources showing a less pronounced cut-off at high z when compared with the less luminous ones. ",
        "introduction": "During the past years, studies of the properties of Active Galactic Nuclei (AGNs) and in particular of their cosmological evolution have become of major relevance within the more general field of galaxy evolution. In fact, it has been found that the properties of the central black hole (BH) are tightly related to those of the host galaxy (e.g. Magorrian et al. 1998; Ferrarese \\& Merritt 2000; Tremaine et al. 2002; McLure \\& Dunlop 2004), so that the energetic feedback that AGN activity can release is now a fundamental ingredient in many theoretical models of galaxy formation (Silk \\& Rees 1998; Fabian 1999; Granato et al. 2001, 2004; Cavaliere \\& Vittorini 2002; Di Matteo et al. 2005; Cirasuolo et al. 2005b).  Previous studies in the optical and X-rays -- bands which are thought to trace the accretion processes onto the central BH -- have established quasar and powerful AGN activity to peak at $z \\sim 2$, with a rapid decline at lower redshifts (e.g. Boyle et al. 2000; Ueda et al. 2003; Croom et al. 2004); on the other hand, the less powerful sources have their major shining phase at $z \\simlt 1$ (e.g. Ueda et al. 2003; Hasinger et al. 2005). The behaviour of the AGN evolution at higher redshifts has been very uncertain for a long time as the relevant observations were biased by selection effects and only considered very small numbers of objects.  The advent of the recent Sloan Digital Sky Survey (SDSS; Fan et al. 1999, York et al. 2000) has allowed to properly explore the high redshift Universe up to $z \\sim 6$ (Fan et al. 2001a,b; 2004). These studies have confirmed the presence of a cut off in the space density of quasars at $z\\sim 2$.  Even though AGNs that show radio emission are only a small fraction of the total population (Sramek \\& Weedman 1980; Condon et al. 1981; Marshall 1987; Miller, Peacock \\& Mead 1990; Kellermann et al. 1989), they represent an important subsample as the radiation at centimetre wavelengths is unaffected by dust obscuration and reddening. Therefore, studies of the evolution of radio-active AGNs provide a less biased view of the behaviour of massive BHs and accretion processes onto them as a function of cosmic time.  Several studies have addressed the evolutionary trend of radio loud sources from the local Universe up to high z (Dunlop \\& Peacock 1990; Toffolatti et al. 1998; Jackson \\& Wall 1999; De Zotti et al. 2005). Shaver et al. (1996, 1999) argued for a drop in the space density of flat spectrum radio quasars by more than a factor 10 between $z\\sim 2.5$ and $z\\sim 6$.  However, a re-analysis of such sources (Jarvis \\& Rawlings 2000) suggests a more gradual (factor $\\sim 4$) decline, decline which is backed up by the work of Jarvis et al. (2001) on steep spectrum sources in the same redshift interval. A luminosity dependent cut-off, with a decrease in space density less dramatic for the most luminous radio sources, has been claimed by Dunlop (1998) and confirmed by other recent studies (i.e. Vigotti et al. 2003; Cirasuolo et al. 2005a).  Unfortunately, the process(es) responsible for the formation of radio jets that mark the class of radio loud objects are still poorly understood.  The mass of the central BH could play an important role in shaping the transition between the population of radio loud (RL) and radio quiet (RQ) AGNs. As recently pointed out by many studies, RL sources seem to have the BHs confined to the upper end of the BH mass function, whereas the BHs in RQ quasars appears to span the full range in BH mass (Laor 2000; McLure \\& Dunlop 2002; Dunlop et al. 2003; Marziani et al. 2003; McLure \\& Jarvis 2004; Metcalf \\& Magliocchetti 2006). Furthermore, the analysis of a large sample of local low luminosity AGNs drawn from SDSS showed the fraction of galaxies hosting a radio-loud AGN to be a strong function of BH and stellar mass (Best et al. 2005).  However, the point is still controversial, and some authors claim no evidence for any relation between radio power and mass of the central BH (Oshlack et al. 2002; Ho 2002; Woo \\& Urry 2002a,b; but see the dissenting view of Jarvis \\& Mclure 2002 and McLure \\& Jarvis 2004, who ascribe the lack of correlation reported by these latter authors as due to selection effects such as Doppler Beaming and orientation).  In the light of the above discussion, the present work is aimed at exploring the radio properties of the highest redshift quasars. The main goal is to investigate if the physical conditions in the early stages of galaxy formation can favour or prevent the formation of relativistic jets and also to test if the radio loudness in quasars exhibits some dependence on cosmic epoch.  For this purpose, we performed deep radio observations of a sample of high redshift quasars selected from SDSS. The optical sample is presented in Section 2, while radio observations and the radio properties of the sample are respectively described in Sections 3 and 4. By  exploiting this unique sample we derive the radio luminosity function in Section 5 and investigate the behaviour of the space density of high redshift QSOs as a function of redshift in Section 6. Our discussion and conclusions are presented in Section 7.  Throughout this paper we adopt the ``concordance'' cosmology, consistent with the Wilkinson Microwave Anisotropy Probe data (Bennett et al. 2003), i.e.: $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70~{\\rm km~s^{-1}}$. ",
        "conclusions": "We have presented deep radio observations of the most distant complete quasar sample drawn from the Sloan Digital Sky Survey. Combining our deep VLA observations with the ones performed by Carilli et al. (2001) and also with the 5 detections from FIRST, we obtained $\\sim 100$ per cent completeness down to $S_{\\rm 1.4 GHz} = 60 \\; {\\rm \\mu Jy}$ over the redshift range $3.8 \\le z \\le 5$.  The fraction of radio detections is relatively high ($\\sim 43$ per cent), similar to what observed locally for bright optical surveys such as the Palomar Bright Quasar Survey (Kellermann et al. 1989) and the LBQS (Hewett et al. 2001). A comparison between the $R^*_{1.4}$ distribution of these high redshift radio quasars with the one derived at $z \\simlt 2$ by Cirasuolo et al. (2003b) suggests that the combined radio and optical properties of quasars might remain overall unchanged from $z \\sim 5$ to the local Universe. However, even though the shape of the $R^*_{1.4}$ distribution is roughly preserved over cosmic time, there is some marginal evidence for a slight over-abundance of radio loud sources at high z when compared with the low redshift samples (see Figure \\ref{histor}), even though not statistically significant due to the small number of sources. Furthermore, it is worth noting that the adoption of a flatter $\\alpha_{\\rm R}$ for high z objects would reduce their radio power and shift them towards lower values of $R^*_{1.4}$, therefore somehow reducing the fraction of purely RL ($R^*_{1.4}$) sources.  An interesting hint to shed some light on the above issue comes from comparisons with the parent optical population. The space density of bright optical quasars ($M_B < -26.5$) at $z\\sim 2$ and $z \\sim 4.4$ is $\\rho_O \\sim 2 \\times 10^{-7} \\rm Mpc^{-3}$ and $\\rho_O \\sim 1.5 \\times 10^{-8} \\rm Mpc^{-3}$, respectively (see Fan et al. 2001b, 2004).  The ratio between the space densities of radio sources with \\lpow $\\ge 24.4$ \\whs (as plotted in Figure \\ref{cumul}) and the $\\rho_O$ of the total optical population is therefore $0.025 \\pm 0.01$ at $z \\sim 2$ and $0.15 \\pm 0.1$ at $z \\sim 4.4$. These figures have been obtained for a radio spectral index $\\alpha_R=0.5$ but, as shown in section 6, the results are not expected to exhibit great variations by adopting a flatter slope. Again, this suggests that at high redshifts the probability of having RL sources is enhanced with respect to that at lower redshifts. However, we stress once more that the statistics we dealt with in this work is very poor and further data and larger samples are needed in order to have a more robust answer. \\\\ Even though a detailed investigation of this phenomenon is outside the possibility of the present data, in a very qualitative way we can relate the suggested excess of RL sources at high $z$ as compared with the lower redshift regime with changes in the accretion rates with cosmic time. The physical conditions of the primordial massive galaxies hosting quasars and the availability of a larger amount of gas in these early stages could in fact allow super-Eddington accretions and eventually favour the formation of powerful radio jets.  We also attempted the first direct estimate of the radio luminosity function of bright quasars at $z \\simgt 4$. Exploiting the deep radio flux limits obtained through VLA observations, we were able to trace the RLF down to \\lpow $\\sim 23.6$ \\whs. It is worth noticing that the transition region between the FR~I and FR~II population occurs at \\lpow $\\sim 24.4$ \\whs.  Therefore, the unique depth -- both in radio and optical -- of this high redshift quasar sample allows us to completely explore the population of optically bright FR~II quasars up to $z\\sim 5$ and furthermore opens a window on the behaviour of the brightest FR~I sources.  Finally, close investigation of the RL quasar space density at different redshifts is suggestive of a differential evolution for the two populations of optically faint and bright objects.  The more luminous sources in fact show a less pronounced cut-off at high z -- with a drop in their space density of only a factor $\\sim 2$ between $z \\sim 2$ and $z \\sim 4.4$ -- when compared with the less luminous ones. Even though the lack of information on the behaviour of optically faint quasars at $z \\simgt 2.2$ does not allow any definitive conclusion, our results indicate the cosmological evolution of radio activity in quasars to be a function of their optical power."
    },
    "0606/astro-ph0606068_arXiv.txt": {
        "abstract": "The observation of high-energy extraterrestrial neutrinos is one of the most promising future options to increase our knowledge on non-thermal processes in the universe. Neutrinos are e.g.\\ unavoidably produced in environments where high-energy hadrons collide; in particular this almost certainly must be true in the astrophysical accelerators of cosmic rays, which thus could be identified unambiguously by sky observations in ``neutrino light''. To establish neutrino astronomy beyond the detection of single events, neutrino telescopes of km$^3$ scale are needed. In order to obtain full sky coverage, a corresponding detector in the Mediterranean Sea is required to complement the IceCube experiment currently under construction at the South Pole. The groups pursuing the current neutrino telescope projects in the Mediterranean Sea, ANTARES, NEMO and NESTOR, have joined to prepare this future installation in a 3-year, EU-funded Design Study named KM3NeT (in the following, this name will also denote the future detector). This report will highlight some of the physics issues to be addressed with KM3NeT and will outline the path towards its realisation, with a focus on the upcoming Design Study. ",
        "introduction": "\\label{sec-phy}  The energy range accessible to neutrino telescopes is intrinsically limited by the detection method to some $10\\gev$ at its lower end, while at energies beyond roughly $10^{17}\\ev$ the neutrino flux is expected to fade below detection thresholds even for future km$^3$-scale detectors. The lower-energy region is dominated by the flux of {\\it atmospheric neutrinos} (cf.\\ Fig.~\\ref{fig-difl}), produced in reactions of cosmic rays with the atmosphere. There are three approaches to identify cosmic muon signals on top of this background: \\begin{enumerate} \\item Neutrinos from specific astrophysical objects ({\\it point sources}) produce excess signals associated to particular celestial coordinates. \\item Neutrinos not associated to specific point sources ({\\it diffuse flux}) are expected to have a much harder energy spectrum than the atmospheric neutrinos and to dominate the neutrino flux above $10^{14}\\rnge10^{15}\\ev$. \\item Exploitation of coincidences in time and/or direction of neutrino events with observations by telescopes (e.g.\\ in the radio, visible, X-ray or gamma regimes) and possibly also by cosmic ray detectors ({\\it multimessenger method}). \\end{enumerate}  The various astro- and particle physics questions to be addressed with the resulting data have been summarised e.g.\\ in \\cite{astro-ph-0503122} and references therein. Here, we will focus on a few central topics, including a recent development:  \\subsection{Neutrinos from galactic shell-type supernova remnants} \\label{sec-phy-snr}  \\begin{figure}[b] \\begin{center} \\epsfig{file=plots/FoV.eps,width=\\columnwidth} \\end{center} \\caption{Field of view of a neutrino telescope at the South Pole (top) and in the Mediterranean (bottom), given in galactic coordinates. A $2\\pi$-downward sensitivity is assumed; the gray regions are then invisible. Indicated are the positions of some candidate neutrino sources.} \\label{fig-fov} \\end{figure}  The shock waves developing when supernova ejecta hit the interstellar medium are prime candidates for hadron acceleration through the Fermi mechanism. Recent observations of gamma rays up to energies of about $40\\tev$ from two shell-type supernova remnants in the Galactic plane (RX\\,J1713.7-3946 and RX\\,J0852.0-4622) \\cite{astro-ph-0511678,aa:437:l7} with the H.E.S.S.\\ \\v Cerenkov telescope support this hypothesis and disfavour explanations of the gamma flux by purely electromagnetic processes. The detection of neutrinos from these sources would, for the first time, identify unambiguously specific cosmic accelerators. Note that this is only possible with Northern-hemisphere neutrino telescopes which, in contrast to the South Pole detectors, cover the relevant part of the Galactic plane in their field of view (see Fig.~\\ref{fig-fov}).  The expected event rates can be estimated using the rough assumptions that the gamma flux follows a power-law spectrum without high-energy cut-off and the muon neutrino and gamma fluxes are in relation $\\phi_{\\nu_\\mu}/\\phi_\\gamma=1/2$, taking into account the relative production probabilities of charged and neutral pions, their decay chains and neutrino oscillations. Preliminary calculations indicate that the first-generation Mediterranean neutrino telescopes may have a chance to observe a few events, whereas a significantly larger signal is expected in a future cubic-kilometre set-up; a tentative estimate of the neutrino sky map of RX\\,J0852.0-4622 after 5~years of data taking with KM3NeT is shown in Fig.~\\ref{fig-km3hess}.  \\begin{figure}[hbt] \\epsfig{file=plots/km3net_skymap.eps,width=\\columnwidth,% bblly=24,bbllx=19,bbury=552,bburx=570,clip=} \\caption{\\protect\\raggedright A skymap of the simulated neutrino signal from RX\\,J0852.0-4622 as seen by a km$^3$-scale neutrino telescope in the Mediterranean Sea after 5 years of data taking. In the simulation, a power-law gamma spectrum without energy cut-off and the relation $\\phi_{\\nu_\\mu}/\\phi_\\gamma=1/2$ have been assumed. The background of atmospheric neutrinos, not included in the plot, can be efficiently reduced by adjusting the lower energy cut without affecting significantly the signal. The circle in the lower left corner indicates the average angular resolution (point spread function).} \\label{fig-km3hess} \\end{figure}  \\subsection{The diffuse neutrino flux} \\label{sec-phy-dif}  The sensitivity of current and future experiments is compared to various predictions of diffuse neutrino fluxes in Fig.~\\ref{fig-difl} (following \\cite{jp:g29:843,arevns:g29:843}). Whereas some of the models are already now severely constrained by the data, others require km$^3$-size neutrino telescopes for experimental assessment and potential discoveries. The measurement of the diffuse neutrino flux would allow for important clues on the properties of the sources, on their cosmic distribution, and on more exotic scenarios such as neutrinos from decays of topological defects or superheavy particles ({\\it top-down scenarios}).  \\begin{figure}[htb] \\epsfig{file=plots/diff_fluxes.eps,width=\\columnwidth,% bblly=24,bbllx=19,bbury=428,bburx=572,clip=} \\caption{Experimental sensitivity to the diffuse neutrino flux for various current and future experiments (red lines), compared to different models for contributions to the diffuse flux (numbered lines). See \\pcite{arevns:g29:843} for detailed explanations. The flux of atmospheric neutrinos is indicated as blue band. Plot provided by courtesy of C.~Spiering. \\label{fig-difl}} \\end{figure}  \\subsection{Search for dark matter annihilation} \\label{sec-phy-dms} The major part of the matter content of the universe is nowadays thought to be formed by {\\it dark matter}, i.e.\\ non-baryonic, weakly interacting massive particles (WIMPs); an attractive WIMP candidate is the lightest supersymmetric particle, the neutralino. Complementary to direct searches for WIMPs, indirect WIMP observations could be possible by measuring neutrinos produced in WIMP annihilation reactions in regions with enhanced WIMP density. Such accumulations may in particular occur due to gravitational trapping, e.g.\\ in the Sun or the Galactic Centre.  The WIMP signal would be an enhanced neutrino flux from these directions, with a characteristic upper cut-off in the energy spectrum below the WIMP mass, $M_\\text{WIMP}$. Although there is no generic upper constraint on $M_\\text{WIMP}$, supersymmetric theories prefer values below $1\\tev$. Substantial detection efficiency down to order $100\\gev$ is therefore essential for indirect WIMP searches through neutrinos. The expected sensitivity depends strongly on assumptions on the WIMP density profile, on $M_\\text{WIMP}$ and on the energy spectrum of neutrinos from WIMP annihilations. At least for some supersymmetric scenarios this sensitivity is compatible or even better than for direct searches \\cite{astro-ph-0503122}. ",
        "conclusions": "\\label{sec-con}  Neutrino astronomy is an emerging field in astroparticle physics offering exciting prospects for gaining new insights into the high-energy, non-thermal processes in our universe. The current neutrino telescope projects in the Mediterranean Sea are approaching installation and promise exciting first data. They have reached a level of technical maturity allowing for the preparation of the next-generation cubic-kilometre detector to complement the IceCube telescope currently being installed at the South Pole. The interest in this project has been further enhanced by the recent H.E.S.S.\\ observations of high-energy gamma rays from shell-type supernova remnants in the Galactic plane, indicating that these objects could well be intense neutrino sources, which would, however, be invisible to IceCube.  The technical design of the future Mediterranean km$^3$ neutrino telescope will be worked out in the 3-year EU-funded KM3NeT Design Study starting in February 2006. The construction of the KM3NeT neutrino telescope during the first years of the next decade thus appears to be possible.  \\vspace*{3.mm}  \\noindent {\\bf Acknowledgement:} The author wishes to thank the organisers of the VLVnT2 Workshop in Catania for their overwhelming hospitality and a very intense, productive and perfectly organised workshop.  {"
    },
    "0606/hep-th0606110_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent works have indicated that in brane inflation models cosmic strings are copiously produced during the brane collision \\cite{tye1,tye2}. This has led to a renewed interest in the physics of cosmic strings and to consider the exciting possibility that there could be long-lived fundamental strings of cosmic size (for reviews, see \\cite{kibble1,polchinski,tye3,kibble2} and references therein). Finding such objects  could constitute a test of string theory.  The dynamics of cosmic strings could lead to interesting astrophysical events such as gravitational waves or black hole formation. Some aspects of the dynamics of cosmic string interactions were studied in \\cite{myers,copeland} (for Abrikosov-Nielsen-Olesen strings, see recently \\cite{achucarro}).  Here we will develop in full detail the classical formalism of string splitting, joining and intercommutation. Our formulas (Appendix) provide explicit expressions for the outgoing string solution starting with an arbitrary initial string configuration before interaction. Understanding the dynamics and the different features of splitting and joining processes is of interest, since these processes are the basis of the interaction rules in string theory.  As an application, we will study the process of possible gravitational collapse as a result of the collision of cosmic fundamental strings. Surprisingly, we will find that gravitational collapse is a quite common phenomenon ensuing  the encounter of strings of equal and opposite maximal angular momentum, which classically are rotating straight strings, and folded in the case of the closed string. If the initial strings just touch at the end points, then they can join forming one single string. If they meet at some intermediate point, then they can interconnect giving rise to two new strings.  We then study the evolution of the resulting strings by the standard flat spacetime dynamics. We find that they typically contract in a finite time to a minimum size, which sometimes is smaller than the gravitational radius $R_s$.  If the strings meet with zero relative transverse momentum, we find that, for generic values of the intersecting positions and angle, a finite fraction of the mass of the resulting interconnected strings collapses into a mathematical point.  In any of these situations, gravitational forces become very strong when the string size approach $R_s$ and should enhance the evolution towards the collapse, ensuing in the formation of a horizon and hence a black hole \\cite{hawking,polmarev,vilenkin} (other  discussions can be found e.g. in \\cite{DV,matsuda}). In our computation the mass (proportional to the length) of the strings appear as an overall scale, therefore this phenomenon can occur for arbitrarily large values of the mass.  If the transverse momentum is not zero, then the resulting strings will be stretched in the transverse direction. The size is of the order of the length times the relative transverse velocity $v$. For non-relativistic relative motions with $v$ much less than the product of the gravitational constant times the string tension, this size will be much smaller than the gravitational radius. We conclude that also in this case, the  interconnection of our strings generically leads to the formation of a black hole.  We also consider a long-lived version of the string with maximum angular momentum. This is a closed string with some component in extra dimensions, whose motion in $3+1$ dimensions is the same as the open (or folded) string with maximum angular momentum. We discuss for which  range of the parameters and for which magnitudes of cross sections  black hole formation is to be expected. ",
        "conclusions": ""
    },
    "0606/astro-ph0606657_arXiv.txt": {
        "abstract": "Recently using Particle-In-Cell simulations i.e. in the kinetic plasma description Tsiklauri et al. and G\\'enot et al. reported on a discovery of a new mechanism of parallel electric field generation, which results in electron acceleration. In this work we show that the parallel (to the uniform unperturbed magnetic field) electric field generation can be obtained in much simpler framework using ideal Magnetohydrodynamic (MHD) description, i.e. without resorting to complicated wave particle interaction effects  such as ion polarisation drift and resulting space charge separation which seems to be an ultimate cause of the electron acceleration. In the ideal MHD the parallel (to the uniform unperturbed magnetic field) electric field appears due to fast magnetosonic waves which are generated by the interaction of weakly non-linear Alfv\\'en waves with the transverse density inhomogeneity. Further, in the context of the coronal heating problem  a new {\\it two stage mechanism} of the plasma heating is presented by putting emphasis, first, on the generation of parallel electric fields within {\\it ideal MHD} description directly, rather than focusing on the enhanced dissipation mechanisms of the Alfv\\'en waves and, second, dissipation of these parallel electric fields via {\\it kinetic} effects. It is shown that a single Alfv\\'en wave harmonic  with frequency ($\\nu = 7$ Hz), (which has longitudinal wavelength  $\\lambda_A = 0.63$ Mm for putative Alfv\\'en speed of 4328 km s$^{-1}$) the generated parallel electric field could account for the 10\\% of the necessary coronal heating requirement. We conjecture that  wide spectrum (10$^{-4}-10^3$ Hz) Alfv\\'en waves, based on observationally constrained spectrum, could provide necessary coronal heating requirement. It is also shown that the amplitude of generated parallel electric field exceeds the Dreicer electric field by about four orders of magnitude, which implies realisation of the run-away regime with the associated electron acceleration. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606182_arXiv.txt": {
        "abstract": "During the last two years we have used the Palomar Testbed Interferometer to observe several explosive variable stars, including V838 Monocerotis, V1663 Aquilae and recently RS Ophiuchi. We observed V838 Monocerotis approximately 34 months after its eruption, and were able to resolve the ejecta. Observations of V1663 Aql were obtained starting 9 days after peak brightness and continued for 10 days. We were able to resolve the milliarcsecond-scale emission and follow the expansion of the nova photosphere. When combined with radial-velocity information, these observations can be used to infer the distance to the nova. Finally we have resolved the recurrent nova RS Oph and can draw some preliminary conclusions regarding the emission morphology. ",
        "introduction": "\\label{sect:intro}  %  With the improving sensitivity limits of ground-based optical interferometers it is over time becoming more likely that eruptive transients such as novae will be bright enough to be observed. The high angular resolution measurements that can be made with such systems should allow observers to directly probe some of the aspects of these explosions. Here we discuss three cases where such observations have been made.  V838 Monocerotis is an explosive variable star that underwent a nova-like event in early 2002 \\cite{b02,mun02}, with a peak magnitude of $m_V \\sim 6.8$ (Fig. \\ref{fig:vband}).  However, the eruption was unlike classical novae in that the effective temperature of the object dropped and the spectral type evolved into a very late M and L type \\cite{evans03}. The eruption mechanism of V838 Mon is not well understood, but is probably a new type of explosive variable. There have been many models proposed: the merger of a main-sequence binary star \\cite{sk03}, a He-flash on a post-AGB star\\cite{vl04}, or even the accretion of several planets \\cite{rm03}.  Classical novae are energetic stellar explosions that occur in systems containing a white dwarf (WD) accreting mass from a late-type stellar companion \\cite{hernanz05}. When the amount of accreted material on the surface of the white dwarf reaches some critical value a thermonuclear-runaway is ignited, giving rise to the observed nova outburst in which material enriched in heavy elements is ejected into the surrounding medium at high velocities.  Nova Aquilae 2005 (ASAS190512+0514.2, V1663 Aql) was discovered on 9 June 2005 by G. Pojmanski \\& A. Oksanen\\cite{iauc8540}.  At the time of discovery the magnitude was $m_V$ = 11.05; the source reached $m_V \\sim 10.8$ the following day, and declined in brightness thereafter. The time to decay 2 magnitudes ($t_2$) was $\\sim16$ days, making V1663 Aql a ``fast'' nova\\cite{payne57}. Soon after discovery M. Dennefeld \\& F. Ricquebourg\\cite{iauc8544} obtained an optical spectrum with features indicating a heavily reddened nova. The H-$\\alpha$ emission lines exhibited P Cygni line profiles and indicated an expansion velocity in the range of $700$ km s$^{-1}$ (Dennefeld, personal communication) to $1000$ km s$^{-1}$ \\cite{iauc8640}.  Direct observations of the expansion of the nova shell provide an opportunity to accurately determine the distance to the nova. Such observations are usually only possible many years after the outburst, when the expanding shell can be resolved.  We have used the Palomar Testbed Interferometer (PTI) to resolve the $2.2 \\mu$m emission from V1663 Aql and measure its apparent angular diameter as a function of time. We were able to follow the expansion starting $\\sim 9$ days after the initial explosion; when combined with radial velocities derived from spectroscopy we are able to infer a distance and luminosity of the object.  Recurrent novae are thought to consist of massive white dwarfs orbiting late-type giants; material is pulled from the giant and accreted onto the WD. As in the case of classical novae, a thermonuclear runaway reaction will on occasion blow off large amounts of matter in a bright, rapid explosion. There are a small number of systems which have been observed through multiple such outbursts, including RS Oph, which showed outbursts in 1898, 1933, 1958, 1967, 1985 and 2006. It has been argued\\cite{hk01} that recurrent novae systems are the progenitors of Type Ia supernovae, as the short duration of the explosions indicate that the WD should be very close to the Chandrasekhar limit.   The Palomar Testbed Interferometer (PTI) was built by NASA/JPL as a testbed for developing ground-based interferometry and is located on Palomar Mountain near San Diego, CA \\cite{colavita99}. It combines starlight from two out of three available 40-cm apertures and measures the resulting interference fringes (See Fig. \\ref{fig:uv} for representative $uv$-plane coverage). The high angular resolution provided by this long-baseline (85-110 m), near infrared ($2.2 \\mu$m) interferometer is sufficient to resolve emission on the milli-arcsecond scale. ",
        "conclusions": "In the past two years we have used the Palomar Testbed Interferometer to observe three eruptive variable stars: V838 Mon (peculiar variable), V1663 Aql (classical nova) and RS Oph (recurrent nova). In each case we were able to resolve the emission and measure its angular size. In the case of V1663 Aql we were able to follow the expansion and derive a geometric distance. These observations of V1663 Aql are only the second time  a classical nova has been resolved by optical interferometry."
    },
    "0606/astro-ph0606461_arXiv.txt": {
        "abstract": "The Unification Model for active galactic nuclei posits that Seyfert 2s are intrinsically like Seyfert 1s, but that their broad-line regions (BLRs) are hidden from our view. A Seyfert 2 nucleus that truly lacked a BLR, instead of simply having it hidden, would be a so-called ``true'' Seyfert 2. No object has as yet been conclusively proven to be one. We present a detailed analysis of four of the best ``true'' Seyfert 2 candidates discovered to date: IC 3639, NGC 3982, NGC 5283, and NGC 5427. None of the four has a broad H$\\alpha$ emission line, either in direct or polarized light. All four have rich, high-excitation spectra, blue continua, and \\textit{Hubble Space Telescope} (\\textit{HST}) images showing them to be unresolved sources with no host-galaxy obscuration. To check for possible obscuration on scales smaller than that resolvable by \\textit{HST}, we obtained X-ray observations using the \\textit{Chandra X-ray Observatory}. All four objects show evidence of obscuration and therefore could have hidden BLRs. The picture that emerges is of moderate to high, but not necessarily Compton-thick, obscuration of the nucleus, with extra-nuclear soft emission extended on the hundreds-of-parsecs scale that may originate in the narrow-line region. Since the extended soft emission compensates, in part, for the nuclear soft emission lost to absorption, both absorption and luminosity are likely to be severely underestimated unless the X-ray spectrum is of sufficient quality to distinguish the two components. This is of special concern where the source is too faint to produce a large number of counts, or where the source is too far away to resolve the extended soft X-ray emitting region. ",
        "introduction": "Seyfert Galaxies are relatively low-luminosity Active Galactic Nuclei (AGNs) that are traditionally divided into two types, 1 and 2, following an empirical optical spectroscopic classification scheme proposed over 30 years ago by \\citet{kh74}.  Seyfert 1s are characterized by strong, broad (widths of several thousand km\\,s$^{-1}$) permitted emission lines of Hydrogen, Helium, and \\ion{Fe}{2} and narrow (typical widths of $500-1000$\\,km\\,s$^{-1}$) forbidden lines such as [\\ion{O}{3}], [\\ion{N}{2}], and [\\ion{S}{2}]. Seyfert 2s have only narrow permitted and forbidden lines of comparable width and are characterized by a very high [\\ion{O}{3}]$\\lambda5007$/H$\\beta$ line ratio.  Both classes also often exhibit a blue featureless continuum, often manifested as an ``ultraviolet excess,'' though this is not a necessary criterion.  The leading explanation for the spectral differences between the two classes of Seyferts is the Unification Model originally motivated by observations of polarized broad emission lines in a number of Seyfert 2 galaxies \\citep[and references therein]{ski93}.  Unification posits that all Seyfert galaxies have the same underlying structure: a central continuum source arising from an accretion disk around a central supermassive black hole, a dense broad-line region (BLR) comprised of high-velocity gas located on light-day scales, and a low-density narrow-line region (NLR) on parsec scales.  The BLR is nestled inside a torus of dusty obscuring material surrounding the nucleus, and whether we see a Type 1 or 2 spectrum depends on the orientation of this torus with respect to our line of sight \\citep[see][Fig 7.1]{peterson97}.  In Seyfert 1s we are looking down the opening of the torus and see the BLR and central continuum source, while in Seyfert 2s the torus blocks the central regions from view and we only see the narrow-line spectrum of the NLR.  In some Seyfert 2s a circumnuclear cloud of gas in a favorable location outside the torus opening acts as a polarizing ``mirror'' that affords an oblique view down the unobscured axis of the torus, and we see the BLR and underlying continuum in polarized light.  A large fraction of the Seyfert 2s observed with spectropolarimetry to date have polarized broad lines \\citep{moran00,tran01,l01,lah04}.  Additional support for the Unification scenario comes from X-ray studies that have demonstrated that many Seyfert 2s are highly absorbed at soft X-rays by very large column densities \\citep[$\\log N_H> 23$, e.g.,][]{mea94,maiolino98,g01}.  Indeed, this idea has become so compelling that many obscured AGNs, especially those that are X-ray bright but have no or only very weak optical counterparts, are called ``Type 2 AGN'', often regardless of whether or not their optical spectra satisfy the original \\citeauthor*{kh74} criteria.  While there is little doubt that a large fraction of Seyfert 2s are indeed obscured Seyfert 1s, there is still no good answer to the question of whether or not there exists a subset of spectroscopically-classified Seyfert 2s that intrinsically lack a BLR. Such BLR-free Seyfert 2s have variously been called ``true Seyfert 2s'' \\citep[e.g.,][]{bd86} or ``pure Seyfert 2s'' \\citep[e.g.,][]{h95} to distinguish them from the Seyfert 2s that are hidden Seyfert 1s in the Unification picture.  We shall adopt the term ``true Seyfert 2s'' in this investigation.  Our motivation in taking up this question is not to attempt to disprove the Unification picture, but rather to try to find the interesting exceptions to it.  Identifying and studying objects that lie at the extremes of the parameter space for the accretion process are critical for informing us about the important physical details of that process.   Why might there be no BLR in some AGN?  Recent observations and theoretical work have suggested that the gas responsible for the BLR is related, at least in part, to the accretion process \\citep[e.g., a high-velocity outflow from the disk,][]{mc98}.  One possibility is that as an AGN becomes active, there is a brief period where conditions conducive to formation of a BLR have not yet been established.  Similarly, when an AGN turns off there might come a time during the shutdown when the BLR can no longer be sustained.  The high densities and small scale (few light days) of the BLR means it responds very rapidly to changes in the ionizing continuum, whereas the low-density, extended scale (10s to 100s of parsecs) of the NLR leads to a slower response by many orders of magnitude: the BLR sees the central continuum as it was days ago, whereas the NLR would see it as it was decades to centuries ago.  Historically, at least a few Seyfert 1s have been observed during extremely low continuum states in which the broad components of the Balmer lines have practically vanished for a brief time, e.g., NGC\\,4151 \\citep{pp84}, NGC\\,1566 \\citep{a85}, and Mrk\\,1018 \\citep{c86}.  Another possibility is that there are natural limits to the existence of a BLR.  \\citet[also \\citealt{n03}]{n00} has proposed a scenario in which the broad lines arise in a disk wind that requires a minimum accretion rate, and the BLR naturally vanishes in low-luminosity objects that fall below the minimum accretion-rate threshold.  \\citet{laor03} has likewise proposed that the BLR may vanish in low-luminosity AGNs, based on extrapolations of empirical scaling relations for luminosity, BLR size scale, and line widths that have been discovered in studies of AGN spectral variability \\citep[see][]{peterson04}.  Whatever the root cause, at first sight a true Seyfert 2 would be virtually indistinguishable from a Seyfert 2 with a hidden BLR and thus hard to identify without a fair amount of supplementary data.  If the picture that they are related to low accretion-rate objects is correct, they would be expected to be relatively rare and primarily identifiable only among nearby AGN.  Thus far, really convincing true Seyfert 2s have proven elusive. This paper presents an attempt to confirm true Seyfert 2 candidates from the local, spectroscopically-identified Seyfert 2 population by using X-ray measurements to determine if an obscuring medium is present.  This paper is organized as follows: \\S\\ref{sec:sampsel} describes how the \\textit{Chandra} targets were selected; \\S\\ref{sec:datan} describes the observations and analysis of the data, with individual targets discussed in \\S\\ref{sec:ic3}--\\ref{sec:n54}; the results are discussed in \\S\\ref{sec:discus}; a summary is presented in \\S\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:discus}  Although targeted as candidate true Seyfert 2s, or AGNs without BLRs, the four objects studied here --- IC\\,3639, NGC\\,3982, NGC\\,5283, and NGC\\,5427 --- all show evidence of obscuration when observed in X-rays. They are therefore not good ``true'' Seyfert 2 candidates, though of course the presence of obscuration does not imply the existence of a BLR behind the obscuration. The picture that emerges is of moderate to high, but not necessarily Compton-thick, obscuration of the nucleus, $N_H \\sim $ few $\\times 10^{21}$ -- $10^{22}$ cm$^{-2}$, and extended soft emission on hundreds-of-parsecs scales that may possibly originate in the NLR.  The spectrum of IC\\,3639 shows moderate absorbing column densities and possibly a reflection component, with an Fe K$\\alpha$ line detected with low significance. The spectrum of NGC\\,5283 clearly shows a power law which is strongly absorbed below $\\sim 4$ keV. The spectra of both objects require a separate component to fit the soft emission. This component is probably a result of the blending of multiple emission lines from a photoionized gas \\citep{bgc05}; however, the CCD spectra analyzed here are not of high enough quality to constrain a multi-parameter model like APEC or MEKAL, and the soft emission was equally well fit by a thermal bremsstrahlung model and a pure blackbody spectrum. In either case the characteristic temperature turns out to be a few tenths of a keV. That the soft and hard components originate in physically distinct regions is suggested by the hardness ratio maps themselves (Figs.~\\ref{fig:ic3post} and \\ref{fig:n52post}), which show harder emission surrounded by extended softer emission. In the case of NGC\\,5283, which has sufficient counts, this was shown more definitively by fitting the spectra of the nucleus and the surrounding emission separately. The harder, power law component clearly is stronger closer to the nucleus, while the extended emission is dominated by the soft component. The extended, soft component also has a lower obscuring column than the nuclear emission.  The two fainter objects, NGC\\,3982 and NGC\\,5427, while too faint to make spectral fitting possible, appear to be consistent with the above picture. The observed emission is soft: NGC\\,3982 has a hardness ratio of $-0.7$ and NGC\\,5427 does not even have a statistically significant detection in the hard band. The faintness is not due to distance --- these two objects are in fact nearer than the two brighter ones. All four nuclei observed here have comparable [\\ion{O}{3}] luminosities \\citep{w92a,nw95} and thus are expected to have comparable $L_X$ \\citep{hphk05}. These two fainter objects may therefore be even more heavily obscured and we are primarily detecting soft X-ray emission from hot, extra-nuclear gas. Is this gas itself responsible for part of the absorption of the nucleus? While not clear in IC\\,3639, in NGC\\,5283 the gas spectrum is consistent with only Galactic absorption, whereas the nuclear spectrum requires an additional column density of $N_H \\sim 10^{22}$ cm$^{-2}$. In the \\citet{laor03} model, the innermost part of an NLR would become a warm absorber, with the effect getting stronger for lower-luminosity AGN.  The gas emitting soft X-rays has the right physical scale (few hundred parsecs) to be the gas that constitutes the NLR, and \\citet{bgc05} find that in their sample of Seyfert 2s the morphology of the X-ray gas matches that of the NLR as determined by [\\ion{O}{3}] $\\lambda 5007$ imaging. Figures~\\ref{fig:ic3hstxcont} and \\ref{fig:n52hstxcont} show structure maps \\citep{pm02} made from wide-band (F606W) \\textit{HST} images of IC 3639 and NGC 5283, respectively, with X-ray contours overlaid. The NLRs are identifiable as bright regions near the centers of the images, and the contours show that the X-ray emission is extended in the same orientation as the NLRs. NGC\\,3982 and NGC\\,5427, which do not show prominent, extended X-ray emission, also do not show extended NLR emission in structure maps \\citep{pm02}. If the soft X-ray-emitting gas is really the NLR, then the source of ionizing radiation is presumably the central engine, and not a starburst. The soft X-ray (0.5--2 keV) luminosity of IC\\,3639 inferred from the spectral fit, $L_{0.5-2\\mathrm{keV}} \\sim 10^{42}$ erg s$^{-1}$, would imply a star formation rate (SFR) of $\\sim 100$ $M_\\sun$ yr$^{-1}$ \\citep{hhptc05}. This is an order of magnitude higher than the known SFR of IC\\,3639 derived from its infrared luminosity \\cite[$\\sim\\! 9\\, M_\\sun$ yr$^{-1}$,][]{dpkc02} and favors the central engine as the ionizing source.  For NGC\\,5283, $L_{0.5-2\\mathrm{keV}} \\sim 10^{40}$ erg s$^{-1}$, requiring an SFR of only a few $M_\\sun$ yr$^{-1}$, which is consistent with what is known based on its infrared luminosity \\citep{k98,pgre01}, and star formation cannot be ruled out, but the AGN luminosity of $\\sim 10^{42}$ \\es\\ could explain the ionized gas without star formation.  For IC\\,3639, the [\\ion{O}{3}] to soft X-ray flux ratio $F_{\\lambda 5007}/F_{0.5-2\\mathrm{keV}} \\approx 5-6$, while the ratio for NGC\\,5283 is $F_{\\lambda 5007} / F_{0.5-2\\mathrm{keV}} \\approx 0.6$, a range similar to that seen in the \\cite{bgc05} sample. Higher quality spectra of similar, absorbed Seyferts indicate that the soft X-ray emission consists of blended emission lines from a photoionized gas \\citep[e.g.][]{bgc05,lhkwz06}. Thus, morphology, energetics, and spectra all suggest that the NLR is the source of the soft X-ray emission seen in at least some Seyfert 2s.  Though obscuration prevents any deductions about the presence or absence of BLRs in these objects, we apply the \\citet{laor03} model to our objects to determine if the unobscured luminosities are still consistent with his model. This model predicts a minimum bolometric luminosity (as a function of black hole mass) needed to sustain a BLR. We obtained published stellar velocity dispersions from \\citet{w92a}, \\citet{nw95} and \\citet{grea05}, [\\ion{O}{3}] fluxes from \\citet{w92a} and \\citet{nw95}, and used the conversion from \\citet{hea04} to change [\\ion{O}{3}] luminosities to bolometric luminosities. A comparison with Fig.~1 of \\citet{laor03} shows that all four of our nuclei have luminosities ($\\log (L_\\mathrm{bol}/\\mathrm{erg}\\,\\mathrm{s}^{-1}) \\sim 44,43,44,43$ for IC\\,3639, NGC\\,3982, NGC\\,5283, NGC\\,5427, respectively) well above the minimum needed to have a BLR. In the case of NGC\\,5283, the absorption-corrected 2--10 keV luminosity that we obtain from the X-ray spectral fit would be by itself greater than the minimum required. Therefore none of the four is a true Seyfert 2 according to the \\citet{laor03} model.  Observations of these four objects show that it can be misleading if we depend only on the hardness ratio to estimate absorption for sources with a small number of counts. For example, NGC\\,3982 has an HR = $-0.7$, which is the HR expected with ACIS-S for a canonical unabsorbed AGN spectrum. What is being measured, however, is the HR of the NLR, with some hard-band contamination from the partially-obscured nuclear source. This is of special concern in cases where very few counts are detected and the hardness ratio is the only quantity that can be determined with any certainty. If any extended soft emission is not taken into account, the true absorption will be underestimated, and the calculated X-ray luminosity will be too low. Even when spectral fitting is possible, if the spectrum is not of sufficient quality to distinguish the two components, an absorbed AGN may mimic an unobscured AGN of lower luminosity. Several, or all, of the apparently unobscured Seyfert 2s that have been discovered so far \\citep[e.g. ][]{pgsz01,pb02,bcc03,ghkgl03,gz03} are probably examples of exactly this kind of mistaken identity. None of these objects has observations with the angular resolution and high signal-to-noise spectrum necessary to rule out the scenario we propose for our sample. Faint AGN at moderate redshift, such as those being found in the deep X-ray surveys, are particularly susceptible to such misclassification. These sources are extremely important because some intrinsically high-luminosity AGN may be masquerading as apparently low-luminosity Seyfert 1s. These obscured AGN are exactly the sources invoked to explain the spectrum of the cosmic X-ray background \\citep{bh05}. Also affected is the estimation of the AGN luminosity function."
    },
    "0606/astro-ph0606711_arXiv.txt": {
        "abstract": "We present $VI$ photometry of the metal-poor inner halo globular clusters NGC~6293 and NGC~6541 using the planetary camera of the WFPC2 on board {\\em Hubble Space Telescope} (HST).  Our color-magnitude diagrams of the clusters show well-defined blue horizontal branch (BHB) populations, consistent with their low metallicities and old ages. NGC~6293 appears to have blue straggler stars in the cluster's central region. We discuss the interstellar reddening and the distance modulus of NGC~6293 and NGC~6541 and obtain $E(B-V)$ = 0.40 and $(m-M)_0$ = 14.61 for NGC~6293 and $E(B-V)$ = 0.14 and $(m-M)_0$ = 14.19 for NGC~6541. Our results confirm that NGC~6293 and NGC~6541 are clearly located in the Galaxy's central regions ($R_{GC} \\leq$ 3 kpc). We also discuss the differential reddening across NGC~6293. The interstellar reddening value of NGC~6293 appears to vary by $\\Delta E(B-V)$ $\\approx$ 0.02 -- 0.04 mag within our small field of view.  The most notable result of our study is that the inner halo clusters NGC~6293 and NGC~6541 essentially have the same ages as M92, confirming the previous result from the HST NIC3 observations of NGC~6287. ",
        "introduction": "Understanding the formation of our Galaxy has always been one of the key quests in modern astrophysics for decades. During the last decade, a tremendous amount of observational data have accumulated regarding the formation of our Galaxy, in particular due to the advent of HST. However, questions associated with the formation and early evolution of our Galaxy remain unanswered.  One of the earliest models for the formation of the Galaxy was that of Eggen, Lynden-Bell, \\& Sandage (1962), who postulated a rapid, monolithic ``collapse\" of a proto-Galaxy. Using similar data, Isobe (1974) and Saio \\& Yoshii (1979) suggested the formation process could have been much longer than a free-fall timescale, occurring over several billion years. An alternate view of the halo formation process was presented by Searle \\& Zinn (1979). Under the assumption that variations in the numbers of blue and red horizontal branch stars in globular clusters is an age indicator, they argued for somewhat more youthful clusters at larger Galactocentric distances. If true, they argued that the Galactic halo may have formed from accretion of ``proto-Galactic fragments\", implying a more chaotic view of the halo's formation. Carrying the argument further, Zinn (1993) argued that the bulk of the old halo globular clusters formed during a monolithic dissipative collapse while the young halo globular clusters formed independently of the Galaxy and later accreted into our Galaxy. The longer timescale of a monolithic formation process and the fragmentary accretion model both suggest a variation in the ages of globular clusters, and a natural assumption is that the Galactic halo may have formed ``inside-out\", with star formation beginning earlier in center due to the smaller free collapse time scale ($\\tau \\propto \\rho^{-1/2}$).  The discoveries that the Universe's expansion and self-gravity are dominated by dark energy and dark matter have led to the basic ``$\\Lambda$CDM\" model, and the formation of galaxies within this model is a much more elaborate yet tractable variant of the original Searle \\& Zinn (1978) hypothesis, and that of White \\& Rees (1978). Numerous sub-galactic mass concentrations form rapidly in the early Universe, and undergo mergers to form the larger galaxies we see today (see Navarro, Frenk, \\& White 1997). Mergers are still going on in the current epoch, in our own Milky Way Galaxy, and in many others as well. Qualitatively, this ``hierarchical assembly\" model implies an ``inside-out\" formation of our Galaxy, and detailed calculations show this to be the case for the dark matter (Helmi et al.\\ 2003) and the stars as well (Robertson et al.\\ 2005; Font et al.\\ 2006).  What do we know about how when star formation began and how rapidly it proceeded in these fragments and the young Galaxy? Two techniques are available to answer these questions. Element-to-iron (or other elemental) ratios provide valuable clues due to different elements' different nucleosynthesis sites. For example, light elements such as magnesium, calcium, and silicon are manufactured more readily in short-lived stars that explode as Type II supernovae, whereas iron is more readily synthesized in Type Ia supernovae, which are thought to arise from mass transfer or mergers of white dwarfs, a process which is thought to require a longer period of time, perhaps $10^{9}$ years or more. Thus in a closed system of stars and gas, stars with high [$\\alpha$/Fe] abundance ratios were likely born prior to significant contributions from Type Ia supernovae. The discovery that the Galaxy's dwarf spheroidal galaxy neighbors have unusually low [$\\alpha$/Fe] ratios, even at very low [Fe/H] values (Shetrone et al.\\ 2001, 2003; Venn et al.\\ 2004) thereby presented a challenge to the idea that the Galaxy's halo was assembled from proto-Galactic fragments, since, apparently, some of the surviving fragments have not shared the same chemical enrichment history as the Galactic halo. The difficulty with relying on element-to-iron ratios is that they do not monitor ages directly. Small galaxies producing stars very slowly will experience enrichment from Type Ia supernovae before the overall [Fe/H] value has risen significantly. Further, incorporation of supernovae ejecta in subsequent generations of stars depends on the host galaxy's ability to retain gas. Font et al.\\ (2006) have shown that the $\\Lambda$CDM model is able to explain the differences in [$\\alpha$/Fe] vs.\\ [Fe/H] patterns in the local dwarf spheroidal galaxies vs.\\ the Galactic halo due to the different rates of star formation in the systems that were absorbed by the Galaxy compared to the survivors.  The second method of age dating is the most direct: comparisons of main sequence turn-offs with stellar model isochrones. While the derived ages do not depend on the history of nucleosynthesis, the method is difficult to apply with the same level of precision. For example, if the timescale for significant contributions of Type Ia supernovae to manifest their presence in subsequent generations of stars, a change in [$\\alpha$/Fe] can be expected to be revealed in stars whose ages differ by less than $10^{9}$ years {\\em in a closed system}\\footnote{It is even more problematical in ensembles constructed from merger fragments, each that may have experienced a different star formation history.}. Measurement of even relative ages with such precision using main sequence data is very challenging. Nonetheless, much careful work has already been done in this area.  We are interested in the particular question of how rapidly did star formation begin throughout the Galaxy and its then more numerous proto-Galactic fragment. The HST observations of one of the most remote metal-poor halo clusters NGC~2419 ($R_{GC}$ $\\approx$ 90 kpc) of Harris et al. (1997) have shown that NGC~2419 and M92 ($R_{GC}$ = 9.6 kpc) essentially have the same age within $\\pm$2 Gyr, suggesting that globular-cluster formation must have started at everywhere at about the same time in our Galaxy (Richer et al. 1996; see also the counterargument of Chaboyer, Demarque, \\& Sarajedini 1996). Moreover, the local dwarf galaxies Carina (Mighell 1997), Draco (Grillmair et al. 1998), Leo I \\& II (Held et al. 2000; Mighell \\& Rich 1996), and Ursa Minor (Mighell \\& Burke 1999) and the globular clusters in the Fornax dwarf galaxy (Buonanno et al. 1998), LMC (Johnson et al. 1999), and the Sagittarius dwarf galaxy (Montegriffo et al. 1998; Layden \\& Sarajedini 2000) appear to be coeval with typical globular clusters of similar metallicity in our Galaxy. These similar ages suggest that globular cluster formation may have begun everywhere at the same time not only in our Galaxy but also in the local dwarf galaxies, despite the differences in the initial physical environments.  The ``inside-out\" model led van den Bergh (1993) to suggest that the most metal-poor globular cluster  near the Galactic center NGC~6287 ([Fe/H] = $-$2.01, Lee \\& Carney 2002; $R_{GC}$ = 1.6 kpc, Lee et al.\\ 2001) may be the oldest globular cluster in our Galaxy, based on its metallicity, horizontal branch (HB) population, and its spatial location in our Galaxy. The recent HST NIC3 observations of the cluster (Lee et al.\\ 2001) have shown that NGC~6287 and M92 essentially have the same age within $\\pm$2 Gyr, suggesting that the metal-poor globular cluster formation must have been triggered roughly everywhere at the same time in our Galaxy. It should be noted that, however, HST NIC3 photometry can be vulnerable to the variation in the intrapixel sensitivity (Lauer 1999) and the temperature dependence of the interstellar reddening law in the HST NICMOS F110W/F160W photometric system (Lee et al.\\ 2001). Therefore, a photometric study with an expanded sample, preferably employing detector plus filter systems not suffering such effects as can be seen in HST NIC3, is necessary to rank the formation epoch of metal-poor inner halo globular clusters in comparison to the intermediate or the outer halo globular clusters.   NGC~6293 ($\\alpha$ = 17:10:10.42; $\\delta$ = $-26$:34:54.2; $\\ell$ = 357.62; $b$ = +7.83) and NGC~6541 ($\\alpha$ = 18:08:02.20; $\\delta$ = $-43$:42:19.7; $\\ell$ = 349.29; $b$ = $-11.18$) are the second and the third most metal-poor globular clusters within 3 kpc from the Galactic center. NGC~6293 and NGC~6541 are located $\\approx$ 1.4 kpc and 2.2 kpc from the Galactic center and $\\approx$ 1.2 kpc and $-$1.4 kpc from the Galactic plane, suffering interstellar reddening of $E(B-V)$ = 0.41 and 0.14, respectively (Harris 1996). In Table~\\ref{tab1}, we list clusters' properties. Since both clusters are thought to be post core-collapsed, as can be more frequently found in the inner halo region than other parts of our Galaxy (Barbuy, Bica, \\& Ortolani 1998), ground-based photometric study of these two clusters has been limited.  Janes \\& Heasley (1991) presented $BV$ photometry of NGC~6293. Their color-magnitude diagram (CMD) showed that the red-giant branch (RGB) and the BHB morphologies of NGC~6293 are similar to those of M92, indicating that the cluster is both old and metal-poor. By comparing with M92, they suggested that $E(B-V)$ = 0.46 and $(m-M)_0 \\approx$ 16.0. Since their observations reached just to the main-sequence turnoff (MSTO), however, they were not able to address the age of the cluster. They also made an interesting point that the RGB-tip magnitude of NGC~6293 appears to be a full magnitude fainter than that of M92. They suggested that this may reflect that the cluster has undergone core collapse. Trager et al.\\ (1995) found that the cluster has indeed undergone core collapse, as may NGC~6541. More recent work by Noyola \\& Gebhardt (2006), relying on HST WFPC2 data, confirms that both clusters show sharply rising surface brightness levels into the innermost regions.  Alcaino (1971, 1979) obtained $BV$ photographic photometry and Alcaino et al.\\ (1997) presented multi-color CCD photometry of NGC~6541. Alcaino (1979) noted that NGC~6541 has BHB stars and its CMD appears to be similar to that of M13. He also noted that NGC~6541 appears to be deficient in bright RGB stars. With deeper CCD photometry, Alcaino et al.\\ (1997) claimed that the CMD of NGC~6541 is similar to those of M13 and M79, finding $E(B-V)$ = 0.15 and [Fe/H] $\\approx$ $-$1.8 for NGC~6541. They noted a discrepancy in the position of NGC~6541 BHB stars in $U - (B - V)$ and $(U - B) - (B - V)$ diagrams, in the sense that $U$ magnitude of BHB stars in NGC~6541 appear to be $\\approx$ 0.3 -- 0.4 mag brighter than those in M79. Using an isochrone fitting method, they claimed that NGC~6541 is very old, one of the oldest clusters they had studied.  Lee \\& Carney (2002) presented high-resolution echelle spectroscopy of both clusters. They noted that both clusters appear to be silicon enhanced and titanium depleted compared to the intermediate halo clusters. In particular, they suggested that [Si/Ti] ratios appear to be related to Galactocentric distances, in the sense that [Si/Ti] ratios decrease with Galactocentric distance, and they proposed that these [Si/Ti] gradients with Galactocentric distance may be due to the different masses of the SNe II progenitors. The high [Si/Ti] ratios toward the Galactic center may be due to higher-mass SNe II contributions. This is an additional hint that both clusters may have formed particularly early. Lee \\& Carney (2002) also derived the radial velocity of the clusters (see also Table~\\ref{tab1}). Their results suggest that the radial velocities of NGC~6293 and NGC~6541 are small enough that they are most likely genuine inner halo clusters unless their tangential motions are extremely large. Unfortunately, neither cluster has a measured proper motion as yet.  In this paper, we explore the relative ages of inner halo clusters NGC~6293 and NGC~6541 using HST PC2 photometry. As mentioned earlier, both clusters have high central concentrations so that only HST can provide the necessary high angular resolution to study main-sequence photometry of these clusters. In section 2, data acquisition and data reduction are discussed. We present the CMDs of the clusters and discuss the main sequence fittings of the clusters to M92 and NGC~2419 in section 3. The interstellar reddening and the distance modulus are also discussed in this section. Finally, we discuss the formation epoch of the metal-poor inner halo clusters in section 4. ",
        "conclusions": "We have presented HST PC2 photometry for the inner halo globular clusters NGC~6293 and NGC~6541. Our CMD for NGC~6293 shows a strong BHB population, consistent with its low metallicity and old age. It also appears to have blue straggler stars and a future study related to these objects is desirable. We could not investigate the RGB-tip luminosity, which was claimed by Janes \\& Heasley (1991) to be abnormally faint, since the bright RGB stars were saturated in our images. Our magnitudes and colors for the NGC~6541 main-sequence TO are in good agreement with the ground-based $VI$ photometry of Alcaino (1997). Our CMD for NGC~6541 appears to be deficient in bright RGB stars, confirming the finding of Alcaino (1997).  We have discussed the interstellar reddening and the distance modulus of NGC~6293 and NGC~6541 with respect to those of M92. For NGC~6293, our interstellar reddening estimate is consistent with previous results, while our distance modulus is $\\approx$ 0.1 mag smaller than the previous estimate by Harris (1996). We also discussed the differential reddening across NGC~6293. It appears that the interstellar reddening value of NGC~6293 varies by $\\Delta E(B-V)$ $\\approx$ 0.02 -- 0.04 mag. For NGC~6541, our age-dating method, which makes use of M92 as a template, appears to suffer from modest metallicity difference effects. Our interstellar reddening and distance modulus of NGC~6541 are $E(B-V)$ = 0.17 and $(m-M)_0$ = 14.24 without the correction for the metallicity effect, and $E(B-V)$ = 0.14 and $(m-M)_0$ = 14.19 with such a correction. Nevertheless, NGC~6293 and NGC~6541 are clearly located in the Galaxy's central regions ($R_{GC} \\leq$ 3 kpc).  The most interesting result of our study is that the inner halo clusters NGC~6293 and NGC~6541 essentially have ages that are indistinguishably different from one of the oldest globular clusters in our Galaxy M92, consistent with the previous result of NGC~6287 by Lee et al.\\ (2001), and the large study of De~Angeli et al.\\ (2005). Furthermore, they appear to have the same ages as the most remote metal-poor globular cluster NGC~2419 ($R_{GC} \\approx$ 90 kpc). Our results strongly support the idea that the globular cluster formation must have begun everywhere at the same time to within $\\approx$ 0.5 -- 1 Gyr in our Galaxy."
    },
    "0606/astro-ph0606527_arXiv.txt": {
        "abstract": "Using the near-infrared integral-field spectrograph SPIFFI on the VLT, we have studied the spatially-resolved dynamics in the $z=3.2$ strongly lensed galaxy 1E0657-56 ``arc$+$core'' by observing the rest-frame optical emission lines [OIII]$\\lambda$5007 and H$\\beta$. The lensing configuration suggests that the high surface brightness ``core'' is the ${\\cal M}\\sim 20$ magnified central $\\sim 1$ kpc of the galaxy, whereas the fainter ``arc'' is the more strongly magnified peripheral region of the same galaxy at about a half-light radius, which otherwise appears to be a typical $z\\sim 3$ Lyman break galaxy.  The overall shape of the position-velocity diagram resembles the ``rotation curves'' of the inner few kpcs of nearby $\\sim{\\cal L}^*$ spiral galaxies. For ${\\cal M}=20$, our data have a spatial resolution of $\\sim$200 pc in the source plane. The projected velocities $v_{rot,proj}$ rise rapidly to $\\sim 75$ \\kms\\ within radii $\\sim 0.5$ kpc from the center, and asymptotically reach a velocity of $\\sim 190$ \\kms\\ within the arc, at a projected radius of a few kpc radius. The rotation curve implies a dynamical mass of $\\log{M_{dyn}/M_{\\odot}}\\sim 9.3$ within the central kpc, and suggests that in this system the equivalent of the mass of a present-day $\\sim{\\cal L}^*$ bulge at the same radius was already in place by $z\\gtrsim 3$. Approximating the circular velocity of the halo by the measured asymptotic velocity of the rotation curve, we estimate a dark matter halo mass of $\\log{M_{halo}/M_{\\odot}}\\sim 11.7\\pm 0.3$, in good agreement with large-scale clustering studies of Lyman break galaxies. The baryonic collapse fraction is low compared to $z\\sim 2$ actively star-forming ``BX'' and low-redshift galaxies, perhaps implying comparatively less gas infall to small radii or efficient feedback. Even more speculatively, the high central mass density might indicate highly dissipative gas collapse in very early stages of galaxy evolution, in approximate agreement with what is expected for ``inside-out'' galaxy formation models. ",
        "introduction": "Dynamical mass estimates of high redshift galaxies are now starting to play a significant role in our developing understanding of galaxy assembly in the early universe, a trend that will likely become even more important in the near future. Directly measuring the dynamical masses from the spatially-resolved spectra of high-redshift galaxies is observationally very challenging, but dynamical masses are less prone to degeneracies and evolutionary bias than mass estimates based solely on photometry. Fascinatingly, they may also allow us to directly probe the baryonic and dark matter content and concentration of galaxies in the early universe, and to measure their angular momenta \\citep{NMFS06}. Ultimately, they will enable us, for example, to directly compare the growth of galaxy mass and angular momentum with model predictions as a function of redshift. Accurately measuring the kinematics of high-redshift galaxies is therefore a major step forward.  To realize these goals, we must show convincingly that the kinematics we measure in a high redshift galaxy have a simple proportionality to the mass distribution and rule out that they are dominated by the orbit or angular momentum loss of mergers or by hydrodynamical processes like, e.g., starburst-driven ``superwinds'' \\citep{lehnert96a}.  The Lyman-break technique has led to the largest sample of spectroscopically confirmed galaxies from z$\\sim$2.7 to 6.4. Despite our rapidly growing understanding of their ensemble properties, such as their luminosity function, clustering, and star-formation history \\citep[e.g.,][]{steidel96, adelberger98}, our knowledge of their detailed intrinsic properties remains rather rudimentary.  LBGs at z$\\sim$3 have typical radii of $r_{e} \\sim 0.3\\arcsec$ \\citep{giavalisco96} so that the spatially resolved kinematics are often difficult to obtain. Thus, dynamical mass estimates for individual LBGs at z$\\sim$3 \\citep[e.g.,][]{pettini01} are based mostly on line widths and only in a handful of cases have velocity gradients been observed. Spatially resolved LBGs are large compared to the overall population, and might be biased towards the strongest line emitting galaxies (e.g., vigorous starbursts) or early-stage mergers and perhaps are not representative for the overall population.  The only way to properly address these issues is to resolve the dynamics of an LBG on fine scale.  Strongly gravitationally lensed LBGs are a promising way to probe small physical scales even with seeing-limited data. The ideal target would be a strongly-lensed, highly inclined LBG, where the kinematic major axis is roughly along a caustic.  Such a configuration would allow several patches of the same galaxy, but at different radii, to be highly magnified to include the intrinsically low surface-brightness periphery. Probing non-lensed LBGs in this way is impossible given their generally small radii, faint magnitudes, and low surface brightnesses.  Unfortunately, strongly lensed LBGs with a favorable lensing geometry are exceedingly rare.  Two cases at z$>$1 have been studied in detail so far, MS1512-cB58 at z=2.8 \\citep[e.g.,][]{teplitz00, pettini02, baker04} and AC114-S2 at z=1.9 \\citep{L-B03}. However, the underlying dynamical mechanism is not conclusively revealed in either case. MS1512-cB58 is magnified by a factor $\\sim 30$, but it is compact and has no apparent velocity gradient \\citep{teplitz00} -- most likely because of an unfavourable lensing geometry. From the mm CO emission line width, \\citet{baker04} measure $M_{dyn} \\sim 10^{10}$ \\msun, not corrected for inclination. AC114-S2 has a velocity gradient \\citep{L-B03}, but is a merger with complex morphology, and its nature is not well constrained. That spatially-resolved spectroscopy of giant arcs can provide valuable constraints on the internal dynamics of the lensed galaxies has recently been shown by \\citet{swinbank06} for a sample of 6 giant arcs at lower redshift, z$\\sim$1. \\citeauthor{swinbank06} found regular kinematics in 4 of the 6 galaxies, consistent with quiescently rotating disks, while in 2 galaxies, they observed complex line profiles of varying widths and irregular velocity structure suggestive of either mergers or outflows.  The strongly lensed z=3.24\\footnote{We adopt a flat concordance cosmology with $\\Omega_{\\Lambda}$=0.7 and H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, in which D$_L$=27.9 Gpc and D$_A$=1.5 Gpc at $z=3.24$. The size scale is 7.5 kpc/\\arcsec. The age of the universe at this redshift and cosmological model is 1.9~Gyrs.} ``arc$+$core'' galaxy behind the z=0.3 X-ray cluster 1E0657-56 \\citep{tucker98} appears to be different from these well-studied high-redshift gravitational arcs. Its lensing configuration suggests the simultaneous magnification of a high surface brightness region at the south-eastern tip of the source that may be associated with the ``core'' of the galaxy as well as a more highly magnified, lower surface brightness region outside the core \\citep[``arc'';][]{mehlert01}.  The total extent of the arc is $\\sim$14\\arcsec, and it has a complex substructure: \\citet{mehlert01} identify 3 faint knots of similar surface brightness within the arc, each separated by a few arcseconds. They propose that the central highest surface brightness region of the lensed galaxy, lying near, but outside the cusp-caustic, is seen as the bright core, whereas a fainter outer region on one side of the same galaxy, which touches the cusp-caustic and is split into three merging images, constitutes the full extent of the arc.  Thus the asymmetric magnified image comprising the near-nuclear region and peripheral patches originating on one side of the galaxy.  The high magnification (${\\cal M}\\ga20$) presents an excellent opportunity to investigate the properties of a z$\\sim$3 galaxy at different radii with high physical and spatial resolution.  The paper is organized as follows: After presenting observations and data reduction in \\S\\ref{sec:obsred}, we turn in \\S\\ref{sec:acs} to the rest-frame UV and optical properties, highlighting the high-resolution ACS morphology. We discuss the spatially-resolved rest-frame emission line kinematics extracted from three--dimensional data cubes obtained with the integral-field spectrograph SPIFFI in \\S\\ref{sec:results}. This includes a detailed discussion of the 146 km s$^{-1}$ velocity gradient in the core and its continuation in the arc. In \\S\\ref{sec:spiral}, we present a detailed comparison with the internal dynamics of low-redshift spiral galaxies. In \\S\\ref{sec:evolution}, we estimate the evolution and angular momentum of the arc$+$core, before investigating the halo mass and the baryonic ``collapse fraction'' in \\S\\ref{sec:dmhalo}. We summarize our results in \\S\\ref{sec:summary} and draw a likely evolutionary scenario for LBGs. ",
        "conclusions": "\\label{sec:summary}  We presented an analysis of the strongly lensed (${\\cal M} = 20$) Lyman break galaxy 1E0657-56 arc$+$core galaxy at redshift $z=3.24$, based on SPIFFI integral-field rest-frame optical spectroscopy, complemented with rest-frame UV imaging and spectroscopy. This galaxy is an excellent target for studying the fine spatial details of a $z\\sim 3$ Lyman break galaxy. We extracted the rest-frame UV colors of the arc$+$core from the deep FORS spectroscopy of \\citet{mehlert01}, and measured directly that the galaxy fulfills the Lyman break criterion, including the $R=25.5$ mag limit, imposed on spectroscopically identified sources, for an unlensed source. The arc$+$core is near the peak of the LBG redshift distribution, and its unmagnified size, optical emission line properties, mass-to-light ratio, and stellar age are within the range estimated for the overall population. Therefore, it is particularly well suited for a detailed analysis of its small-scale properties. We find a slightly lower star-formation rate than average, most likely due to our uncertain extinction estimate, likely underlying \\hb\\ absorption, or missing flux by not accounting for unlensed or multiply lensed regions of the galaxy.  Through studying magnified high surface brightness regions of an LBG at $z\\sim 3$, we can investigate the structure and nature of LBGs at high physical resolution. The dynamical mass within 500 pc is about $2\\times 10^9$ \\msun, while at about 2 kpc radius \\citep[approximately the half-light radius of a typical $z\\sim 3$ LBG, e.g.,][]{giavalisco96}, the mass is similar to the average stellar mass $\\sim 10^{10}$ \\msun of LBGs. Stellar masses derived from SED modelling include light at larger radii and make assumptions about the initial mass function that may be unwarranted and generally lead to higher masses \\citep{NMFS06}, so any discrepancy is not totally unexpected. However, our estimated mass of the core is also typical for the bulges of local $\\sim {\\cal L}^{\\star}$ spiral galaxies. In addition, the arc$+$core has a specific angular momentum similar to that of local spiral galaxies.  The combination of mass surface density, metallicity and the dynamical time perhaps suggests an interesting evolutionary picture for the 1E0657-56 arc$+$core. Since the properties of the 1E0657-56 arc$+$core are well within the typical range of LBGs, this outline of the evolution of the 1E0657-56 arc$+$core might, with some caution, be applicable to $z\\sim 3$ LBGs generally.  Compared to local spirals, most of the mass within 500 pc for the 1E0657-56 core appears to be already in place. However, the metallicity in the nuclei of low-redshift spirals is approximately solar, while we have found that the 1E0657-56 arc$+$core has at most about half solar gas-phase abundances, and this estimate is clearly dominated by the emission from the core. We do not know the gas fraction of the 1E0657-56 arc$+$core, however, CO observations of the lensed z$=$2.7 LBG MS1512-cB58 by \\citet{baker04} suggest that LBGs might be gas rich, with gas fractions of possibly up to 50\\%. For a simple closed box model with such high gas fractions, the metallicity will double within several 10 to 100~Myrs. This time estimate likely increases by factors of a few if including outflows or inflows with the metallicity of the intergalactic medium at $z\\sim 3$. Our orbital time estimate suggests that the intense star-formation is likely to last long enough to increase the metallicity to about solar.  Because $z \\sim 3$ LBGs have similar co-moving densities as local luminous (${\\cal L}\\gtrsim {\\cal L}^{\\star}$) galaxies, \\cite{steidel96} suggested that they represent the formation of the spheroidal component of massive galaxies. Our analysis suggests that this hypothesis is plausible since we measure a mass and mass surface density similar to local $\\sim{\\cal L}^{\\star}$ spiral galaxies, and also fulfill the metallicity constraint, after allowing for further evolution in the on-going episode of intense star-formation. Moreover, the low baryon collapse fraction within $\\sim r_e$ might hint that a substantial amount of gas resides on larger scales within the halo (maybe gas blown out during intense star-formation or pre-enriched material from the IGM).  However, we also find $v_c/\\sigma\\lesssim 3$, significantly larger than the $v_c/\\sigma \\sim 0-1.2$ of bulges \\citep[][and references therein]{kormendy04}. Thus, if the arc$+$core is representative of the overall population, LBGs will have to lose factors of a few in their circular velocities to have $v_c/\\sigma$ ratios consistent with local bulges, and certainly substantially more angular momentum to evolve into massive ellipticals.  Currently, very few models address the evolution of individual disk galaxies within the context of the hierarchical model in detail, which makes a quantitative comparison rather difficult. Overall, models of the formation of large scale structure and the evolution of galaxies within a $\\Lambda$CDM cosmology favour ``inside-out'' galaxy evolution, where the inner regions of galaxies form earlier than the peripheries \\citep[e.g.,][]{samland03, abadi03}. Such a scenario quite naturally explains observations at low redshift, such as metallicity and stellar population (age) gradients observed in local galaxies.  Although these models produce inner regions of galaxies that collapse relatively early, unfortunately they also predict that only a relatively small amount of mass will be in place by $z\\sim 3$ compared to the final mass of the galaxy. Most of the mass at small radii is acquired rather late, more likely around redshifts of-order $z=1$ \\citep{samland03}. However, as emphasized by \\citet{immeli04}, the timing and spatial distribution of the star-formation ``history'' depends crucially on the infall history and on how efficiently the kinetic energy gained from dynamical and mechanical heating during collapse is dissipated. They show that the gas in galaxies with large dissipation efficiency will strongly fragment, and interactions between individual subclumps and dynamical friction will make the fragments coalesce to the central regions more rapidly.  In other words, the efficiency of dissipation and fragmentation may essentially be a free parameter which could be constrained observationally.  In comparison to local ${\\cal L}^*$ spiral galaxies, we find in the arc$+$core at $z=3.2$ a significant and comparable mass surface density, while the overall relative mass is rather low.  In light of the models already discussed, this might indicate highly dissipative gas collapse during the earliest phases of galaxy evolution.  The later evolution and perhaps the formation of the disk might either be driven by infall of material \\citep[e.g.,][]{samland03} or by the merger of gas rich galaxies supported by strong feedback to prevent the baryons from collapsing into the central regions \\citep[e.g.,][]{robertson05}.  The latter of these hypotheses might explain the apparent inefficiency of the baryon collapse of the z$\\sim$3 LBGs compared to galaxies at lower redshift."
    },
    "0606/astro-ph0606241_arXiv.txt": {
        "abstract": "s{ Galaxy clusters furnish extremely rich information on the contents and structure of our universe.  The potential of galaxy cluster studies to constrain dark energy, for example, motivates a number of ambitious cluster surveys.  Among these, surveys based on the Sunyaev-Zel'dovich (SZ) effect are particularly powerful for their ability to cleanly select clusters out to redshifts $z>1$.  Now poised to begin surveying substantial areas of sky, dedicated interferometers, bolometer cameras and the Planck satellite will soon produce large cluster catalogs that will provide a precise measure of the cosmic expansion rate over a range of redshifts and precipitate a new understanding of structure and galaxy formation.  I review the science potential of these surveys and examine some issues of SZ cluster catalog construction. } ",
        "introduction": "Detailed observations of the cosmic microwave background anisotropies \\cite{cmb_gen}, distant SNIa \\cite{snIa_gen} and of the galaxy distribution \\cite{lss_gen} have driven the tremendous advance in recent years leading to the development of a standard cosmological model. Ongoing research aims to test the model's coherence and to answer outstanding fundamental questions: What is the nature of dark matter and of the mysterious dark energy accelerating the present expansion of the universe? What is the physics of the early universe, in particular of the inflation epoch? How can we use detailed cosmological observations to probe fundamental physics, such as the neutrino sector? And how do galaxies form and evolve in the cosmic web of large--scale structure.  These exciting questions inspire large observational programs centered on the CMB, SNIa searches and wide--field multiband surveys.  Galaxy clusters offer a unique avenue of attack on several of these key questions: \\begin{enumerate} \\item Their abundance and evolution with redshift are highly sensitive to the statistics and growth rate of the density perturbations, which in turn depends on the cosmological parameters; \\item They are ideal tracers of the largest scale structures.  With a mean separation of $\\sim 50$~Mpc, they efficiently sample structures of wavelength $\\sim 100$~Mpc and larger, such as the baryon acoustic peaks in the matter power spectrum (see Figure~\\ref{fig:Veff}) \\cite{bao}; \\item They provide a well--defined, quasi--closed environment for galaxy formation studies.  In the cluster environment we directly observe the stellar, diffuse gas and dark matter components of the cosmic fluid; \\item Combined X--ray and millimeter cluster observations (the Sunyaev-Zel'dovich effect; see below) permit distance measurements and hence the construction of a Hubble diagram with which to measure the matter density and dark energy abundance and equation--of--state (like SNIa distance measurements). \\end{enumerate} ",
        "conclusions": "Future galaxy cluster surveys will provide a wealth of information on dark energy, dark matter and structure and galaxy formation.  Among these surveys, those based on the SZ signal will profit from its intrinsic ability to find clusters at high redshift and its expected tight correlation to halo mass.  Over the next 5 years, these surveys will provide large, well--defined (in terms of mass) catalogs containing hundreds to thousands of clusters at redshifts beyond unity or, in other words, multiplying by 10--100 the number of known clusters at these redshifts\\footnote{The new WMAP3 results were published at the time of writing.  The lower value of $\\sigma_8$ favored by the new release can lower the predicted cluster counts by up to a factor $\\sim 2$. We are evaluating the changes for each experiment in detail.}.  This will give us a new view of galaxy formation in dense environments and a measure of the expansion rate at this crucial epoch marking the transition between matter and dark energy domination.  Many of the surveys will begin this year (2006), and a number of important surveying issues require further study.  These include accurate evaluation of survey selection functions and observational errors.  I have shown two results from our studies based on a matched filter detection algorithm and simulations of different surveys.  We find that high resolution ground--based surveys select clusters not simply on flux, but on a combination of flux and angular size, and that this must be properly accounted for when modeling cluster counts.  We also find that observational scatter on measured cluster flux $Y$ (e.g., $\\sim 40\\%$ in the case of SPT) largely exceeds the intrinsic scatter predicted by numerical simulations~\\cite{motl_etal05}.  Furthermore, confusion with primary CMB anisotropy severely compromises photometry in single frequency observations; follow--up at a second frequency (or in X--rays) will therefore be necessary for these surveys."
    },
    "0606/astro-ph0606288_arXiv.txt": {
        "abstract": "To detect Earth-like planets in the visible with a coronagraphic telescope, two major noise sources have to be overcome: the photon noise of the diffracted star light, and the speckle noise due to the star light scattered by instrumental defects. Coronagraphs tackle only the photon noise contribution. In order to decrease the speckle noise below the planet level, an active control of the wave front is required. We have developed analytical methods to measure and correct the speckle noise behind a coronagraph with a deformable mirror. In this paper, we summarize these methods, present numerical simulations, and discuss preliminary experimental results obtained with the High-Contrast Imaging Testbed at NASA's Jet Propulsion Laboratory. ",
        "introduction": "\\label{sec:intro}  In coronagraphs like any other optical system, aberrations give rise to scattered light in the form of speckles in the focal plane, thus preventing very high-contrast imaging. A wave front sensing and control system can remedy this, and make the detection of faint objects possible\\cite{Malbet95}. Such a system is absolutely critical to the success of NASA's Terrestrial Planet Finder Coronagraph (TPF-C)\\cite{Coulter04} that aims not only at the detection, but also at the spectroscopy of terrestrial planets in the visible range. This is a very challenging task as the contrast\\footnote{The contrast is defined here as ratio of the luminosity of the star to that of the planet at the wavelength of interest.} between a terrestrial planet and its Sun-like star amounts to $\\sim 10^{10}$ in the visible. With that goal in mind, we have devised a theory\\cite{Borde06} to measure and correct wave front aberrations, so as to clear out of speckles an area of the focal plane, referred to as dark hole (DH). This area is then suitable for very faint object detection.  In the following, we first summarized the theory (\\S\\ref{sec:theory}), then we present simulations (\\S\\ref{sec:simulations}), and finally we discuss preliminary experimental results (\\S\\ref{sec:experiments}). ",
        "conclusions": "\\label{sec:conclusion} Although our first experimental results are modest, they prove that our approach is sound and should be pursued. Future work will include an experimental demonstration of the energy minimization approach in monochromatic light, as well as a generalization of the theory to polychromatic light with corresponding experiments. Our wave front sensing and control theory is general and can be applied to other types of coronagraph, as demonstrated in Ref.~\\citenum{Giveon06} for shaped-pupil coronagraphs.  Note that these experiments are by no means representative of the best result achieved to date with the HCIT. For these, we refer the reader to Ref.~\\citenum{Trauger04} and to the paper by Trauger et al. in these proceedings."
    },
    "0606/astro-ph0606594_arXiv.txt": {
        "abstract": "\\noindent We consider the excitation of radial and non-radial oscillations in low-mass B stars by the iron-bump opacity mechanism. The results are significant for the interpretation of pulsations in subdwarf B stars, helium-rich subdwarfs and extreme helium stars, including the EC14026 and PG1716 variables.  We demonstrate that, for radial oscillations, the driving mechanism becomes effective by increasing the contrast between the iron-bump opacity and the opacity from other sources. This can be achieved either by increasing the iron abundance or by decreasing the hydrogen abundance. The location of the iron-bump instability boundary is found to depend on the mean molecular weight in the envelope and also on the radial order of the oscillation.  A bluer instability boundary is provided by increasing the iron abundance alone, rather than the entire metal component, and is required to explain the observed EC14026 variables. A bluer instability boundary is also provided by higher radial order oscillations. Using data for observed and theoretical period ranges, we show that the coolest EC14026 variables may vary in the fundamental radial mode, but the hottest variables {\\it must} vary in modes of higher radial order.  In considering non-radial oscillations, we demonstrate that g-modes of high radial order and low spherical degree ($l<4$) may be excited in some blue horizontal branch stars with near-normal composition ($Z=0.02$). Additional iron enhancement extends the g-mode instability zone to higher effective temperatures and also creates a p-mode instability zone. The latter is essentially contiguous with the radial instability zone. With sufficient iron, the p-mode and g-mode instability zones overlap, allowing a small region where the EC14026 and PG1716-type variability can be excited simultaneously. The overlap zone is principally a function of effective temperature and only weakly a function of luminosity. However its location is roughly 5000{\\rm K} too low compared with the observed boundary between EC14026 and PG1716 variables. The discrepancy cannot be resolved by simply increasing the iron abundance. ",
        "introduction": "% \\label{intro}  The discovery of a significant contribution to stellar opacity from iron-group elements at temperatures around 200\\,000\\,K \\citep{OPAL92, OP95} had major consequences for studies of pulsation in stars. First, it solved the long-standing Cepheid mass problem \\citep{Mos92,Kan94}. Second, it provided a natural explanation for hitherto unexplained stellar variability, notably with respect to the $\\beta$ Cepheids \\citep{Dzi93} and some extreme helium stars \\citep{Sai93}. Third, it initiated searches for variability (predicted and observed) in stars hitherto thought to be stable. There are now known to be several classes of variable in which pulsations are driven by the opacity (or $\\kappa-$) mechanism excited by the Fe- or Z- opacity bump. It is one goal of this paper to explore some of the connections between these groups and hence to understand better the nature of Fe-bump driven pulsations.  \\begin{figure} \\begin{center} \\epsfig{file=sdbv_obs.ps,width=8cm,angle=0} \\caption[Fe-bump variables] {The loci of low-mass variable stars believed to pulsate due to Fe-bump instability, including EC14026 variables (upright triangles), PG1716 variables (inverted triangles) and helium-rich variables (squares). A selection are labeled. HS0702+6843 and Balloon 090100001 exhibit both EC14026 and PG1716-type behaviour \\citep{Sch06,Ore05}. Straight diagonal lines represent $\\log g=6.0$ (below) and 4.0 (above). These features are reproduced in Figs.~\\ref{Z-stable} to \\ref{Fe-modes}. } \\label{sdbv_obs} \\end{center} \\end{figure}  An early success was the explanation of radial pulsations in the extreme helium star V652\\,Her \\citep{Sai93}, together with the prediction \\citep{Sai94} and subsequent discovery of radial pulsations in a second helium star, BX\\,Cir \\citep{Kil95}. In these stars, both with periods of $\\sim 9\\,300$s and with masses 0.59 and 0.47 \\Msolar\\ respectively \\citep{Jef01,Woo02} (Fig.~\\ref{sdbv_obs}), the cause of pulsation instability is clearly the Fe- opacity bump, exaggerated by a reduction of the background opacity due to hydrogen at these temperatures. The failure to detect any variability in a third helium star, HD144941, of similar temperature and luminosity but very low metallicity \\citep{Jef96,Jef97}, confirmed the r\\^ole of the Fe-opacity in driving pulsations in these low-mass early-type stars.  Another notable success has been the prediction and discovery of short-period pulsations in subluminous B (sdB) stars \\citep{Cha96, Cha97, Kil97, Bil97}. Briefly, multi-periodic oscillations with periods between 90 and 600 seconds are observed in approximately 10\\% of hot subdwarf B stars; such stars are variously known as EC14026 variables (after the class prototype EC\\,14026--2647) or sdBVs. These pulsations are successfully explained if the stars are ``extreme horizontal branch stars'', that is they consist of a helium-burning core of some 0.5\\Msolar\\ overlaid by a very thin layer of hydrogen. Because of the high effective temperature and the high surface gravity, theoretical models have shown that  radiative forces on the ions in the stellar envelope act differentially such that substantial chemical gradients are established over a diffusion time scale $\\sim 10^5$y \\citep{Mic89}. The consequent levitation and accumulation of iron in layers at around  200\\,000\\,K \\citep{Cha95} enhances the Fe-opacity bump sufficiently that radial and non-radial p-mode oscillations are excited. The theory successfully explained the observed distribution  (Fig.~\\ref{sdbv_obs}) of pulsating and non-pulsating sdB stars in effective temperature and surface gravity \\citep{Cha01}. Although consequential more on frequency eigenvalues than stability criteria, the theory is also  sufficiently well developed that it has been possible to compare predicted and observed pulsation frequencies in some pulsating sdB stars \\citep{Bra01, Cha05a, Cha05b}.  The unexpected discovery of oscillations with periods of a few hours in many sdB stars lying red-ward of the EC14026 instability domain presented a new challenge \\citep{Gre03}. While radiative levitation of iron is still operative in these stars, sometimes known as PG1716 variables (after the prototype PG\\,1716+426), p-modes are reported to be stable in the chemically stratified models \\citep{Cha01}. On the other hand, non-radial g-modes of high radial order ($k\\geq10$) and high spherical degree ($l\\geq3$) were found to be unstable, but only in models of stars cooler than those in which variability had been detected \\citep{Fon03}. While the observed periods imply a g-mode origin, modes of such high degree are not  generally thought  to be observable as variations in total flux due to geometric cancellation. The challenge is therefore to shift the g-mode blue edge to higher effective temperatures and to excite modes of lower spherical degree.  A further challenge has been provided by the detection of variability in the helium-rich sdB star LS\\,IV$-14^{\\circ}116$ with periods of $\\sim1800 - 3000$s \\citep{Ahm05}. While the canonical sdB stars mostly have surface helium abundances one tenth of the solar value ($\\sim10\\%$ by number), the much rarer He-rich sdB stars comprise a spectroscopically distinct and highly inhomogeneous group with surface helium abundances ranging from some 30\\% to nearly 100\\% by number, and surface gravities over a much broader range than seen in normal sdB stars \\citep{Ahm03}. The natural response to the discovery of pulsations in LS\\,IV$-14^{\\circ}116$ was to try to explain them in terms of Fe-bump instability in either a helium star like V652\\,Her, or in some sort of mutant sdB star. With $T_{\\rm eff}$ similar to the EC14026 stars, but with a lower surface gravity (higher $L/M$ ratio), the pulsation periods were too long to be explained by Fe-bump driven p-modes \\citep{Ahm05}. One question is whether g-modes might be excited in such a star. This  challenge is important  because a viable picture of the star that explains both the $L/M$ ratio and the high surface abundance of helium has yet to be established. Any model that successfully reproduced the observed oscillations would assist this process.  In order to address these questions, we need to develop our understanding of instability in low-mass stars. By gaining a general insight into what affects the type of pulsation that can be driven, it should be possible to determine what model properties need to be modified in order to reproduce observed oscillations.  The important theme is clearly Fe-bump instability. We have already investigated the r\\^ole of envelope hydrogen abundance in the excitation of Fe-bump pulsations \\citep{Jef98}. We recall that pulsational instability exists in all stars of sufficiently high $L/M$ ratio due to the presence of strange-mode instability \\citep{Gau90}. Fe-bump instability sets in for lower $L/M$ stars with $T_{\\rm eff}\\sim 20 - 30\\,000$K when either the iron-abundance is raised sufficiently or the hydrogen abundance is reduced sufficiently. The second occurs because reducing the hydrogen opacity increases the contrast between the iron opacity and the background opacity due to other sources, effectively increasing the opacity gradients $\\kappa_T = (\\partial\\ln\\kappa/\\partial\\ln T)_\\rho, \\kappa_{\\rho}=(\\partial\\ln\\kappa/\\partial\\ln\\rho)_T$ in the driving zone.  Therefore there should be a physical connection between, for example, the radial pulsations seen in the extreme helium star V652\\,Her (essentially normal iron, reduced hydrogen) and the p-mode oscillations in EC14026 variables (normal hydrogen, enhanced iron). Note that radial modes are special cases of p-modes with spherical degree $l=0$; such modes are generally present in EC14026 stars alongside modes of higher spherical degree. However, the EC14026 instability region reported by \\citet{Cha01} lies considerably blue-ward of the Fe-bump instability finger described by \\citet{Jef98}. Understanding this paradox represents an important step towards understanding Fe-bump instability in general.  Our initial goal is therefore to explore Fe-bump instability for radial modes in comparatively simple (i.e. homogeneous) stellar envelopes as a function of composition (section II). With success here, it becomes possible to restrict the volume of model space required to explore the stability of modes of higher degree, in particular g-modes (section III), and also to identify what is required in more detailed models in order to reproduce the observations accurately. ",
        "conclusions": "% \\label{conclusion}  Our original goals were to understand oscillations observed in the helium-rich subdwarf LS\\,IV$-14^{\\circ}116$, and in the PG1716 variables. We therefore carried out a broad review of Fe-bump driven pulsational instability in low-mass stars\\footnote{Oscillations in main-sequence stars were deliberately excluded from our study.}.  By first considering radial modes, we have demonstrated the essential r\\^ole of chemical composition such that instability increases with the contrast between the iron-bump opacity and other opacity sources. This increased contrast may be achieved either by increasing the iron abundance, confirming earlier work by \\citet{Cha01} or by reducing the hydrogen abundance \\citep{Jef98}. At least one of these is necessary to excite oscillations in all of the {\\it low-mass} Fe-bump pulsators discovered to date.  We have further demonstrated that the blue-edge for radial instability is affected by the mean molecular weight in the stellar envelope, so that increasing the iron abundance alone provides a bluer instability region than increasing the iron abundance in concert with all elements heavier than helium. The former is required to explain the locus of the EC14026 instability region, making the general assumption that the properties of non-radial p-modes modes are closely linked to the corresponding radial mode of the same radial order.  Furthermore, the blue-edge also depends on the radial order of the oscillations, so that higher-order modes may be found in hotter stars. By comparing theoretical and observed periods for EC14026 stars, we have shown that low-order or fundamental mode radial oscillations are only likely to be seen in the coolest EC14026 stars and that the hottest stars must oscillate in higher-order modes.  However, we have been unable to explain the oscillations in the helium-rich star  LS\\,IV$-14^{\\circ}116$; the observed periods are simply too long compared with the $L/M$ ratio obtained from spectroscopy. A lower surface gravity would definitely help.  Considering non-radial pulsations, we have focused on g-mode instability and in particular on instability in modes with spherical degree $l<4$. We have discovered a small g-mode instability island on the blue horizontal branch which does not require iron enhancement or hydrogen depletion. However, since it requires $Z=0.02$ and a partial reflection of the wave at the interface between the hydrogen-rich envelope and the helium core, it will take some observational effort to verify whether there are any real pulsating horizontal-branch stars which correspond with these models. Effective temperatures are expected to be between 16\\,000 and 20\\,000 K, and pulsation periods between 1 and 4 hours.  With a factor of 10 enhancement of iron on a background metallicity $Z=0.02$,  a large instability region develops, even for $l=1$, which extends from $<13\\,000$ to $\\sim 24\\,000$ K ($l=3$). With such iron enhancement there is always an overlap between p- and g-mode instability regions so that stars near the boundary would be expected to exhibit both modes simultaneously. Depleting the envelope hydrogen abundance tends to shift both the g-mode blue edge and the radial/p-mode red edge to lower temperatures.  Many of these results have been demonstrated individually \\citep{Sai93,Jef98,Cha96,Fon03}; this investigation has placed them in a more general context, as well as delivering some new results. Of these, the r\\^ole that mean molecular weight plays in determining the instability boundaries for radial and p-mode oscillations has implications for understanding the oscillations in EC14026 variables. We have made no effort to justify the adopted abundances on physical grounds, we have simply sought parametric solutions which satisfy the observations. Radiative levitation \\citep{Cha95} is known to operate in extreme horizontal-branch stars and provides a natural explanation for the required iron enhancements. The fact that it operates selectively accords well with our deduction that only iron should be enhanced.  Our discovery that g-mode oscillations may be excited in blue horizontal branch stars with envelopes of ``normal'' composition has two-fold consequences. In addition to the observational question already posed, it may be less hard to excite g-mode pulsations in PG1716 stars than hitherto supposed; maybe we don't have to work so hard to find the necessary chemical structure in the stellar envelope as we do in the case of EC14026 variables.  Nonetheless, we have been unable to resolve the problem of why the observed boundary between the EC14026 and PG1716 stars occurs at $\\sim 29\\,000$ K. The theoretical p-mode/g-mode boundary remains persistently at $\\sim 24\\,000$ K over a wide range of envelope compositions. It is possible that a sharp discontinuity in composition immediately below the iron bump could help; this and other experiments with stratified envelopes remain to be investigated."
    },
    "0606/astro-ph0606077_arXiv.txt": {
        "abstract": "Estimates of the bulk metal abundance of the Sun derived from the latest generation of model atmospheres are significantly lower than the earlier standard values.  In Paper I we demonstrated that a low solar metallicity is inconsistent with helioseismology if the quoted errors in the atmospheres models (of order 0.05 dex) are correct.  In this paper we undertake a critical analysis of the solar metallicity and its uncertainty from a model atmospheres perspective, focusing on CNO.  We argue that the non-LTE corrections for abundances derived from atomic features are overestimated in the recent abundance studies, while systematic errors in the absolute abundances are underestimated.  If we adopt the internal consistency between different indicators as a measure of goodness of fit, we obtain intermediate abundances [C/H] = 8.44 +/- 0.06, [N/H] = 7.96 +/- 0.10 and [O/H] = 8.75 +/- 0.08.  The errors are too large to conclude that there is a solar abundance problem, and permit both the high and low scales.  However, the center-to-limb continuum flux variations predicted in the simulations appear to be inconsistent with solar data, which would favor the traditional thermal structure and lead to high CNO abundances of (8.52, 7.96, 8.80) close to the seismic scale.  We argue that further empirical tests of non-LTE corrections and the thermal structure are required for precise absolute abundances.  The implications for beryllium depletion and possible sources of error in the numerical simulations are discussed. ",
        "introduction": "The uncertainty in the absolute chemical composition of stars is the limiting factor in our ability to do high precision stellar astrophysics.  Traditionally, we have had to rely on a small database of fundamental stellar parameters such as mass, distance, and radius.  However, current and upcoming space missions promise a wealth of astrometric and photometric data.  Large surveys undertaken primarily for other purposes (microlensing, planet searches and cosmology) have discovered thousands of eclipsing binaries, yielding numerous precise mass estimates.  The rapidly developing field of optical interferometry has also permitted a growing number of direct radius estimates.  Asteroseismology is also growing in importance, and missions such as COROT promise a wealth of detailed information on the pulsational properties of solar-like stars.  Our stellar interiors models have become highly sophisticated and successful when compared with observational diagnostics.  In particular, the resolution of the solar neutrino problem in favor of the solar model predictions and the agreement between theoretical predictions and helioseismic data are both encouraging signs.  The combination of better observations and theory has opened the prospect of a new era of precision stellar astrophysics, which could have broad consequences for diverse subfields of astronomy.  Stellar atmospheres theory has traditionally employed a series of approximations when deriving abundances.  Classical models assume an ad hoc turbulent velocity field adjusted to yield abundances independent of excitation potential and line strength.  Convection is usually treated in an approximate fashion, with the mixing length theory.  Horizontal temperature fluctuations (granulation) are not included.  The models also typically assume that the molecular and atomic levels are described by local thermodynamic equilibrium (LTE), e.g. by the local temperature alone.  The compilations of solar abundances used for theoretical solar models \\citep{ag1989,gn1993,gs1998} employed model atmospheres with approximations at the level described above.  The mean thermal structure employed was semi-empirical \\citep[hereafter HM74]{hm1974}.  When these approximations are relaxed, different conclusions about the abundances are obtained.  Departures from LTE are expected at a modest level for solar conditions, and have been investigated by a number of authors \\citep[for example][]{c1986,sh1990,k1993,stb2001,w2001}.  Numerical simulations of convection have matured to the level where they can be used to predict velocity fields, temperature fluctuations, and changes in the mean thermal structure of the upper atmosphere \\citep{sn1998}.  Abundances derived from these simulations yield a very different pattern, which has been developed in a series of papers \\citep{pla2001,pla2002,asplund2000a,asplund2000b,agsak2004,agsak2005}; papers by Lodders (2003) and \\citet{ags2005} summarize the revised abundance scale.  The net effect is in the sense of systematically lower metal abundances.  The downward revisions for the heavier elements (e.g. Fe and Si) are small, while the claimed reduction in the abundances of lighter species (especially CNO) is more dramatic.  Models employing different treatments of granulation and non-LTE corrections \\citep{h2001,sh2002} predict smaller abundance reductions.  The central temperature predicted by interiors models is sensitive to the abundances of the heavier elements, but not the lighter ones.  As a result, the new abundance scale does not disturb the agreement between interiors models and observational data for purposes such as the mass-luminosity relationship and solar neutrino fluxes.  However, the inferred solar sound speed profile, and the radii of interiors models, is sensitive to the bulk metallicity. Serious problems have emerged when comparing interiors models with the revised abundance scale.  These discrepancies are evidence for problems in our understanding of stellar interiors, stellar atmospheres, or both.  In Paper I \\citep{dp2006} we investigated the errors in solar abundances predicted by the combination of stellar interiors models and helioseismic data.  In this paper we examine the uncertainties in the abundance predictions from stellar atmospheres theory.  We begin with a brief summary of the results from Paper I, and follow with a discussion of the motivation and main results from the revised stellar atmospheres models.  In Section 2, we perform a critical analysis of the precision of the solar CNO abundances and discuss the implications for Be.  We demonstrate in that section that the errors in the abundances are larger than previously estimated, and that there is evidence that the ``best'' current solar CNO abundances are intermediate between the new and old scales, with errors permitting both.  We discuss the implications of our finding and future tests in Section 3. In particular, we argue that inconsistencies between the solar thermal structure and that predicted by the simulations would favor a higher abundance scale closer to the seismic value and discuss uncertainties in the numerical convection simulations.   \\subsection{Constraints from Helioseismology}  Helioseismology provides two powerful constraints on the solar composition: diagnostics of the internal solar temperature gradient and diagnostics of the equation of state.  Inversions of the observed solar pulsation frequencies yield accurate measures of the sound speed as a function of depth.  In turn, the gradient in the sound speed can be directly tied to the temperature gradient.  Since the temperature gradient is related to the opacity, and thus the composition, information on the solar abundances is encoded in the seismic data for the radiative interior.  One can even obtain meaningful constraints on the solar age from the helium abundance profile deduced in the deep interior.  In addition to the vector information on the sound speed profile, there are also precise scalar quantities that can be extracted.  The thermal structure at the base of the solar convection zone is nearly adiabatic, while the temperature gradient in the interior is radiative.  As discussed in Paper I, Section 2.2 the resulting discontinuity in $\\nabla$ generates a distinct signal that can be used to precisely localize the base of the solar convection zone \\citep[$R_{cz}= 0.7133 \\pm 0.0005 R_{sun}$][]{ba2004}.  Seismology also sets strict limits on convective overshooting \\citep[$< 0.05 H_P$,][]{cdmt1995}.  The depth of the solar convection zone is sensitive to the light metal abundances in the Sun but insensitive to most of the other uncertainties in solar interiors models (see Paper I for a detailed error budget).  Ionization also induces a depression in the adiabatic temperature gradient, and the absolute abundances of the species in question can be inferred from the magnitude of the perturbation in the surface convection zone.  An extremely precise surface helium abundance can be deduced from this effect ($Y_{surf}=0.2483 \\pm 0.0046$, see Paper I, Section 2.3 for the sources used in this estimate).  More recently, \\citet{ab2006} have demonstrated that the ionization signal of metals in the convection zone can be detected in the seismic data, leading to a bulk metallicity $Z=0.017 \\pm 0.002$.  Because the majority of the solar metals are in the form of CNO, this is primarily a constraint on their abundance.  In principle one might be able to use this technique to solve for individual heavy element abundances by fitting the strength of distinct ionization stages.  However, it is not yet clear that there is sufficient spatial resolution in the seismic data to permit such a detailed analysis.  In Paper I we demonstrated that the combination of the surface convection zone depth and surface helium abundance constraints was a powerful diagnostic of the solar heavy element abundances.  The surface helium abundance is tied to the initial solar helium abundance with a correction for gravitational settling.  The initial helium is sensitive to the central opacity and the abundances of the heavier metals (especially iron).  The convection zone depth is sensitive to the opacity at temperatures ~ 2 million K, where bound-free opacity from light metals (CNONe) is an important contributor.  The most significant new finding in Paper I was that the combination of the two scalar constraints could be used to rule out some abundance combinations with high statistical significance.  The detailed sound speed profile adds additional information; we found that models consistent with the scalar constraints could be constructed with low oxygen and very high neon, but such models exhibited substantial sound speed deviations relative to solar data in the deep interior.  The inferred solar oxygen and iron abundances ($[O/H]=8.86 \\pm 0.05, [Fe/H]= 7.50 \\pm 0.05$) are consistent with the \\citet{gs1998} absolute abundances, but strongly inconsistent with the new abundance scale within its quoted errors \\citep{l2003,ags2005}.  Although there are potentially positive chemical evolution consequences for the revised abundance scale \\citep{turck2004}, it is not easy to generate interiors models that are consistent with both seismology and the low abundance scale.  The most commonly cited possible explanations on the interiors side (high neon, enhanced gravitational settling, and errors in the high temperature radiative opacities) are all strongly disfavored.  As previously mentioned, solutions with high neon degrade agreement with the sound speed profile, and are also problematic from a stellar atmospheres perspective \\citep{s2005}.  An increase in the degree of gravitational settling increases the convection zone depth but decreases surface helium, trading improved agreement with one diagnostic for worse agreement in another.  Enhanced differential settling of metals with respect to helium is inconsistent with the underlying physics and would have to be extreme \\citep{gwc2005}.  Three independent quantum mechanical calculations yield extremely similar Rosseland mean opacities at the temperatures of interest for the base of the solar convection zone.  As discussed in Paper I, both the atomic physics and equation of state are relatively simple in this regime, and the concordance between different calculations is thus not a surprise.  Physical processes neglected in classical stellar models (such as rotational mixing and radiative acceleration) can be independently constrained by other data, and in any case they would tend to induce higher rather than lower surface abundances.  The scalar constraints are insensitive to the other theoretical ingredients in standard solar models (e.g. convection theory, surface boundary conditions, low temperature opacities, equation of state, and nuclear reaction rates).  The considerations above indicate that it is extremely challenging to reconcile a low solar metal abundance with current stellar interiors models and seismic data.  This does not imply that a metal-poor Sun is impossible, but it certainly motivates an investigation of the uncertainty in the atmospheres models used to derive the abundances.  \\subsection{Model Atmosphere Ingredients}  The new solar abundance estimates are derived from a variety of changes in the model atmospheres.  Changes in oscillator strengths and equivalent widths of spectral lines contribute for some diagnostics, and as discussed below we largely concur with the revised values.  The magnitude of non-LTE corrections depends on the atomic model and the relative importance of photo-excitation and collisions on the level populations in the model atmosphere.  The comprehensive re-examinations of the solar oxygen \\citep[hereafter AGSAK04]{agsak2004} and carbon \\citep[hereafter AGSAB05]{agsab2005} adopted a particular model for NLTE corrections, and we assess its accuracy and uncertainties by comparison with limb darkening data and other published calculations.  In \\citet{pla2001,pla2002} the abundances derived from forbidden lines have been reduced by the application of blending corrections; AGSAK04 and AGSAB05 used the revised equivalent widths for C and O abundance studies.  Both the uncertainty in these corrections and their central value is incorporated in our error analysis.  Three other coupled changes in the atmosphere models are the treatment of convective velocity fields (``macro/microturbulence''), horizontal temperature fluctuations (granulation), and the impact of convective overshooting on the mean thermal stratification.  All of these features are derived from numerical convection simulations; a good discussion can be found in \\citet{sn1998}.  Their combined impact can be deduced by the comparison of the results in the 3D case in the published studies with the results from the semi-empirical 1D HM74 thermal structure and the theoretical 1D MARCS models.  Since the initial 3D model atmosphere is derived with physics similar to that in the MARCS code, the impact of the convection treatment can be indirectly inferred by comparing MARCS and 3D abundances.  A comparison of HM74 with MARCS and 3D is a measure of the impact of different choices of the thermal structure.  The numerical convection simulations predict line profiles in excellent agreement with the data for iron and silicon lines \\citep{asplund2000a,asplund2000b}.  There are some trends with excitation potential that may be related to NLTE corrections \\citep{stb2001} and issues with the initial generation of simulations when compared with line profiles in the outer solar photosphere \\citep{scott2006}.  The concordance of the predicted amplitude of horizontal temperature fluctuations with the solar granulation pattern is encouraging \\citep{sn1998, asplund2000a}.  The validity of the steep solar temperature gradient in the simulations is not as clear; there is an apparent conflict between the simulations and the mean thermal structure of the solar atmosphere, as well as the degree of convective penetration for the upper atmospheric layers \\citep{ay2006}.  Since these effects are all tied together in the abundance studies, we focus on the agreement between different diagnostics of abundance as a valuable test of the precision of the results obtained from different atmospheric models. ",
        "conclusions": "Our basic conclusion is simple: the difference between the solar CNO abundances as derived from model atmospheres and model interiors considerations is not statistically significant.  The systematic errors in photospheric abundance indicators will have to be reduced before a ``solar abundance problem'' can be established (or ruled out) with confidence. However, the disagreement between the solar thermal structure and that of the simulations would favor the higher abundance scale, and there is some recently published evidence to that effect.  If this is confirmed, it switches the nature of the problem from being a question of the correct abundance scale to a question of the uncertainties in numerical convection simulations. We begin with a synthesis and explanation of our findings.  We then divide our conclusions into two parts. We recommend steps to more firmly establish the photospheric abundance scale, and contend that accurate solar abundances require tests of the thermal structure of the models and the magnitude of non-LTE abundance corrections.  In our final subsection we then gather together evidence that the atmospheric abundance scale problem may be tied to the limited resolution in the convection simulations or errors in the underlying model atmosphere treatment.  The consequences for the solar beryllium abundance, which is a useful diagnostic of internal mixing, are also explored.  The two main justifications for the superiority of the 3D hydro atmospheres are the treatment of line broadening and the inclusion of granulation.  Both of these represent genuine improvements in the atmospheric physics.  However, neither of these effects is actually primarily responsible for the difference in the solar abundance scale.  Many of the abundance indicators are insensitive to the effective microturbulence.  If temperature fluctuations are imposed on a semi-empirical Holweger-Mueller atmosphere, the resulting granulation corrections are usually smaller than the 3D convection effects reported by Asplund and collaborators, and frequently opposite in sign \\citep{h2001}.  The main driver behind the systematic reductions in abundance derived from the 3D models is a theoretically predicted change in the thermal structure, coupled with large assumed non-LTE corrections for atomic features.  Neither of these changes is directly supported by observational tests.  Instead, the argument for the superiority of the abundances derived from the newer model atmospheres is an indirect one, focused on the concordance of abundances derived from different indicators.  A consistent chain of logic emerges from the comprehensive studies of oxygen (AGSAK04) and carbon (AGSAB05). Classical LTE model atmospheres tend to yield internally consistent, and high, carbon and oxygen abundances for atomic and molecular indicators. The application of a different thermal structure in the 3D hydro atmospheres drastically reduces the abundances inferred from highly temperature sensitive molecular indicators, but has a smaller effect on atomic features. Large NLTE corrections are then applied to the abundances derived from permitted atomic features for both 1D and 3D models. The net result is that the abundance estimates from 1D models become internally inconsistent (atomic indicators yield lower abundances than molecular ones), while abundances derived from the 3D models are internally consistent.  The abundances derived from forbidden lines are insensitive to NLTE effects, but they are reduced in the newer generation of models by the inclusion of blending features.  As a secondary argument, the fits to individual indicators are argued to be superior in the 3D models when compared to the fits to individual indicators in the 1D models.  This approach is appealing on the surface, but when examined in detail the picture is decidedly more ambiguous.  If anything, the hints from the data would lean towards the opposite conclusion.  The abundances derived from forbidden lines have the smallest systematic errors, but errors in both the theoretical oscillator strengths and the treatment of blending features result in non-negligible random errors.  More to the point, the internal consistency of abundances derived from forbidden and molecular lines is actually similar in the 3D and 1D cases.  From Table 1, the forbidden and molecular oxygen abundances are (8.71, 8.64) for 3D and (8.77, 8.84) for the HM; the differences are identical.  Given the errors, neither discrepancy is statistically significant with high confidence.  The abundances reported for permitted atomic features in AGSAK04 and AGSAB05 are significantly lower for 1D models than the corresponding molecular abundances, while the reported 3D results are in agreement.  In the case of oxygen, this rests completely on the assignment of large NLTE corrections.  These corrections were obtained under the assumption that hydrogen collisions were unimportant.  Detailed studies of the response of the triplet to limb-darkening indicate that models including hydrogen collisions are favored, and the inferred NLTE corrections decrease.  As a result, the internal consistency of the oxygen indicators is comparable for the different classes of atmospheres.  Nitrogen is consistent for HM74 models and inconsistent (but at less than $2 \\sigma$) for the 3D case.  In the case of carbon, the situation is made more complex by significant zero-point offsets between earlier studies of carbon abundances that are not explained.  Again, the assignment of larger NLTE corrections is uncertain (and, unlike the case of oxygen, not directly tested against limb-darkening data).  A clean distinction between models on the basis of consistency is not obtained.  However, the 3D models do yield different molecular and atomic abundances for both N and O, and might also do so for C.  One might then hope to find distinct differences in the quality of the fits to different molecular indicators.  The usual patterns, unfortunately, manifest themselves as simple zero-point shifts.  For every case where there are issues with the 1D models (e.g. small trends with excitation potential in the [O/H] derived from (v,r) OH transitions in the HM model) there are comparable or larger effects for the 3D models (e.g. substantial trends in the [O/H] derived from (r,r) OH transitions).  In a recent preprint, \\citet{scott2006} examined CO indicators, and the resulting pattern is illustrative.  The 3D models yielded similar results for two of the three features studied, while the 1D models performed better in a different pair of indicators.  The $C^{12}/C^{13}$ ratio from the 1D models ranges from 69 to 84, while the same ratio for the 3D models ranges from 83 to 108.  These values should be contrasted with the expected terrestrial ratio of 89.  Scott et al. (2006) choose comparisons that favor the 3D models, while an advocate of the traditional models might reasonably stress the other cases.  In our view, the best choice of models is not clearly distinguishable from the CNO abundance studies.  We recommend caution when extrapolating these model results to other stars, where the differential effects can be even more drastic.  \\subsection{Establishing the Absolute Photospheric Abundance Scale}  The single most important test that is required for atmospheres theory is a discriminant between the different proposed thermal structures of the solar atmosphere.  The recent paper by Ayres et al. (2006) makes an important contribution by making direct comparisons of solar data with the thermal properties of the simulations.  They present evidence that the solar center-to-limb variations in continuum flux are inconsistent with the predictions of the 3D hydro simulations.  They also note that the predicted magnitude of fluctuations in the upper atmosphere from the simulations appears to be larger than the observed pattern.  Ayres et al. then construct an empirical model of the atmosphere and derive a high oxygen abundance (8.85) from CO molecular features under the assumption of a fixed C/O ratio.  In retrospect this conclusion is not surprising.  The HM model is not a purely theoretical exercise; it was constructed to reproduce the mapping of the source function as a function of optical depth inferred from limb darkening studies of continuum flux and strong lines \\citep[see also][]{arg1998}.  The relative trends we have inferred from atomic and molecular abundance indicators support the conclusions of Ayres et al., but the current errors make our evidence in this matter suggestive but not conclusive.  It would also be highly beneficial to repeat the HM exercise with the full 3D models as opposed to the restricted form of them that Ayres et al. (2006) had available to them.  Ultimately, the absolute accuracy of photospheric abundances is directly tied to the absolute accuracy of the thermal structure.  This suggests that an approach similar to that of \\citet{sh2002} may be the optimal one.  In their paper they examined the impact of temperature fluctuations around an assumed mean empirical thermal structure, which in their case was the HM74 model.  Interestingly, the abundance corrections that they derive would act in the sense of increasing the concordance between abundance indicators.  Oxygen abundances from atomic indicators would be slightly increased; although they did not consider molecular features directly, the net effect would certainly have the same sign as that obtained from 3D hydro models, namely a decrease in the inferred abundance.  In such a differential approach, deviations between the mean structure of the simulations and the empirical data would be used as guidance concerning the underlying physics.  In contrast, the 3D model abundances assume that the ab initio profile is correct.  A similar approach could be employed for the velocity field that replaces the microturbulence and macroturbulence in traditional 1D atmospheres.  A second ingredient that must be tested empirically, rather than by theoretical assertion, is the magnitude of NLTE corrections.  The available evidence suggests that NLTE corrections are in general small for the Sun, but for the level of precision required in the absolute abundance scale these small corrections are significant.  Studies of different spectral features yield different conclusions about the physical model employed in NLTE studies.  This implies that there are significant uncertainties in absolute theoretical calculations.  Fortunately, NLTE corrections can be constrained by the response of line strength to limb darkening in the Sun.  It should be possible to develop improved theoretical models with a sufficient database of information developed in this fashion.  One other stringent test of NLTE effects may be to focus on the species whose relative abundances can be reliably inferred from meteoritic data.  For example, NLTE effects may be significant for iron \\citep{stb2001} but less so for Si \\citep{w2001}.  \\citet{h2001} noted that there may be a conflict between the photospheric and meteoritic Fe/Si ratio, albeit one of marginal significance.  A similar situation may exist for Na \\citep{ags2005}.  Another tractable problem is the absolute error for the forbidden C and O lines.  In these cases, uncertainties in the line profiles and continuum levels should be included.  Better atomic data (such as oscillator strengths for both the lines and the blending features) would also be useful.  The accuracy of the theoretically predicted turbulent velocity field as a function of optical depth should also be subjected to a more rigorous analysis.  \\citet{scott2006} present evidence that the generation of simulations used for the abundance analysis yielded poor fits to the line bisectors of CO lines.  Higher resolution simulations gave better line profile fits, but for (unspecified) unrealistically high C/O abundances.  The higher resolution simulations were not employed in the CO abundance analysis in that paper.  It is worth keeping in mind that line profiles are integral quantities, and as a result the uniqueness of the solutions is not established by individual cases of good fits.  This is particularly true when the abundance itself is treated as a free parameter.  It would be extremely useful if future papers on abundances derived using numerical simulations illustrated individual line fits, as well as quantifying the actual impact of the ``effective microturbulence'' on the abundance estimates.  It is useful to separate out the impact of velocity broadening from the effect of granulation and temperature gradient changes.  This can be done by using the mean thermal structure and temperature fluctuations from the simulations and a more traditional micro/macroturbulence model to infer abundances, and comparing the results with the full 3D models.  \\citet{scott2006} constructed such a test case (their 1DAV model), and found only small abundance offsets, of order 0.04 dex for oxygen derived from IR OH lines.  They also inferred carbon abundances from CO; in this case O was held fixed and the carbon was adjusted to fit different molecular indicators.  The deviations in the derived carbon abundances relative to the 3D case ranged from small (0.01 dex for the LE lines) to modest (0.06 dex for the weak $\\Delta \\nu = 1$ lines) to large (0.14 dex for the $\\Delta \\nu = 2$ lines). These deviations may explain the changes in excitation potential that \\citet{ay2006} needed to obtain consistent abundances within a 1D framework. This exercise implies that the impact of the improved microphysics varies substantially for different indicators, and is worth quantifying across the board.  An alternate exercise (using the revised velocity field and relative temperature fluctuations while adopting a HM74 mean thermal structure) might also be illuminating.  \\subsection{Uncertainties in Numerical Convection Simulations}  First-principles theoretical model atmosphere calculations have undeniable strengths.  The ability to naturally reproduce line widths and include granulation is a powerful addition to our ability to reliably interpret stellar and solar spectra.  The principal difficulty with such models is that errors in the input physics generate absolute errors in the inferred atmospheric structure that cannot be calibrated away in the absence of explicit free parameters.  This phenomenon is the major reason why numerical convection simulations have not replaced the simple mixing length theory in stellar interiors calculations.  Interiors models that can reproduce observed stellar radii are simply more useful for most purposes than models with a better physical treatment of convection that fail to do so.  Before the results from such models are adopted as the new abundance standard, it will be necessary to perform an extensive theoretical error analysis and to compare the models with the strongest observational constraints.  We believe that accurate solar abundance calculations must reproduce the observed solar thermal structure, and from the Ayres et al. (2006) paper the Asplund models employing numerical convection simulations appear to yield a temperature gradient steeper than the real Sun.  This could be caused by errors in the background (1D) stellar atmospheres treatment; for example, uncertainties in the equation of state and continuous opacities will induce absolute errors in the thermal structure.  An approach similar to that employed in interiors models would be useful for assessing the uncertainties in the thermal structure and abundance predictions, and this should be included in the error budget for abundances.  It is more likely, however, that the major error source in 3D hydro model atmospheres is related to uncertainties in the numerical convection simulations. The approximations in hydro simulations of giant planet atmospheres are have been demonstrated to be strongly affected by the quality of the assumed physics \\citep{eg2004}. \\citet{zs2006} also provides a good summary of the uncertainties in the related problem of terrestrial and solar dynamo models.  Another phenomenon that could be related is the issue of convective overshooting below surface convection zones.  Numerical simulations have tended to favor extensive overshooting, and the early models had a substantial nearly adiabatic overshoot region, in conflict with the stringent limits set by seismology (less than $ 0.05 H_p$).  More recent 3D \\citet{bct2002} and 2D \\citep{rg2005,rg2006} calculations found that the filling factor for plumes is smaller than previously thought, which led to an overestimate in earlier models of the changes induced by overshooting in the thermal structure.  The newer simulations predict strongly subadiabatic overshooting (effectively, overmixing without changing the thermal structure), which is consistent with the seismic limits.  However, they still produce a substantial mixed region below the surface convection zone of order $0.4 H_p$.  Since even a small overmixing of $0.05 H_p$ drastically increases pre-MS lithium depletion \\citep{mhp1997}, which is already too efficient relative to stellar data \\citep{ptc2002}, it is likely that even this reduced overshooting is too large to be compatible with stellar constraints.  We argue that there is a common pattern in both \"undershooting\" and \"overshooting\" above and below convective regions.  In both cases, the numerical simulations may be overestimating the degree of mixing and the impact on the thermal structure of convection outside the formal bounds set by the Schwartzschild criterion.  There are two plausible error sources that should be investigated.  The treatment of heat transfer in the atmosphere convection simulations is necessarily simplified, and this may be leading to an artificial inhibition in energy transport between turbulent cells projected into the radiative atmosphere and their surroundings.  Resolution effects, however, may be even more important.  Even the highest resolution simulations available today are many orders of magnitude away from being able to reproduce the characteristic Reynolds numbers in the Sun.  Scott et al. (2006) found significant changes in line bisectors for the outer layers of their solar model when they increased their resolution, and these changes were in the sense of reducing the temperature contrast in the upper atmosphere and improving the shape of the bisectors relative to data.  Numerical tests with substantially increased resolution may shed some interesting light on the sensitivity of the predictions to the underlying numerics; 2D convection simulations may be useful in this regard.  We are optimistic that the net effect of such testing will be a greatly improved understanding of the strengths and weaknesses of theoretical atmospheres models, just as we are confident that the net result of the solar abundance controversy will be a far more secure knowledge of stellar abundances."
    },
    "0606/astro-ph0606307_arXiv.txt": {
        "abstract": "{We report on the results of an EPIC XMM-Newton observation of the faint source SAX~J1748.2$-$2808 and the surrounding field. This source was discovered during the BeppoSAX Galactic center survey performed in 1997-1998. A spatial analysis resulted in the detection of 31 sources within the EPIC field of view. SAX J1748.2$-$2808 is clearly resolved into 2 sources in EPIC images with the brighter contributing almost 80\\% of the 2--10~keV flux. Spectral fits to this main source are consistent with an absorbed power-law with a photon index of $1.4\\pm 0.5$ and absorption equivalent to $14 ^{+6} _{-4}\\times 10^{22}$~cm$^{-2}$ together with an iron line at $6.6 ^{+0.2} _{-0.2}$~keV with an equivalent width of $780 ^{+620}_{-380}$~eV. The significantly better statistics of the \\xmm\\ observation, compared with \\sax, allows to exclude a thermal nature for the X--ray emission. A comparison with other observations of SAXJ1748.2$-$2808 does not reveal any evidence for spectral or intensity long-term variability. Based on these properties we propose that the source is a low-luminosity high-mass X-ray binary located in the Galactic center region. ",
        "introduction": "\\label{sect:intro}   \\src\\ is an X--ray source discovered with the Narrow Field Instruments on-board \\sax\\ during a survey of the Galactic Center region (hereafter GC) performed in September 1997 (Sidoli~\\cite{s:00}, Sidoli et al.~\\cite{s:01}). The \\sax\\ spectrum was severely absorbed and displayed an intense Fe~K emission line. The spectrum was poorly constrained, making both thermal and non-thermal nature for the X--ray emission possible. The source, unresolved at the angular resolution of the MECS instrument  (FWHM~$\\sim$1$'$), is located in the direction of the giant molecular cloud Sgr~D.  At the time of its discovery, the nature of {\\mbox \\src} was uncertain and its intense Fe~K line emission, together with its highly absorbed spectrum, made it a unique object in the GC region which could well represent the bright tail of a distribution of similar unresolved objects significantly contributing to the diffuse Fe line emission (at 6.7~keV) from the galactic ridge (Koyama et al.~\\cite{k:89}; Ebisawa et al.~\\cite{e:01}).   Interestingly, \\src\\ displays properties very similar to a class of sources subsequently discovered with the INTEGRAL satellite (see e.g., Walter et al.~\\cite{w:04}, and references therein): these objects show strong photoelectric absorption, hard 2--10 keV spectra,  and often display intense Fe line emission. Most of them also show X--ray pulsations, thus indicating that they are likely high mass X--ray binaries embedded in a local absorbing gas.  Here we report the results of an \\xmm\\ observation the region of sky around \\src, performed with the main goal of unveiling the nature of this intriguing source.   \\begin{figure*}[ht!] \\centering \\includegraphics[angle=0,height=8cm]{5418fig1.ps} \\caption{{\\em Left panel:} Combined (pn+MOS1+MOS2) EPIC image (2--10 keV) centered on \\src. Contours (at 5, 10, 20 and 30 counts/pixel) mark the two sources (the ``main'' and the ``faint'') resolved with \\xmm. {\\em Right panel:} Optical image of the source field, from the ``Second Epoch Survey\" of the southern sky made by the Anglo-Australian Observatory (AAO) with the UK Schmidt Telescope (digitized plates available from  STScI at {\\em http://archive.stsci.edu/}). The star positionally coincident with the faint source is 0600-28834001 in the USNO-A2.0 catalog.} \\label{fig:resolved} \\end{figure*} ",
        "conclusions": "\\label{sect:discussion}   In Sidoli et al. (\\cite{s:01}) we reported the discovery of a new X--ray source in the direction of the Sgr~D region, \\src. Our new  \\xmm\\ observation allows to resolve it into two sources (sources 3 and 12 in Table~\\ref{tab:cat}), with a brighter ``main'' source contributing almost 80\\% of the source flux in the 2--10~keV energy range.  The fainter source is harder (detected only above 5~keV) than the ``main'' one. A possible optical counterpart is the star 0600-28834001 of the USNO-A2.0 catalog (B=18.1, R=13.4), which is listed as [RHI84]10-672  in the   Raharto et al. (1984) catalog of M-type stars. The derivation of log(f$_{X}$/f$_{V}$) is highly uncertain, but assuming, e.g., a blackbody emission at kT$\\sim$1~keV, absorbed with a column density of 10$^{24}$~cm$^{-2}$, the 5--10 keV flux translates into a 0.3-3.5 keV flux $\\sim$4$\\times$10$^{-11}$~erg~cm$^{-2}$s$^{-1}$ (corrected for the absorption), and to a log(f$_{X}$/f$_{V}$)$\\sim$1.8, clearly not stellar. Thus, the hardness of the X--ray emission excludes a coronal emission for the fainter source.  The refined sky position of the brighter source allows to reject all the possible associations discussed in  Sidoli et al. (\\cite{s:01}). The \\sax\\ spectrum was affected by a high interstellar absorption, $N_{\\rm H}$$\\sim$$10^{23}$~cm$^{-2}$, suggesting that the source is probably  located at the GC distance (in this case the luminosity in the 2--10 keV energy band is $\\sim$10$^{34}$~erg~s$^{-1}$). A strong Fe~K line was present (with a line centroid of $6.62\\pm0.30$~keV), and a good fit was obtained both with a power-law plus a Gaussian line, and with a hot thermal plasma model with a temperature, kT, of $6^{+35} _{-4}$~keV. Thus, the \\sax\\ spectrum was consistent with both thermal and non-thermal models.  The significantly better statistics of the \\xmm\\ spectrum and the smaller uncertainties in the spectral slope, favor a non-thermal nature for the X--ray emission of the ``main'' source. A hard power-law ($\\Gamma \\sim$ 1.4) is a good fit to the data, with an iron line and a high photoelectric absorption (\\nh = 10--20~$\\times$10$^{22}$~cm$^{-2}$). The absorption is probably not intrinsic, since the source is located within about 1~degree from the direction of the Galactic center. Thermal models do not fit the X--ray spectrum as well, and result in very high temperatures (for example, a thermal plasma should be hotter than 15~keV). Among the thermal models tried, the blackbody is the best in fitting the spectrum, but results in a high temperature ($\\sim$2~keV) and in an emitting region of less then 0.1~km at the galactic center distance. Thus, the X--ray spectral shape favors a non-thermal nature for the X--ray emission.   The X--ray emission appears to be stable; the ``main'' source has been detected at large off-axis angle during two previous \\xmm\\ observations performed in September 2000 and March 2003 (both pointed at the SNR G0.9+0.1). In both occasions, \\src\\ did not show evidence for any strong flux variability. Moreover, the total emission from ``main'' plus ``faint'' source, is compatible with that observed with  \\sax\\ in 1997 (see Fig.~\\ref{fig:sax}).   These properties are suggestive of three possibilities: a binary system containing a compact object, a background AGN, or reflection from a molecular cloud core (e.g., similar to the X--rays emission and fluorescent iron line produced from the molecular cloud Sgr~B2; Revnivtsev et al.~\\cite{r:04}). This third possibility, already discussed in Sidoli et al.~\\cite{s:01}, seems now unlikely, based on the high spatial resolution of the \\xmm\\ observation. The compact cores contained in the giant molecular cloud Sgr~B2, for example, are about 1~pc in size (Lis \\& Goldsmith~\\cite{lg:91}), while the \\xmm\\ spatial resolution (FWHM$\\sim$6$''$) allows us to exclude a source with a size larger than $\\sim$0.25~pc at a distance of 8.5~kpc.  The shape of the X--ray spectrum and the parameters of the iron line are consistent with a background AGN. It should be a nearby object, since the Fe line is not red-shifted. Assuming a distance of 5~Mpc, the 2--10~keV unabsorbed flux corresponds to a luminosity of $\\sim$3.4$\\times$10$^{39}$~erg~s$^{-1}$, which is quite low, but still compatible with a low-luminosity Seyfert galaxy (Terashima et al.~\\cite{t:02}). Note that no radio counterpart is present in the NED catalogue within 30$''$ of the X--ray position, and \\src\\ does not show evidence for X--ray variability on timescales of years, while X--ray temporal variability and presence of radio emission are typical properties of AGNs.  The X--ray spectral properties of \\src\\ are reminiscent of the soft gamma-ray sources discovered with the INTEGRAL satellite (see e.g., Kuulkers~\\cite{k:05} for a review). Several of their X--ray counterparts display hard and heavily absorbed spectra, together with intense fluorescent Fe line emission, indicative of dense gaseous envelopes around the compact object, illuminated by the central source. In few of them, the association with OB optical counterparts and the detection of X--ray pulsations, suggest that they are highly absorbed HMXRBs, not detected in previous surveys at soft X--rays. The derived luminosity of these INTEGRAL sources is around 10$^{36}$~erg~s$^{-1}$, although there is a large uncertainty in the distance estimates, and  the true luminosity could be much less than this. \\src\\ lies in the direction of SgrD molecular cloud, near to SgrB2, which is an important site of star formation, so it is not unlikely that \\src\\ is indeed a HMXRB. The low X--ray luminosity suggests that it belongs to a class of massive X--ray binaries with low persistent emission (in the range 10$^{34}$--10$^{35}$~erg~s$^{-1}$), wind-accreting and with no outbursts (e.g., 4U~2206+54, Masetti et al.~\\cite{m:04}). On the other hand, these sources typically show temporal variability on different timescales (sometimes with flares), which has not been observed in \\src\\ (perhaps because of the poor coverage). However, wind-fed HMXRBs are usually quite stable X--ray emitters on long timescales (months or years). Low luminosity wind-accreting neutron stars has been predicted by Pfahl et al. (2002), who proposed that most of the faint sources detected in the $Chandra$ survey of the GC (Wang et al. 2002) could be of this kind. A search for hard unidentified sources from ROSAT PSPC observations seems to confirm that a new class of fainter wind-fed X--ray binaries exists in our Galaxy (Suchkov \\& Hanisch 2004).   Other kinds of galactic X--ray binaries, containing neutron stars or black-holes, seem to be unlikely; the luminosity ($\\sim$10$^{34}$~erg~s$^{-1}$) suggests an object in quiescence: but low-mass X-ray binaries (LMXRBs) in quiescence (soft X--ray transients) typically have much softer spectra (blackbody temperatures $\\sim$0.1--0.3~keV; e.g., Verbunt \\& Lewin~\\cite{vl:04}), while black-hole X--ray novae in quiescence have much lower luminosities (Kong et al.~\\cite{k:02}).   In conclusion, among the different hypotheses discussed above, the spectral shape (hard, non--thermal), X--ray luminosity, the presence of Fe line emission, seem to favor a low luminosity HMXRB."
    },
    "0606/physics0606083_arXiv.txt": {
        "abstract": "A first detailed study of the ground state of the H$_3^+$ molecular ion in linear configuration, parallel to a magnetic field direction, and its low-lying $\\Si,\\Pi,\\De$ states is carried out for magnetic fields $B=0-4.414 \\times 10^{13}\\,$G in the Born-Oppenheimer approximation. The variational method is employed with a single trial function which includes electronic correlation in the form $\\exp{(\\ga r_{12})}$, where $\\ga$ is a variational parameter. It is shown that the quantum numbers of the state of the lowest total energy (ground state) depend on the magnetic field strength. The ground state evolves from the spin-singlet ${}^1\\Si_g$ state for weak magnetic fields $B \\lesssim 5 \\times 10^{8}\\,$G to a weakly-bound spin-triplet ${}^3\\Si_u$ state for intermediate fields and, eventually, to a spin-triplet $^3\\Pi_u$ state for $5 \\times 10^{10}\\,\\lesssim B \\lesssim 4.414 \\times 10^{13}\\,$G. Local stability of the linear parallel configuration with respect to possible small deviations is checked. ",
        "introduction": "The behavior of atoms, molecules and ions placed in a strong magnetic field has attracted a significant attention during the last two decades (see, in particular, review papers \\cite{Liberman:1995,Lai:2001,Turbiner:2006}). It is motivated by both pure theoretical interest and by possible practical applications in astrophysics and solid state physics. In particular, the knowledge of the energy levels can be important for interpretation of the spectra of white dwarfs (where a surface magnetic field ranges in $B\\approx 10^{6}-10^{9}$\\,G) and neutron stars where a surface magnetic field varies in $B\\approx 10^{12}-10^{13}$\\,G, and even can be $B\\approx 10^{14}-10^{16}$\\,G for the case of magnetars.  Recently, it was announced that in a sufficiently strong magnetic field $B \\gtrsim 10^{11}$\\,G  the exotic molecular ion H$_3^{2+}$ can exist in linear configuration with protons situated along the magnetic line \\cite{Turbiner:1999} (for discussion see a review \\cite{Turbiner:2006}). In general, it is a metastable long-living system which decays to H$_2^+ + p$. However, at $B \\gtrsim 10^{13}$\\,G the ion H$_3^{2+}$ becomes stable. This system does not exist without or for weak magnetic fields. The ion H$_3^{2+}$ constitutes the simplest one-electron polyatomic molecular ion in a strong magnetic field. The H$_3^{2+}$ ion has been proposed as being the most abundant chemical compound in the atmosphere of the isolated neutron star 1E1207.4-5209 \\cite{Turbiner:2004m}. A detailed review of the current status of one-electron molecular systems, both traditional and exotic, that might exist in a magnetic field $B \\geq 10^{9}$\\,G can be found in \\cite{Turbiner:2006}.  The molecular ion $H_3^+$ is the simplest stable two-electron polyatomic molecular ion. It has a long history since its discovery by J.J.~Thomson \\cite{thomson}. Its exceptional importance in astrophysics related to interstellar media explains the great interest in this ion from astronomy, astrophysics and chemistry communities (for a detailed review, see, \\cite{tennyson}). For all these reasons, there have been extensive theoretical and experimental works on this molecular ion since the pioneering (semi-quantitative) work by Coulson \\cite{coulson}.  The first variational calculations \\cite{Hirschfelder} of the total energy of the molecular ion H$_3^+$ showed that the equilibrium configuration might be either linear or equilateral triangular. However, this was not well-established until 1964 \\cite{Christoffersen} when it was shown that the equilibrium configuration for the state of the lowest total energy is an equilateral triangular configuration, while the linear configuration of the H$_3^+$ ion may occur in excited state(s). Since that time a large number of excited states has been studied \\cite{schaad} (for a general review, see \\cite{tennyson}). In particular, it has been found that there exists a single spin-triplet state which appears in a linear configuration ${}^3\\Si_u$. This is also the unique known state of H$_3^+$ in the linear configuration. No spin-triplet states have been found for a triangular (spacial) configurations so far.  Although the molecular ion H$_3^+$ is characterized by the equilateral triangular configuration as being the optimal in field-free case, it is expected that in a magnetic field $B \\approx 0.2$\\,a.u. (see below) a linear configuration, parallel to a magnetic field direction, gives the lower total energy and becomes the optimal configuration. Somehow, a similar phenomenon already happened for the one-electron exotic molecular ion H$_3^{2+}$ \\cite{Turbiner:2002} where the optimal configuration is triangular at $10^8 \\lesssim B \\lesssim 10^{11}\\,$G and becomes linear parallel at $B \\approx 10^{11}\\,$G. It is worth noting that for H$_3^+$ in field-free case the difference between the total energy of the ground state (triangular configuration) and of the lowest linear configuration is very small, $\\approx 0.13$\\,Ry, in comparison to characteristic energies in a magnetic field.  To the best of our knowledge there exists a single attempt to explore the molecular ion H$_3^+$ in a magnetic field \\cite{warke}. We repeated all numerical calculations of this work following its guidelines with use of its formulas (see below, Tables~I, V, VI) - in fact, no single number from \\cite{warke} was confirmed. However, in \\cite{warke} it was made a qualitative statement that with a magnetic field increase the transition from equilateral stable equilibrium configuration to linear equilibrium configuration may occur. This statement we confirm. We predict that this transition takes place at a magnetic field $\\approx 0.2\\,$a.u. A detailed study of a triangular configuration and of this transition will be published elsewhere \\cite{tg}.  Atomic units are used throughout ($\\hbar$=$m_e$=$e$=1), although energies are expressed in Rydbergs (Ry). The magnetic field $B$ is given in a.u. with a conversion factor $B_0 = 2.35 \\times 10^9$\\,G. ",
        "conclusions": "We study the low-lying energy states of H$_3^+$ molecular ion in linear configuration parallel to a magnetic field from 0  up to $4.414 \\times 10^{13}\\,$G using the variational method in the Born-Oppenheimer approximation. The total energy curves display a well pronounced minimum at finite internuclear distances at $R_+=R_-$ for the lowest states with magnetic quantum numbers $m=0,-1,-2$, total spins $S=0,1(m_s=-1)$ and parity $p=\\pm 1$. A level distribution for several magnetic field strengths is shown on Fig.~6. If in field-free case there exist only two eigenstates in a linear configuration, but many more states in linear parallel configuration can appear when a magnetic field is imposed.   \\begin{figure}[htb] \\begin{center} \\includegraphics*[height=9.5in,width=7.in,angle=0]{Elevels.ps} \\vskip -70pt \\caption{Total energy of  the low-lying levels for $B=0, 1, 10, 100\\ \\mbox{and}\\ 1000\\,$ a.u. (energy scale is kept the same for all presented magnetic fields but the reference points depend on them)} \\end{center} \\label{fig:levels} \\end{figure}  In general, for all studied states, as the magnetic field increases the equilibrium internuclear distances $R_{eq}$ decreases and the system becomes more compact, while the total energies of spin-singlet states increase whereas that of spin-triplet states decrease.  The state of the lowest total energy in linear parallel configuration depends on the magnetic field strength. It evolves from spin-singlet (unstable towards a deviation from linearity) $^1\\Si_g$ for weak magnetic fields $B\\lesssim 0.2\\,$a.u. to spin-triplet (stable towards a deviation from linearity) ${}^3\\Si_u$ for intermediate fields and eventually to spin-triplet ${}^3\\Pi_u$ state for $B \\gtrsim 20 $\\,a.u. which remains the ground state until the Schwinger limit $B = 4.414 \\times 10^{13}\\,$G.  It is worth emphasizing that for weak magnetic fields, $B\\lesssim 0.2\\,$a.u., the global ground state is given by a triangular configuration \\cite{tg} and then, for larger magnetic fields, the global stable ground state corresponds to a linear parallel configuration \\footnote{In order to make such a claim that the state of the lowest energy corresponds by a linear parallel configuration we make a very natural physically assumption that there are no any other spacial configuration which may provide a lower total energy. But in order to be rigorous we must investigate the total energy surface for all possible spacial configurations.}. The H$_3^+$ ion in the ${}^3\\Si_u$ state is weakly bound. For all studied magnetic fields the total energy surface well corresponding to the ground state contains at least one longitudinal vibrational state.  It is interesting to compare the evolution of the ground state for H$_3^+$ with magnetic field change with that of other two-electron systems (see \\cite{Turbiner:2006London} and references therein). For atomic type H$^-$ and He systems there is no domain of magnetic field where the spin-triplet, $m=0$ state is the ground state: for weak fields the ground state is the spin-singlet, $m=0$ state and then it becomes the spin-triplet, $m=-1$ state for large fields. For the hydrogen molecule the ${}^3\\Si_u$ state is unbound for all magnetic fields unlike the case of H$_3^+$. It implies that the H$_2$ molecule does not exist as a bound system for $0.18 \\lesssim B \\lesssim 15.6$\\,a.u., where the unbound state ${}^3\\Si_u$ has the lowest total energy at infinitely-large distance between protons. A similar situation occurs for the He$_2^{2+}$-ion: it does not exist as a bound system for $0.85 \\lesssim B \\lesssim 1100$\\,a.u. \\cite{Turbiner:2006He2}.  What is the lowest-lying excited state for weak magnetic fields $B\\lesssim 0.2\\,$a.u. is not clear yet. This question, and also the whole domain $B\\lesssim 0.2\\,$a.u., will be studied elsewhere. In the domain of magnetic fields $0.2~\\le~B~\\le~5\\,$a.u. the lowest-lying excited state is ${}^3\\Si_g$, then for $B \\gtrsim 5\\,$a.u. the lowest-lying excited state is ${}^3\\Pi_u$. For $B \\gtrsim 20 \\,$a.u., where the ${}^3\\Pi_u$ state becomes the ground state, the lowest-lying excited state is ${}^3\\Si_u$. However, at $B \\gtrsim 1000 \\,$a.u. until the Schwinger limit the lowest-lying excited state is ${}^3\\De_g$.  It is interesting to note that at $B = 1000\\,$a.u. the H$_3^+$ ion exists with ${}^3\\Pi_u$ as the ground state ($E_T=-44.54\\,$a.u.) with two possible excited states: ${}^3\\De_g$ ($E_T=-40.38\\,$a.u.) and ${}^3\\Si_u$ ($E_T=-35.99\\,$a.u.) with energies below the threshold of dissociation to H$_2({}^3\\Pi_u) + p$ ($E_T=-35.44\\,$a.u.). For larger magnetic fields the situation becomes different. For instance, at $B = 10000 \\,$a.u. for the H$_3^+$ ion ($E_T=-95.21\\,$a.u.) only one excited state, ${}^3\\De_g$ ($E_T=-87.45\\,$a.u.), exists with energy below the dissociation threshold to H$_2({}^3\\Pi_u) + p$ ($E_T=-71.39\\,$a.u.). Similar situation holds up to the Schwinger limit $B=4.414 \\times 10^{13}\\,$G: a single excited state ${}^3\\De_g$ lies below the dissociation threshold.  It is found that many states in linear configuration which do not exist for $B=0\\,$ begin to be bound at relatively small magnetic field $B \\approx 0.2\\,$a.u. A study of the existence of the bound states which might appear in a spacial configuration is our goal for a future study. Another goal is related to a study of transition amplitudes for different electronic states.  Present consideration is based on the use of a simple variational trial function (\\ref{psi}). This function can be easily generalized and extended in the same way as was done in a variational study of various one-electron systems in a strong magnetic field (see \\cite{Turbiner:2006}). This will allow to improve the present results and might be done in future. However, we are not sure that such a study is crucially important. It is related to a fact that typical accuracies in astronomical observations of neutron star radiation would not be higher $10^{-3} - 10^{-4}$ unlike to spectroscopical accuracies in laboratory where they can be by several orders of magnitude higher."
    },
    "0606/astro-ph0606131_arXiv.txt": {
        "abstract": "Clio is an adaptive-optics camera mounted on the 6.5 meter MMT optimized for diffraction-limited L' and M-band imaging over a $\\sim15''$ field. The instrument was designed from the ground up with a large well-depth, fast readout thermal infrared ($\\sim 3-5\\:\\mu m$) 320 by 256 pixel InSb detector, cooled optics, and associated focal plane and pupil masks (with the option for a coronograph) to minimize the thermal background and maximize throughput. When coupled with the MMT's adaptive secondary AO (two warm reflections) system's low thermal background, this instrument is in a unique position to image nearby warm planets, which are the brightest in the L' and M-band atmospheric windows. We present the current status of this recently commissioned instrument that performed exceptionally during first light. Our instrument sensitivities are impressive and are sky background limited: for an hour of integration, we obtain an L'-band 5 $\\sigma$ detection limit of of 17.0 magnitudes (Strehl $\\sim 80\\%$) and an M-band limit of 14.5 (Strehl $\\sim 90\\%$). Our M-band sensitivity is lower due to the increase in thermal sky background. These sensitivities translate to finding relatively young planets five times Jupiter mass (M$_{Jup}$) at 10 pc within a few AU of a star. Presently, a large Clio survey of nearby stellar systems is underway including a search for planets around solar-type stars, M dwarfs, and white dwarfs. Even with a null result, we can place strong constraints on planet distribution models. ",
        "introduction": "Despite the exhaustive search for extrasolar planets, only one confirmed exoplanet orbiting around the brown dwarf 2M1207 has been imaged through VLT/NACO observations~\\cite{2005A&A...438L..25C}. The first detection of exoplanet photons resulted from Spitzer observations of a secondary eclipse in a known planetary system; however, the system was not resolved~\\cite{2005ApJ...626..523C}. The contrast levels needed to discern Jupiter-sized planets visibly (or even in the near-IR) from their bright stars typically require next generation hardware. Thus, the $>100$ \\emph{indirect} planet discoveries have resulted mostly from high-precision radial velocity (RV) surveys where some have been followed up with transit studies to determine the physical properties of a planetary system~\\cite{2005PThPS.158...24M}.  A handful, including Jupiter-sized planet TrES-1~\\cite{2004ApJ...613L.153A}, have turned up in transit studies (e.g. TrES, OGLE) and been confirmed. However, RV and transit surveys are inherently biased towards large, close-in planets, leaving lower-mass planets on orbits with wider semi-major axes largely unexplored and their distributions unknown. Direct imaging is well-poised to uncover planets at greater separations from stars and therefore complements current RV techniques. At present, direct imaging can uncover the distribution and properties of extrasolar giant planets (EGP) and place strong constraints on current evolutionary models. The $3-5\\:\\mu m$ imaging offers a new window for tackling these scientific questions. ",
        "conclusions": "The instrument is in a state where it can be used for routine observations and is able to set very interesting new limits in planet detection. We are finally able to do a systematic survey of a wide variety of nearby systems and place mass constraints down to 5 M$_{Jup}$ even for moderate age systems. In the future, we hope to extend the capabilities of our instrument through the addition of L' and M-band PSF shaping phase plates and a grism for crude spectroscopy of our objects. We plan to investigate other coronographic techniques to improve 3-5 micron contrasts."
    },
    "0606/astro-ph0606125_arXiv.txt": {
        "abstract": "Hollerbach and R\\\"udiger have reported a new type of magnetorotational instability (MRI) in magnetized Taylor-Couette flow in the presence of combined axial and azimuthal magnetic fields.  The salient advantage of this ``helical'' MRI (HMRI) is that marginal instability occurs at arbitrarily low magnetic Reynolds and Lundquist numbers, suggesting that HMRI might be easier to realize than standard MRI (axial field only). We confirm their results, calculate HMRI growth rates, and show that in the resistive limit, HMRI is a weakly destabilized inertial oscillation propagating in a unique direction along the axis. But we report other features of HMRI that make it less attractive for experiments and for resistive astrophysical disks.  Growth rates are small and require large axial currents.  More fundamentally, instability of highly resistive flow is peculiar to infinitely long or periodic cylinders: finite cylinders with insulating endcaps are shown to be stable in this limit.  Also, keplerian rotation profiles are stable in the resistive limit regardless of axial boundary conditions. Nevertheless, the addition of toroidal field lowers thresholds for instability even in finite cylinders. ",
        "introduction": "\\label{sec:intro} The magnetorotational instability (MRI) is probably the main source of turbulence and accretion in sufficiently ionized astrophysical disks \\cite{bh98}. MRI was first discovered theoretically \\cite{ve59,chan60,bh91}, then later supported numerically \\cite{hgb95,bnst95,mt95}, but has never been directly observed in astronomy.  No laboratory study of MRI has been completed except for that of \\citet{sisan04}, whose experiment proceeded from a background state that was not in MHD equilibrium.  We and others therefore have proposed experimental demonstrations of MRI \\cite{jgk01,gj02,npc02}. The experimental geometry planned by most groups is a magnetized Taylor-Couette flow: an incompressible liquid metal confined between concentric rotating cylinders, with an imposed background magnetic field sustained by currents external to the fluid.  The challenge for experimentation, however, is that liquid-metal flows are very far from ideal on laboratory scales.  While the fluid Reynolds number $Re\\equiv \\Omega_{1}r_{1}(r_{2}-r_{1})/\\nu$ can be large, the corresponding \\emph{magnetic} Reynolds number $\\Rm\\equiv\\Omega_{1}r_{1}(r_{2}-r_{1})/\\eta$ is modest or small, because the magnetic Prandtl number $Pr_{\\rm m}\\equiv\\nu/\\eta\\sim 10^{-5}-10^{-6}$ in liquid metals; here $\\nu\\lesssim 10^{-2}\\cm^2\\s^{-1}$ is the kinematic viscosity and $\\eta$ is the magnetic diffusivity.  Standard MRI modes will not grow unless both the rotation period and the Alfv\\'en crossing time are shorter than the timescale for magnetic diffusion.  This requires both $\\Rm\\gtrsim 1$ and $S\\gtrsim 1$, where $S\\equiv V_{A}(r_{2}-r_{1})/\\eta$ is the Lundquist number, and $V_{A}=B/\\sqrt{\\mu_0\\rho}$ is the Alfv\\'en speed. Therefore, $Re\\gtrsim 10^6$ and fields of several kilogauss must typically be achieved.  Recently, Hollerbach and collaborators have discovered that MRI-like modes may grow at much reduced $\\Rm$ and $S$ in the presence of a helical background field, a current-free combination of axial and toroidal field \\citep{hr05,rhss05}. \\begin{equation} \\label{eq:backgroundfield} \\b{B}^{(0)}=B_z^{(0)}\\left(\\b{e}_z+\\beta\\frac{r_1}{r}\\b{e}_\\theta\\right) \\end{equation} in cylindrical coordinates $(r,\\theta,z)$, where $B_z^{(0)}$ and $\\beta$ are constants.  (When it will not cause ambiguity, we will omit the superscript $(0)$ from $\\b{B}$ and $B_z$ hereafter.) Henceforth, ``standard MRI'' (SMRI) will refer to cases where the $\\beta=0$, and ``helical MRI'' (HMRI) to modes that require $\\beta\\ne0$.  In centrifugally stable flows---meaning that $d(r^2\\Omega)^2/dr>0$, where $\\Omega=V_\\theta^{(0)}/r$ is the background angular velocity---SMRI exists only when $\\Rm$ and $S$ exceed thresholds of order unity \\cite{jgk01,gj02}.  Remarkably, however, HMRI may persist in such flows even as both parameters tend to zero, though not independently: more precisely, the thresholds are $\\ll 1$ and would vanish if the fluid were inviscid ($\\nu=0$).  In a fixed geometry and flow profile, the resistive limit may be approached theoretically by increasing $\\eta$ with all other parameters held constant. The growth rate of inviscid HMRI is then $\\propto\\eta^{-1}$ so that the hydrodynamic case is approached continuously. The special case of toroidal-only magnetic field ($\\beta =\\infty$) is stable \\cite{hs06}.   A novel feature of the background state for HMRI is that there is a uniform axial flux of angular momentum carried by the field, $rT_{\\varphi z}^{(\\rm mag)}=-rB_\\theta B_z/\\mu_0$ and an associated axial Poynting flux $\\Omega$ times this.  In an infinite or periodic cylinder, the question of the sources and sinks of these axial fluxes need not arise, but in an experimental device, a torque is exerted by the axial field on the radial sections of the coil that complete the circuit containing the axial current.  Related to this perhaps, the dispersion relation for linear modes is sensitive to the sign of the axial wavenumber ($k_z$), and the instabilities of axially infinite or periodic cylinders are travelling rather than standing waves, as noted by Knobloch \\citep{ke92,ke96}.  This begs the question what should happen to the modes in finite cylinders, a question that has motivated much of our analysis.  Even the analysis for periodic cylinders implies two practical difficulties for an HMRI experiment.  First, as will be seen, the typical growth rates tend to be smaller than those of SMRI except in regimes where SMRI would also be unstable.  Secondly, the axial current needed for the required toroidal fields tend to be quite large: $I[\\unit{kA}]=5 B_\\theta r[\\mbox{kG-cm}]$.  In Section \\ref{sec:theory} we analyze the linear stability of HMRI using complementary approximations, some for infinite/periodic cylinders and others for finite ones.  The results are compared with one another and with fully nonlinear axisymmetric simulations. Our conclusions are summarized in Section \\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis} We have analyzed the linear development of helical magnetorotational instability in a non-ideal magnetohydrodynamic Taylor-Couette flow, paying particular attention to the effects of the axial boundary conditions.  A number of complementary approximations and numerical methods have been used.  For infinitely long or periodic cylinders, we confirm that there is an axisymmetric MHD instability that persists to smaller magnetic Reynolds number and Lundquist number in the presence of \\emph{both} axial and toroidal background magnetic field than the standard MRI that exists for axial field alone. The new mode is an overstability and propagates axially in the direction of the background Poynting flux $-r\\Omega B_\\theta B_z/\\mu_0$.  In highly resistive flows, the new mode is a weakly destabilized hydrodynamic inertial oscillation. Growth depends also on the ratio of shear to rotation, \\emph{i.e.} Rossby number: for all aspect ratios $r_2/r_1$ that we have explored, and certainly for narrow gaps, the keplerian Rossby number is stable. Even for those profiles that permit growth, the rate tends to be rather small, except in flows that are sufficiently ideal to permit growth of standard MRI (axial field only).  We have also considered finite cylinders with insulating endcaps, which are closer to experimental reality but which do not permit traveling modes that propagate indefinitely along the axis. Astrophysical disks also, of course, have limited vertical thickness. These boundary conditions reduce the growth rate of the helical mode, and stabilize highly resistive flows entirely. The small growth rate may make it difficult to detect the instability in the face of Ekman circulation and other experimental imperfections. In addition, the new mode requires toroidal fields at least as large as the axial field, and therefore large axial currents, which introduces additional engineering challenges.  For all of these reasons, the experimental advantage of helical MRI over standard MRI is open to question, as is the relevance of HMRI to astrophysical disks, although it may be relevant to stellar interiors and jets, where the magnetic geometry and the Rossby number may be more favorable. Also, HMRI may have theoretical significance that goes beyond its direct applications.  It is not understood why linearly and axisymmetrically stable rotating flows are often also nonlinearly and nonaxisymmetrically unstable, especially since subcritical transition does occur at some Rossby numbers \\citep{ll05}. The fact that even a very poorly coupled magnetic field can sometimes linearly destabilize such flows hints that it might also affect nonlinear transition."
    },
    "0606/astro-ph0606639_arXiv.txt": {
        "abstract": "With the excellent angular resolution of the {\\em Chandra X-ray Observatory}, it is possible to geometrically determine the distance to variable Galactic sources, based on the phenomenon that scattered radiation appearing in the X-ray halo has to travel along a slightly longer path than the direct, unscattered radiation. By measuring the delayed variability, constraints on the source distance can be obtained if the halo brightness is large enough to dominate the point spread function (PSF) and to provide sufficient statistics. The distance to Cyg X-3, which has a quasi-sinusoidal light curve, has been obtained with this approach by Predehl et al. Here we examine the feasibility of using the delayed signature of type I X-ray bursts as distance indicators. We use simulations of delayed X-ray burst light curves in the halo to find that the optimal annular region and energy band for a distance measurement with a grating observation is roughly 10--50\\arcsec~and 1--5 keV respectively, assuming \\chan's effective area and PSF, uniformly distributed dust, the input spectrum and optical depth to GX 13+1, and the Weingartner \\& Draine interstellar grain model. We find that the statistics are dominated by Poisson noise rather than systematic uncertainties, e.g., the PSF contribution to the halo. Using {\\em Chandra}, a distance measurement to such a source at 4 (8) kpc could be made to about 23\\% (30\\%) accuracy with a single burst with 68\\% confidence. By stacking many bursts, a reasonable estimate of systematic errors limit the distance measurement to about 10\\% accuracy. ",
        "introduction": "Type I X-ray bursts potentially provide a delayed signature in the halo that can be used to measure the distance to variable X-ray sources with a method proposed by Tr\\\"{u}mper \\& Sch\\\"{o}nfelder (1973; hereafter TS73), prior to the detection of the first X-ray halo around GX 339-4 by {\\em Einstein} (Rolf 1983). The TS73 method is based on the phenomenon that grain-scattered photons from an X-ray source travel along slightly longer paths to the telescope than the direct, unscattered photons. Temporal variations in the intensity are therefore delayed and smeared when they appear in the X-ray halo. Using the single scattering approximation, it is possible to predict the delayed light curve as a function of halo angle, energy, and distance to an accuracy that is limited by knowledge of the distribution of dust along the line of sight and by the accuracy of the interstellar grain model. Hu et al. (2004) developed a similar method to measure distances with X-ray halo variability using the frequency domain.  The feasibility of the TS73 distance measurement method has been demonstrated by Predehl et al. (2000), who derived a distance of 9$^{+4}_{-2}$ kpc to Cyg X-3 using the source's quasi-sinusoidal variability. In this paper, we examine the applicability of the TS73 method to the abrupt increase in flux provided by type I X-ray bursts, which result from unstable thermonuclear ignition of accreted material on the surface of the neutron star. Sharp and temporary intensity changes, such as that produced by type I bursts or gamma-ray bursts (GRBs), produce a qualitatively different response in the X-ray halo than the response from smooth intensity variations like those of Cyg X-3. For example, consider a smooth change in point source flux over an hour or so. Eventually the halo brightness will reach a new base level corresponding to the new point source intensity, and information on the source distance and dust distribution is contained in the rate of response and the resulting phase shift of the light curve. With X-ray bursts a new base level is not obtained; rather, the burst appears as a slight and temporary increase in halo flux after a time delay that depends on the scattering geometry. The intrinsically small signal-to-noise ($S/N$) ratio provided by bursts makes their potential application as distance indicators particularly difficult, requiring the physical parameters governing the delayed light curve in the X-ray halo be measured to a high degree of accuracy, and requiring careful treatment of the point spread function (PSF).  Gamma-ray bursts produce a similar halo light curve, although the problem is much simpler: With spectroscopic determination of the distances to host galaxies of GRBs, the degeneracy between the distribution of dust and the distance to the source is removed. This has been exploited by both the {\\em Swift} and {\\em XMM-Newton} observatories to determine the distances to Galactic interstellar dust clouds along various sightlines (Vaughan et al. 2004; Tiengo \\& Mereghetti 2006; Vaughan et al. 2006).  The physical parameters governing the delayed light curve in the halo are the peak and persistent flux, the spectral shape, the scattering optical depth, the scale time of the approximately exponential burst light curve, and the distribution and composition of interstellar dust along the line of sight. Two of the most widely used interstellar grain models are the Mathis, Rumpl, \\& Nordsieck (1977; MRN) model and the Weingartner \\& Draine (2001; WD01) model. The MRN model is composed of silicate and graphite grains with a size distribution of $n(a) \\propto a^{-3.5}$, which reproduces the observed extinction of starlight. The WD01 grain model additionally accounts for the observed infrared and microwave emission from the diffuse ISM by including sufficient small carbonaceous grains. In this paper we use the WD01 interstellar grain model and we assume that the dust is distributed uniformly along the line of sight.  The structure of this paper is as follows: We begin by describing the \\chan~observation that was used in this analysis and we briefly discuss pile-up (\\S~2). In \\S~3 we present two analytic equations from Draine (2003) describing the scattering of X-rays from interstellar grains and the resulting halos, from which we derive an equation describing the predicted counts spectrum for an arbitrary range of angles and energies. We also discuss our treatment of the PSF. The derived equation is used to estimate the optical depth to GX 13+1 assuming uniformly distributed dust along the line of sight by fitting it to the observed surface brightness distribution (SBD). In \\S~4 we demonstrate our method for calculating the expected delayed halo light curve of singly-scattered photons from a burst. In \\S~5 we show that, given a particular set of assumptions, the 1--5 keV energy range and the 10--50\\arcsec~annular region yield the highest time-averaged $S/N$ for viewing bursts in the halo. With this binning, simulations are carried out to test the feasibility of applying the TS73 method to type I X-ray bursts. We finish in \\S~6 by summarizing our results. ",
        "conclusions": "In this paper, we studied the applicability of the TS73 method to the delayed signature of type I X-ray bursts in the halo. The signal produced by an X-ray burst in the halo is invariably small. Regular type I bursts last $\\sim$5--150 s, but the scattering geometry, which governs the distribution of time delays, assures that these photons will be spread out over time delays greater than $\\sim$10 ks for uniformly distributed dust. Nevertheless, if the physical parameters governing the light curve of the burst in the halo are well-known, the statistics provided by \\chan's effective area can yield a distance measurement to about 25\\% accuracy with a single burst. By stacking many bursts, systematic errors likely limit the potential accuracy to roughly 10\\%. Although non-grating observations have improved statistics due to larger effective area, pile-up precludes accurate flux measurements (although the transfer streak can be used) and require the use of larger halo angles, decreasing the $S/N$ as the evolution of the delayed burst light curve occurs on longer time scales. Source variability on 200 s scales with 40\\% rms amplitude increase the distance uncertainty by $\\sim$15\\%. The quoted errors assume very favorable source characteristics: (1) a bright source with a  flux $\\sim$0.6 photons/cm$^{2}$/s during persistent emission from 1--5 keV, (2) a large optical depth to scattering ($\\tau_{\\rm sca}\\approx$1.4 at 1 keV), both of which increase the $S/N$ ratio of the delayed burst light curve in the halo, (3) a burst with large total fluence, characterized by $t_{\\rm b} = 60$ s and $F_{\\rm peak}/F_{\\rm persistent} = 7$. In the general case, the uncertainty in the inferred distance will be somewhat larger. Finally, note that the interstellar grain model represents an additional source of systematic error.  In many cases the distance to Galactic X-ray sources has been estimated using other methods, e.g., using the requirement that the inferred peak luminosity is sub-Eddington, or using optical photometry of the (low-mass) binary companion, and the potential accuracy of the TS73 method applied to bursters may only provide additional distance constraints for a small fraction of all type I bursters. Moreover, if the distribution of dust differs significantly from the assumed or inferred distribution, the distance measurement can be unreliable. A future observatory with a larger effective area and arcsec resolution would provide superior statistics, however, making the TS73 method an attractive means of measuring the distance to a larger subset of the population of absorbed X-ray bursters."
    },
    "0606/astro-ph0606063_arXiv.txt": {
        "abstract": "We present spectroscopic confirmation of nine moderate redshift galaxy groups and poor clusters selected from the \\textit{ROSAT} Deep Cluster Survey. The groups span the redshift range z $\\sim$ 0.23 --0.59 and have between 4 and 20 confirmed members. The velocity dispersions of these groups range from $\\sim$ 125 to 650 km s$^{-1}$. Similar to X-ray groups at low redshift, these systems contain a significant number of early-type galaxies. Therefore, the trend for X-ray luminous groups to have high early-type fractions is already in place by at least z $\\sim$ 0.5.  In four of the nine groups, the X-ray emission is clearly peaked on the most luminous early-type galaxy in the group. However, in several cases the central galaxy is composed of multiple luminous nuclei, suggesting that the brightest group galaxy may still be undergoing major mergers.  In at least three (and possibly five) of the groups in our sample, a dominant early-type galaxy is not found at the center of the group potential.  This suggests that many of our groups are not dynamically evolved despite their high X-ray luminosities.  While similar systems have been identified at low redshift, the X-ray luminosities of the intermediate redshift examples are one to three orders of magnitude higher than those of their low redshift counterparts. We suggest that this may be evidence for group downsizing: while massive groups are still in the process of collapsing and virializing at intermediate redshifts, only low-mass groups are in the process of forming at the present day. ",
        "introduction": "Groups of galaxies constitute the most common galaxy associations, containing as many as 50--70\\% of all galaxies \\citep{tur76,gel83,eke06}. They are, therefore, an important laboratory for studying the processes associated with galaxy formation and evolution.  In recent years, optical and X-ray studies of groups at low redshift have provided new insights into these important systems. In particular, there are strong correlations between the morphological composition of the luminous galaxies, the velocity dispersion, and the presence of X-ray emission \\citep{zab98,mul98, mul03,osm04}.  Specifically, diffuse X-ray emission is found almost exclusively in those groups dominated by early-type galaxies.  In turn, the early-type fraction is strongly correlated with the group velocity dispersion and, thus, the group mass.  In the most luminous X-ray groups, the brightest group galaxy (BGG) is always a very massive elliptical, located at the peak of the X-ray emission \\citep{ebe94,mul96, mul98,hel00,mul03,osm04}.  As the peak of the X-ray emission is likely coincident with the center of the group, this implies that the BGG lies at the center of the group potential. Indeed, the position of the BGG is also indistinguishable from the center of the group potential as defined by the mean velocity and projected spatial distribution of the galaxies \\citep{zab98}. The fact that the BGG is located at the center of the potential suggests the formation of the BGG is intimately linked to the formation and evolution of the group itself.  Given their relatively low velocity dispersions, groups of galaxies provide ideal sites for galaxy-galaxy mergers \\citep{bar85,aar80,mer85,mil04,tay05,tem06}. This implies that significant changes in the star formation rates and morphological appearance of galaxies may be occurring in groups. To better understand how galaxies evolve in the group environment, groups must be observed over a wide range of cosmic time.  However, observations of groups at even moderate redshifts have been limited because of the difficulty of finding groups given their low galaxy densities.  \\citet{all93} photometrically selected a sample of groups at intermediate redshifts by targeting known radio galaxies. Their study suggested a progressive bluing in the galaxy population. Small samples of groups at higher redshifts have also been found in deep redshift surveys \\citep{lub98,coh00} and around lensed quasars \\citep{rus01,fas02,nai02,ray03,gra04,fau04,wil06}.  The recent completion of very large redshifts surveys now allow large group samples to be kinematically-defined from moderate redshifts up to z $\\sim$ 1 \\citep{car01,ger05}.  Wilman et al. (2005a,b) studied a large sample of groups at moderate redshifts selected from the CNOC2 survey and found that the fraction of group members undergoing significant star formation increases strongly with redshift out to z $\\sim$ 0.5. Therefore, there is evidence for some evolution in the group environment in the last $\\sim$ 5 billion years at least among optically-selected group samples.  X-ray emission from the hot intragroup medium provides another way to identify candidate groups at high redshifts. \\textit{ROSAT} was the first X-ray telescope capable of finding such systems and large numbers of group candidates at intermediate redshifts were found in deep surveys with this telescope \\citep{ros95,sca97,bur97,ros98,jon98,vik98,rom00,ada00,per02,jon02,bur03}. The ROSAT surveys suggest there is little or no evolution of the X-ray luminosity function of groups and poor clusters out to at least z=0.5. More recently, \\citet{wil05} arrived at a similar conclusion using twelve groups and clusters from the early data taken as part of the \\textit{XMM} Large-Scale Structure (LSS) Survey.  Upon completion, the XMM-LSS survey will provide a large sample of X-ray selected groups and poor clusters out to redshifts of z $\\sim$ 0.6 or higher.  In this paper, we provide the first results from an extensive multi-wavelength study of nine X-ray selected galaxy groups and poor clusters in the redshift range 0.2 $<$ z $<$ 0.6 selected from deep \\textit{ROSAT PSPC} pointings.  Our data allow the first detailed look at the morphological composition of X-ray groups at intermediate redshifts.  A detailed study of the X-ray properties of six of these systems based on \\textit{XMM-Newton} observations is provided in a companion paper (Jeltema et al. 2006; hereafter Paper II). We assume a $\\Lambda$ cold dark matter cosmology with $\\Omega$$_{m}$=0.27, $\\Lambda$=0.73, and H$_{\\rm o}$=70 km s$^{-1}$ Mpc$^{-1}$ throughout this paper. ",
        "conclusions": "\\subsection{Group Membership}  In most cases, the identification of the group in redshift space is trivial as the spatial distribution of galaxies on the sky coincide with the X-ray emission.  However, for a few of these fields, there are several different galaxy systems superposed along the line of sight.  This leads to some ambiguity about the true redshift of the X-ray system in two of the nine systems studied here. The RXJ1205.9+4429 field contains two significant galaxy systems, one at z $\\sim$ 0.35 and another z $\\sim$ 0.59. The \\textit{XMM-Newton} observation of this field shows that the X-ray emission is clearly centered on a luminous early-type galaxy that is part of the z=0.59 system (Paper II). Thus, we adopt this value as the redshift of this group. We note that the preliminary RDCS redshift corresponded to the lower redshift system. Therefore, the X-ray luminosity is actually considerably higher than originally reported in the RDCS. The true X-ray luminosity of this system is high enough that it does not meet our original selection criterion, suggesting this is likely a much richer system than the other objects in our sample. \\citet{ulm05} have recently published a detailed study of the X-ray and optical properties of this system and conclude that it is a fossil group. However, our spectroscopy and imaging data indicate that the magnitude difference between the brightest and second brightest confirmed member is $\\sim$ 1.2 mag. in the R-band. Thus, this system is not a fossil group by the standard definition usually adopted in the literature \\citep{jon03}.   In the case of RJX0341-4459 field, we measure redshifts for five galaxies in a system at z $\\sim$ 0.41. The five galaxies have a spatial distribution similar to the X-ray emission. However, there are also three foreground galaxies distributed over the same area. As these three objects have very different redshifts, they are not part of a single galaxy system. Thus, we believe the X-ray emission is most likely associated with the system at z $\\sim$ 0.41, although higher resolution X-ray images are required to be sure. Examples like this demonstrate the difficulty sometimes encountered when trying to identify low galaxy density systems (i.e. groups) at high redshift even when X-ray emission is present.  We determine group membership for each system using the ROSTAT package \\citep{bee90}. We start by considering all galaxies within $\\pm$3000 km s$^{-1}$ of the group's mean velocity. This is a large enough range to include all potential group members. We then calculate the biweight estimators of location (mean velocity) and scale (velocity dispersion). Objects with velocities greater than three times $\\sigma$$_{\\rm biwt}$ are then removed from the sample and a new mean location and scale are calculated. This process is repeated until there are no more objects to be clipped. This procedure resulted in the removal of one galaxy from three of the groups and none from the remainder. Figure 4 shows the velocity distributions of each member relative to the final mean velocity of the group. The final mean velocity and velocity dispersion are given in Table 2. For all of the systems studied here, the classical velocity dispersion (i.e.  $\\sigma_{\\rm Gauss}$, the Gaussian estimator) is in good agreement with the biweight velocity dispersion estimate. For approximately half of our sample, the velocity dispersions are based on only $\\sim$ 5 velocity measurements. These velocity dispersions are rather uncertain. Studies of low redshift X-ray groups suggest velocity dispersions based on a small number of galaxies can be significantly underestimated \\citep{zab98}.  \\subsection{The L$_{\\rm X}$-$\\sigma$ Relationship}   Figure 5 shows the L$_{\\rm X}$-$\\sigma$ relationship for our nine groups along with the sample of nearby groups from \\citet{osm04} and the moderate redshift X-ray groups from \\citet{wil05}. As can be seen from the figure, several of the groups fall significantly off the relationships found for nearby groups and clusters.  The two most deviant points in our sample correspond to the RXJ1334.0+3750 and RXJ1648.7+6019 groups.  These are the two groups from our \\textit{XMM-Newton} survey where the X-ray emission is not centered on an early-type BGG (Paper II).  In both cases, these groups have very low velocity dispersions for their given X-ray luminosities.  There are several possible explanations for why these groups fall so far off the relationships found for low redshift groups and clusters.  First, our velocity dispersions estimates for these groups may be artificially low because they are based on relatively small numbers (6 and 8 members, respectively).  \\citet{zab98} find that velocity dispersions estimated from the five brightest galaxies can be underestimated by as much as a factor of three. A similar factor would bring our two most deviant groups in agreement with the relationship found for nearby groups and clusters. This idea can be tested by obtaining more velocity measurements.  Second, the X-ray luminosities of these groups may have been enhanced or contaminated in some way. For example, the observed X-ray emission may be dominated by galaxy emission that is unresolved with our \\textit{XMM-Newton} observations. While we believe this is very unlikely given the high X-ray luminosities of our groups and the extent and morphology of the X-ray gas, we cannot rule this possibility out without higher resolution X-ray images. Thirdly, the velocity dispersions may have been reduced in some way. \\citet{hel05} have studied several nearby groups that fall off the L$_{\\rm X}$-$\\sigma$ relationship in a similar manner to our groups (although the X-ray luminosities of the nearby groups are nearly two orders of magnitude lower than the present examples). They suggest several physical mechanisms that could reduce the velocity dispersions including dynamical friction and tidal heating. They also suggest that orientation effects can lead to an artificially low observed velocity dispersion.  Finally, the low velocity dispersions could be an indication that these groups are in the process of collapsing for the first time and therefore the measured velocity dispersions do not yet accurately reflect the depth of the group potential (see Section 4.5).   \\subsection{Morphological Composition}  Studies of X-ray groups at low redshift have revealed a very strong tendency for these systems to be dominated by early-type galaxies \\citep{ebe94,pil95,hen95,mul96,zab98}. Table 2 lists the early-type fraction for our groups (based on the \\textit{HST} morphological classifications, where possible).  For all but one of our objects, the early-type fractions are in the range $\\sim$ 0.4--0.8. For the four groups with just four to six members known, these fractions could be somewhat over-estimated as our analysis is restricted to the brightest group members, which tend to be ellipticals \\citep{zab98}. However, even for the groups with many more members identified, the early-type fractions are comparable to those of rich clusters sampled out to similar radii \\citep{whi93}. Thus, the trend for X-ray groups to contain a large number of early-type galaxies appears to be in place out to at least z $\\sim$ 0.5. The one exception in our sample is the RXJ0210.4-3929 group, which based on the HST imaging is dominated by spiral galaxies. The large number of spirals in this system make it very unusual among X-ray groups at both low and moderate redshifts.  A correlation between early-type fraction and velocity dispersion has been noted for nearby group samples \\citep{hic88,zab98,osm04}, suggesting that galaxy morphology is related to the depth of the group potential. For groups with well-determined membership, the relationship is surprisingly robust \\citep{zab98}. In Figure 6, we plot these quantities for our sample along with the low redshift data from \\citet{zab98}. Among our moderate redshift groups, there is considerable scatter and no indication of a trend between early-type fraction and velocity dispersion. We suspect much of this scatter is an indication that neither quantity is well-determined for most of our groups. However, we note that the two groups in our sample with membership data comparable to that of the \\citet{zab98} sample (i.e. $\\sim$ 20 known members) do appear to follow the trend found at low redshift. In fact, these two groups suggest that the relationship found by \\citet{zab98} extends to the range of poor clusters. As noted by \\citet{zab98}, the relationship cannot be the same as for rich clusters as it would predict an unphysical early-type fraction for clusters with velocity dispersions above $\\sim$ 800 km s$^{-1}$.   \\subsection{The Brightest Group Galaxy}  Previous work on low redshift X-ray groups indicates that the X-ray emission is usually centered on a luminous elliptical galaxy \\citep{ebe94,mul98,hel00,mul03,osm04}.  In almost every case, this elliptical is the most luminous galaxy in the group.  As the peak in the X-ray emission is likely coincident with the center of the group potential, this implies that the brightest group galaxy (BGG) lies at the center of the potential.  Unfortunately, it is difficult to define the peak of the X-ray emission for the present sample from the low signal-to-noise \\textit{ROSAT} images.  However, six of the nine groups have now been observed by \\textit{XMM-Newton} and four of these are consistent with the X-ray peak being coincident with the brightest group elliptical (Paper II).  In all four groups with a central BGG, the radial velocity of the BGG is consistent with the mean velocity of the group within the velocity errors. Thus, similar to the case found for nearby X-ray groups \\citep{zab98}, the BGG is likely at or near the center of the group potential in these systems.  However, for three of the four groups where we find a dominant BGG, the central object appears to be composed of multiple nuclei (see Figure 7).  In the two most spectacular cases (RXJ0720.8+7109 and RXJ1256.0+2556), the central object has three components.  Although multiple nuclei in CD galaxies in clusters are fairly common \\citep{hoessel80,schneider83,lauer88}, in the majority of cases there is a large magnitude difference between the various components. In contrast, for both of our three nuclei systems, the second nuclei have R-band magnitudes within $\\sim$ 0.5 mag. of the brightest component.  In the case of the RXJ0720.8+7109 group, the second nucleus is the second brightest galaxy in the group.  Previous studies of multiple nuclei in clusters have found large velocity offsets between the various components in some cases, indicating these systems are not bound and are thus not in the process of merging \\citep{merritt84,Tonry85,smith85}.  For the RXJ0720.8+7109 group, we obtained a spectrum of the two brightest components and find a radial velocity difference of $\\sim$ 200 km s$^{-1}$. Given the typical errors on our velocity measurements ($\\sim$ 100 km s$^{-1}$), our data are consistent with a similar radial velocity for the two components. Thus, the two components may be bound. For the other triple system, RXJ1256.0+2556, we only obtained a velocity for the central component, so we cannot infer anything further about the nature of the multiple nuclei.  As noted above, in only four of the nine groups in our sample is the X-ray emission clearly centered on an early-type galaxy. In two of the other groups the most luminous galaxy is an elliptical, but the existing X-ray data are not sufficient to determine an X-ray center (RXJ0341.9-4459 and RXJ1347.9+0752).  Thus, we cannot draw strong conclusions for these two groups as to whether the X-ray emission is peaked on the dominant elliptical galaxy or not.  For both groups, the brightest elliptical is offset in velocity from the mean velocity of the group by several hundred kilometers per second. However, we have very few velocity measurements for both systems, so the mean velocity of the group is not well-determined and we cannot draw strong conclusions regarding a potential offset of the BGG from the group center.  However, the remaining three groups in our sample deviate strongly from the low redshift trend for there to be a dominant elliptical galaxy at the center of the group potential.  The RXJ1334.0+3750 group contains a dominant elliptical, but the peak of the X-ray emission is offset significantly from this galaxy (Paper II).  The velocity of this galaxy is consistent with the mean velocity of the group within the errors on each measurement. However, given the very low velocity dispersion of this system ($\\sigma$=121$^{+58}_{-45}$) and the small number of known members (6), we cannot draw strong conclusions with the existing velocity data.  The most luminous galaxy in the RXJ1648.7+6019 group is also an elliptical, although the group contains several galaxies of comparable luminosity.  Furthermore, the \\textit{XMM-Newton} data suggest the X-ray emission is not centered on any particular galaxy, but is instead distributed in a chain morphology similar to the distribution of galaxies near the group center (Paper II). The brightest elliptical is also offset in velocity from the mean velocity of the group by more than 200 km s$^{-1}$. This provides further evidence that this galaxy is not at the center of the group potential.  A chain-like morphology is also found in the RXJ0210.4-3929 group. In this case, the brightest galaxies in the chain are spirals. Unfortunately, we do not have \\textit{XMM-Newton} data for this system, so we cannot be sure where the X-ray emission peaks. Regardless, the most luminous early-type galaxy near the group center is nearly a magnitude fainter than the brightest spirals and has a velocity offset by nearly 400 km s$^{-1}$ from the mean velocity of the group. Thus, this group also lacks a dominant central early-type galaxy.  \\subsection{Evidence for Group Downsizing}  As discussed in the last section, at least three (and potentially five) of the nine groups in the present sample do not appear to have a central dominant early-type galaxy.  Furthermore, in three of the four groups where the X-ray emission is centered on a BGG, the central galaxy is not a single object, but rather is composed of multiple components. These observations suggest that most of the groups in our sample are not dynamically evolved. Instead, we appear to be catching them in the process of virialization.  The global X-ray properties of these groups are consistent with the properties of more dynamically relaxed systems, however (Paper II). This suggests that the X-ray properties of groups are already in place early in the formation of these systems.  Specifically, the intragroup medium properties appear to be largely set prior to the BGG experiencing its last major merger and settling at the center of the group potential.  This scenario would also explain the lack of evolution observed in the X-ray luminosity function of groups out to z $\\sim$ 0.5 \\citep{ros95,jon02,wil05} despite the morphological peculiarities we find over the same redshift interval.  If true, the temperature of the hot gas component may provide a better indication of the global group potential early on than the velocity dispersion of the galaxies. This might explain the very low velocity dispersions observed for the RXJ1334.0+3750 and RXJ1648.7+6019 groups: The X-ray temperatures of these systems reflect the massive group potentials, but the velocity dispersions of the galaxies do not yet accurately probe the group mass. Cosmological simulations of groups that include both the intragroup gas and galaxies may be able to test this idea.  The late assembly of the BGG in groups is consistent with the results of simulations in hierarchical cosmological models \\citep{dub98}. The groups in our sample appear to cover a range in dynamical state and can therefore provide some clues into the formation process of the BGG. The RXJ0210.4-3929 and RXJ1648.7+6019 groups do not yet contain a dominant early-type galaxy and thus they are likely at the earliest stages of group formation. The morphological compositions of these groups are very different, with RXJ0210.4-3929 consisting mostly of spirals and RXJ1648.7+6019 mostly of early-type galaxies. This suggests that both early and late type galaxies can be the dominant contributor to the final merger product. The groups with a multiple component BGG are likely much further along in the virialization process.  In fact, the BGG in these groups is probably undergoing its final major merger. Finally, only one of the groups in our sample is consistent with being a relaxed, virialized system (RXJ0329.0+0256). In this case, the BGG is at the center of the group potential as determined from both the X-ray emission and the velocity distribution of the group members and is unlikely to undergo any more major mergers.  The fact that many of our intermediate redshift groups do not have a dominant central elliptical is somewhat surprising given that such groups appear to be very rare among local X-ray group samples \\citep{mul03,osm04}.  One potential concern in comparing our intermediate redshift groups to local samples is that the best studied local samples were not selected in a similar manner. In fact, the largest ROSAT surveys of groups were performed with very heterogeneous samples of groups mostly drawn from optical catalogs \\citep{mul00,mah00,hel00,mul03,osm04}. To allow a better comparison to low redshift systems, we have selected a sample of nearby X-ray groups and poor clusters from two surveys based on the \\textit{ROSAT} All-Sky Survey: the NORAS \\citep{boh00} and REFLEX \\citep{boh04} group and cluster samples.  From each survey, we have selected all of groups and clusters with X-ray luminosities between $\\sim$ 2 $\\times$ $10^{42}~h_{70}^{-2}~{\\rm ergs~s^{-1}}$ and $\\sim$ 2 $\\times$ 10$^{43}~h_{70}^{-2}~{\\rm ergs~s^{-1}}$ in the \\textit{ROSAT} band (i.e. the corresponding selection criterion for our intermediate redshift sample) with z $\\le$ 0.05. Eliminating duplicate entries from the two catalogs produces a sample of 74 X-ray luminous groups and clusters. Unfortunately, the vast majority of these systems have not been previously studied in detail in either the optical or X-ray bandpasses. However, a literature search reveals that 19 of the 74 systems have deeper X-ray images published. The existing data for this subset of groups suggests these X-ray selected systems follow the trends found among the more heterogeneously selected nearby group samples. In particular, in all 19 of the groups, the X-ray emission is centered on a luminous early-type galaxy. Furthermore, we find no multiple-nuclei examples among the 19 nearby BGGs.  This suggest that the differences we find between our intermediate redshift systems and local samples are not the result of a selection effect. Rather, the intermediate redshift groups appear to be less dynamically evolved than present day luminous X-ray groups.  A closer examination of low redshift samples suggests that there are some local examples of X-ray groups without a central BGG \\citep{mul03,osm04,ras06}. However, the X-ray luminosities of these systems are one to three orders of magnitude lower than the X-ray luminosities of our moderate redshift groups.  Among the $\\sim$ 60 low redshift X-ray groups that have been studied in detail with \\textit{ROSAT}, the most luminous example of a system without a early-type BGG is the NGC 5171 group (L$_{\\rm X}$ $\\sim$ 3 $\\times$ 10$^{42}$ erg/s; \\citet{osm04}).  Thus, among the most X-ray luminous (L$_{\\rm X}$ $>$ 5 $\\times$ 10$^{42}$ erg s$^{-1}$) groups in the nearby universe, there appear to be no known counterparts to the systems we find at intermediate redshifts.  The failure to find nearby examples of such systems suggests that the most X-ray luminous groups have largely reached virialization by z $\\sim$ 0. This suggest that we are witnessing group downsizing: While the most luminous (and thus most massive) groups are still in the process of virializing at intermediate redshifts, this process is restricted to much less luminous (and thus less massive) systems at the present day."
    },
    "0606/astro-ph0606580_arXiv.txt": {
        "abstract": "Observational and theoretical evidence in support of metallicity dependent winds for Wolf-Rayet stars is considered. Well known differences in Wolf-Rayet subtype distributions in the Milky Way, LMC and SMC may be attributed to the sensitivity of subtypes to wind density. Implications for Wolf-Rayet stars at low metallicity include a hardening of ionizing flux distributions, an increased WR population due to reduced optical line fluxes, plus support for  the role  of single WR stars as Gamma Ray Burst progenitors. ",
        "introduction": "Wolf-Rayet (WR) stars represent the final phase in the evolution of very massive stars prior to core-collapse, in which the H-rich envelope has been stripped away via either stellar winds or close binary evolution, revealing products of H-burning (WN sequence) or He-burning (WC sequence) at their surfaces, i.e. He, N or C, O (Crowther 2007).  WR stellar winds are significantly denser than O stars, as illustrated in Fig.~\\ref{WRross}, so their visual spectra are dominated by broad emission lines, notably HeII $\\lambda$4686 (WN stars) and CIII $\\lambda$4647-51, CIII $\\lambda$5696, CIV $\\lambda$5801-12 (WC stars). The spectroscopic signature of WR stars may be seen individually in Local Group galaxies (e.g. Massey \\& Johnson 1998), within knots in local star forming galaxies (e.g. Hadfield \\& Crowther 2006) and in the average rest frame UV spectrum of Lyman Break Galaxies (Shapley et al. 2003).  In the case of a single massive star, the strength of stellar winds during the main sequence and blue supergiant phase scales with the metallicity (Vink et al. 2001). Consequently, one expects a higher threshold for the formation of WR stars at lower metallicity, and indeed the SMC shows a decreased number of WR to O stars than in the Solar Neighbourhood. Alternatively, the H-rich envelope may be removed during the Roche lobe overflow phase of close binary evolution, a process which is not expected to depend upon metallicity.   \\begin{figure} \\centerline{\\psfig{file=test.eps,width=4.0in}} \\caption{Comparisons between stellar radii at Rosseland optical depths of 20 (= $R_{\\ast}$, black) and 2/3 (= $R_{2/3}$, grey) for HD~66811 (O4\\,If), HD~96548 (WN8) and HD~164270 (WC9), shown to scale, together with the wind region corresponding to the primary optical wind line forming region, $10^{11} \\leq n_{e} \\leq 10^{12}$ cm$^{-3}$ (hatched) in each case, illustrating the highly extended winds of WR stars with respect to O stars (Crowther 2007).}\\label{WRross} \\end{figure}   WR stars represent the prime candidates for Type Ib/c core-collapse supernovae and long, soft Gamma Ray Bursts (GRBs).  This is due to their immediate progenitors being associated with young massive stellar populations, compact in nature and deficient in either hydrogen (Type Ib) or both hydrogen and helium (Type Ic).  For the case of GRBs, a number of which have been associated with Type Ic hypernovae (Galama et al. 1998; Hjorth et al. 2003), a rapidly rotating core is a requirement for the collapsar scenario in which the newly formed black hole accretes via an accretion disk (MacFadyen \\& Woosley 1999).  Indeed, WR populations have been observed within local GRB host galaxies (Hammer et al. 2006).  At solar metallicity, single star models predict that the core is spun down either during the red supergiant (via a magnetic dynamo) or Wolf-Rayet (via mass-loss)  phases. The tendency of GRBs to originate from metal-poor environments (e.g. Stanek et al. 2006) suggests that stellar winds from single stars play a role in their origin since Roche lobe overflow in a close binary evolution would not be expected to show a strong metallicity dependence.  In this article, evidence in favour of a metallicity dependence for WR stars is presented, of application to the observed WR subtype distribution in Local Group galaxies, plus properties of WR stars at low metallicity including their role as GRB progenitors.  \\begin{figure} \\centerline{\\psfig{file=wrpop.eps,width=4.8in}} \\caption{Subtype distribution of Milky Way ($<$3kpc), LMC and SMC WR stars, in which known binaries are shaded (Crowther 2007).}\\label{wrpop} \\end{figure} ",
        "conclusions": "Observational and theoretical evidence supports reduced wind densities and velocities for low metallicity WR stars, which addresses the relative WR subtype distribution in the Milky Way and Magellanic Clouds, plus the reduced WR line strengths in the SMC with regard to the Galaxy and LMC. The primary impact at low metallicity is as follows; (a) an increased WR population due to lower line fluxes from individual stars, of particular relevance to I~Zw\\,18 and SBS0335-052E; (b) harder ionizing fluxes from WR stars, potentially responsible for the strong nebular HeII $\\lambda$4686 seen in low metallicity HII galaxies; (c) responsible for the reduced density of GRB environments with respect to Solar metallicity WR counterparts.  Finally, a metallicity dependence for WR winds may help to reconcile the relative number of WN to WC stars observed in surveys (e.g. Massey \\& Johnson 1998) with evolutionary predictions. Evolutionary models for which rotational mixing is included yet metallicity dependent WR winds are not (Meynet \\& Maeder 2003) fail to predict the high N(WC)/N(WN) ratio observed at high metallicities (Hadfield et al. 2005), whilst models which account for the Vink \\& de Koter (2005) WR wind dependence compare much more favourably with observations (Eldridge \\& Vink 2006), in spite of the neglect of rotational mixing."
    },
    "0606/astro-ph0606549_arXiv.txt": {
        "abstract": "It is shown that in strongly magnetized neutron stars, there exist upper limits of magnetic field strength, beyond which the self energies for both neutron and proton components of neutron star matter become complex in nature. As a consequence they decay within the strong interaction time scale. However, in the ultra-strong magnetic field case, when the zeroth Landau level is only occupied by protons, the system again becomes stable against strong decay. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606255_arXiv.txt": {
        "abstract": "A newly identified kiloparsec-scale X-ray jet in the high-redshift $z$=3.89 quasar 1745+624 is studied with multi-frequency Very Large Array, Hubble Space Telescope, and {\\it Chandra} X-ray imaging data.  This is only the third large-scale X-ray jet beyond $z>$ 3 known and is further distinguished as being the most luminous relativistic jet observed at any redshift, exceeding $10^{45}$ erg/s in both the radio and X-ray bands. Apart from the jet's extreme redshift, luminosity, and high inferred equipartition magnetic field (in comparison to local analogues), its basic properties such as X-ray/radio morphology and radio polarization are similar to lower-redshift examples.  Its resolved linear structure and the convex broad-band spectral energy distributions of three distinct knots are also a common feature among known powerful X-ray jets at lower-redshift. Relativistically beamed inverse Compton and `non-standard' synchrotron models have been considered to account for such excess X-ray emission in other jets; both models are applicable to this high-redshift example but with differing requirements for the underlying jet physical properties, such as velocity, energetics, and electron acceleration processes. One potentially very important distinguishing characteristic between the two models is their strongly diverging predictions for the X-ray/radio emission with increasing redshift. This is considered, though with the limited sample of three $z>$ 3 jets it is apparent that future studies targeted at very high-redshift jets are required for further elucidation of this issue. Finally, from the broad-band jet emission we estimate the jet kinetic power to be no less than $10^{46}$ erg/s, which is about 10$\\%$ of the Eddington luminosity corresponding to this galaxy's central supermassive black hole mass ${\\cal M}_{\\rm BH}$ \\simgt $10^9 \\, {\\cal M}_{\\odot}$ estimated here via the virial relation.  The optical luminosity of the quasar core is about ten times over Eddington, hence the inferred jet power seems to be much less than that available from mass accretion. The apparent super-Eddington accretion rate may however suggest contribution of the unresolved jet emission to the observed optical flux of the nucleus. ",
        "introduction": "}   \\begin{figure*} \\epsscale{0.75} \\plotone{f1.eps} \\figcaption[f1.eps]{\\label{figure-1} VLBI 2.3 GHz map of the parsec-scale jet in quasar 1745+624 from averaging 6 images from the USNO database (4 mas circular beam plotted on bottom left).  The quasar core is the peak (239.0 mJy/bm) toward the upper left. The lowest contour level is 0.3 mJy/bm (2 times the measured rms in the image) increasing by factors of $\\sqrt{2}$. } \\end{figure*}   \\begin{figure*} \\epsscale{1.75} \\plotone{f2.eps} \\figcaption[f2.eps]{\\label{figure-2} Multi-frequency VLA images of 1745+624. The $\\sim$2.5$\\arcsec$ long radio jet extending to the southwest of the core (placed at the origin) is seen clearly. The lowest contour levels are 0.60, 0.33, 0.40, and 0.45 mJy/beam (3 times the measured off source rms in the images), at 1.5, 4.9, 8.5, and 15 GHz, respectively. The positive levels (solid contours) are spaced by factors of $\\sqrt{2}$ up to the image peaks of 484, 441, 464, and 577 mJy/beam. The elliptical restoring (naturally weighted) beams are plotted at the bottom left corner: their dimensions are 0.530\\arcsec$\\times$0.340\\arcsec\\ at PA=--72$^{\\circ}$, 0.490\\arcsec$\\times$0.305\\arcsec\\ at PA=--67$^{\\circ}$, 0.255\\arcsec$\\times$0.204\\arcsec\\ at PA=--29$^{\\circ}$, and 0.182\\arcsec$\\times$0.128\\arcsec\\ at PA=--74$^{\\circ}$.} \\end{figure*}  \\begin{figure*} \\plotone{f3.eps} \\figcaption[f3.eps]{\\label{figure-3} VLA 4.9 GHz total intensity [I; left] and polarized intensity [P; right] images of 1745+624 at 0.35\\arcsec\\ resolution (beam plotted in bottom left). Contour levels begin at 0.35 (I) and 0.16 (P) mJy/bm, and increase by factors of $\\sqrt{2}$ up to peaks of 441.4 (I) and 18.1 (P) mJy/bm. The tick marks show the orientation of the electric vector position angles, with a correction of --6\\deg\\ applied corresponding to the integrated rotation measure of 28.4$\\pm$0.5 radians m$^{-2}$ \\citep{ore95}. Selected fractional polarization levels are indicated.} \\end{figure*}   Quasars form a class of objects that frequently warrant superlatives in their descriptions:  1745+624 (4C+62.29) is no exception. It was identified as one of the highest redshift X-ray quasars at the time of its discovery \\citep[$z$=3.87;][]{bec92} from spectroscopic followup of statistically significant sources from the {\\it Einstein} X-ray Observatory. Further optical spectroscopy by \\citet{sti93} refined the redshift to $z$=3.89, affirmed by \\citet{hoo95}; the latter value is adopted here. After the {\\it Einstein} detection, it was established as a bona fide X-ray source by {\\it ROSAT} \\citep{fin93}, {\\it ASCA} \\citep{kub97}, and {\\it BeppoSAX} \\citep{don05}.  By virtue of its high-redshift, it is one of the most radio-luminous quasars known \\citep[cf. Figure~1 of][]{jes03} -- its observed 1.4 GHz luminosity\\footnote{We adopt H$_{0}=71~$km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$ and $\\Omega_{\\rm \\Lambda}=0.73$ and have converted quoted literature values of luminosities and proper motions to this cosmology. If we assume the cosmology adopted in \\citet{jes03}, $L_{\\rm 1.4~GHz}=10^{28.4}$ W Hz$^{-1}$ (1 Watt = 10$^{7}$ erg s$^{-1}$).}, $L_{\\rm 1.4 GHz}=10^{29}$ W Hz$^{-1}$, is over an order of magnitude greater than that of the archetype luminous quasar 3C~273 \\citep{con93}. It is probably no coincidence that its 2.5\\arcsec\\ long jet \\citep[18 kpc projected;][]{bec92} is also the most luminous radio jet currently known \\citep[that we are aware of; cf. ][]{liu02,che05}, accounting for $\\sim$1/4th of the source's total 1.4 GHz luminosity.  This quasar is of interest to us because of this prominent radio jet, visible on both milli-arcsecond \\citep[][Figure~\\ref{figure-1}]{tay94,fey00} and arcsecond-scales (Figures~\\ref{figure-2} \\& \\ref{figure-3}).  We were prompted to identify such high-redshift systems with large-scale jets after the {\\it Chandra} X-ray Observatory detection of a $\\sim$2.5\\arcsec\\ long X-ray extension \\citep{sie03,yua03,che04} in the $z$=4.3 quasar GB~1508+5714 \\citep{hoo95}. Subsequent VLBA imaging of GB~1508+5714 has revealed a parsec-scale jet that can be traced out to $\\sim$100 milli-arcseconds \\citep[0.7 kpc projected;][]{che06}, aimed in the general direction of the kpc-scale structure, supporting the jet interpretation. These are the two most distant quasars with a {\\it kiloparsec-scale} jet\\footnote{Adopting the conventional requirement of a jet being at least 4 times longer than it is wide \\citep{bri84}. A single feature associated with a jet was discovered in the more distant $z$=5.47 blazar Q0906+6930 by \\citet{rom04}, but on VLBI scales only.} detected at any wavelength \\citep{che05}.  The identification of such systems in the early Universe is extremely interesting for a number of reasons. Jets are signposts for ``active'' black hole/accretion disk systems \\citep[e.g.,][]{beg84}, requiring central nuclei of high-$z$ galaxies to be sufficiently well-developed to sustain gravitational collapse. At such early epochs, it is remarkable that the jet production process is apparently efficient and persists long enough ($>$Myr) to have produced such large (10's--100's kpc-scale) structures. In this context, jets may even serve as tracers of the directionality of the BH/disk axis in these early accreting systems, as chronicled in the ``bumps and wiggles'' in our jet images.  Also, these jets are an efficient means to shock heat ambient gas thus triggering early star formation \\citep[e.g.,][]{ree89}.  Among other important issues, such high-redshift systems are potentially very important to our understanding of the emission processes responsible for the production of the broad-band jet emission, and the X-rays in particular. Although the number of quasars with kpc-scale jets detected in the X-rays is growing \\citep{har06}\\footnote{See a current on-line census at:  http://hea-www.harvard.edu/XJET/}, the origin of this radiation is still in active debate.  \\citet{sch02} noted early on that a comparison of cosmologically distant jets with local examples could in principle provide crucial arguments to the debate. This is because the two main contending models for the X-ray emission of large-scale quasar jets -- synchrotron radiation and inverse Compton scattering of the cosmic microwave background (IC/CMB) -- have, in their simplest versions, markedly different predictions of the X-ray jet properties with redshift (see \\S~\\ref{section-z}).  A high-resolution {\\it Chandra} image of the $z$=3.89 quasar 1745+624 was obtained in an early observing cycle (Table~\\ref{table-1}) and X-ray emission from the arcsecond-scale radio jet is in fact detected and is quite prominent also. Apart from the GB~1508+5714 case mentioned above, evidence has been presented recently for at least one more high-$z$ quasar with extended X-ray emission \\citep[J2219--2719 at $z$=3.63;][]{lop06} with a radio counterpart \\citep[][and manuscript in preparation]{che05}. This makes 1745+624 one of only three very high-redshift ($z>$3) systems with extended X-rays associated with a radio jet.  With the scarcity of such high-redshift systems and their importance in current discussions of X-ray jet emission models, we have independently analyzed the archival {\\it Chandra} data. It was also realized that 1745+624 hosts an exceptional jet making it conducive to study. This is because while both the GB~1508+5714 and J2219--2719 cases each show a single faint ($\\sim$1 mJy at 1.4 GHz) arcsecond-scale feature separated from their bright nuclei, the 1745+624 jet is 2 orders of magnitude brighter than the other two cases and shows {\\it resolved structure}. In addition, it is the most luminous X-ray jet observed thus far \\citep[cf.,][]{har06}.  To fully exploit the potential of the {\\it Chandra} observation, we have also analyzed archival multi-frequency imaging data from the Very Large Array (VLA), and a Hubble Space Telescope (HST) image to map the spectral energy distributions of the extended components in the quasar jet. These data, along with a new deep very long baseline interferometry (VLBI) map of the extended subarcsecond-scale jet are described in \\S\\S~\\ref{section-vlbi}--\\ref{section-chandra}, with a critical assessment of the derived spectra of the kpc-scale components in \\S~\\ref{section-spectra}. The physical properties of these extended components are then discussed (\\S~\\ref{section-discuss}), with comparisons made to the other two quasars with X-ray detected knots at comparably high-redshift ($z\\simgt$3). These distant examples are compared to other currently known X-ray jets in quasars at lower redshift ($z\\simlt$2) in light of expectations from synchrotron and inverse Compton emission models. Our results are summarized in \\S~\\ref{section-summary}. ",
        "conclusions": "}  At $z$=3.89, the 1745+624 system displays a fully formed jet as observed by us when the Universe was only about 12$\\%$ of its present age.  For its projected length $L = 18$ kpc, and assuming a constant hot spot advance velocity $\\beta_{\\rm adv} \\sim 0.1-1$ \\citep[see in this context][but also comments at the end of \\S~\\ref{section-hotspot}]{sch95}, the lifetime of the radio source is roughly $t_{\\rm life} \\sim L/ (\\beta_{\\rm adv} c \\, \\sin \\theta) \\sim$ Myr (the jet viewing angle $\\theta \\sim 10\\deg$ is taken here for illustration; see below). This structure signals long-term active accretion at a very early cosmic epoch, along with a roughly stable black hole/accretion disk axis \\citep[e.g.,][]{beg84}.  Such large-scale jets (deprojected $\\sim$100 kpc in this case) at high-redshifts are not necessarily rare, though they have not been systematically searched for and studied. \\citet[][and manuscript in preparation]{che05} has only recently identified $\\sim$10 other jets on a comparable scale in a VLA study of a sample of $z>$3.4 (out to $z$=4.7) flat-spectrum-core radio sources; 1745+624 remains the most radio-luminous example observed (\\S~\\ref{section-intro}).  Despite the extreme luminosity of the jet in 1745+624, its extended radio and X-ray emissions appear quite similar to other more local examples in regard to its morphology, spectral energy distribution, and radio polarization.  This is quite surprising as higher-$z$ jets have propagated through the (supposedly different) intergalactic medium of the early Universe \\citep[e.g.,][]{ree89}. In the following, we first consider available evidence for relativistic beaming in the 1745+624 jet on both parsec and kpc-scales (\\S~\\ref{section-describe}).  Then, we discuss the emission properties of the kpc-scale jet (\\S~\\ref{section-discussjet}) and the hot spot (\\S~\\ref{section-hotspot}), confronting the multiwavelength data with expectations from different emission models. Along the way, we make comparisons to other X-ray jets at low and high redshift. Finally, the broad-band data allow us to estimate the energetics of the jet, which we compare to the accretion parameters of the quasar nucleus (\\S~\\ref{section-energetic}).   \\subsection{Radio Constraints on Relativistic Beaming of the Jet\\label{section-describe}}  Several lines of evidence suggest that relativistic beaming in the 1745+624 jet is significant on the observed scales.  Both the radio (VLBI and VLA-scale) and X-ray structures appear one-sided to our detection limits.  The best constraints on the sidedness comes from the presented 2.3 GHz USNO map (jet to counter-jet ratio $>$50; \\S~\\ref{section-vlbi}) and proper motion studies with VLBI. Twelve epochs of higher resolution 8 GHz data from the same USNO program, obtained over $\\sim$2.5 years \\citep{pin04}, show a basically stationary inner ($\\sim$1--1.5 mas distant) component, although velocities of a few $c$ are allowed by the formal fit (B.G. Piner, 2005, private communication).  Maps from the Caltech-Jodrell Bank survey obtained over a longer time baseline at 5 GHz reveal three superluminal components presumably further down the jet \\citep[][this is the highest redshift object in the sample -- see Figure~1 in the latter]{tay94,ver03}. Converting to our adopted cosmology, their proper motions correspond to apparent speeds ranging from $\\sim$6$c$ up to $\\sim$16$c$. This restricts the maximum angle of the small-scale jet to our line of sight to be $<$7--19\\deg.  The superluminal VLBI jet is oriented at PA$\\sim$205--208\\deg\\ in the sky \\citep{tay94}, increasing gradually to $\\sim$216\\deg\\ out to $\\sim$110 mas (Figure~\\ref{figure-1}; \\S~\\ref{section-vlbi}). This further increases out to the visible arcsecond-scale structure (the hot spot PA=227\\deg).  The maximal (projected) changes in position angle between the different features is 6\\deg\\ (cf. Table~\\ref{table-2}). The required intrinsic bend in the jet to account for these changing PAs is $\\sim$6\\deg\\ $\\times \\sin(20\\deg)$$\\sim$2\\deg, or even less (supposing the jet inclination $<$20\\deg\\ as indicated by the proper motions).  Since we can trace the VLBI-scale emission out to the observed {\\it Chandra} scales, the VLBI proper motions (and only small intrinsic bends) constrain the kiloparsec-scale jet to be aligned within $\\simlt$25\\deg\\ to our line-of-sight, with a likely range of $\\sim$10--20\\deg, or less. Radio asymmetry studies of large scale quasar jets require jets like this to be relativistic \\citep[$\\Gamma$$>$1, with strict upper limits of $\\Gamma$$<$2--3;][]{war97}; these constraints have bearing on the following discussion of the origin of the X-ray emission in the 1745+624 jet.  \\subsection{X-ray Emission Mechanisms in the Large-scale Jet\\label{section-discussjet}}  The absence of optical emission and the steep radio spectra of the jet knots in 1745+624 preclude a straightforward power-law extrapolation of the low energy (synchrotron) emission into the X-ray band (\\S~\\ref{section-spectra}). The resultant convex broad-band spectral energy distribution (SED; Figure~\\ref{figure-7}) is a common characteristic of X-ray emitting knots in the jets of powerful quasars \\citep[e.g.,][]{schwartz00,sam04,mar05,har06}. As widely discussed, (unbeamed) inverse Compton emission also can not account for the observed excess of X-rays, and it is evident that this applies also to the present example at very high-redshift (\\S~\\ref{section-basic}). We explore the hypothesis that the bright X-ray emission is produced by {\\it relativistically beamed} inverse-Compton scattering of the CMB photons \\citep{tav00,cel01} as advocated by many workers. The main feature of this model is that in order to fulfill energy equipartition, significant beaming of the large-scale quasar jet emission is required (\\S~\\ref{section-beamed}).  ``Non-standard'' synchrotron models have also been proposed to reproduce the overall convex radio-to-X-ray SEDs and we discuss this possibility in \\S~\\ref{section-synchrotron}. The main difference between these models is the expected diverging redshift-dependence of the X-ray emission, and we discuss issues and prospects for distinguishing such a dependence from the observations (\\S~\\ref{section-z}).   \\subsubsection{Basic Considerations\\label{section-basic}}  Generally, one can calculate the magnetic field in a synchrotron emitting radio source if we assume magnetic field energy equipartition with the radiating particles (corresponding to the minimum total energy of the system). In the 1745+624 radio jet, we use the observed radio properties of K1.4 and K1.8 (assuming negligible contribution from relativistic protons and no beaming, $\\delta$=1) to calculate fairly high values of $B_{\\rm eq, \\, \\delta=1} \\approx 180$ $\\mu$G and $200$ $\\mu$G, respectively.  These values are more than an order of magnitude higher (due to the exceptionally luminous nature of this radio jet; \\S~\\ref{section-intro}) than those computed in lower-redshift quasar jets \\citep{kat05}, but are comparable to values inferred for lower power FRI radio galaxies \\citep{sta06}. The corresponding magnetic field energy densities of the knots, $U_{\\rm B} = B_{\\rm eq, \\, \\delta=1}^2/ 8 \\pi \\approx 1.3-1.6 \\times 10^{-9}$ erg\\,cm$^{-3}$, are $\\sim$6--7 times greater than the CMB energy density at the quasar's redshift, $U_{\\rm CMB} = 4 \\times 10^{-13} \\, (1+z)^4$ erg\\,cm$^{-3}$ $\\approx 2.3 \\times 10^{-10}$ erg\\,cm$^{-3}$. The latter quantity is, however, comparable to the energy densities of the knots' synchrotron photons (in the case of sub-relativistic jet velocity),  \\begin{equation} U_{\\rm syn} \\approx {L_{\\rm syn} \\over 4 \\pi \\, R^2 \\, c} , \\label{equation-usyn} \\end{equation}  \\noindent namely, $U_{\\rm syn} \\approx$ 2.8 and 3.8 $\\times 10^{-10}$ erg\\,cm$^{-3}$ for K1.4 and K1.8, respectively. In this estimate, we approximated the knots as spheres with radii, $R = 0.2\\arcsec / 2 \\approx 0.7$ kpc (\\S~\\ref{section-radio}), and took the synchrotron luminosities, $L_{\\rm syn} \\approx 10 \\times L_{\\rm 5 \\, GHz} \\approx 5-7 \\times 10^{44}$ erg\\,s$^{-1}$, as appropriate for the radio continuum extending over a few decades in frequency with spectral index $\\alpha_{\\rm r} \\approx 1$.  For the spectral index $\\alpha_{\\rm r} \\approx 1$, the expected X-ray flux densities ($F_{\\rm X}$) from the synchrotron self-Compton (SSC) plus IC/CMB processes are simply:  \\begin{equation} F_{\\rm X}^{\\rm ssc} + F_{\\rm X}^{\\rm IC/CMB} \\approx \\left({\\nu_{\\rm r} \\over \\nu_{\\rm X}}\\right) \\, \\left({U_{\\rm syn} + U_{\\rm CMB} \\over U_{\\rm B}}\\right) \\, F_{\\rm r}. \\label{equation-ssc}\\end{equation}  \\noindent Using the 5 GHz radio flux measurements ($F_{\\rm r}$), this severely underpredicts ($\\sim$0.054 nJy (K1.4) and $\\sim$0.074 nJy (K1.8) at 1 keV) the observed X-ray jet emission by factors of 87 and 42 in knots K1.4 and K1.8, respectively (Table~\\ref{table-2}).   \\subsubsection{Relativistic Beamed X-ray Emission?\\label{section-beamed}}  The bulk velocity of kpc-scale quasar jets are likely to be at least mildly relativistic \\citep[$\\Gamma > 1$, ][]{bri94b,war97}. This can significantly influence the expected inverse Compton X-ray fluxes calculated above. Specifically, relativistic beaming ($\\delta > 1$) increases the IC/CMB emission, and decreases the SSC one.  In 1745+624, the superluminal VLBI jet can be traced out to kpc-scales (\\S~\\ref{section-describe}), so the beamed IC/CMB X-ray emission will dominate over that produced by SSC in the visible jet. Assuming energy equipartition as an additional constraint, the expected IC/CMB flux is (for $\\alpha\\approx$1):  \\begin{equation} F_{\\rm X}^{\\rm IC/CMB} \\approx \\left({\\delta \\over \\Gamma}\\right)^2 \\left({\\nu_{\\rm r} \\over \\nu_{\\rm X}}\\right) \\, \\left({\\Gamma^2 \\, U_{\\rm CMB} \\over \\delta^{-10/7} \\, U_{\\rm B}}\\right) \\, F_{\\rm r} \\, \\, \\, \\propto \\, \\, \\, \\delta^{24/7}, \\label{equation-iccmb} \\end{equation}  \\noindent since in the jet rest frame (denoted by primes) $U'_{\\rm CMB} \\approx \\Gamma^2 \\, U_{\\rm CMB}$, and the equipartition magnetic field $B'_{\\rm eq} = B_{\\rm eq, \\, \\delta=1} \\, \\delta^{-5/7}$. Here, $\\Gamma$ is the jet bulk Lorentz factor, and $U_B = B_{\\rm eq, \\, \\delta=1}^2 / 8 \\pi$. This gives the required $\\delta = 4.6$ and 3.9 for the knots K1.4 and K1.8, respectively, to explain the observed X-rays via inverse Compton emission.  In this interpretation, the (small) decrease of the required Doppler factor along the jet is a direct consequence of the fading X-rays with the corresponding brightening of the radio emission (Figure~\\ref{figure-5}; cf. similar multi-wavelength morphologies in other well-studied jets like, 3C~273 \\citep{mar01,sam01}, 0827+243 \\citep{jor04}, PKS~1127--145 \\citep{sie02}, and PKS~1136--135 \\citep{sam06}).  Although there is evidence that the 1745+624 jet is highly relativistic on parsec scales (\\S~\\ref{section-describe}), there are few independent constraints (the one-sidedness of the radio and X-ray jets) to directly support the fairly high Doppler factors on these large scales, as implied in this interpretation. The inferred Doppler factors require minimum Lorentz factors, $\\Gamma\\geq (\\delta + \\delta^{-1})/2$ = 2.4 (K1.4) and 2.1 (K1.8). These approach the strict upper limits on the speeds of large-scale quasar jets deduced from radio asymmetry studies \\citep[$\\Gamma$$<$2--3;][]{war97}. If we assume the 1745+624 jet flow is moving near these maximal values, the derived Doppler factors require that the jet is inclined within $\\simlt$12--15\\deg\\ to our line of sight, in agreement with the values inferred in \\S~\\ref{section-describe}.  Given the discussed Lorentz factors of $\\sim$2--3, the Doppler factor is fairly insensitive to jet angles considered; for instance, $\\delta$=5.5--3.6 ($\\Gamma$=3) and $\\delta$=4.0--3.2 ($\\Gamma$=2.1) for jet angles, $\\theta$=0\\deg--15\\deg. Thus we can {\\it just} plausibly explain the level of X-ray jet emission in this jet via the IC/CMB mechanism if we push the speeds to be consistent with upper limits inferred from previous studies, together with the small viewing angles.  It should be stressed that these estimates assumed a single-zone emitting region, and there are generally more free parameters in both Compton and synchrotron models invoking substructure in the jet \\citep[e.g.,][]{cel01,sta02,sie06}.  Related to the beaming issue, we should remark that the \\emph{observed} radio power of the 1745+624 jet (excluding the hot spot region) is very large, $L_{\\rm 5 \\, GHz} \\approx 10^{44}$ erg\\,s$^{-1}$, suggesting its synchrotron luminosity of order $\\sim 10^{45}$ erg\\,s$^{-1}$ or greater. Such enormous radio powers from \\emph{large-scale jets} (exclusive from extended radio lobes) are quite extreme, so invoking even moderate beaming factors would reduce the intrinsic jet frame luminosity (by 1--2 orders of magnitude) to more comfortable levels.  This relativistically beamed IC/CMB model also requires that there is a sufficient number density of low energy electrons to upscatter CMB photons into the observed {\\it Chandra} band.  Radio emission would be radiated by these electrons at very low-frequencies (10's MHz in this case) and subarcsecond-resolution radio imaging capabilities at these energies are not widely available.  Curiously, \\citet{sti93} noticed that the {\\it integrated} radio spectrum of 1745+624 contains a steep-spectrum excess at low-frequencies; this quasar would be catalogued as a steep-spectrum radio quasar in low-frequency surveys and/or if it were local (the $1+z$ shift in observed frequencies). In our own compilation of radio data from the literature (Figure~\\ref{figure-7}), we find $\\alpha$=0.55$\\pm$0.01 from 38 MHz to 1.4 GHz and $\\alpha$$\\sim$0 at the higher frequencies where our imaging observations were taken. If we extrapolate the observed individual spectra of the (dominant $\\alpha$$\\sim$0) core and steep-spectrum extended component to lower frequencies, it implies that the latter dominates the total low-frequency emission. The well-studied $z$=0.158 quasar 3C~273 shows a similar low-frequency excess apparent down to $\\sim$10's MHz \\citep{cou98} and resolved low-frequency imaging shows its extended X-ray emitting \\citep[e.g.,][]{mar01,sam01,jes06} radio jet does indeed dominate the total flux at low frequencies \\citep[cf. figure~A1 in][]{con93}.  Further, although this low-frequency emission in 1745+624 is probably dominated by the (brighter and steeper spectrum at cm-wavelengths;  \\S~\\ref{section-hotspot}) hot spot, this does not preclude a low-frequency contribution from the jet. Interestingly, the X-ray spectral index of the jet (Table~\\ref{table-2}) matches the low-frequency integrated radio one, as would be expected in the IC/CMB scenario.  Future low-frequency imaging observations can clarify this by showing the relative contributions of the jet and hot spot in this and similar systems \\citep[see][for future prospects]{harris06}.  \\subsubsection{Synchrotron Models for the X-ray Emission\\label{section-synchrotron}}  In lieu of the beamed IC/CMB interpretation, synchrotron models with non-standard and/or multiple electron components are able to account for the ``excess'' X-ray emission \\citep{der02,sta02,sta04,jor04}. These models are attractive because in nearby low-power FRI jets, it is widely believed that synchrotron X-ray emission is produced \\citep[e.g.,][]{har06}, although they do not show the same convex SEDs, and emit at much smaller (few kpc) scales than the discussed quasar jets (10's--100's kpc scales). The challenge remains to explain the strong upturn in the X-ray band (i.e., the small values of $\\alpha_{\\rm ox}$) observed in powerful jets such as in 1745+624.  While it is true that an increased energy density of the CMB photons at $z > 1$ will result in a stronger radiative cooling of the ultra-relativistic electrons, the maximum electron energy available assuming efficient (although realistic) continuous acceleration process is still high enough to allow for production of an appreciable level of keV synchrotron photons.  In the case of the 1745+624 jet, the comoving energy density of the magnetic field (by assumption close to equipartition with the radiating electrons, as given above) is lower than the comoving energy density of the CMB photons, $U'_{\\rm B} < U'_{\\rm cmb}$, if only $\\delta^{5/7} \\Gamma > 2.5$. Thus, assuming even very moderate bulk velocity and beaming, the jet electrons cool mainly by inverse-Comptonization of the CMB photon field.  In the equipartition derived magnetic field of $B$=200 $\\mu{\\rm G}$, the electrons emitting the highest energy photons detected ($\\sim$6 keV, observed) have electron energies, $\\gamma = E_{\\rm e} / m_{\\rm e} c^2 = (\\nu_{\\rm Hz} (1+z) / \\delta B_{\\rm \\mu{\\rm G}})^{1/2}\\sim 10^8$. Since their cooling are dominated by inverse Compton losses, the appropriate timescale can be evaluated roughly as $t_{\\rm cool}' = 3\\,m_e\\,c / 4 \\sigma_{\\rm T} \\gamma U_{\\rm CMB}' \\sim 4.4$ yrs. This evaluation is rough, because in fact the considered high energy electrons are expected to radiate in a transition between Thomson and Klein-Nishina cooling regimes, since  \\begin{equation} \\left({\\Gamma \\, \\varepsilon_{\\rm cmb} \\over m_{\\rm e} c^2}\\right) \\, \\left({E_{\\rm keV} \\over m_{\\rm e} c^2}\\right) \\sim 1 \\, , \\end{equation}  \\noindent with the energy of the CMB photon $\\varepsilon_{\\rm cmb} \\approx (1+z)$ milli-eV, and the jet bulk Lorentz factor $\\Gamma$ of the order of a few \\citep[see in this context][]{der02}. For such parameters and the considered electron energies, the `optimistic' electron acceleration timescale $t'_{\\rm acc} \\sim \\zeta \\, r_{\\rm g} / c \\sim 0.01$ yr, where $r_{\\rm g}$ is the electron gyroradius and $\\zeta \\sim 10$ is the efficiency factor depending on the jet plasma conditions \\citep[see][]{aha02}. This is still much shorter than the timescale for radiative losses, while the latter is order of magnitudes longer than the timescale for the electron escape from the emission region.  The postulated stochastic acceleration mechanism can produce a flat-spectrum high-energy electron component, with the total energy density not exceeding the magnetic field energy density \\citep[see][for a discussion]{sta02}. The \\emph{observed} synchrotron spectrum of such particles (from the unresolved emitting region) is expected to show a spectral index modified (increased) by the radiative losses. With the condition of continuous injection of freshly accelerated electrons (characterized by a power-law energy spectrum $n'(\\gamma) \\propto \\gamma^{-p}$), such a spectral index would be $\\alpha = 0.5$ for any electron index $p < 2$ \\citep[see][]{hea87}, which is very close to the observed X-ray spectral index ($\\alpha_{\\rm X} \\approx 0.62^{+0.16}_{-0.17}$) of the 1745+624 jet. Moreover, the total luminosity of such a high-energy synchrotron component, if indeed limited only by the energy equipartition requirement, should be comparable to or less than the luminosity of the low-energy (radio) synchrotron component. That is in fact the case, since the \\emph{total} $1$ keV luminosity of the 1745+624 jet (excluding the terminal feature discussed in \\S~\\ref{section-hotspot}) is $\\sim 3 \\times 10^{45}$ erg s$^{-1}$ while, as we argue in \\S~\\ref{section-basic} above, the radio jet luminosity is $> 10^{45}$ erg s$^{-1}$.  Note, that in the framework of this synchrotron interpretation, the X-ray-to-radio jet luminosity ratio is not expected to change systematically with redshift in contrast to expectations in an IC/CMB interpretation, and this issue is discussed in the next subsection.  \\subsubsection{Redshift Dependence of the X-ray Emission?\\label{section-z}}  The simplest versions of synchrotron and IC/CMB models give significantly different predictions for the X-ray jet emission at high redshift because of the strong increase of the CMB energy density $\\propto (1+z)^4$ \\citep[e.g.,][]{sch02}.  With all else equal, they predict divergences in the observed X-ray to radio monochromatic luminosity ($f_{\\nu}=\\nu\\,F_{\\nu}$) ratio, $f_{\\rm x}$/$f_{\\rm r}$ exceeding a factor of 100 or more at $z$$>$2 (cf. Equation~\\ref{equation-iccmb}). This difference should be readily apparent in observations of $z\\sim$4 jets like in 1745+624 if the X-rays were dominated by IC/CMB emission. X-ray jets detected so far have observed $f_{\\rm x}$/$f_{\\rm r}$ values which range over 4 orders of magnitude and there is no obvious sign of a $(1+z)$-dependence in the latest large compilation of \\citet{kat05}.  However, there are complicating factors which will skew such correlations, in both models, such as source-to-source (and knot-to-knot within a single jet) variations in jet magnetic fields, electron energy densities, and speeds.  Particularly in current samples biased toward ``blazars'', the X-ray jet emission is particularly sensitive to the (very small) jet angles to our line of sight if there is significant relativistic beaming on these large scales.  For an IC/CMB origin, \\citet[][cf. figure~4 therein]{che04} outlined such a scenario to account for the large $f_{\\rm X}/ f_{\\rm r}$ ratio of the $z$=4.3 GB~1508+5714 quasar jet ($>$ 100; among the largest value found thus far) in comparison to other known, lower redshift ($z\\simlt$2) jets due mostly to its extreme redshift\\footnote{$f_{\\rm X}/ f_{\\rm r}$ can be equivalently expressed as the radio-to-X-ray power-law slope; $\\alpha_{\\rm rx}$=0.73 for GB~1508+5714 \\citep{che04} while $\\alpha_{\\rm rx} \\sim 1.1-0.8$, translating to $f_{\\rm X}/ f_{\\rm r} \\sim 0.1-30$ at lower redshifts \\citep[see figures 5 and 9 in][respectively]{sam04,kat05}.\\label{footnote-x}}.  This scenario can be extended to the $z$=3.63 core-dominated quasar J2219--2719 where 6 total extended X-ray counts were recently detected in an $\\sim$8 ksec exposure \\citep{lop06}, and can be attributed to a $\\sim$1 mJy radio feature 2\\arcsec\\ south of the quasar \\citep[thus $f_{\\rm X}/ f_{\\rm r}>100$ also; ][and manuscript in preparation]{che05}. Two quasars showing particularly low $f_{\\rm x} / f_{\\rm r}$ ratios, despite their fairly high-redshifts ($z$=1.4 and 2) also stood out -- these are especially lobe-dominated quasars in comparison to the other examples, and a connection to beaming was suggested \\citep[see][for a discussion]{che04}.  One caveat to note is that the Doppler factors required for the knots in the 1745+624 jet (in the minimum-energy + IC/CMB framework) are rather moderate, $\\delta=4-5$, and similar to that found in the GB~1508+5714 case \\citep[][\\S~\\ref{section-intro}]{che04}. The required $\\delta$ in these two $z\\sim$4 quasar jets are at the low end of the corresponding values calculated for lower-redshift objects ($\\delta\\sim$4--25). In fact, these few critical data points at high-$z$ help to define the trend discussed by \\citet[figure 10 therein]{kat05} that the maximal values of $\\delta$ required to explain the X-ray jet emission as IC/CMB {\\it decreases with redshift}. Also, in contrast to the GB~1508+5714 and J2219--2719 cases, the $f_{\\rm x}$/$f_{\\rm r}$ values of the 1745+624 jet are closer to those observed at lower-redshifts (Table~\\ref{table-2} and footnote~\\ref{footnote-x}).  Although second-order effects from the intrinsic jet properties should be factored in, the low implied jet Doppler factors and lack of a very obvious redshift-dependence of $f_{\\rm x}$/$f_{\\rm r}$, can be taken as evidence against an IC/CMB interpretation. In particular, if there were more extreme Doppler factors than in the $z$$\\sim$4 quasars observed so far, the large-scale jets would be orders of magnitude brighter and will outshine the active cores in X-rays via the IC/CMB process \\citep{sch02}.  Currently however, no such sources have been found in a growing number of high-$z$ radio-loud quasars imaged with subarcsecond-resolution by {\\it Chandra} \\citep{bas04,lop06}, so where are the very highly beamed high-$z$ X-ray jets? There is of course a natural Malmquist bias inherent in studying very high-redshift objects, so this and other selection biases must be addressed in future samples.  On the other hand, if these trends persist with further observations, it may instead reflect intrinsic differences between large-scale quasar jets located at different redshifts and/or their different environments, which are then probed with observations. For example, in the IC/CMB interpretation, the trend may imply that high-$z$ jets of this kind are intrinsically slower than their lower-redshift counterparts. This may be due to a more disturbed environment (evidenced by the ``alignment effect'' of high-$z$ radio galaxies; e.g., Rees 1989) through which high-$z$ jets are propagating. Such an scenario, in fact, was suggested early on to explain the distorted morphologies observed in a large samples of high-$z$ ($\\sim$1.5--3) quasars and radio galaxies \\citep[e.g.,][]{bar88}.  As current X-ray studies have tended toward known well-studied radio jets \\citep[e.g.,][]{sam04,mar05} that are relatively local, it is important to test if this and other trends persist and can be clarified with larger, more homogenous samples of distant ($z>$2--4) jets.  In analogy to VLBI proper motion studies, the earliest superluminal motions detected tended to be the highest \\citep{ver94} and larger ensembles of source measurements now tend toward lower values \\citep{ver03,kel04}. This may mean that beaming factors of kpc-scale X-ray jets as inferred from IC/CMB calculations are more typically smaller than the highest ones found thus far.    \\subsection{The Terminal ``Hot Spot'' \\label{section-hotspot}}  Recent {\\it Chandra} studies have aimed to distinguish the X-ray detected terminal jet features (i.e., the hot spots) from the jets \\citep{hard04,kat05,tav05}. Many of these knots show much the same problem as in the jet, i.e. the inability to extrapolate the radio-to-optical power-law slopes smoothly in the X-ray band, and the underprediction of not-beamed inverse-Compton (SSC and IC/CMB) emission to the observed. In this context, it is useful to discuss the terminal feature K2.5 of the 1745+624 jet as a terminal hot spot.  Knot K2.5 follows the conventional empirical definition of a hot spot only loosely \\citep{bri94b}, as applying these strict criteria to such high-redshift, small angular-size jets is quite restrictive.  Hot spots are strong terminal shocks formed at the tips of powerful jets where the strong compression of the jet magnetic field leads to a transverse configuration and relatively high radio polarization. This expectation is in good agreement with observations of sources located at low redshifts. K2.5 is situated at the edge of the radio source and it is indeed compact with a spatial extent of $\\sim 0.15\\arcsec \\times 0.075\\arcsec \\approx 1 \\times 0.5$ kpc (\\S~\\ref{section-radio}), which is typical for the known X-ray detected hot spots in quasars and FR II radio galaxies.  However, the parallel configuration of the magnetic field with respect to the jet axis, as well its low degree of linear polarization relative to the rest of the jet (Figure~\\ref{figure-3}), are not expected, nor typical of the hot spots resolved in more nearby radio sources \\citep{bri84}. {\\it This may betray the dominance of an underlying jet flow in this feature.}  The 5 GHz flux of the hot spot in 1745+624 implies a monochromatic luminosity, $L_{\\rm 5 \\, GHz} \\approx 2.42 \\times 10^{44}$ erg\\,s$^{-1}$. This is higher, although still comparable to the 5 GHz luminosities of hot spots at lower redshifts \\citep[e.g.,][]{kat05}.  Assuming energy equipartition as in the jet, and a cylindrical geometry for the emitting region with radius $0.15\\arcsec/2 \\approx 0.5$ kpc (\\S~\\ref{section-radio}) and similar length, we calculate the hot spot magnetic field, $B_{\\rm eq, \\, \\delta=1} \\approx 660$ $\\mu$G. This also is in rough agreement (though again slightly higher) with the minimum energy magnetic fields derived for the other knots \\citep{kat05}.  The equipartition magnetic energy density of the hot spot, $U_{\\rm B} \\approx 1.7 \\times 10^{-8}$ erg\\,cm$^{-3}$, is more than one order of magnitude greater than the energy density of the CMB radiation at the quasar's redshift. However, the energy density of the synchrotron photons thereby, $U_{\\rm syn} \\approx 4.7 \\times 10^{-9}$ erg\\,cm$^{-3}$, is roughly comparable to $U_{\\rm B}$. In this evaluation, we took the hot spot radio spectral index, $\\alpha_{\\rm r} \\approx 1$, its synchrotron luminosity $L_{\\rm syn} \\approx 10 \\times L_{\\rm 5 \\, GHz}$, the geometry as described above, and neglected relativistic corrections. The implied 1 keV flux of the hot spot due to the SSC plus IC/CMB processes is then $\\sim$0.18 nJy (Equation~\\ref{equation-ssc}). As in the jet emission (\\S~\\ref{section-basic}), this is inconsistent with the observed X-ray flux, though here the deviation is less drastic (factor of $\\sim$14; Table~\\ref{table-2}). However, one should keep in mind many of the underlying assumptions in such calculations before claiming a true disagreement between model predictions and the observations. In particular, the hot spot structure is in fact unresolved, hence a slightly smaller volume than considered, together with, e.g. small deviations from energy equipartition and spectral curvature (see below), could foreseeably bring the measurement into better agreement with the model prediction.  However, if we take the observed excess in X-rays over the estimated SSC flux at face value, this would not be the first such case for a hot spot. \\citet{hard04} has argued for a synchrotron origin of this X-ray excess in many local hot spot sources, though in the case of 1745+624, the optical upper limit precludes a straightforward extrapolation of the (very steep) radio spectrum up to the X-ray band. On the other hand, \\citet{tav05} has argued that there is an IC/CMB contribution from an underlying relativistic portion of the jet terminus (unresolved by our observations) to the few X-ray detected hot spots studied in several $z\\sim 1$ quasars.  The latter possibility is quite attractive in the case of 1745+624: as we noted earlier, the terminal jet feature in this source does not have typical hot spot-like radio polarization properties, and there could very well be an unresolved portion of the jet mixed in.  If the multi-band radiative output of K2.5 is dominated by a relativistic jet flow, this would imply $\\delta$=5 for an IC/CMB origin of the X-ray flux (using Equation~\\ref{equation-iccmb} with all the parameters discussed above). This beaming factor is quite high, though comparable to values derived in lower-redshift quasars \\citep{tav05}. It is also slightly higher than the values obtained for the proper jet (knots K1.4 and K1.8; see \\S~\\ref{section-beamed}), though again, there are the usual uncertainties in these calculations to keep in mind.  Regardless, the problem of terminal X-ray/radio knots in this and many quasar jets can be viewed as a simple extension of that in the jet knots.  One relevant issue in this discussion is the apparent lack of a hot spot on the side opposite of the nucleus to the prominent jet. If the advance speed of the jet is sub-relativistic, then some emission should have been visible on the counter-jet side. This can be reconciled by the fact that the usually inferred (sub-relativistic) advance speeds of hot spots in powerful radio sources assume that the jet thrust is balanced with the ram-pressure of the ambient material within a typical galaxy group/poor cluster environment \\citep{beg84}.  However, if the jets propagate into the radio cocoon formed in a previous epoch of radio activity, i.e. in an environment with a substantially decreased number density and pressure, the advance velocity of the jet may be close to $c$.  In such a case, combined Doppler and light-travel effects may result in an apparent lack of a visible counter hot spot. Such a scenario was proposed by \\citet{staw04} for the 3C~273 jet. Since the observed properties of 1745+624 and 3C~273 radio quasars are so similar, one could also apply this argument to explain the lack of emission visible on the counter-jet side in 1745+624.  Let us finally discuss yet another issue regarding this dominant radio feature, namely the extrapolation of its radio spectrum to lower frequencies. As discussed in \\S~\\ref{section-beamed}, an extrapolation of its observed $\\sim$$1-10$ GHz continuum $\\propto \\nu^{-1.3}$ to lower frequencies exceeds the extrapolation of the flat-spectrum core component $\\propto \\nu^0$ at about $\\sim 500$ MHz, and joins smoothly with the observed integrated low frequency emission at $\\sim 200$ MHz (see Figure~\\ref{figure-7} [left]). In other words, one can attribute the $\\sim$30--300 MHz integrated emission to the extended radio structure. Since the extended emission is probably dominated by the brightest feature, the hot spot (we can not exclude a contribution from the fainter inner jet; \\S~\\ref{section-beamed}), this would imply a broken power-law character of its synchrotron spectrum, $\\propto \\nu^{-0.5}$ for $\\nu < 300$ MHz and $\\propto \\nu^{-1.3}$ for $\\nu > 300$ MHz. In the standard continuous-injection model of terminal features in powerful jets \\citep{hea87}, the spectral break would correspond to the energies, $E_{\\rm e} / m_{\\rm e} c^2 \\sim 10^3$. This low-frequency break is expected in face of the large computed hot spot equipartition magnetic field, $B_{\\rm eq, \\, \\delta=1} \\approx 660$ $\\mu$G \\citep[see a discussion in][]{bru03,cheung05}.   \\subsection{The Jet Energetics \\label{section-energetic}}  The multiwavelength emissions of the 1745+624 jet allow us to estimate the kinetic power of the outflowing plasma in this source. The kinetic power of the ultrarelativistic electrons and magnetic field characterizing the proper jet, under the minimum-power hypothesis, is roughly $L_{\\rm j} \\sim L_{\\rm B} + L_{\\rm e} \\sim 2 \\times \\pi R^2 \\, c \\Gamma^2 \\, U'_{\\rm B}$, where $R \\approx 0.7$ kpc is the jet radius and $U'_B = 1.5 \\times 10^{-9} \\, \\delta^{-10/7}$ erg s$^{-1}$ is the jet comoving magnetic field energy density, as discussed in \\S~\\ref{section-discussjet}. Assuming for illustration, a jet viewing angle $\\sim 10\\deg$ (see \\S~\\ref{section-describe}) and jet bulk Lorentz factor at the kpc-scale $\\Gamma \\sim 3$ (leading to the `comfortable' value of the jet Doppler factor $\\delta \\sim 4.7$), one obtains a kinematic factor $\\delta^{-10/7} \\, \\Gamma^2 \\sim 1$, and therefore a jet kinetic luminosity (regarding solely ultrarelativistic jet electrons and the jet magnetic field) $L_{\\rm j} \\sim 10^{45}$ erg s$^{-1}$.  If we regard the terminal feature as a ``true'' hot spot, we can estimate the kinetic energy of the jet. We approximate the broad-band synchrotron spectrum of the hot spot by a broken power-law: $\\propto \\nu^{-0.5}$ between the assumed minimum synchrotron frequency $\\nu_{\\rm min} = 10$ MHz and the break frequency $\\nu_{\\rm br} = 300$ MHz, followed by $\\propto \\nu^{-1.3}$ between $\\nu_{\\rm br}$ and the assumed maximum synchrotron frequency $\\nu_{\\rm max} = 10^{12}$ Hz (see \\S~\\ref{section-hotspot}). The radiative efficiency factor, i.e. the ratio of the power emitted by the hot spot electrons via synchrotron radiation, and the total power injected to these electrons at the terminal shock, is:  \\begin{equation} \\eta_{\\rm rad} \\sim {\\ln (\\nu_{\\rm max} / \\nu_{\\rm br}) \\over \\ln (\\nu_{\\rm max} / \\nu_{\\rm min})} \\sim 0.7 \\, . \\end{equation}  \\noindent The reconstructed radio continuum of the hot spot implies also its total synchrotron luminosity $L_{\\rm syn} \\approx 10 \\times L_{\\rm 5 \\, GHz} \\sim 2.5 \\times 10^{45}$ erg s$^{-1}$. Thus, the total kinetic power of the jet transported from the active center to the jet terminal point is roughly:  \\begin{equation} L_{\\rm kin} \\sim {L_{\\rm syn} \\over \\eta_{\\rm rad} \\, \\eta_{\\rm e}} \\, , \\end{equation}  \\noindent where $\\eta_{\\rm e}$ is the fraction of the jet kinetic energy transformed in the terminal shock to the internal energy of ultrarelativistic (i.e.  synchrotron emitting) electrons. In the case of energy equipartition between the electrons and the magnetic field reached at the hot spot, one should expect very roughly $\\eta_{\\rm e} \\sim 0.5$ \\citep[see][for a wider discussion]{sik01}. This gives $L_{\\rm kin} \\sim 0.75 \\times 10^{46}$ erg s$^{-1}$. In fact, $\\eta_{\\rm e} = 0.5$ should be regarded as an upper limit, keeping in mind a possible proton contribution and work done by the outflow in pushing out the ambient medium, making $10^{46}$ erg s$^{-1}$ a safe lower limit for the kinetic power of the 1745+624 jet. The jet power estimated at the beginning of this subsection is an order of magnitude lower than its total kinetic power constrained from the hot spot emission. This may indicate a dynamical role of non-radiating jet particles \\citep[either cold electrons or protons; see in this context,][]{sik00}.  Let us compare the kinetic power of the 1745+624 jet estimated above with the accretion parameters of the quasar core. Due to the large distance of the quasar, and hence related observational difficulties, these accretion parameters cannot be evaluated precisely thus should be treated with some caution. Nevertheless, we note \\citet{sti93} and \\citet{sk93} found a broad emission CIV line in the spectrum of the 1745+624 quasar at the observed wavelength $\\lambda_{\\rm CIV}^{\\rm obs} = 7553$ \\AA, and an observed ${\\rm FWHM_{CIV}^{\\rm obs}} = 64.9$ \\AA. This gives the intrinsic velocity dispersion:  \\begin{equation} v_{\\rm CIV} = c \\, {{\\rm FWHM_{CIV}} \\over \\lambda_{\\rm CIV}} = c \\, {{\\rm FWHM_{CIV}^{\\rm obs}} \\over \\lambda_{\\rm CIV}^{\\rm obs}} \\approx 2576 \\, {\\rm km \\, s^{-1}} \\, . \\end{equation}  \\noindent \\citet{sti93} reported also a steep-spectrum ($\\alpha_{\\rm opt} = 1.3$), non-variable and low-polarized optical continuum of the 1745+624 quasar, which we consider below as being produced by the circumnuclear gas solely. In other words, we assume that the host galaxy and the nuclear portion of the jet do not contribute significantly to the observed optical emission of the discussed object. At the observed wavelength $6600$ \\AA, corresponding to the source rest frame wavelength $1350$ \\AA, the detected continuum flux is $F_{6600} \\approx 0.08$ mJy. Thus, the intrinsic continuum luminosity at the emitted frequency $1350$ \\AA, is simply $L_{1350} = 4 \\pi d_{\\rm L}^2 \\, \\left[\\lambda F_{\\lambda}\\right]_{6600} \\approx 5.4 \\times 10^{46}$ erg s$^{-1}$. These estimates allow us to find the bolometric luminosity of the quasar core (assumed to be about ten times the $V$-band core luminosity), $L_{\\rm bol} \\sim 10 \\times L_{5500} \\sim 10 \\, \\left(\\nu_{5500} / \\nu_{1350}\\right)^{1-\\alpha_{\\rm opt}} \\, L_{1350} \\sim 8.2 \\times 10^{47}$ erg s$^{-1}$, as well as the mass of the central supermassive black hole via the virial relation \\citep{ves02,ves06}:  \\begin{equation} {{\\cal M}_{\\rm BH} \\over {\\cal M}_{\\odot}} = 5.4 \\times 10^6 \\, \\left({v_{\\rm CIV} \\over 1000 \\, {\\rm km/s}}\\right)^2 \\, \\left({L_{1350} \\over 10^{44} \\, {\\rm erg/s}}\\right)^{0.53} \\simgt 1 \\times 10^9. \\end{equation}  \\noindent This can be taken as a lower limit in this case because the relation was calibrated for the broad component of the CIV line, which was not measured independently from the narrow one \\citep{sti93,sk93}; including the latter underestimates the measured broad-line width \\citep{ves02} which could be $>$3,000 km s$^{-1}$.  The corresponding Eddington luminosity is $L_{\\rm Edd} \\simgt 10^{47}$ erg s$^{-1}$ and the Eddington ratio, $\\Lambda \\sim L_{\\rm bol} / L_{\\rm Edd} \\sim 1-10$. The latter is quite large and may indicate contribution of the emission due to an unresolved portion of the jet to the nuclear optical continuum produced by the accreting matter.  Finally, the total observed $1.4$ GHz flux of the 1745+624 radio source \\citep[780.6 mJy; ][]{con98}, gives the total monochromatic radio luminosity $L_{\\rm R} \\approx 1.6 \\times 10^{45}$ erg s$^{-1}$ thus a radio loudness parameter, ${\\cal R} \\sim 10^5 \\, L_{\\rm R} / L_{\\rm V} \\sim 10^{3} - 10^{4}$. Note that the discussed object, which possesses the most powerful observed radio jet known, is extremely `radio-loud', as the standard division between radio-quiet and radio-loud sources \\citep{kel89} is ${\\cal R} = 10$. With the obtained values of the Eddington ratio and radio loudness, 1745+624 fits into the `radio-loud' trend formed on the $\\Lambda - {\\cal R}$ plane by local ($z < 0.5$) FR I radio galaxies and radio selected broad-line AGNs, as discussed in detail by \\citet{sik06}.   }  We have analyzed multi-frequency radio, optical and X-ray imaging data for the kpc-scale jet in the $z$=3.89 quasar 1745+624. This quasar hosts the most powerful large-scale radio and X-ray jet yet observed, with monochromatic radio and X-ray luminosities, $L_{\\rm 5 \\, GHz} \\approx 1.3 \\times 10^{44}$ erg s$^{-1}$ and $L_{\\rm 1 \\, keV} \\approx 2.8 \\times 10^{45}$ erg s$^{-1}$, respectively (excluding emission from the powerful terminal ``hot spot''). Aside from its large radiative output, and related large inferred equipartition magnetic field (in comparison to local powerful jet sources), its properties (multi-wavelength morphology, radio polarization, and broad-band spectral energy distribution) are broadly similar to other lower-redshift quasar jets. This is unexpected in light of the dramatic increase in the CMB energy density and strong cosmic evolution of the intergalactic medium at such high-redshifts;  this should have manifested in a dramatically different appearance of such large-scale outflows.  The spectral energy distributions of the resolved linear structures are discussed.  In this jet, inverse-Compton scattered emission on the CMB photons can just account for the X-ray emission if the jet is inclined close to our line of sight ($\\simlt$10\\deg) and if it is also moving relativistically, $\\Gamma$=2--3. Several indirect arguments from the multi-scale radio observations of 1745+624 support such values of the large-scale jet kinematic parameters. The data are also consistent with a synchrotron interpretation provided the electron acceleration timescale is much shorter than that of the radiative losses, which we deem likely. The main distinguishing feature of the two models should manifest in drastically different X-ray/radio emission properties of such high-redshift jets though with only three current $z>$3 examples, no trends are yet apparent.  Via the virial relation, we estimate that a ${\\cal M}_{\\rm BH} \\simgt 1 \\times 10^9 \\, {\\cal M}_{\\odot}$ supermassive black hole resides at the center of this galaxy.  The broad-band emission of the extended components allow us to additionally estimate a jet kinetic power of $L_{\\rm kin} \\sim 10^{46}$ erg\\,s$^{-1}$, which is a small fraction of the Eddington luminosity ($L_{\\rm kin}/L_{\\rm Edd} \\sim 0.1$) corresponding to this black hole mass. As the bolometric luminosity of the quasar is $L_{\\rm bol} / L_{\\rm Edd} \\sim 10$, the inferred jet kinetic power seems to be much less than the estimated accretion power. Despite this, the 1745+624 quasar, hosting the most powerful radio jet known, is extremely `radio-loud' in comparison to more local radio selected quasars with a radio-loudness parameter ${\\cal R} \\geq 10^{3}$, and is surprisingly similar to local radio selected broad line AGNs regarding the accretion parameters."
    },
    "0606/astro-ph0606019_arXiv.txt": {
        "abstract": "The concordance model of cosmology and structure formation predicts the formation of isolated very massive stars at high redshifts in dark matter dominated halos of $10^5$ to $10^6\\Msun$. These stars photo-ionize their host primordial molecular clouds, expelling all the baryons from their halos. When the stars die, a relic \\ion{H}{2} region is formed within which large amounts of molecular hydrogen form which will allow the gas to cool efficiently when gravity assembles it into larger dark matter halos. The filaments surrounding the first star hosting halo are largely shielded and provide the pathway for gas to stream into the halo when the star has died. We present the first fully three dimensional cosmological radiation hydrodynamical simulations that follow all these effects. A novel adaptive ray casting technique incorporates the time dependent radiative transfer around point sources. This approach is fast enough so that radiation transport, kinetic rate equations, and hydrodynamics are solved self-consistently. It retains the time derivative of the transfer equation and is explicitly photon conserving. This method is integrated with the cosmological adaptive mesh refinement code {\\sl enzo}\\, and runs on distributed and shared memory parallel architectures. Where applicable the three dimensional calculation not only confirm expectations from earlier one dimensional results but also illustrate the multi--fold hydrodynamic complexities of \\ion{H}{2} regions. In the absence of stellar winds the circumstellar environments of the first supernovae and putative early gamma--ray bursts will be of low density $\\sim 1$ cm$^{-3}$. Albeit marginally resolved, ionization front instabilities lead to cometary and elephant trunk like small scale structures reminiscent of nearby star forming regions. ",
        "introduction": "\\footnote{Visualizations and animations of the simulations presented here can be found at {\\tt http://www.slac.stanford.edu/} {\\tt $\\sim$jwise/research/RT1/}} With the fundamental cosmological parameters pinned down by recent observations, questions related to first structure formation have now, in principle, no free parameters. Hydrodynamical simulations that start with cosmological initial conditions have shed light on the nature of the first luminous objects in the universe \\markcite{1998ApJ...508..518A, 1999AIPC..470...58N, 2000ApJ...540...39A, 2002Sci...295...93A, 2002MNRAS.330..927H, 2003ApJ...592..645Y}({Abel} {et~al.} 1998; {Norman}, {Abel}, \\&  {Bryan} 1999; {Abel}, {Bryan}, \\&  {Norman} 2000, 2002; {Hutchings} {et~al.} 2002; {Yoshida} {et~al.} 2003). \\markcite{2002Sci...295...93A}{Abel} {et~al.} (2002) found that the first luminous objects are isolated massive stars with masses somewhere in the range between 30 and 300 solar masses depending on how much feedback from the proto-star affects the accretion rate. The stellar build up from predicted accretion rates in the absence of feedback have been confirmed in smooth particle hydrodynamic simulations~\\markcite{2004NewA....9..353B}({Bromm} \\& {Loeb} 2004). There are physical reasons from proto-stellar evolution that accretion may come to an end when the star reaches $\\sim 100\\Msun$ \\markcite{2003ApJ...589..677O}({Omukai} \\& {Palla} 2003) but full radiation hydrodynamical simulations, even in one dimension, have not been possible to follow primordial stars up to the zero age main sequence. The early proto-stellar evolution, however, is understood in some detail \\markcite{1998ApJ...508..141O, 2002MNRAS.334..401R, 2004MNRAS.348.1019R}({Omukai} \\& {Nishi} 1998; {Ripamonti} {et~al.} 2002; {Ripamonti} \\& {Abel} 2004).  The immediate relevance of primordial massive stars to cosmological reionization has prompted a number of investigations even though radiative transfer effects could only be accounted for crudely \\markcite{2000ApJ...535..530G, 2001NewA....6..437G, 2004MNRAS.350...47S, 2005ApJ...628L...5O, 2006ApJ...639..621A}({Gnedin} 2000; {Gnedin} \\& {Abel} 2001; {Sokasian} {et~al.} 2004; {O'Shea} {et~al.} 2005; {Alvarez}, {Bromm}, \\&  {Shapiro} 2006) or in non-cosmological two-dimensional calculations~\\markcite{2004MNRAS.348..753S}({Shapiro}, {Iliev}, \\&  {Raga} 2004). In spherical symmetry, the evolution of the first \\ion{H}{2} regions was studied by \\markcite{2004ApJ...610...14W}{Whalen}, {Abel}, \\&  {Norman} (2004) and \\markcite{2004ApJ...613..631K}{Kitayama} {et~al.} (2004). These authors found that when the ionization front slows to become D-type it accelerates all the baryonic material to ten times the escape velocity of the dark matter halos that host these first stars. However, neither the stability of the ionization front, nor the impact of photo-ionization on the surrounding filaments could be addressed in the simplified one dimensional models. Significant amounts of molecular hydrogen are formed in relic \\ion{H}{2} regions facilitated by their large initial electron fractions \\markcite{2002ApJ...575...33R, 2005ApJ...628L...5O}(e.g. {Ricotti}, {Gnedin}, \\&  {Shull} 2002; {O'Shea} {et~al.} 2005). Whether, photo-ionization leads to a net increase, decrease, or a delay in star formation remains controversial \\markcite{1996ApJ...467..522H, 2000ApJ...534...11H, 2001ApJ...551..599H, 2000MNRAS.314..611C, 2003MNRAS.346..456O, 2005ApJ...628L...5O, 2006ApJ...639..621A}({Haiman}, {Rees}, \\&  {Loeb} 1996; {Haiman}, {Abel}, \\&  {Rees} 2000; {Haiman}, {Abel}, \\&  {Madau} 2001; {Ciardi} {et~al.} 2000; {Oh} \\& {Haiman} 2003; {O'Shea} {et~al.} 2005; {Alvarez} {et~al.} 2006). These outstanding questions will be resolved by direct numerical simulations that accurately follow the radiation transport as well as the cosmological hydrodynamics. Such calculations are presented in this {\\sl Letter}\\ for the first time. ",
        "conclusions": "\\label{conc}  Early \\ion{H}{2} regions have a profound impact on the earliest structures. They displace a large amount of baryons, raise the gas entropy, and enable the formation of large amounts of molecular hydrogen after the star dies. If Pop III stars indeed have little mass loss before they die \\markcite{2001ApJ...550..890B}({Baraffe}, {Heger}, \\&  {Woosley} 2001) our results illustrate that putative gamma ray bursts and supernovae will occur in a low density environment with densities $\\sim 1$ cm$^{-3}$ \\markcite{2004ApJ...604..508G}({Gou} {et~al.} 2004). The mixing of the first heavy elements occurs outside of dark matter halos in the intergalactic medium in regions that eventually will re-collapse to form galaxies. Consequently, the elements necessary for life may have been spread early and widely.  The rich subject of ionization front instabilities \\markcite{1979ApJ...233..280G,1996ApJ...469..171G}(cf. {Giuliani} 1979; {Garcia-Segura} \\& {Franco} 1996) is important in early as well as galactic star formation. It can now be addressed in three dimensions at high spatial resolution using our radiation hydrodynamical adaptive mesh simulation techniques."
    },
    "0606/astro-ph0606533_arXiv.txt": {
        "abstract": "A wide-field galaxy redshift survey allows one to probe galaxy clustering at largest spatial scales, which carries invaluable information on horizon-scale physics complementarily to the cosmic microwave background (CMB).  Assuming the planned survey consisting of $z\\sim 1$ and $z\\sim 3$ surveys with areas of $2000$ and 300 deg$^2$, respectively, we study the prospects for probing dark energy clustering from the measured galaxy power spectrum, assuming the dynamical properties of dark energy are specified in terms of the equation of state and the effective sound speed $c_{\\rm e}$ in the context of an adiabatic cold dark matter dominated model. The dark energy clustering adds a power to the galaxy power spectrum amplitude at spatial scales greater than the sound horizon, and the enhancement is sensitive to redshift evolution of the net dark energy density, i.e. the equation of state. We find that the galaxy survey, when combined with CMB expected from the Planck satellite mission, can distinguish dark energy clustering from a smooth dark energy model such as the quintessence model ($c_{\\rm e}=1$), when $c_{\\rm e}\\simlt 0.04$ (0.02) in the case of the constant equation of state $w_0=-0.9$ ($-0.95$). An ultimate full-sky survey of $z\\sim 1$ galaxies allows the detection when $c_{\\rm e}\\simlt 0.08$ (0.04) for $w_0=0.9$ ($-0.95$). These forecasts show a compatible power with an all-sky CMB and galaxy cross-correlation that probes the integrated Sachs-Wolfe effect.  We also investigate a degeneracy between the dark energy clustering and the non-relativistic neutrinos implied from the neutrino oscillation experiments, because the two effects both induce a scale-dependent modification in the galaxy power spectrum shape at largest spatial scales accessible from the galaxy survey. It is shown that a wider redshift coverage can efficiently separate the two effects by utilizing the different redshift dependences, where dark energy clustering is apparent only at low redshifts $z\\simlt 1$. ",
        "introduction": "Various cosmological probes such as supernovae in distant galaxies \\cite{Perlmutter99,Riess99}, the cosmic microwave background (CMB) sky \\cite{WMAP3}, and the galaxy redshift surveys \\cite{Tegmark04,2dF,Scranton,Eisenstein} have given strong evidence that a dark energy component, such as the cosmological constant, constitutes approximately $70\\%$ of the total energy density of the universe, which derives the accelerating cosmic expansion at low redshifts.  Because there is no plausible theoretical explanation for its existence and magnitude (e.g. see \\cite{weinberg89,Ratra}), observational exploration of the nature of dark energy is one of the most important issues in modern cosmology as well as particle physics.  An observational dark energy task we should first explore to tackle this fundamental problem would be to determine whether the accelerating expansion is as a consequence of the cosmological constant.  Relaxing this assumption leads to a generalized dark energy with dynamically evolving energy density, which can be characterized by a time-dependent equation of state $w(a)=p_{\\rm de}/\\rho_{\\rm de}$ (the cosmological constant corresponds to $w=-1$). The current level of accuracy in constraining the constant $w$-parameter is $\\sigma(w)\\sim 0.1$ at $1\\sigma$ level with the best-fit value containing a model with $w=-1$ (e.g. see \\cite{Eisenstein,SNLS,WMAP3}). Future prospects aimed at pinning down the constraint on $w$ by a factor of 10 have been extensively investigated to address the usefulness of various cosmological experiments based on massive galaxy surveys such as tomographic weak lensing experiment (e.g. \\cite{Huterer,TJ}), the baryon oscillation experiment (e.g.  \\cite{SE,HuHaiman,taka05}) and the cluster abundance experiment (e.g.  \\cite{Laura}).  Another important consequence of a generalized dark energy is the spatial perturbation, providing an independent clue to resolving the nature of dark energy from the equation of state.  There are many previous efforts made to study how an inclusion of the dark energy perturbations modify the cold dark matter (CDM) structure formation scenarios: theoretical predictions on the modifications in the CMB power spectra and the galaxy power spectra \\cite{Caldwell,Chiba,Garriga,Dave,Ma,Hu02,Moroi,Huterer,Dedeo,Uzan} and the observational exploration of the dark energy clustering signal from the WMAP data at low multipoles \\cite{WMAP3,Bean,Lewis,GordonHu}. The dark energy clustering is relevant only for structure formation at low redshifts $z\\simlt 1$, where the net energy density is dominant to the cosmic expansion.  In addition, a reasonable model including the dark energy perturbation predicts that the dark energy can cluster together with matter components at large spatial scales (inevitably at super-horizon scales), whilst the dark energy is smooth at small scales so that the additional component does not largely change the small-scale structure formations such as galaxy formation. For these reasons, the dark energy clustering effect on the CMB observables is likely to appear only via the integrated Sachs-Wolfe effect (ISW) at low multipoles, where the Sachs-Wolfe effect generated at the recombination epoch is significant contamination to separate the ISW effect from the measured power spectrum.  The transition scale to divide the dark energy clustering and smooth regimes can be usefully modeled by the effective sound speed of dark energy \\cite{Hu98,Hu02}. In this model, smoothness of dark energy can be tested by searching for the signature of the sound speed from cosmological observables. Hu and Scranton \\cite{HuScranton} carefully investigated a prospect of how the ISW effect measured via the angular cross-correlation between CMB and galaxy distribution can be used to probe the dark energy clustering, assuming an all-sky, deep multi-color imaging galaxy survey out to $z\\sim 2$.  A galaxy redshift survey offers an alternative means for probing the dark energy clustering, through the measured statistical properties of three-dimensional gravitational clustering at largest scales.  Compared to the angular correlations, from a cosmological point of view, the redshift survey carries more information on the underlying mass distribution due to a gain of the modes along the line-of-sight or redshift direction.  Therefore, the purpose of this paper is to, for the first time, investigate the ability of a galaxy redshift survey for testing the smoothness of dark energy from the measured galaxy power spectrum.  In fact, there are several future plans for high-redshift galaxy surveys that are already being constructed or seriously under consideration: the Fiber Multiple Object Spectrograph (FMOS) on Subaru telescope \\cite{fmos}, its significantly expanded version, WFMOS \\cite{wfmos}, the Hobby--Ebery Telescope Dark Energy eXperiment (HETDEX) \\cite{hetdex}, and the Cosmic Inflation Probe (CIP) mission \\cite{cip}. These surveys probe galaxies at higher redshifts $z\\simgt 0.5$ than the existing surveys such as SDSS and 2dF surveys. Such a high-redshift survey has several advantages over the lower redshift ones. First, given a fixed solid angle, the comoving volume in which we can observe galaxies is larger at higher redshifts than in the local universe, thereby reducing the sample variance. This would make it more straightforward to obtain a well-behaved survey geometry that can help measure largest-scale perturbations. Second, density perturbations at smaller spatial scales are still in the linear regime or only in the weakly non-linear regime at higher redshift, which gives us more leverages on measuring the shape of the linear power spectrum to break the parameter degeneracies. In this paper, we will consider the survey design close to the proposed WFMOS survey, which consists two types of surveys different in redshift coverage and survey area: $0.5\\le z\\le 1.3$ with 2000 degree$^2$ and $2.5\\le z\\le 3.5$ with $300$ degree$^2$, respectively. The redshift coverage of WFMOS is suitable to probe the dynamical dark energy whose effects are apparent only at low redshifts $z\\simlt 1$, as we will show below.  The structure of this paper is as follows. In Sec.~\\ref{prel} we start with writing down background cosmological equations, the Hubble expansion and the angular diameter distance, in terms of cosmological parameters. In Sec.~\\ref{de}, we define the effective sound speed parameter to model dynamical properties of dark energy clustering, and review how the dark energy clustering leads to a scale-dependent modification in the linear power spectrum shape assuming the adiabatic initial condition. Sec.~\\ref{gps} defines the galaxy power spectrum in terms of the primordial power spectrum, the transfer functions and the scale-dependent growth rate of mass clustering. In Sec.~\\ref{method}, we first define survey parameters intended to resemble a future survey being planned, and describe a methodology to model the galaxy power spectrum observed from a redshift survey that includes the two-dimensional nature in the line-of-sight and transverse directions due to the cosmological and redshift distortion effects.  We then present the Fisher information matrix formalism that is used to estimate the projected uncertainties in the cosmological parameter determination provided the measured galaxy power spectrum. In Sec.~\\ref{results} we show the prospects of the future survey for probing the dark energy clustering.  In addition, we carefully study how a degeneracy between the dark energy clustering and massive neutrinos can be lifted by utilizing the redshift information of galaxy clustering.  Finally, we present conclusion and some discussion in Sec.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  In this paper we have investigated the ability of a future galaxy redshift survey to test the smoothness of dynamical dark energy, which is another important consequence of a generalized dark energy with $w(z)\\ne-1$ and provides independent information on the nature of dark energy from that carried by the equation of state. The dark energy clustering signatures can be measured via a scale-dependent transition in the power of the galaxy power spectrum appearing at scales comparable with the dark energy sound horizon (see Fig.~\\ref{fig:pk}). It was shown that, for WFMOS survey, the sound speed can be detected at more than a 1-$\\sigma$ level, if the sound speed is small enough as $c_{\\rm e}\\simlt 0.04 $ (0.02) when $w_0=-0.9$ ($-0.95$) (see Figs.~\\ref{fig:ce} and \\ref{fig:ce-w}).  An effective way to improve the ability is to enlarge the survey volume especially for low redshift slices at $z\\sim 1$. An ultimate full-sky survey could improve the lower bound on the detectable $c_{\\rm e}$ by a factor of 2.  Another interesting possibility of the future galaxy survey is the use of the galaxy power spectrum to weigh the neutrino mass, as investigated in \\cite{TK}. We carefully investigate possible degeneracies between the dark energy clustering and the non-relativistic neutrinos in the galaxy power spectrum (see Fig.~\\ref{fig:pk_fnu-ce} and also see \\cite{Hannestad} for the related discussion on the degeneracy between the neutrino mass and the dark energy equation of state). We showed that having a wider redshift coverage can efficiently separate the two effects by utilizing the different redshift dependences; the dark energy is prominent only at low redshifts $z\\simlt 1$. In addition, modified gravity theories have received much recent attention as an alternative explanation of cosmic acceleration without dark energy, where gravity is weaker than the Einstein gravity at scales comparable to the horizon (e.g., \\cite{Shirata,Raul,Peacock,Fritz}). An important test to discriminate this possibility from the dark energy model is to simultaneously explore the cosmic expansion history and the growth of structure formation from cosmological experiments, because the two possibilities predict different growth rate even for an identical expansion history. Once again, wider redshift coverage of a galaxy redshift survey would be powerful to make a robust test of discriminating between effects of dark energy clustering, massive neutrinos, and modified gravity, which warrants a detailed study (also see \\cite{Peacock} for the discussion on the interplay between the massive neutrinos and the modified gravity in the shape of mass power spectrum). We hope that the results shown in this paper will give a useful guidance to designing a future survey if the dark energy clustering is desired to pursue as one of the science goals.  We have assumed throughout this paper that dynamical properties of the dark energy are specified in terms of the equation of state and the effective sound speed, with the adiabatic initial condition.  A merit of this approach is we need not assume a specific form of the Lagrangian of dark energy sector.  However, the limitation is there is no guarantee this modeling can be applied to a full range of generic dark energy models. For example, we have ignored a contribution from the trace-free stress perturbation of dark energy, which is valid if the dark energy is a scalar field (e.g. see \\cite{Uzan} for the extension).  Since we have little idea of what dark energy is, it is worth exploring a more general modeling of dark energy properties and investigate the resulting effect on structure formation.  Unless such a theoretical model is available, we cannot extract all the cosmological information inherent in the future high-precision cosmological data sets in an unbiased way.  Another non-trivial assumption we have made in this paper is the adiabatic initial conditions for dark energy perturbations, without any rigorous reasoning. If the dark energy field had existed since the inflationary epoch in the early universe, the dark energy perturbations naturally contained iso-curvature modes, and the effect just emerges at low redshifts. Hence, a galaxy redshift survey could by its own open up a new window for exploring the iso-curvature modes in structure formation to probe the physics in the early universe, complementarily to the CMB or the CMB-galaxy cross correlation \\cite{GordonHu,HuScranton}. Furthermore, if the primordial non-Gaussianity (e.g. see \\cite{Komatsu} for a thorough review) is dominantly imprinted onto the sector of dark energy, the galaxy survey allows us to explore the signal via, e.g., the bispectrum measurement of galaxy clustering. These are our future study and will be presented elsewhere.  The linear perturbation theory makes secure predictions on structure formation, which allows an accurate interpretation of the cosmological data sets from the detailed comparison. A main obstacle that a galaxy redshift survey could contain in this procedure is uncertainties involved in the galaxy biasing, where unknown, non-linear gastrophysics may pollute the linear information even at large scales.  Although we have assumed a scale-independent, linear bias, we do not think that a more realistic bias such as a scale-dependent bias is a significant contamination to the dark energy clustering constraint.  Because the dark energy clustering induces characteristic scale-dependent modification and strong redshift evolution for $P(k)$ (see Fig.~\\ref{fig:pk}), these properties can be very robust to discriminate the dark energy clustering signals from the galaxy bias contamination or more generally other systematics.  Finally, we comment on the non-linear gravitational clustering, which could also provide a contamination to the linear theory predictions. Encouragingly, however, the non-linear effects can be to some extent corrected based on our refined knowledge of the structure formation scenarios, based on the perturbation theory as well as the N-body simulations \\cite{Suto,Jain94,Jeong,Fritz}. Furthermore, as an interesting possibility, coupling of different Fourier modes induced by non-linear gravitational clustering could induce a transfer of the dark energy clustering effect at large scales to the power of $P(k)$ at smaller scales.  If this is the case to be apparent on the measured power spectrum, an inclusion of the non-linearity in the model predictions could improve a signal-to-noise to test smoothness of dark energy, from the smaller-scale clustering information. This issue will be worth exploring in detail, and will be presented elsewhere.   \\bigskip {\\it Acknowledgments:} The author thanks E.~Komatsu for valuable discussions and also thanks for T.~Futamase, O.~Lahav, D.~Nitta, P.~Norberg and Y.~Suto for helpful discussions. He also acknowledges the warm hospitality of Caltech where this work was partly done.  This work was supported in part by a Grand-in-Aid for Scientific Research (17740129) of the Ministry of Education, Culture, Sports, Science and Technology in Japan as well as the COE program at Tohoku University.  We acknowledge the use of the publicly-available CMBFAST code \\cite{cmbfast}."
    },
    "0606/astro-ph0606705_arXiv.txt": {
        "abstract": "{\\hd} is a unique system containing one massive star ({\\stA}) that is apparently entering the luminous blue variable phase, and an eclipsing companion ({\\stB}) that may have already evolved beyond this phase to become a Wolf-Rayet star. In this paper we present the results from {\\fuse} observations obtained in 1999, 2000, and 2002 and one far-UV observation obtained by {\\orfeus}/BEFS in 1993 shortly before the first eruption of {\\hd}. The eight phase-resolved spectra obtained by {\\fuse} in 2002 are analyzed in the context of a wind-eclipse model. This analysis shows that the wind of the eruptor obeyed a very fast velocity law in 2002, which is consistent with the line-driving mechanism. Large amplitude line-profile variations on the orbital period are shown to be due to the eclipse of {\\stB} by the wind of {\\stA}, although the eclipse due to gas flowing in the direction of {\\stB} is absent. This can only be explained if the wind of {\\stA} is not spherically symmetric, or if the eclipsed line radiation is ``filled-in\" by emission originating from somewhere else in the system, e.g., in the wind-wind collision region. Except for a slightly lower wind speed, the {\\orfeus}/BEFS spectrum is very similar to the spectrum obtained by {\\fuse} at the same orbital phase: there is no indication of the impending eruption. However, the trend for decreasing wind velocity suggests the occurrence of the ``bi-stability\" mechanism, which in turn implies that the restructuring of the circumbinary environment caused by the transition from ``fast, rarefied wind\" to ``slow, dense wind\" was observed as the eruptive event. The underlying mechanism responsible for the long-term decrease in wind velocity that precipitated this change remains an open issue. ",
        "introduction": "{\\hd} (Sk\\,78, AV\\,229) is the most luminous eclipsing binary system in the Small Magellanic Cloud (SMC) and a member of the young stellar cluster NGC\\,346. It is an ideal object for studying interacting winds and the evolution of very massive stars in binaries. The system achieved notoriety in 1994, when its primary star (hereafter {\\stA}; M$\\sim$50 \\msun) was discovered to be undergoing an eruptive event of a yet-undetermined nature. Although the system had been monitored since the early 1980s, the sudden brightening by $\\sim$2 magnitudes in late 1993 went unrecorded except for the visual estimates of A. Jones \\citep{Bateson93}. An intensive ultraviolet (UV) monitoring campaign was initiated in mid-1994, and this database provides a unique record of the declining phase of the outburst. However, the most intriguing question that {\\hd} raises involves the mechanism responsible for the eruption. We know that between 1987 and 1995 its UV spectrum went successively through all the ``late\" nitrogen Wolf-Rayet (WR) classes -- WN6 $\\rightarrow$ WN7 $\\rightarrow$ WN8 $\\rightarrow$ WN11 $\\rightarrow$ WN6 -- while also exhibiting significant changes in wind velocity and visual brightness; see \\citet{Koenigsberger04} for a general review. Although some of the characteristics of the eruption are similar to those observed in Luminous Blue Variables (LBVs), both its extraordinary luminosity ($\\sim3 \\times 10^6$~{\\lsun}) and its increase in luminosity during eruption \\citep[$\\sim10^7$~\\lsun;~][]{Drissen01} place it in a category currently shared only with $\\eta$ Carinae.  The second component of the eclipsing system (hereafter {\\stB}; M$\\sim$28 \\msun) also appears to be a WR star. It has been difficult to isolate its spectral characteristics from the combined spectrum; see \\citet{Koenigsberger04} for a discussion of this issue, which is further complicated by the possibility that some emission at line frequencies arises in a wind-wind colliding region \\citep{Moffat98}. If {\\stB} is indeed a WNE star, then it is the evolved remnant of the originally more massive star of the binary system, and is therefore a good pre-supernova candidate. Alternately, it could simply be a less massive star whose outer layers have been significantly modified by a mass-transfer process, which has led to the current state of {\\stA}.   The spectrum of {\\hd} indicates the presence of a third star. It can be seen in the continuum \\citep{Breysacher91}, in the optical spectra obtained in the 1980s \\citep{Niemela88}, and in the ``stationary\" UV photospheric lines that are visible in {\\it HST/STIS} spectra \\citep{Koenigsberger02}. This component (hereafter {\\stC}) is likely to be an O4--6 supergiant that could be gravitationally bound to the {\\stA}$+${\\stB} pair with a very long orbital period. Thus, it might be responsible for the instability of the system \\citep{Koenigsberger94}. Alternately, {\\stC} could simply be a line-of-sight object whose light contaminates the spectrum of the {\\hd} system. The presence of this third star complicates the interpretation of an already complicated system. However, because the evolution of massive stars plays such an important role in many astrophysical phenomena, it is important to understand the mechanisms responsible for the peculiar behavior of {\\hd}.   In this paper, we analyze a set of phase-resolved spectroscopic observations of {\\hd} obtained with the {\\it Far Ultraviolet Spectroscopic Explorer} (\\fuse) satellite. These data were obtained over an interval spanning many orbital cycles, but include a subset of 8 observations taken over $\\sim$6 orbits, which permit orbital changes in the structure of the wind of {\\stA} to be assessed. In addition, we use an archival spectrum obtained by {\\orfeus} shortly before {\\hd} erupted. This fortuitous observation constrains the time scale for the development of the instability leading to the eruptions. Our analysis is organized as follows. In \\S\\ref{obs}, we describe the observational material. The {\\fuse} light curve is presented in \\S\\ref{fuselc}, while the far-UV spectrum of {\\hd} and its orbital variability are discussed in \\S\\ref{fusesp}. In \\S\\ref{model} we analyze the orbital variations in terms of a wind-eclipse model, and discuss the implications. Additional discussion and our conclusions are presented in \\S\\ref{concl}. ",
        "conclusions": "\\subsection{The Wind Structure of Star~A}  We have used phase-resolved wind profiles of the {\\ion{P}{5}} resonance doublet obtained by {\\fuse} in 2002 together with a simple wind eclipse model to constrain the velocity law of {\\stA}, the relative continuum intensities of {\\stA} and {\\stB}, and the physical dimensions of the {\\hd} system. We find that, to a first approximation, the {\\ion{P}{5}} wind profiles require that the wind of {\\stA} is expanding with a ``fast\" velocity law, as is typical of outflows driven by radiation pressure in spectral lines. Systematic variations of profile strength and morphology as a function of orbital phase are well reproduced by a combination of ``wind eclipses\" of {\\stB} and physical eclipses of {\\stA} by {\\stB}.  However, there are four important discrepancies between the observations and the synthetic line profiles obtained from our simplified model. The first consists of the presence of an emission ``spike\" near the rest velocity in the components of the {\\ion{P}{5}} resonance doublet that is not predicted by the model. We show in \\S\\ref{emC} that this problem can be resolved by assuming that the excess emission arises from the third component of the system, {\\stC}. The presence of a {\\ion{P}{5}} wind feature is consistent with the classification of {\\stC} as an O4--O6 supergiant on the basis of its photospheric absorption lines \\citep{Koenigsberger01}. Unfortunately, solutions for the three remaining discrepancies are less straightforward.   \\noindent{\\bf Absence of a wind eclipse at $v \\geq 0:$} The wind eclipse around phase $\\sim$0.0 produces reduced emission over the range of velocities associated with the column of material projected against the disk of {\\stB}. We have shown that the wind of {\\stA} expands rapidly, so that it achieves its terminal speed before it reaches the radius of its relative orbit about {\\stB}. Thus, at phases immediately before and after $\\phi=$0.00, the entire emission lobe should be significantly reduced in intensity by the wind eclipse. However, the observations show the expected effect only at $v \\leq 0$~{\\kms}; i.e., only on the blue side of the P~Cygni profile.  Thus, it seems as if only the portion of the wind of {\\stA} that is approaching the observer has the geometry and kinematics assumed by the model. We see two possible explanations for this discrepancy: \\begin{enumerate} \\item There is an extended region of {\\ion{P}{5}} emission associated with {\\stB} that compensates for the absorption produced by the wind eclipse. This emission could come from the surface of the shock cone associated with the WWC region, which effectively truncates the spherical symmetry of the outflow.  \\item The wind of {\\stA} is not spherically symmetric, but is significantly perturbed in the direction of {\\stB}. However, the asymmetry must be present in the wind acceleration region, which we have shown lies very close to the surface of {\\stA}, far from the expected location of the WWC stagnation point. Hence, if this explanation is correct, we conclude that the wind structure of {\\stA} in the direction of its companion differs from that of other directions, as opposed to a wind structure that is truncated by the WWC region.  This difference could either be {\\em intrinsic} to {\\stA} or could be the result of ``sudden radiative braking\" \\citep{Gayley97}. Note that if the wind structure towards the companion is non-standard, then the WWC surfaces that are drawn in Fig.~\\ref{wwcgeom} are not valid. For example, a slower wind towards the companion would have the effect of moving the stagnation point towards {\\stA}.  \\end{enumerate}   \\noindent{\\bf Strength of P~Cygni Absorption at $\\phi=$0.36:} If {\\ion{P}{5} $\\lambda$1117} is formed exclusively in the wind of {\\stA}, then the eclipse model predicts that the absorption trough of its P~Cygni profile should exhibit significant weakening during secondary eclipse, because {\\stB} occults the column of material responsible for producing this absorption. However, only a modest degree of weakening is observed at slow velocities; see, e.g., Fig~\\ref{p5montage} and Fig.~\\ref{p5montmSk80}. At large velocities, $v \\geq v_\\infty$, the absorption seems to be even more pronounced during secondary eclipse. A natural explanation for this behavior is that the wind of {\\stB} {\\em also emits and absorbs} {\\ion{P}{5}}; i.e., that our assumption concerning the unique association of {\\ion{P}{5}} with {\\stA} is not correct. Unfortunately, the present version of the wind-eclipse model is not able to test this idea, since a rigorous calculation must include detailed radiative transfer through two stellar winds and the WWC region. Instead, we simply note that if the spectrum of {\\stB} does posses a P~Cygni wind profile in the {\\ion{P}{5}} resonance doublet, then it differs significantly from the spectrum of the WR star in {Sk\\,108}.  \\noindent{\\bf Extent of the P~Cygni absorption edge:} The final discrepancy is the extent of the P~Cygni absorption  edge. Although {\\vinf $=$ 1750~\\kms} for {\\stA}, the edge velocity is $v_{edge} \\geq -2200$~{\\kms}. The greater extent of $v_{edge}$ is a well-known phenomenon in O-type and WR stars, so it is not too surprising to find it in {\\hd} as well. What is surprising, however, is that the extended, ``soft\" blue edge of the P~Cygni absorption trough does not vary with orbital phase. Since {\\stB} also drives a strong stellar wind, it is possible that its radiation field further accelerates the outflow from {\\stA} that has already achieved {\\vinf}. However, since this speculation cannot be tested directly with the present data set, the issue must remain open for now.  \\subsection{New Insights into the Eruption Mechanism}  The changes in the wind of {\\hd} prior to the eruptions in 1993--1994 were characterized by a progressive decrease in terminal velocity, which was accompanied by a systematic increase in density \\citep{Koenigsberger98b}. The first sign that a significant change was occurring in {\\hd} can be found in its {\\iue} spectra of 1986 \\citep{Koenigsberger04}, although peculiarities were already present in the early 1980s \\citep{Niemela88}. The amplitude of the perturbation seems to have grown gradually over the subsequent $\\sim$13 years until some kind of critical state was reached. This critical state produced the sudden eruption in 1993, which was followed by a second, stronger eruption in 1994. According to the visual data of Albert Jones \\citep[reproduced in][]{Koenigsberger04}, the eruption started around HJD 2,449,299 (1993 November 7), which is only 51 days (2.6 orbital cycles) after the {\\orfeus}/BEFS spectrum was acquired.   Figure~\\ref{cforfeus} compares the {\\ion{P}{5} $\\lambda$1117} wind profile in the {\\orfeus}/BEFS spectrum with its counterpart in the {\\fuse} P2230102 spectrum. Both spectra were obtained at essentially the same orbital phase, though they are separated by 168 orbital periods. In addition to slightly stronger P~Cygni absorption around $-500$~{\\kms}, the only morphological difference between the two spectra is the location of the discrete absorption component that marks the position of terminal speed: in 1993 {\\vinf $=$ 1530~\\kms}, while {\\vinf $=$ 1750~\\kms} in 2002. Hence, the trend for decreasing wind velocity persisted from the early 1980s up to $\\sim$51 days before the start of the eruption.   Evidently, the overall structure of the wind of {\\stA} was very similar just prior to the first eruption in 1993 and in 2002, when {\\hd} was in a quiescent state. Hence, we infer that the wind must have been driven by the same mechanism at both epochs; and we have argued on the basis of the rapid acceleration required to model the {\\ion{P}{5}} wind profile that in 2002, this mechanism was radiation pressure in spectral lines. Consequently, in 1993 the conditions in the star must have changed from supporting a ``normal\" radiatively driven outflow to a state of significantly enhanced mass loss in less that 51 days; i.e., whatever set of critical conditions initiated the eruption occurred within a remarkably short time interval. This rapid reconfiguration of the properties of the wind provides a strong constraint on the unknown mechanism responsible for the eruption.  The fact that the wind velocity decreased systematically for many years before the eruption suggests that some critical limit was reached in the autumn of 1993. If the decreasing wind velocity is associated with an increasing wind density, as appears to be the case from the growth of emission line intensities over the 1980--2000 timescale \\citep{Koenigsberger04}, it is possible that the critical conditions are related to the ``bi-stability limit\" \\citep{Lamers95,Vink99}. In this case, the eruptions would be interpreted as the observable manifestation of the transition from the ``fast, rarefied wind\" side of the bi-stability limit to the ``slow, dense wind\" side and the concomitant re-structuring of the circumbinary environment. Since two major eruptions occurred within $\\sim$1.5 years, it appears that {\\stA} remained very close to its ``bi-stability limit\" for this long before beginning to relax to its pre-outburst state.  Of course, the crucial unanswered question is what underlying evolutionary process caused {\\stA} to move towards its ``bi-stability limit\" in the first place, e.g., by triggering enhanced mass-loss rate or decreasing the velocity of the radiatively driven outflow. Identifying this trigger mechanism in {\\hd} is certain to improve our understanding of eruptive phenomena that occur in other massive stars and binary systems. Since these outbursts may drive significantly greater mass loss than possible via stellar winds, they have great potential to alter both the evolutionary histories of these stars and the yields of chemically enriched material they provide to their local environments."
    },
    "0606/astro-ph0606475_arXiv.txt": {
        "abstract": "We cross-correlate the new 3 year Wilkinson Microwave Anistropy Probe (WMAP)  cosmic microwave background (CMB) data with the NRAO VLA Sky Survey (NVSS) radio galaxy data, and find further evidence of late integrated Sachs-Wolfe (ISW) effect taking place at late times in cosmic history. Our detection makes use of a novel statistical method \\cite{Baldi et al. 2006a, Baldi et al. 2006b} based on a new construction of spherical wavelets, called needlets. The null hypothesis (no ISW) is excluded at more than 99.7\\% confidence. When we compare the measured cross-correlation with the theoretical predictions of standard, flat cosmological models with a generalized dark energy component parameterized by its density, $\\omde$, equation of state $w$ and speed of sound $\\cs2$, we find $0.3\\leq\\omde\\leq0.8$ at 95\\% c.l., independently of $\\cs2$ and $w$. If dark energy is assumed to be a cosmological constant ($w=-1$), the bound on density shrinks to $0.41\\leq\\omde\\leq 0.79$. Models without dark energy are excluded at more than $4\\sigma$. The bounds on $w$ depend rather strongly on the assumed value of $\\cs2$. We find that models with more negative equation of state (such as phantom models) are a worse fit to the data in the case $\\cs2=1$ than in the case $\\cs2=0$. ",
        "introduction": "The most outstanding problem in modern cosmology is understanding the mechanism that led to a recent epoch of accelerated expansion of the universe. The evidence that we live in an accelerating universe is now compelling. The luminosity distance at high redshift ($z\\sim 1$) measured from distant type Ia supernovae is consistent with a negative deceleration parameter ($q_0<0$ at $\\sim 3\\sigma$) and shows strong evidence of a recent transition from deceleration to acceleration \\cite{Riess et al. 2004, Riess et al. 1998, Perlmutter et al. 1999}. The amount of clustered matter in the universe, as detected from its gravitational signature through a variety of large scale probes (redshift surveys, clusters of galaxies, etc.)  cannot be more than $\\sim 1/3$ of the total content of the universe \\cite{Springel et al. 2006}. Observations of the cosmic microwave background (CMB) anisotropy have constrained the value of cosmological parameters with outstanding precision. The recent WMAP data (\\cite{Bennett et al. 2003, Spergel et al. 2003, Spergel et al. 2006}) have shown that the total density of the universe is very close to its critical value. Taken together, these results are a strong indication in favor of a non-null cosmological term, which would at the same time explain the accelerated expansion of the universe and provide the remaining $\\sim 2/3$ of its critical density.  The precise nature of the cosmological term which drives the accelerated expansion, however, remains mysterious. The favoured working hypothesis is to consider a dynamical, almost homogeneous component (termed {\\em dark energy}) with negative pressure (or, equivalently, repulsive gravity) and an equation of state $w\\equiv p/\\rho<-1/3$ \\cite{Peebles & Ratra 1988, Caldwell et al. 1998, Wang et al. 2000, Peebles & Ratra 2003}. Such a framework helps alleviating a number of fundamental problems arising when a constant cosmological term is interpreted as the energy density of the vacuum \\cite{Weinberg 1989}.  One key indication of an accelerated phase in cosmic history is the signature from the integrated Sachs-Wolfe (ISW) effect \\cite{Sachs & Wolfe 1967} in the CMB angular power spectrum. This is directly related to variations in the gravitational potential: in particular, it traces the epoch of transition from a matter-dominated universe to one dominated by dark energy. This effect (which is usually called {\\em late ISW}, as opposed to a {\\em early ISW} generated during the radiation-matter transition), shows up as a contribution in the low multipole region of the CMB spectrum. A detection of a late ISW signal in a flat universe is, in itself, a direct evidence of dark energy. Furthermore, the details of the ISW contribution depend on the physics of dark energy, and are therefore a powerful tool to better understand its nature. Unfortunately, the low multipole region of the angular power spectrum is also  the most affected by cosmic variance, making the extraction of the ISW signal a difficult task.  A useful way to separate the ISW contribution from the total signal is to cross-correlate the CMB anisotropy pattern (imprinted during the  recombination epoch at $z\\sim 1100$)  with tracers of the large scale structure (LSS) in the local universe \\cite{Crittenden & Turok 1996}. Detailed predictions of the ability to reconstruct the ISW using this technique were obtained by a number of authors \\cite{Cooray  2002, Hu & Scranton 2004, Afshordi 2004, Corasaniti et al. 2003, Pogosian et al. 2005}. This kind of analysis has been performed several times during the past few years, using different CMB data sets and various tracers of clustering. The first detection of the ISW \\cite{Boughn & Crittenden 2004, Boughn & Crittenden 2005} was obtained  by combining the WMAP 1st year CMB data with the hard X-ray background observed by the High Energy Astronomy Observatory-1 satellite (HEAO-1 \\cite{Boldt 1987}) and with the radio galaxies of the NRAO VLA Sky Survey (NVSS \\cite{Condon et al. 1998}). The positive correlation with NVSS was later confirmed by the WMAP team \\cite{Nolta et al. 2004}. Other large scale structure tracers that led to similar positive results were the APM galaxy survey \\cite{Maddox et al. 1990}, the Sloan Digital Sky Survey (SDSS \\cite{York et al. 2000}) and the near infrared 2 Micron All Sky Survey eXtendend Source Catalog (2MASS XSC \\cite{Jarrett et al. 2000}) \\cite{Fosalba et al. 2003, Scranton et al. 2003, Fosalba & Gaztanaga 2004, Afshordi et al. 2004, Padmanabhan et al. 2005, Cabre et al. 2006}.  A somewhat different strategy to attack the problem was recently adopted by other authors, who attempted to seek the ISW signal in spaces other than the pixel space of the maps or the harmonic space of the angular power spectrum \\cite{Vielva et al. 2006, McEwen et al. 2006}. This approach relies on spherical wavelets as a tool to exploit the spatial localization of ISW (at large angular scales) in order to get a more significant detection of the effect.  The purpose of the present paper is twofold. On one side, we want to perform a further analysis of the CMB-LSS cross-correlation, in order to obtain an independent check on previous results. We combine the recent 3rd year release of WMAP CMB sky maps with the radio galaxy NVSS catalogue, and carry out our investigation in wavelet space. We make use of a new type of spherical wavelets, the so-called {\\em needlets} \\cite{Narcowich et al. 2006, Baldi et al. 2006a, Baldi et al. 2006b} which have a series of advantages over previously used wavelets, as will be described in detail later. This then represents at the same time a check on previous results \\cite{Vielva et al. 2006, McEwen et al. 2006} and a significant improvement of the statistical and technical aspects of the problem. On the other side, we follow a rather general approach to dark energy modelization, as first proposed in \\cite{Hu 1998}. Within this framework the phenomenology of dark energy is characterized by three physical parameters: its overall density $\\omde$,  its equation of state $w$, and the sound speed $c_s^2$. This parameterization has the advantage of being model independent, allowing one to encompass a rather broad set of fundamental models, and of giving a more realistic description of the dark energy fluid, for example accounting for its clustering properties, a feature that was shown to have quite a strong effect on theoretical predictions \\cite{Weller & Lewis 2003}. As shown in \\cite{Hu & Scranton 2004, Bean & Dore 2004, Corasaniti et al. 2003} the ISW signature can in principle be able to set constraints on the parameters of this generalized dark energy scenario: however, previous analyses of the ISW from CMB-LSS cross-correlation made a number of unrealistic simplyfing assumptions on the dark energy component and were only able to either found confirmations for its existence by constraining its density, or to set limits on its equation of state under restrictive hypotheses on its clustering properties (one notable exception being the analysis performed in \\cite{Corasaniti et al. 2005} which applied a parameterization similar to ours to make a likelihood analysis of the cross-correlation data points estimated in \\cite{Gaztanaga et al. 2004}). Our approach is more ambitious, as we attempt a more realistic description of dark energy and derive constraints on the combined set of three above mentioned parameters.  The paper is organized as follows. In Section \\ref{data} we describe the dataset we use to perform our analysis. In Section \\ref{techniques} we  outline the theoretical predictions of the expected ISW signal and discuss the statistical techniques we apply to extract it from the data. In Section \\ref{results} we present our results and the derived constraints on dark energy. Finally, in Section \\ref{conclusions} we discuss our main findings and conclusions. ",
        "conclusions": "\\label{conclusions}  We have analyzed the WMAP 3 year CMB temperature data, in conjunction with the NVSS radio galaxy survey, and found further evidence of a correlation between the CMB fluctuation pattern and the local distribution of matter, consistent with an ISW effect taking place at a late epoch of cosmic evolution. When a flat universe is assumed (as suggested by CMB observation) the detection of a late ISW signature is a strong evidence in favour of a dark energy component.  Our findings are based on a new construction of spherical wavelets that has a number of advantages with respect to previous studies. The presence of a correlation between the CMB and the LSS is established with a high level of confidence.  We have also improved the treatment of the dark energy component, introducing a more general parameterization than those used is similar earlier analysis. Quite interestingly, we find that although the case for a non zero dark energy contribution to the total density is compelling and robust, the constraints on $w$ do depend on the assumed clustering properties of the dark energy component, namely its sound speed $\\cs2$. Phantom models, and also the ordinary cosmological constant case $w=-1$, perform worse when a quintessence behaviour $\\cs2=1$ is assumed. This is due to the fact that there exist models with $w\\sim -0.4$ which predict more correlation at smaller angular scales ($\\theta\\sim 2^\\circ$). This is an intriguing result, that could imply a ISW effect taking place at redshifts as high as $z\\sim 1$, earlier than expected in the cosmological constant case.  A similar preference for larger values of $w$ in quintessence models was also found in \\cite{McEwen et al. 2006}.  Whether this is an indication of interesting physics taking place between the dark energy and dark matter components is a subject that requires further investigation. Clearly, the observation of ISW is proving quite promising as a tool to answer the questions arising from the mysterious nature of dark energy. While the CMB data have reached a great degree of accuracy on the angular scales that are more relevant for the detection of ISW, deeper redshift surveys and better catalogues can, in the future, improve the tracing of the local matter distribution, thus allowing to reduce the errors on the cross-correlation determination."
    },
    "0606/astro-ph0606196_arXiv.txt": {
        "abstract": "We experimentally and numerically tested the separability of two independent equally--luminous monochromatic and white light sources at the diffraction limit, using Optical Vortices (OV), related to the Orbital Angular Momentum (OAM) of light. The diffraction pattern of one of the two sources crosses a phase modifying device (fork--hologram) on its center generating the Laguerre--Gaussian (L--G) transform of an Airy disk. The second source, crossing the fork--hologram in positions different from the optical center, acquires different OAM values and generates non--symmetric L--G patterns. We formulated a criterion, based on the asymmetric intensity distribution of the superposed L--G patterns so created, to resolve the two sources at angular distances much below the Rayleigh criterion. Analogous experiments carried out in white light allow angular resolutions which are still one order of magnitude below the Rayleigh criterion. The use OVs might offer new applications for stellar separation in future space experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606643_arXiv.txt": {
        "abstract": "{\\sl Planck\\/} will be the first mission to map the entire cosmic microwave background (CMB) sky with mJy sensitivity and resolution better than $10'$. The science enabled by such a mission spans many areas of astrophysics and cosmology.  In particular it will lead to a revolution in our understanding of primary and secondary CMB anisotropies, the constraints on many key cosmological parameters will be improved by almost an order of magnitude (to sub-percent levels) and the shape and amplitude of the mass power spectrum at high redshift will be tightly constrained. ",
        "introduction": "{\\sl Planck\\/} will be the first mission to map the entire cosmic microwave background (CMB) sky with mJy sensitivity and resolution better than $10'$ \\cite{BlueBook}. The science enabled by such a mission spans many areas of astrophysics and cosmology, but in this short proceedings I can focus on only a few.  (Further discussion of the cosmological science enabled by {\\sl Planck\\/} was covered by Lloyd Knox in his talk at this meeting.)   In particular I want to focus on the dramatic revolution {\\sl Planck\\/} will represent in the study of primary CMB anisotropies and the universe at $z=10^3$, with its implications for low-$z$ studies such as those of dark energy.  I also want to make a push for a CMB-centric view of structure formation which emphasizes the exquisite constraints on large-scale structure that we already have from the CMB at high-$z$.  Before I begin with these science topics, it is important to remind ourselves how revolutionary {\\sl Planck\\/} will be.  In addition to wider frequency coverage (crucial for control of foregrounds) and better sensitivity than {\\sl WMAP\\/}, {\\sl Planck\\/} has the resolution needed to see into the damping tail of the anisotropy spectrum.  In fact {\\sl Planck\\/} will be the first experiment to make an almost cosmic variance limited measurement of the temperature anisotropy spectrum around the $3^{\\rm rd}$ and $4^{\\rm th}$ acoustic peaks.  \\begin{figure} \\begin{center} \\resizebox{2.7in}{!}{\\includegraphics{irvine06_1a.eps}} \\resizebox{2.7in}{!}{\\includegraphics{irvine06_1b.eps}} \\end{center} \\caption{Forecast measurements of the temperature anisotropy power spectrum {}from 4 years of {\\sl WMAP\\/} or 1 year of {\\sl Planck\\/} data assuming nominal sensitivities.  We have chosen the same binning scheme to show the advantage that higher resolution and sensitivity confers on {\\sl Planck\\/} for high-$\\ell$ science.  Figures from \\protect\\cite{BlueBook}.} \\label{fig:tt} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{2.7in}{!}{\\includegraphics{irvine06_2a.eps}} \\resizebox{2.7in}{!}{\\includegraphics{irvine06_2b.eps}} \\end{center} \\caption{Forecast measurements of the polarization anisotropy power spectrum {}from 4 years of {\\sl WMAP\\/} or 1 year of {\\sl Planck\\/} data assuming nominal sensitivities.  One can see clearly the advantage that higher sensitivity confers on {\\sl Planck\\/} for polarization science. Figures from \\protect\\cite{BlueBook}.} \\label{fig:ee} \\end{figure}  What does this dramatic increase in our knowledge of the temperature and polarization anisotropy spectra tell us about cosmology, fundamental physics and the formation of structure?  Here I will highlight just a few areas where we expect a large impact. ",
        "conclusions": "{\\sl Planck\\/} will provide a dramatic advance in our knowledge of primary and secondary CMB anisotropies.  The constraints on many key cosmological parameters will be dropped to percent, or sub-percent, levels and the shape and amplitude of the mass power spectrum at high redshift will be tightly constrained.  Beyond our desire to know the basic parameters of the universe accurately, and to perform truly precision tests of our cosmological model, the increase in precision will be important for a host of low redshift experiments, including those that aim to constrain the nature of the dark energy.  I would like to thank the organizers of this conference for a pleasant and productive meeting, and Daniel Eisenstein and Wayne Hu for conversations and collaborations upon which some of this work rests.  I am grateful to the many members of the {\\sl Planck\\/} collaboration who have labored tirelessly to make {\\sl Planck\\/} a reality, and especially to Charles Lawrence for his leadership and tireless enthusiasm -- and the chocolate donuts. MJW was supported in part by NASA."
    },
    "0606/astro-ph0606207.txt": {
        "abstract": "The Low Mach Number Approximation (LMNA) is applied to 2D hydrodynamical modeling of Type I X-ray bursts on a rectangular patch on the surface of a non-rotating neutron star. Because such phenomena involve decidedly subsonic flows, the timestep increase offered by the LMNA makes routine simulations of these deflagrations feasible in an environment where strong gravity produces significant stratification, while allowing for potentially significant lateral differences in temperature and density. The model is employed to simulate the heating, peak, and initial cooling stages in the deep envelope layers of a burst. During the deflagration, B\u00e9nard-like cells naturally fill up a vertically expanding convective layer. The Mach number is always less than 0.15 throughout the simulation, thus justifying the low Mach number approximation. While the convective layer is superadiabatic on average, significant fluctuations in adiabaticity occur within it on subconvective timescales. Due to convective layer expansion, significant compositional mixing naturally occurs, but tracer particle penetration through the convective layer boundaries on convective timescales is temporary and spatially limited. Thus, mixing occurs on the relatively slow burst timescale through thermal expansion of the convective layer rather than from mass penetration of the convective layer boundary through particle convection. At the convective layer boundaries where mixing is less efficient, the actual temperature gradient more closely follows the Ledoux criteria. ",
        "introduction": "In the thermonuclear flash model, Type I X-ray bursts (henceforth referred to as \\textit{bursts}) are understood to be caused by the explosive ignition of hydrogen and/or helium gas, which have accreted onto the outer surface of neutron stars from relatively low-mass, binary companion donors. Extensive theoretical calculations using diffusional-thermal and 1D hydrodynamical models have successfully reproduced many of the general observational features of these bursts, such as the energies involved ($\\sim$10$^{38}$-10$^{39}$ ergs), their rise times (seconds), durations ($\\sim$ 10-100 seconds), spectral softening, and recurrence intervals (several hours). (For reviews, see Taam 1985, Lewin et al. 1995, Cumming 2004, and Strohmayer \\& Bildsten 2006.) However, multidimensional hydrodynamic modeling of bursts has been much more limited, partly due to limitations in computational resources. Using FLASH (Fryxell et al. 2000) based on the Piecewise Parabolic Method (PPM) (Colella \\& Woodward 1984), Zingale et al. (2001) simulated bursts as 2D detonations, albeit assuming very low mass accretion rates. Also using FLASH, Zingale et al. (2002) simulated bursts as 2D deflagrations using an artificial temperature perturbation in lieu of a realistic burning network. Spitkovsky et al. (2002) used the shallow water approximation to examine how flames would propagate during bursts around a two-layer neutron star surface, but their method is incompressible, assumes an ideal gas law, and does not account for thermal diffusion. In a different context, Dearborn et al. (2006) studied the related helium flash problem in the 3D cores of evolved giant stars using Djehuty (Baz\u00e1n 2003), which is based on an explicit Lagrange-Eulerian hydrodynamic method.  Our current modeling effort simulates the heating, peak, and initial cooling stages of the deep envelope layers undergoing a burst with 2D hydrodynamics using the Low Mach Number Approximation (LMNA) and remedies the shortcomings of previous work in several respects: (1) compressibility effects are included; (2) potentially significant lateral fluctuations in temperature and density are allowed; (3) the essential input physics are considered, including thermal diffusion, a realistic equation of state, and a burning network (a $3\\alpha$ nuclear burning network is reported here for simplicity, and more complete networks can be included); (4) no assumption is made regarding the nature of convection; (5) sound waves are naturally excluded from the domain, thereby eliminating both the acoustic timestep restriction and acoustic boundary reflection problems which can plague fully compressible simulations; and, (6) the computational time is reduced by a factor of 10-100 compared to fully compressible methods, due to the corresponding increase in the timestep.  Other approximation methods which eliminate the acoustic timestep restriction include the Boussinesq approximation (Spiegel \\& Veronis 1959; Miralles 2000), the anelastic approximation (Ogura \\& Phillips 1962; Glatzmaier 1984), and implicit methods (e.g., Deupree 2000). However, at present only the LMNA can successfully model the compressibility effects and significant lateral fluctuations in temperature and density which characterize deflagrations, as well as offer substantial increases in timestep and avoid acoustic boundary complications. While the LMNA is routinely used to model terrestrial combustion (e.g., Bayliss et al. 1992; McGrattan et al. 2004), only recently have efforts begun to adapt it to the astrophysical setting. Alternative astrophysical LMNA models have been recently developed by Bell et al. (2004a) and Almgren et al. (2006a, 2006b). Thus far, they have been used to examine 2D Landau-Darrieus planar flame instabilities during the early development of a Type Ia supernova (Bell et al. 2004b), Rayleigh-Taylor unstable flames in 2D and 3D (Bell et al. 2004c; Zingale 2005), and tested against the anelastic approximation and other fully compressible methods in a regime where all are valid (Almgren et al. 2006a, 2006b). The model by Bell et al. neglects background stratification, because the domain on which it was applied was much smaller than one pressure scale height. The model by Almgren et al. does include background stratification, and is thus more comparable to ours. However, our LMNA model differs from Almgren et al.'s in several important respects: (1) they evolve the density via the continuity equation, whereas we evolve the temperature via the energy equation; (2) they neglect thermal diffusion, because it is expected to be unimportant in their modeling of Type Ia supernova, whereas we include it in our model; (3) they allow a time-dependent background state, while presently, we assume it is time-independent, a condition which is easily relaxed and is reserved for future work; (4) we reformulate (using a novel function which will be described in detail below) the perturbative pressure gradient and buoyancy forces to accurately calculate the net vertical force in the presence of significant cancellation for large gravity.  In $\\mbox{\\S}$2, the motivation and essence of our LMNA model are presented. The results of applying the LMNA to simulate a Type I X-ray burst in 2D are described in $\\S$3. Next, $\\S$4 briefly describes verification studies which were performed. Finally, a discussion of the key findings and limitations of the present model, along with ideas for future development and applications, is presented in $\\S$5. ",
        "conclusions": "The major accomplishments and findings of this project are summarized as follows: (1) The low Mach number approximation has been developed, verified, and implemented to study astrophysical deflagrations where large vertical pressure variations exist and the Mach number $M$ is small. (2) When applied to subadiabatic initial conditions representing the pre-burst peak stage of a Type I X-ray burst, a vertically expanding convective layer of B\u00e9nard-like cells naturally develops, and the vertical extent of the larger cells matches that of the convective layer. The convective layer expands to two pressure scale heights during the burst progression. (3) Even at their maximum values, convective flow speeds are substantially subsonic ($M_{peak}<0.15$), while the deviation of the pressure from the hydrostatic base state is always at most of order $M^{2}$. (4) As the convective layer expands, fuel is naturally mixed into the convective layer, and mixing within the layer is very efficient. However, at the convective layer boundaries, less efficient mixing results in significant composition gradients, such that $\\grad$ more closely follows $\\grad_{L}$ there. Penetration on convective-timescales is limited and temporary. (5) Both sub- and superadiabaticity are found within the convective layer, but it is slightly superadiabatic on average. (6) Convection significantly affects energy transport.  In the present results, convection develops naturally as a consequence of superadiabatic gradients arising from heat inputted into the system by nuclear burning in a bursting layer. No model for convection is assumed; indeed, we have not been able to establish agreement with the predictions of mixing-length theory. Throughout the burst, the average values of the actual gradient $\\grad$ are best described as generally between $\\grad_{ad}$ and $\\grad_{L}$, while the instantaneous values are closer to $\\grad_{L}$. Moreover, some penetration occurs at the convective layer boundaries, where on average, $\\grad$ clearly deviates from $\\grad_{ad}$ to conform better to $\\grad_{L}$. Whether the Schwarzschild or Ledoux criteria is satisfied in regions where composition gradients exist in massive stars is still an open question in astrophysics. Semiconvection, a relatively slow mixing caused by composition gradients, is poorly understood, and models which examine the Schwarzschild vs. Ledoux gradients yield conflicting results (Merryfield 1995; Canuto 2000). Moreover, Canuto (2000) demonstrates the Schwarzschild criteria necessarily implies convective overshooting, and the Ledoux criteria also necessarily implies overshooting if convection is non-local. The current model includes all the key elements of convective and semiconvective processes, and it provides an opportunity to further examine this important issue.  Several limitations to our current model should be noted. (1) The contribution of energy generation $\\dot{s}_{12,\\alpha}$ due to subsequent $\\alpha$ captures on $_{6}^{12}C$ is neglected in the current computation. Calculations show $\\dot{s}_{12,\\alpha}$ becomes important post-burst peak due to the rise in ash concentration $Z$ as a result of burning, and its inclusion is expected to increase the duration and magnitude of the simulated burst. (2) At more advanced stages of the burst ($log$ EGR$_{max}$ $\\geq$ 19), the upper convective boundary has reached the upper domain boundary. Comparison with a model which has an extended height shows that dynamical results at these stages may be subject to upper boundary influences. However, these differences were relatively minor and did not affect the qualitative behavior described here. Moreover, because the current model does not include the true surface of the star, observable light curves can not be rigorously calculated. (3) While most of the computational domain remains degenerate throughout the burst sequence, calculations show that degeneracy decreases as the burst progresses, and by the end of the calculation, it begins to be lifted in the upper third of the domain where the densities are smallest. Thus, expansion effects may become important at more advanced stages of the burst, and the current assumption that the base state is time-independent may need to be relaxed to better model the dynamics which arise at these later times. (4) The current model neglects rotation and magnetism. We note that these effects can be incorporated within the Low Mach number formalism, and we expect to address them in future work. Nevertheless, the results presented here describe reasonable qualitative behavior of the flow field as a burst progresses. The substantially subsonic flows arising during the burst, the convective layer filling up with larger cells of roughly the same vertical extent, and the better agreement of $\\grad$ with $\\grad_{L}$ near the convective layer boundaries are likely to describe the behavior when these additional effects are included.  Presently, the LMNA method is a powerful computational tool and has been successfully applied to routinely simulate 2D Type I X-ray burst deflagrations, a problem which has thus far been intractable with other methodologies. Continuing to develop the algorithm will be a vital aspect of future work. Computationally, enhancements may include relaxing the time-independence of the hydrostatic base state, extending the model to 3D, implementing different coordinate systems, incorporating adaptive gridding techniques, and improving the input physics, such as incorporating rotation, more complete nuclear burning networks, and a sub-grid turbulence model. The LMNA model is well-suited to routinely model astrophysical deflagrations which occur during Type I X-ray bursts, the pre-ejection stage of classical novae, the pre-detonation stage of supernovae, and the hydrodynamics and burning in the cores of main sequence stars. The LMNA method represents a useful tool which enables the routine investigation of a wide variety of interesting and important astrophysical questions."
    },
    "0606/astro-ph0606332_arXiv.txt": {
        "abstract": "We present results from our two year study of ground-layer turbulence as seen through the 6.5-meter Magellan Telescopes at Las Campanas Observatory. The experiment consists of multiple, moderate resolution, Shack-Hartmann wavefront sensors deployed over a large 16 arcminute field. Over the two years of the experiment, the ground-layer turbulence has been sampled on eleven nights in a variety of seeing and wind conditions. On most nights the ground-layer turbulence contributes 10\\% to the total visible-band seeing, although a few nights exhibit ground-layer contributions up to 30\\%. We present the ground-layer turbulence on the sampled nights as well as a demonstration of its strength as a function of field size. This information is combined with data from a MASS-DIMM seeing monitor adjacent to the Magellan Telescopes to infer the annual ground-layer contribution to seeing at Las Campanas. ",
        "introduction": "\\label{sect:intro} The concept of ground-layer adaptive optics correction (GLAO) is to make a modest improvement in image quality over a large field of view (many arcmin) by compensating for the turbulence in the lowest layers of the atmosphere.  GLAO is complementary to full-atmosphere adaptive optics correction (AO), which delivers diffraction-limited imaging, but only over small fields of view ($<$ 1 arcmin).  Recent studies of the atmosphere at astronomical sites have been conducted using a variety seeing monitors, and also with balloons carrying micro-thermal sensors \\cite{1994A&A...284..311V,1998A&AS..130..141K,2001A&A...369..364A,MASS_3,2003RMxAC..19...23S,2003SPIE.4839..466W}. A number these studies provide measuremants of both the total seeing and the distribution of turbulence as a function of altitiude. The new results suggest that typically more than half of the turbulent power in the atmosphere occurs at low altitude ($\\lesssim$ 1.0 km). This ground-layer is thought to arise from the interaction of moving air with local topography, which may differ from site to site depending on wind direction and ground contour upstream of the telescope.  The presence of a strong, low altitude turbulence layer provides an opportunity to improve the seeing using an AO system with modest operating parameters \\cite{2000SPIE.4007.1022R}.  Near the ground, the wind speed is low so the crossing time for a turbulent cell is typically on the order of ten milliseconds or more.  Because the effective height of the ground layer is largest for the largest turbulent scales, a GLAO system favors correction of the lowest order modes.  For atmospheric layers near the telescope, the isoplanatic angle will be large. Thus, a GLAO system might operate at low frequency (100 Hz) and low order ($<$50 modes on a 6.5-m telescope) while significantly improving the image quality over a large field of view (many arcmin). Initial GLAO models predicted a gain of a factor of two in image size over fields of 10 arcmin or more, resulting in a factor of four gain in sensitivity for background limited observations \\cite{2000SPIE.4007.1022R,2003SPIE.4839..673T}. The potential of GLAO correction is evident, however most of the work to date has consisted of modeling based on a few atmospheric measurements.  Only limited amounts of data have been taken at the telescope.  Although several on-going studies reveal at least some information about ground-layer turbulence \\cite{2004SPIE.5490..758W,2005ApJ...634..679L}, none of the current experiments provides adequate empirical measurements of the parameters and performance of a GLAO system.  Early laser-tomography experiments, for example, are restricted to smaller fields of view.  Since in any event different observatory sites might be expected to exhibit different ground-layer conditions, we have chosen to conduct an experiment to directly measure the effect of ground-layer turbulence at the Magellan Telescopes.  This experiment was conducted as part the Giant Magellan Telescope (GMT) project, in order to study the implications of GLAO for the design of GMT and its instrumentation. ",
        "conclusions": "We have taken one of the first sets of moderate speed (100 Hz), wide-field (arcminutes), and moderate resolution (59 cm sub-aperture) wavefront sensor data, in order to asses the gains of a GLAO system. We have taken this data on eleven nights which span a range of meteorological, seeing, and reported ground-layer strength conditions. On the best nights, we find a RMS image size reduction of 30\\% over a 7 arcminute separation. On a typical night we find a 10\\% correction to a visible-band image."
    },
    "0606/astro-ph0606274_arXiv.txt": {
        "abstract": "Cataclysmic Variables (CVs) are close binary systems where mass is transferred from a red dwarf star to a white dwarf star via an accretion disk. The flickering is observed as stochastic variations in the emitted radiation both in the continuum and in the emission line profiles.\\\\ The main goal of our simulations is to compare synthetic Doppler maps with observed ones, aiming to constrain the flickering properties and wind parameters.\\\\ A code was developed which generates synthetic emission line profiles of a geometrically thin and optically thick accretion disk. The simulation allows us to include flares in a particular disk region. The emission line flares may be integrated over arbitrary ``exposure'' times, producing the synthetic line profiles. Flickering Doppler maps are created using such synthetic time series. The presence of a wind inside the Roche lobe was also implemented. Radiative transfer  effects in the lines where taken into account in order to reproduce the single peaked line profiles frequently seen in nova-like CVs. ",
        "introduction": "The flickering, or rapid variability, is observed as stochastic fluctuations in the emitted radiation. The typical timescales range from seconds to tenths of minutes, and the amplitudes vary from cents of magnitude to more than one magnitude. The flickering is a defining characteristic of cataclysmic variables, being frequently used to classify an object as a CV. Flickering is also observed in other classes of objects, as in some symbiotic stars \\cite{Miko90}. The first CV where flickering was observed is UX UMa \\cite{Line49}. Since then, many photometric studies were made aiming to locate the flickering source region in many systems. \\inlinecite{Diaz01} proposed a tomographic method to map the flickering source regions using line profile data. The flickering tomography was applied to V442 Oph system, and an isolated flickering source region could not be identified. The objective of our simulations is to generate flickering tomograms from synthetic line profiles and compare these tomograms with the observed maps, aiming to constrain some parameters of the flickering and locate its forming region.\\\\ The presence of winds in cataclysmic variables is noticed by strong wind driven lines in UV, i.e. C IV, and the occurrence of P-Cygni profiles. As some tomograms present emission at low velocities, we have implemented the presence of a optically thick wind, aiming to reproduce this behavior. The comparison of model tomograms with observations may help us to constrain the main wind parameters.\\\\ In this proceeding we present the main physical concepts and parameterizations contained in the simulations and some preliminary results that arise from them. ",
        "conclusions": "Simulations of synthetic accretion disc line profiles including flickering and wind were performed. From the simulations one can see that high S/N and high time-resolution spectra are needed in order to obtain information from flickering tomograms. The information contained in the low amplitude and high frequency flickering is lost first. The flickering information is also lost if long integration times are used. From the wind simulations we conclude that a line optical depth greater than 10 is needed to obtain single peaked wind line profiles.\\\\ As incoming work this code will be used to generate synthetic tomograms with flickering, which will be compared to observed flickering tomograms of V3885 Sgr, aiming to locate the sites of flickering production and constrain the flickering parameters in this system."
    },
    "0606/astro-ph0606568_arXiv.txt": {
        "abstract": "We present parameter estimation forecasts for present and future 3D cosmic shear surveys.  We demonstrate in particular that, in conjunction with results from cosmic microwave background (CMB) experiments, the properties of dark energy can be estimated with very high precision with large-scale, fully 3D weak lensing surveys. In particular, a 5-band, 10,000 square degree ground-based survey of galaxies to a median redshift of $z_m=0.7$ could achieve $1$-$\\sigma$ marginal statistical errors, in combination with the constraints expected from the CMB Planck Surveyor, of $\\Delta w_0=0.108$ and $\\Delta w_a=0.099$. We parameterize the redshift evolution of $w$ by $w(a)=w_0+w_a(1-a)$ where $a$ is the scale factor. Such a survey is achievable with a wide-field camera on a $4$ metre class telescope. The error on the value of $w$ at an intermediate pivot redshift of $z=0.368$ is constrained to $\\Delta w(z=0.368)=0.0175$. We compare and combine the 3D weak lensing constraints with the cosmological and dark energy parameters measured from planned Baryon Acoustic Oscillation (BAO) and supernova Type Ia experiments, and find that 3D weak lensing significantly improves the marginalized errors on $w_0$ and $w_a$ in combination, and provides constraints on $w(z)$ at a unique redshift through the lensing effect.  A combination of 3D weak lensing, CMB and BAO experiments could achieve $\\Delta w_0=0.037$ and $\\Delta w_a=0.099$. We also show how our results can be scaled to other telescopes and survey designs. Fully 3D weak shear analysis avoids the loss of information inherent in tomographic binning, and we also show that the sensitivity to systematic errors in photometric redshift is much less. In conjunction with the fact that the physics of lensing is very soundly based, the analysis here demonstrates that deep, wide-angle 3D weak lensing surveys are extremely promising for measuring dark energy properties. ",
        "introduction": "Our knowledge of cosmology has advanced considerably in recent years. Led by detailed measurements of the microwave background radiation and large-scale structure, many of the most important cosmological parameters are now known with good accuracy. This advance has come about principally through the all-sky maps of the microwave sky taken with the Wilkinson Microwave Anisotropy Probe (WMAP) (Bennett et al., 2003), supplemented by higher-resolution observations of the Arcminute Cosmology Bolometer Array Receiver (ACBAR) and Cosmic Background Imager (CBI) (Kuo et al., 2004; Pearson et al., 2003). When combined with large-scale structure information from the Anglo-Australian 2 degree field galaxy redshift survey (2dFGRS) (Colless et al., 2001; Percival et al., 2001), the Lyman-$\\alpha$ forest (Croft et al., 2002; Gnedin and Hamilton, 2002), and measurements of galaxy bias (Verde et al., 2002), the data establish the concordance model of an accelerating Universe dominated by dark energy and dark matter (Spergel et al., 2003; Spergel et al., 2006).  The acceleration of the Universe is also apparent in observations of distant supernovae (e.g. Riess et al., 2000).  The determination of the density, baryon content and expansion rate of the Universe shifts the major unanswered questions in cosmology to the nature of the dark matter and dark energy. Dark energy in particular can be probed through its cosmological effects on the distance-redshift relation and the growth rate of structure.  The question of the precise nature of the dark energy is a far-reaching one. We use the simplest phenomenological model of dark energy by parameterizing the equation of state of the vacuum, \\begin{equation} w\\equiv p/(\\rho c^2), \\end{equation} where $\\rho$ is its energy-density and $p$ is the dark-energy/vacuum-pressure. $w=w(a)$ may vary with scale factor. If it is found with high precision that $w=-1$, then the dark energy cannot be distinguished with large-scale measurements from a modification to the gravity law along the lines suggested by Einstein with the cosmological constant.  If, however, it can be established with a degree of certainty that $w$ differs from $-1$ at any redshift, then it cannot be associated with such a change to the gravity law, and is most naturally accounted for by a new field. This would be an extremely important discovery, and the time-evolution of the field would be a useful constraint on models. Some possibilities exist in the literature, such as those proposed by Ratra and Peebles (1988), but none is a clear favourite candidate.  Weak lensing is a very attractive proposition for studying dark energy, as it is sensitive to both of these effects, and, equally importantly, the physics of weak lensing is well understood.  A key part of this is that it is sensitive to the distribution of matter in the Universe, regardless of its form.  Furthermore, since weak lensing analysis can be done in a way which is either dependent on the distance-redshift relation alone (see e.g. Taylor et al., 2006; Jain \\& Taylor, 2003) or on both the distance-redshift relation and the growth factor (this paper), it can in principle distinguish between modified gravity models and dark energy models.  The main lesson of the field of microwave background astronomy is that with well-understood physics, robust results can be obtained with high precision.  Weak lensing observations are, however, a technical challenge, as the imaging requirements are severe.  Thus, it has only been in the last five years or so that the first measurements of cosmic shear have appeared (Bacon, Refregier and Ellis, 2000; Kaiser, Wilson and Luppino, 2000; van Waerbeke et al., 2000; Wittman et al., 2000). Weak lensing measurements to date have concentrated on obtaining the matter density parameter $\\Omega_m$ and the amplitude of mass density fluctuations (Hoekstra, Yee and Gladders, 2002; Jarvis et al., 2003; Rhodes et al., 2004; Heymans et al., 2004; Hoekstra et al., 2006; Semboloni et al., 2006). More ambitiously, weak lensing observations have started to put constraints on the equation of state of dark energy (Jarvis et al., 2005; Semboloni at al., 2006).  Theoretically, the prospects for determining dark energy properties (specifically its equation of state $w$) using weak lensing have been explored in a number of papers (e.g. Taylor et al., 2006; Hu and Tegmark, 1999; Huterer, 2002; Heavens, 2003; Refregier, 2003; Simon, King and Schneider, 2004; Takada and Jain, 2004; Song and Knox, 2004; Ishak et al., 2004; Ishak, 2005). The prospects for determining $w$ as a function of redshift $z$ are markedly improved when 3D information on the individual lensed sources is available.  Source distances could come from spectroscopic redshifts, but given the depth and the sky area required, they are more likely to be estimated from photometric redshifts.  With 3D information, the lensing pattern can be analyzed in shells at different distances (e.g. Hu, 1999; Hu and Jain, 2004; Ishak, 2005), or by analyzing the shear pattern as a fully three-dimensional field (Heavens, 2003).  It is the latter possibility which we investigate in this paper.  The statistical properties of the shear pattern are influenced by many cosmological parameters, including $w(z)$.  In this paper we extend the analysis of Heavens (2003) to small-angle surveys as well as computing the expected marginal errors on $w$ (and its evolution), using a Fisher matrix approach.  We investigate issues of depth vs area, and the number of photometric bands which should be used, to determine the dark energy properties as accurately as possible.  The main focus of the paper is in computing the expected statistical errors, but we do consider the impact of some systematics (Ishak et al., 2004; Bernstein, 2005; Huterer et al., 2005).  The layout of the paper is as follows:  in Section \\ref{s:method} we detail the transform method used and compute the covariance matrix of the transform coefficients; in Section \\ref{s:params} we outline how the expected statistical errors on parameters are calculated; in Section \\ref{s:Errors} we present the survey design and how we can scale to other surveys and in Section \\ref{Optimisation For a Wide-Field Lensing Survey} we  present an optimization of survey design and the parameter errors; in Section \\ref{Parameter Forecasts} we consider the synergy of 3D weak lensing with other dark energy probes and discuss future surveys and finally we give our conclusions in Section \\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions} In this paper we have presented a 3D weak lensing spectral method suitable for high-$\\ell$ studies, and investigated how well 3D weak lensing surveys could determine the equation of state of dark energy.  The accuracy which could be achieved if systematic errors can be controlled is impressively high, provided the surveys are analyzed in 3D: marginal statistical errors of $\\Delta w_0=0.108$, on the current value of $w\\equiv p/(\\rho c^2)$, and its evolution $w_a$ constrained to $\\Delta w_a=0.397$ are possible with a 10,000 square degree survey in 5 bands to a median source depth of $z_m=0.7$. At a pivot redshift of $z=0.37$ such an experiment could constrain $w(z)$ to  $\\Delta w(z=0.37)=0.0175$. Such a survey is possible with darkCAM, in conjunction with data from the Planck satellite. Even without Planck, the accuracy from 3D weak lensing alone is still impressively high, and better than any other dark energy probe considered on its own. The fact that the physics of 3D weak lensing is well-understood, combined with the small statistical error forecasts, makes 3D weak lensing a formidable prospect for advancing cosmology in the next decade. The errors on $w$ are comparable to, but a little better than, predictions from tomography (Hu and Jain, 2004; Ishak, 2005). The constraints on $w(z)$ at the pivot redshift and the figure of merit, $\\Delta w(z_{\\rm pivot})*\\Delta w_a$, of the experiments considered were also discussed.  We have investigated optimizing a wide field survey to measure the equation of state parameters $w_0$ and $w_a$ and found an optimal survey strategy of $z_m=1.0$ covering $2400$ square degrees for a $5$ optical band survey. We found that increasing the number of optical bands to $9$ or $17$ makes little difference to the marginal errors when the 3D weak lensing result is combined with a Planck prior. The effect of including infrared bands in a wide field survey was investigated by varying the photometric redshift error, it was found that adding $4$ infrared bands to a $5$ band optical survey improves the marginal constraints on $w_0$ slightly from $\\Delta w_0=0.108$ to $\\Delta w_0= 0.097$.  Three alternative dark energy probes were considered: a Planck CMB experiment; a WFMOS BAO experiment and a SNAP SNIa experiment. All possible combinations of experiments were considered and the figure of merit and pivot redshifts of the combinations shown. In such a competitive environment 3D weak lensing places strong constraints on the dark energy parameters and in combination with other experiments provides a unique degeneracy in the ($w_0$,$w_a$) plane which is manifest as a strong constraint at a particular pivot redshift.  We have addressed the issues of biased photometric redshift estimates (e.g. Ma, Hu and Huterer, 2005) and show that the method is relatively insensitive to this. We also investigated the effect that a sample of outliers, with poor photometric redshift estimates, would have on the predicted marginal errors. The effect of outliers on the marginal error of $w_0$ is small although the way in which such a sample is treated is important.  We have not considered errors due to the intrinsic alignment of galaxies (Heavens, Refregier and Heymans, 2000; Croft and Metzler, 2000; Catelan, Kamionkowski and Blandford, 2001; Crittenden et al., 2001; Jing, 2002), as these may be reduced to a negligible level by removing pairs which are close in photometric redshifts (Heymans and Heavens, 2003; King and Schneider, 2002). This procedure has already been demonstrated in the analysis of the COMBO-17 data (Heymans et al., 2004). We have also not addressed other issues of systematics, such as optical distortions, or possible alignment of foreground galaxies with shear (Hirata and Seljak, 2004), which may be reduced using techniques such as template fitting (King, 2005). Nevertheless, the fact that the statistical errors are very small is very encouraging. Clearly to achieve the accuracies quoted here is going to be a formidable challenge for control of systematics, but at least the statistical error forecasts are small enough that the promise of accurate measurement of the equation of state of dark energy may be realized."
    },
    "0606/astro-ph0606042_arXiv.txt": {
        "abstract": "Neutrino oscillations affect light element synthesis through the $\\nu$-process in supernova explosions. The $^7$Li and $^{11}$B yields produced in a supernova explosion of a 16.2 $M_\\odot$ star model increase by factors of 1.9 and 1.3 in the case of large mixing angle solution with normal mass hierarchy and $\\sin^{2}2\\theta_{13} \\ga 2 \\times 10^{-3}$ compared with those without the oscillations. In the case of inverted mass hierarchy or nonadiabatic 13-mixing resonance, the increment of their yields is much smaller. Neutrino oscillations raise the reaction rates of charged-current $\\nu$-process reactions in the region outside oxygen-rich layers. The number ratio of $^7$Li/$^{11}$B could be a tracer of normal mass hierarchy and relatively large $\\theta_{13}$, still satisfying $\\sin^{2}2\\theta_{13} \\le 0.1$, through future precise observations in stars having strong supernova component. ",
        "introduction": "In the final evolutionary stage of massive stars, most region of the stars except collapsing core explodes as supernova explosions. The collapsing core releases its gravitational energy with gigantic amount of neutrinos. The emitted neutrinos interact with nuclei in the exploding material and new species of nuclei are produced; this synthetic process is called the $\\nu$-process \\citep{de78,wh90}. There are several species produced through the $\\nu$-process. For light elements, $^7$Li and $^{11}$B are mainly produced through the $\\nu$-process (Woosley et al. 1990; Woosley \\& Weaver 1995; Yoshida, Emori, \\& Nakazawa 2000; Rauscher et al. 2002; Yoshida et al. 2004; Yoshida, Kajino, \\& Hartmann 2005). Some $^{19}$F is also produced through the $\\nu$-process \\citep{wh90,ww95,rh02}. For neutron-deficient heavy nuclei $^{138}$La and $^{180}$Ta are also produced through charged-current interactions with $\\nu_e$ \\citep{ga01,hk05}. Neutrino-driven winds from proto-neutron stars are considered to be one of the promoting sites for $r$-process heavy elements \\citep[e.g.,][]{ww94,tw94,ot00,tl04}.  Supernova explosion is one of the important sites for supplying $^7$Li and $^{11}$B as well as Galactic cosmic rays, AGB stars, and novae during Galactic chemical evolution (GCE) \\citep[e.g.,][]{fo00}. In previous studies, we showed that the amounts of $^7$Li and $^{11}$B strongly depend on the neutrino energy spectra and the total neutrino energy \\citep{yt04,yk05}. We also constrained the neutrino energy spectra from the gravitational energy of a proto-neutron star and GCE models \\citep{yk05}. In these studies it has been assumed that the neutrino spectra do not change in the supernova ejecta.  On the other hand, recent remarkable progress in neutrino experiments has confirmed the phenomenon of neutrino oscillations \\citep[e.g.,][]{mv04}. The experiments on atmospheric neutrinos \\citep[e.g.,][]{SK04}, solar neutrinos \\citep[e.g.,][]{SNO04}, and reactor neutrinos \\citep[e.g,][]{ap03,KL05} constrained most of parameter values in neutrino oscillations, such as squared mass differences and the mixing angles. Resultantly, the large mixing angle (LMA) solution turns out to be a unique solution for the 12- and 23-mixings. However, mass hierarchy between 1 and 3 mass eigenstates has not been clarified \\citep[e.g.,][]{SK04} and only upper limit of $\\sin^{2}2\\theta_{13}$ has been determined \\citep{ap03}.  Supernova neutrino is another promoting target for neutrino experiments. When SN 1987A occurred, Kamiokande group and IMB group found eleven and eight events of neutrino detection \\citep{hk87,IMB87}. Owing to the development of neutrino experiments, much larger events of the neutrino detection are expected when a supernova explosion occurs in neighboring galaxies. In order to evaluate the neutrino flux and their energy dependence, neutrino oscillations in supernova explosions have been investigated qualitatively \\citep{ds00} and quantitatively \\citep{tw01}. They showed that in the case of adiabatic resonance for $\\sin^{2}2\\theta_{13}$ (LMA-L in Takahashi et al. 2001) the transition probability from $\\nu_e$ to $\\nu_{\\mu,\\tau}$ changes from 0 to almost 1 in the O/C layer in their supernova model. Finally, the energy spectrum of $\\nu_e$ changes to the one close to the $\\nu_{\\mu,\\tau}$ spectrum emitted from neutrino sphere. In the case of nonadiabatic resonance for $\\sin^{2}2\\theta_{13}$ (LMA-S in Takahashi et al. 2001), the change of neutrino spectra is smaller. The effects from mass hierarchy and from the change of the density profile due to the shock propagation were also investigated \\citep{ts03a,ts03b}.  Since neutrino oscillations change neutrino spectra, it is expected that the amounts of $^7$Li and $^{11}$B change by the effect of neutrino oscillations \\citep{yk06}. During supernova explosions, neutral-current reactions such as $^4$He($\\nu,\\nu'p)^3$H and $^4$He($\\nu,\\nu'n)^3$He are important for $^7$Li and $^{11}$B production \\citep{yt04}. We note here that the total reaction rates of neutral-current reactions do not change by the neutrino oscillations. The energy spectrum summed up in all neutrinos and antineutrinos does not change by the oscillations. On the other hand, the reaction rates of charged-current reactions such as $^4$He($\\nu_e,e^-p)^3$He and $^4$He($\\bar{\\nu}_e,e^+n)^3$H are expected to increase by neutrino oscillations. As shown in \\citet{ds00} and \\citet{tw01}, the mean energies of $\\nu_e$ and $\\bar{\\nu}_e$ increase by the neutrino oscillations. This increase will raise the efficiency of the $^7$Li and $^{11}$B production. If we obtain some clear signals of neutrino oscillations in the abundances of $^7$Li and $^{11}$B, we would constrain the parameter values of neutrino oscillations from observations of light elements. This is a new procedure to constrain neutrino oscillation parameters completely different from the detections of supernova neutrinos.  In the present study, we investigate light element synthesis in supernova explosions taking account of the change of the neutrino spectra due to neutrino oscillations. We also evaluate the dependence of the yields of $^7$Li and $^{11}$B in the supernova ejecta on the mixing angle $\\sin^{2}2\\theta_{13}$ and mass hierarchy.  We set the luminosity and the energy spectrum of each flavor of neutrinos just emitted from a proto-neutron star in \\S 2. We also set the parameter values of neutrino oscillations from the results of recent neutrino experiments. We explain a supernova explosion model and a nuclear reaction network for light element synthesis. We mention the cross sections of charged-current reactions of the $\\nu$-process to evaluate the reaction rates including neutrino oscillations. In \\S 3, we show the transition probabilities of neutrino flavors with different values of $\\sin^{2}2\\theta_{13}$ and mass hierarchy. We also discuss the effect of the oscillations on the reaction rates of charged-current $\\nu$-process reactions. In \\S 4, we show the calculated mass fraction distribution of $^7$Li and $^{11}$B taking account of neutrino oscillations. Then, we show the dependence of the $^7$Li and $^{11}$B yields on $\\sin^{2}2\\theta_{13}$ and mass hierarchy. We also show the dependence on the temperatures of $\\nu_{\\rm e}$ and $\\bar{\\nu}_{\\rm e}$ just emitted from proto-neutron star. In \\S 5, we discuss the neutrino oscillations with supernova shock propagation and show the change of the $^7$Li and $^{11}$B yields by this effect. We also discuss the $^7$Li and $^{11}$B yields related to observations of stars which have traces of supernova explosions and supernova remnants. Finally, we conclude our study in \\S 6. ",
        "conclusions": "We studied light element nucleosynthesis through the $\\nu$-process in supernovae taking account of neutrino oscillations. The parameters of neutrino oscillations were adopted from the evaluations through several neutrino experiments. We used a supernova model corresponding to SN 1987A and investigated the dependence of the $^7$Li and $^{11}$B yields on the mixing angle $\\theta_{13}$ and mass hierarchy. The obtained results are summarized as follows: \\begin{itemize} \\item Neutrino oscillations affect the yields of $^7$Li and $^{11}$B synthesized in supernova explosions. Especially, the $^7$Li yield increases by a factor of 1.9 in the normal mass hierarchy and adiabatic H-resonance ($\\sin^{2}2\\theta_{13} \\ga 2 \\times 10^{-3}$) compared with that without neutrino oscillations. The $^{11}$B yield increases by a factor of 1.3.  \\item In the inverted mass hierarchy, the increase in the $^7$Li and $^{11}$B yields is smaller: the yields of $^7$Li and $^{11}$B increase by factors of 1.3 and 1.2.  \\item Neutrino oscillations in supernovae make the reaction rates of charged-current $\\nu$-process reactions larger. The reaction rates of neutral-current $\\nu$-process reactions do not change. Thus, the final amounts of the $\\nu$-process products increase by the neutrino oscillations. In our study, main important $\\nu$-process reactions for the $^7$Li and $^{11}$B production are $^4$He($\\nu,\\nu'p)^3$H($\\alpha,\\gamma)^7$Li, $^4$He($\\nu,\\nu'n)^3$He($\\alpha,\\gamma)^7$Be, $^{12}$C($\\nu,\\nu'p)^{11}$B, and $^{12}$C($\\nu,\\nu'n)^{11}$C. When we consider neutrino oscillations, the following charged-current $\\nu$-process reactions become also important: $^4$He($\\nu_e,e^-p)^3$He($\\alpha,\\gamma)^7$Li, $^4$He($\\bar{\\nu}_e,e^+n)^3$H ($\\alpha,\\gamma)^7$Be, $^{12}$C($\\nu_e,e^-p)^{11}$C, $^{12}$C($\\bar{\\nu}_e,e^+n)^{11}$B.   \\item The neutrino temperatures also affect the $^7$Li and $^{11}$B yields due to neutrino oscillations. Large difference of the temperatures of $\\nu_e$ and $\\nu_{\\mu,\\tau}$ brings about larger increase in the yields compared with those without neutrino oscillations.  \\item The shock propagation effect on the neutrino oscillations would slightly reduce the increment of the $^7$Li and $^{11}$B yields. In our model, most of neutrinos pass through the He/C layer before the shock wave arrives at the O/C layer, i.e., the resonance density region. When the shock wave is in the O-rich layers,  the density change by the shock does not influence strongly neutrino oscillations.  \\end{itemize}"
    },
    "0606/astro-ph0606724_arXiv.txt": {
        "abstract": "This paper considers the techniques to distinguish normal star forming (NSF) galaxies and active galactic nuclei (AGN) hosts using optical spectra. The observational data base is a set of 20\\,000 galaxies extracted from the Sloan Digital Sky Survey, for which we have determined the emission line intensities after subtracting the stellar continuum obtained from spectral synthesis. Our analysis is based on photoionization models computed using the stellar ionizing radiation predicted by population synthesis codes (essentially Starburst 99) and, for the AGNs, a broken power-law spectrum. We explain why, among the four classical emission line diagnostic diagrams, (\\oiii/\\Hb\\ vs \\oii/\\Hb, \\oiii/\\Hb\\ vs \\nii/\\Ha\\ (the BPT diagram), \\oiii/\\Hb\\ vs \\sii/\\Ha, and \\oiii/\\Hb\\ vs \\oi/\\Ha), the BPT one works best. We show however, that none of these diagrams is efficient in detecting AGNs in metal poor galaxies, should such cases exist. We propose a new divisory line between ``pure''  NSF galaxies and AGN hosts: $ y = (-30.787+1.1358x+0.27297x^2){ \\rm tanh}(5.7409x) -31.093, $ where $y = $ log (\\oiii/\\Hb), and $x$= log (\\nii/\\Ha).  According to our models, the divisory line drawn empirically by Kauffmann et al. (2003) includes among  NSF galaxies objects that may have an AGN contribution to \\Hb\\ of up to 3\\%.  The Kewley et al. (2001) line allows for an AGN contribution of roughly 20\\%.  About 20\\% of the galaxies in our entire sample that can be represented in the BTP diagram are found between our divisory line and the Kauffmann et al. line, meaning that the local Universe contains a fair proportion of galaxies with very low level nuclear activity, in agreement with the statistics from observations of nuclei of nearby galaxies.  We also show that a classification into NSF and AGN galaxies using only \\nii/\\Ha is feasible and useful.  Finally, we propose a new classification diagram, the $DEW$ diagram, plotting $D_n(4000)$ vs max(EW\\oii,EW\\neiii). This diagram can be used with  optical spectra for galaxies with redshifts up to $z = 1.3$, meaning an important progress over classifications proposed up to now.  Since the $DEW$ diagram requires only a small range in wavelength, it can also be used at even larger redshifts in suitable atmospheric windows. It also has the advantage of not requiring stellar synthesis analysis to subtract the stars and of allowing one to see \\emph{all} the galaxies in the same diagram, including passive galaxies. ",
        "introduction": "Until recently, it was believed that active galactic nuclei were found in only a small fraction of all galaxies (Huchra \\& Burg 1992). However, it was already known that a large fraction of galaxies have nuclei with a very low level activity (called LINERs by Heckman 1980, for Low Ionization Nuclear Emission Regions), and that this activity would not be detectable in distant galaxies.  Generally, normal star forming (NSF) galaxies are distinguished from those containing an active galactic nucleus (AGN) using diagrams where are plotted emission line ratios.  The most common diagnostic diagrams are those of Baldwin, Phillips \\& Terlevich (1981, BPT) and Veilleux and Osterbrock (1987, VO). The lines in  NSF galaxies are emitted by HII regions, which are ionized by massive stars, while AGNs are ionized by a harder radiation field. Therefore, for a given \\oiii/\\Hb\\ or \\oiii/\\oii\\ ratio, AGN galaxies will show higher \\oii/\\Hb, \\nii/\\Ha, \\sii/\\Ha, or \\oi/\\Ha\\ ratios \\footnote {In the entire paper  \\oiii\\ stands for \\Oiii, \\oii\\ for \\Oii, \\nii\\ for \\Nii, \\sii\\ for \\Sii, and \\oi\\ for \\Oi.}, than NSF galaxies. The dividing line between NSF galaxies and AGN hosts has slightly changed over the years. In BPT and VO, it was a compromise between what was suggested by a limited number of data points and some coarse grids of crude photoionization models. More recently, Kewley et al. (2001) proposed a theoretical boundary,  defined by the upper envelope of their grid of photoionization models in which the ionizing source was provided by young stellar clusters.  With the advent of the Sloan Digital Sky Survey (SDSS, York et al. 2000, Abazajian et al. 2004), the number of data points increased by orders of magnitude.  Also, techniques to model the stellar component of the spectra and subtract it from the observed spectrum to obtain the pure nebular spectrum became practicable on a large number of objects. As a result, in the \\oiii/\\Hb vs. \\nii/\\Ha\\ diagram, the $\\sim$ 50000 SDSS galaxies having S/N $>3$ in all the four lines and pertaining to a complete sample of about 120000 galaxies clearly outline two wings (Kauffmann et al. 2003) which look like the wings of a seagull. Kauffmann et al. (2003) have defined a purely empirical dividing line between NSF and AGN galaxies. This dividing line is significantly below the line drawn by Kewley et al. (2001).  Interestingly, the SDSS has definitely shown that, in the local Universe, the number of galaxies hosting AGNs is of the same order as that of NSF galaxies (within a factor which depends on selection criteria and definitions). Studies based on other galaxy samples (e.g. Carter et al.  2001) also came to a similar conclusion, but the SDSS results are stronger, being based on a much larger number of objects, a clear selection function, high resolution spectra and elaborate subtraction of stellar features.  There is actually an important difference between the original BPT or VO diagrams and the Kauffmann et al. (2003) diagram. The former were constructed using spectra of known giant HII regions (mainly located in spiral galaxies) and known  nearby active galactic nuclei, while the Kauffmann et al. (2003) plot concerns galaxy spectra obtained  through  3\\arcsec\\ fibres  which, at their $z \\sim 0.1$ typical redshift, corresponds to 6 kpc (for $H_0=70$ km s$^{-1}$ Mpc$^{-1}$) . Hence, in many galaxies, the region covered by the fibre encompasses a significant fraction of volume and light of the entire galaxy.  Thus, galaxies that occupy the same position as LINERs in these diagrams a priori have no reason to be galaxies \\emph{hosting} a LINER, since the emission line flux from the low ionization nuclear emission region is small with respect to the emission line flux from a  region of several kiloparsecs in diameter, at least in galaxies which still form stars.   One of the persistent questions in  astronomy  is what causes or favors non-stellar activity in galaxies (see e.g. the proceedings of  the IAU symposium ``The Interplay among Black Holes, Stars and ISM in Galactic Nuclei'', Storchi-Bergmann et al. 2004). The SDSS is revolutionizing our ways to attack this problem (e.g. Heckman et al. 2004, Kauffmann et al. 2004, Best et al., 2005, Fukugita et al. 2004, Hao et al. 2005a,b, Pasquali et al. 2005), and deeper surveys will follow. In view of this, it is important to revisit the classification criteria of galaxies in order to lay them on sounder ground. This is the purpose of the present paper.   The paper is organized as follows. In Sect. 2, we present the data sample, and the method to measure emission line intensities. In Sect. 3 we show and discuss some classical emission line diagrams. In Sect. 4 we compare the distribution of observational points with the location of photoionization models for giant HII regions.  In  Sect. 5, we propose a simple model to  account for the emission line properties of AGN host galaxies. In Section 6, we present our boundaries to distinguish NSF galaxies and AGN hosts in classical emission-line diagrams, and we propose alternative classifications, including one that can be easily used for high redshift objects. The last section summarizes our results. ",
        "conclusions": "We have considered a sample of 20\\,000 galaxies extracted from the Sloan Digital Sky Survey and constituting a magnitude-limited sample.  We have applied the spectral synthesis technique described in previous papers in this series to the spectra of these galaxies in order to properly subtract the starlight and obtain a pure nebular spectrum. The emission line intensities have been measured with our automated procedure.  These data have been used to revisit the classical diagrams that are used to distinguish normal star forming galaxies from galaxies hosting an AGN, and to propose new diagrams.  We first analyzed the four classical emission line ratio diagrams: \\oiii/\\Hb\\ vs \\oii/\\Hb, \\oiii/\\Hb\\ vs \\nii/\\Ha\\ (the BPT diagram), \\oiii/\\Hb\\ vs \\sii/\\Ha, and \\oiii/\\Hb\\ vs \\oi/\\Ha. From a purely observational point of view, the BPT diagram is the one which best distinguishes two categories of galaxies, as it distributes the galaxies in two wings which look like the wings of a seagull. The left wing, identified with the sequence of normal star forming galaxies, is very narrow. The right wing, which appeared clearly for the first time in the paper by Kauffmann et al. (2003) also based on SDSS galaxies, is constituted of galaxies hosting an AGN. We have computed a series of photoionization models, using as an input the spectral energy distributions from evolutionary stellar population synthesis. We used the population synthesis code Starburst 99 (Leitherer et al. 1999) in the version which incorporates  the most elaborated stellar atmospheres for the massive stars (Smith et al. 2002). Our photoionization models confirm this interpretation and allow us to draw physically based divisory lines in all the four classical diagrams. However, the models are too schematic to reproduce the observed \\sii/\\Ha\\ and \\oi/\\Ha\\ line ratios correctly. Therefore, the model sequence that best divides NSF and AGN galaxies in the \\oiii/\\Hb\\ vs \\oii/\\Hb\\ or \\oiii/\\Hb\\ vs \\nii/\\Ha\\ diagrams cannot be safely used to distinguish NSF and AGN galaxies in the \\oiii/\\Hb\\ vs \\sii/\\Ha, and \\oiii/\\Hb\\ vs \\oi/\\Ha diagrams.  We propose the following divisory line between NSF galaxies and AGN hosts in the BPT diagram: \\begin{eqnarray} y = (-30.787+1.1358x+0.27297x^2){ \\rm tanh}(5.7409x) \\nonumber\\\\ -31.093, \\end{eqnarray} where $y$= log (\\oiii/\\Hb), and $x$= log (\\nii/\\Ha), replaced by log (\\nii/\\Ha)=-0.4 if \\oiii/\\Hb\\ is not available. This line is actually  close to the line drawn empirically by Kauffmann et al. to distinguish NSF galaxies from AGN hosts. We found that the Kauffmann et al. line  includes among NSF galaxies objects that have an AGN contribution to \\Hb\\ of up to 3\\%. Thus, depending on the problem one is interested in, one may want to use either the Kauffmann. et al line, or the line we propose in this paper, in order to segregate NSF galaxies from AGN hosts. The Kewley line is much less restrictive, and  allows for an AGN contribution of roughly 20\\%.  Since the BPT diagram is very populated between the line defined by Eq. (11) and the Kauffmann line (it contains about 11\\% of the galaxies in our sample, including passive galaxies), this means that the local Universe contains a fair proportion of galaxies with very low level nuclear activity, in agreement with the statistics from observations of galactic nuclei eg., Ho, Fillipenko \\& Sargent (1997).   We  point out that emission line ratio diagrams are not efficient in detecting the presence of an AGN in low metallicity galaxies, if such cases exist.   We have shown that  a classification into NSF and AGN galaxies using only \\nii/\\Ha\\ is feasible and useful.  Finally, we propose a new classification diagram (named the $DEW$ diagram), which uses $D_n(4000)$  vs  max(EW\\oii,EW\\neiii). This classification has many advantages: \\begin{itemize} \\item It can be used at much  larger redshifts than the previous emission line classifications. With SDSS spectra, it can be applied to galaxies with redshifts up to $z = 1.3$. \\item It requires only a small range in wavelengths, so it can also be used at even larger redshifts in suitable windows in the near infra-red. \\item It  can be used without a stellar synthesis analysis to subtract the stars. \\item It allows one to see \\emph{all} the galaxies in the same diagram, including passive galaxies (the definition of passive assumes a certain detection limit of emission lines).  Hence, all galaxies can be classified. \\end{itemize}  This method has drawbacks too: \\begin{itemize} \\item  It is not exactly equivalent to the usual BPT classification. But does it matters? \\item  Old galaxies with a recent starburst ($< 10^7$ yr) will be mistaken for AGN hosts. \\item The borderline between NSF and AGN galaxies is somewhat ``porous'' (but this is the case of almost any frontier). Note that in the BPT diagram, the borderline is also not very well defined at the low excitation end. \\end{itemize}  We note that our proposed classification in the $DEW$ diagram is actually more compatible with that based on the  \\oiii/\\Hb\\ vs \\nii/\\Ha\\ diagram (when using our boundary line) than a classification based on the \\oiii/\\Hb\\ vs \\oii/\\Hb\\ diagram which is used in some papers.  With this new classification scheme at hand, it will be possible to investigate the evolution of AGN galaxy populations in a much larger redshift range than has been done so far, and on  firmer grounds."
    },
    "0606/astro-ph0606038_arXiv.txt": {
        "abstract": "We present the results of HST observations of the host star for the first definitive extrasolar planet detected by microlensing. The light curve model for this event predicts that the lens star should be separated from the source star by $\\sim 6\\,$mas at the time of the HST images. If the lens star is a late G, K or early M dwarf, then it will be visible in the HST images as an additional source of light that is blended with the source image. Unless the lens and source have exactly the same colors, its presence will also be revealed by a systematic shift between centroids of the source plus lens in different filter bands. The HST data indicates both of these effects: the HST source that matches the position of the source star is 0.21 magnitudes brighter in the ACS/HRC-F814W filter than the microlensing model predicts, and there is an offset of $\\sim 0.7\\,$mas between the centroid of this source in the F814W and F435W filter bands. We conclude the planetary host star has been detected in these HST images, and this identification of the lens star enables a complete solution of the lens system. The lens parameters are determined with a Bayesian analysis, averaging over uncertainties in the measured parameters, interstellar extinction, and allowing for the possibility of a binary companion to the source star. This yields a stellar mass of $M_\\ast = 0.63{+0.07\\atop -0.09}\\msun$ and a planet mass of $M_p = 2.6 {+0.8\\atop -0.6} M_{\\rm Jup}$ at an orbital separation of $4.3 {+2.5\\atop -0.8}$AU. Thus, the lens system resembles our own Solar System, with a planet of $\\sim 3$ Jupiter-masses in a Jupiter-like orbit around a star of two-thirds of a Solar mass. These conclusions can be tested with future HST images, which should reveal a broadening of the blended source-plus-lens point spread function due to the relative lens-source proper motion. ",
        "introduction": "A new window on the study of extrasolar planets has been opened with the discovery of four planets by the gravitational microlensing method \\citep{bond-moa53,ogle71,ogle390,ogle169}. This method \\citep{mao-pac,gouldloeb} probes some regions of extrasolar planet parameter space that are not accessible with other planet detection methods. For example, microlensing probes a random sample of stars toward the Galactic bulge and has no large selection effects based on stellar type. The sensitivity of microlensing to low-mass planets \\citep{em_planet,wamb} at separations of a few AU is also unique among current methods, as the recent discoveries of OGLE-2005-BLG-390Lb and OGLE-2005-BLG-169Lb have demonstrated \\citep{ogle390,ogle169}. It is also relatively straight-forward to determine the planet detection efficiency as a function of planet mass \\citep{planet-limit,gaudi-planet-lim}, and the detection efficiency for such low-mass planets is more than an order of magnitude lower than for Jupiter-mass planets. The microlensing discovery of two low-mass planets, despite the much lower detection efficiency, indicates that planets of $\\sim 10\\mearth$ are more common than Jupiter-mass planets at separations of a few AU around the most common stars in our Galaxy, and the fraction of stars with planets of $\\sim 10\\mearth$ is estimated to be about 40\\%. This would seem to confirm a key prediction of the core accretion model for planet formation: that Jupiter mass planets are much more likely to form in orbit around G and K-dwarf stars than around M-dwarfs \\citep{laughlin,ida_lin,boss}. In fact, \\citet{laughlin} have argued that the Jupiter-mass planets found in microlensing events \\citep{bond-moa53,ogle71} are more likely to orbit white dwarfs than M-dwarfs. Clearly, the microlensing detections would provide tighter constraints on planet formation theories if the properties of the host stars were known.  In this paper, we present an analysis of images taken with the Hubble Space Telescope's Advanced Camera for Surveys (ACS) of the source and lens stars for the first microlensing event that yielded a definitive planet detection: OGLE-2003-BLG-235/MOA-2005-BLG-53. This event had a planetary lens system with a mass ratio of $q = 0.0039 {+0.0011\\atop -0.0007}$, so the planet is likely to have a mass similar to that of Jupiter or larger. Like the three other definitive microlensing planet detections to date, the light curve of this event has sharp features that resolve the finite angular size of the source star. This implies that the features of the lens system crossed the angular radius of the source star in $t_\\ast = 0.059\\pm 0.007\\,$days. This can be combined with the source star angular radius, $\\theta_\\ast = 0.53\\pm 0.04\\,\\mu$as to yield the angular Einstein radius: $\\theta_E = t_E \\theta_\\ast /t_\\ast = 0.55\\pm 0.07\\,$mas. The angular source star radius has been determined from the source star magnitude and color \\citep{bond-moa53} using the color-color relations of \\citet{bessellbrett} and the empirical relations between angular radius and surface V-K brightness of \\citet{vanbelle} and \\citet{kervella_dwarf}. This corresponds to a linear source radius of $1 \\rsun$, for $D_S = 8.8\\,$kpc, which is consistent with the dereddened color $(V-I)_0 = 0.76$ of a G-dwarf source \\citep{bond-moa53}. Our $\\theta_\\ast$ value differs slightly from the value given in \\citet{bond-moa53} due to the improved zero-point for ground-based photometry given below. $\\theta_E$ is related to the lens system mass by \\begin{equation} M_L = {c^2\\over 4G} \\theta_E^2 {D_S D_L\\over D_S - D_L} \\ , \\label{eq-m_thetaE} \\end{equation} where $D_L$ and $D_S$ are the lens and source distances, and $M_L = M_\\ast + M_p$ is the total lens mass. Since $D_S$ is known (approximately) and $M_\\ast \\gg M_p$, eq.~\\ref{eq-m_thetaE} can be considered to be a mass-distance relation for the lens star.  Eq.~\\ref{eq-m_thetaE} can be combined with a multicolor main sequence star mass-lumonosity relation \\citep{kroupa_tout,gray,allen} to determine the brightness and colors of the lens star as a function of its distance. This allows us to predict the properties of the combined image of the lens and source in the HST frames, as shown in Fig.~\\ref{fig-predict}. The fraction of the combined flux that is due to the lens, $f_{\\rm lens}$, is shown in the top panel. During the microlensing event, the lens and source were separated by $< 0.1\\,$mas, but the light curve indicates a relative proper motion of $\\mu_{\\rm rel} = \\theta_\\ast /t_\\ast = 3.3\\pm 0.4\\,$mas/yr. So, the lens-source separation was $5.9\\pm 0.7\\,$mas when the HST images were taken, 1.78 years after peak magnification, and color difference between the lens and source stars implies that the centroid position for the unresolved images of the lens and source will show a color dependence. This color dependent centroid shift is shown in the bottom panel of Fig.~\\ref{fig-predict}. This figure indicates that the lens star could contribute as much as $\\sim 30$\\% of the total lens+source flux and yield a B-I centroid shift of $\\sim 0.65\\,$mas for a K-dwarf lens star of 0.6-0.8$\\,\\msun$. ",
        "conclusions": "With the detection of the lens star for microlensing event OGLE-2003-BLG-235/MOA-2005-BLG-53 with HST images taken less than 2 years after peak magnification, we have demonstrated that the ambiguities in the interpretation of planetary microlensing light curves can readily be resolved with high angular resolution follow-up observations. The resolution of these ambiguities allows the planetary microlensing events to place more significant constraints on planet formation theories. For example, we can now confirm the core accretion theory prediction of \\citet{laughlin} that the OGLE-2003-BLG-235/MOA-2005-BLG-53 planetary host star is not an M-dwarf. Instead, it is almost certainly a K-dwarf. If other gas giant planets discovered by microlensing, such as OGLE-2005-BLG-71Lb \\citep{ogle71} also orbit stars that are more massive than the typical microlens star, then this would confirm the core accretion theory prediction that gas giant planets form very rarely in orbit about low-mass stars \\citep{ida_lin,laughlin}. However, the suggestion by \\citet{laughlin} that the lens star might be a white dwarf is not supported by the HST data since a white dwarf lens would be much too faint to be detected.  Although there is evidence that the lens star is detected in both the observed flux from the target star and the color dependent centroid shift, both of the signals are relatively weak. However, the centroid shift signal will grow with time, and will be 2.5 times stronger if similar observations are made in late 2007. It should also be possible to detect an additional effect: the broadening of the target star PSF. In late 2007, the lens and source will be separated by $\\sim 14\\,$mas, or about half an HST-ACS/HRC pixel. According to detailed simulation by one of us (JA), an exposure of  $\\sim 2000\\,$sec (about one orbit) should suffice to measure the PSF broadening in the F814W passband.  It should be possible to identify the lens star for most other planetary microlensing events. Because the relative proper motion, $\\mu_{\\rm rel}$, is usually determined from light curve features for planetary microlensing events, we can routinely construct plots similar to Fig.~\\ref{fig-predict} for other planetary events. The $\\mu_{\\rm rel}$ value for this event is smaller than average, so this event is not particularly favorable for such measurements. If the lens star had a lower mass, we might have to wait $\\sim 5$ years after the event or obtain more exposure time to detect the centroid shift. The situation for the host star of the $\\sim 13\\mearth$ planet in event OGLE-2005-BLG-169, is more favorable because the relative proper motion is large, $\\mu_{\\rm rel} = 8.4\\pm 0.6\\,$mas/yr, and the source star is about a magnitude fainter. As a result, the lens star can be detected in HST images 2 years after the event, even if it is at the bottom of the main sequence (Bennett, Anderson \\& Gaudi, in preparation). The only planetary microlensing event to date that may not allow the detection of the lens star is OGLE-2005-BLG-390 \\citep{ogle390}, because the source star is a giant. However, the lens star for this event may be detectable in the future with the Very Large Telescope Interferometer (Beaulieu et al, in preparation). A space-based microlensing survey \\citep{gest-sim}, like the proposed Microlensing Planet Finder \\citep{mpf-spie} would routinely detect the lens stars by the methods discussed in this paper without the need for follow-up observations with a different telescope."
    },
    "0606/astro-ph0606662_arXiv.txt": {
        "abstract": "{ In this paper we calculate the luminosity distance - redshift relation for a special type of flat Friedmann brane with cosmological constant. This special case is singled out by its simplicity, the luminosity distance being given in terms of elementary functions. We compare our analytical result with the expresssion of the luminosity distance for the flat Friedmann-Lemaitre-Robertson-Walker (FLRW) universe and discuss the differences. ",
        "introduction": "[\\citet{RS}] suggested a new model for the gravitational interaction acting in five non-compact dimensions, the fifth dimension being warped. In brane cosmological models, emerging as generalizations of [\\citet{RS}], our observable universe is a four-dimensional space-time hypersurface (the brane), which has cosmological symmetries and is embedded in the warped five-dimensional bulk. The standard model interactions are confined to the brane, but gravitational dynamics is modified as compared with general relativity, at least at high energies (also at late times in the so-called induced gravity models). Consequently, the luminosity - redshift relation is also changed.  The relation between the luminosity distance and redshift is a powerful tool of the cosmology and has a long history of its own [\\citet{perl,pad}]. In general relativity, a milestone was the work of Mattig [\\citet{matt}], in which this relation has been derived for the FLRW universe with vanishing cosmological constant.  In Section 2 we discuss various luminosity distance - redshift relations. Subsection 2.1 contains the definition of the radial coordinate distance. In Subsection 2.2 we use a standard method for the calculation of the luminosity distance [\\citet{star}] in a FLRW universe with cosmological constant. The result cannot be represented by elementary functions as it contains elliptic integrals of the first kind.  In Subsection 2.3 we calculate the luminosity distance - redshift relation for the flat Friedmann brane embeded in $Z_{2}$ symmetrically into the five-dimensional Schwarzschild-anti de Sitter space-time (5D SADS). For a special value of the brane tension, this relation becomes even simpler than in general relativity, containing only elementary functions. We briefly discuss the assumptions which led to this special case. We compare the luminosity distance - redshift relations for flat Friedmann brane and for FRLW universe with cosmological constant in the Concluding Remarks. ",
        "conclusions": "We have derived the analytical expressions of the luminosity distances for both a flat FLRW universe with cosmological constant and a Friedmann brane embedded into 5D ADS bulk. These expressions are substantially different, as they depend on the Friedmann equation. In the case of the Friedmann brane we have imposed two simplifying assumptions yielding the luminosity distance in terms of elementary functions. There are two values of the cosmological constants and of the brane tension, which are in accordance with these assumptions. The higher value of the $\\Omega_{\\Lambda}$ (see: \\textit{Table 1}) is very close to today's preferred value [\\citet{lid}].  The luminosity distances as function of redshift for all three cases is represented in Fig. 1. On the two plots, $d_{L}$ is represented from $z=0$ to $z=10$ and $z=0$ to $z=2$, respectively. The motivation for the second graph is that supernova observations extend nowadays up to $z=2$. On the plots, we see that all three luminosity distances grow monotonically with increasing redshift. The steepest curve belongs to the FLRW universe. The middle curve is for the case of the brane with the higher value of the cosmological constant. This curve, in the range $z=0..2$, is extremely close to one pertinent to a flat FLRW universe.  Since the values of both brane tensions are much below the theoretically predicted limit, our brane model qualifies as a \"toy model\". The constraints on brane tension [see: \\citet{maar}] imply that $\\Omega_{\\lambda}$ should be small. Nowadays, $\\Omega_{d}$ is also small, being proportional to $a_0^4$. Thus a perturbative treatment can give rise to a more realistic solution for the luminosity distance for Friedmann brane models [see: \\citet{KN}]. However, such realistic solutions for the luminosity distance will be more complicated than the correponding expressions in general relativity.  \\begin{acknowledgement} We thank L\\'{a}szl\\'{o} \\'{A}rp\\'{a}d Gergely for raising this problem and guidance in its elaboration. This work was supported by OTKA grants no. T046939 and TS044665. \\end{acknowledgement}"
    },
    "0606/astro-ph0606381_arXiv.txt": {
        "abstract": "We report the detection of variable linear polarization from Sgr A* at  a wavelength of $3.5\\,$mm, the longest wavelength yet at which a detection has been made.  The mean polarization is $2.1 \\pm 0.1$\\% at a position angle of $16 \\pm 2^\\circ$ with rms scatters of 0.4\\% and 9$^\\circ$ over the five epochs.  We also detect polarization variability on a timescale of days.   Combined with previous detections over the range 150--400\\,GHz (750--2000\\,$\\mu$m), the average polarization position angles are all found to be consistent with a rotation measure of $-4.4 \\pm 0.3 \\times 10^5\\,$rad\\,m$^{-2}$.  This implies that the Faraday rotation occurs external to the polarized source at all wavelengths.  This implies an accretion rate $\\sim 0.2 - 4 \\times 10^{-8}\\, {\\rm M}_\\odot \\,$yr$^{-1}$ for the accretion density profiles expected of ADAF, jet and CDAF models and assuming that the region at which electrons in the accretion flow become relativistic is within $10\\,R_{\\rm S}$.  The inferred accretion rate is inconsistent with ADAF/Bondi accretion. The stability of the mean polarization position angle between disparate polarization observations over the frequency range limits fluctuations in the accretion rate to less than $5$\\%.  The flat frequency dependence of the inter-day polarization position angle variations also makes them difficult to attribute to rotation measure fluctuations, and suggests that both the magnitude and position angle variations are intrinsic to the emission. ",
        "introduction": "\\label{Introduction}  Linear polarization can be an important diagnostic of relativistic jets and accretion flows associated with black hole systems.  In the case of the massive black hole in the Galactic Center, Sgr A*, the properties of its mm-wavelength linear polarization probes the accretion environment on scales inaccessible with other techniques (Bower et al.\\,1999a,b; Aitken et al.\\,2001; Bower et al.\\,2001; Bower et al.\\,2003; Marrone et al.\\,2006).  The apparent absence of linear polarization at wavelengths exceeding 2.7\\,mm and sharp rise in polarization fraction at shorter wavelengths sets an upper limit to the rotation measure (RM).  This limits the mass accretion rate to $\\sim 10^{-7} {\\rm M}_\\sun {\\rm\\ yr^{-1}}$ at distances of $10 - 1000$ Schwarzschild radii from the black hole, which eliminates certain classes of accretion flow (Quataert \\& Gruzinov 2000; Agol 2000), but is consistent with CDAF and jet models (e.g. Falcke, Mannheim and Biermann 1993).  The RM measures the accretion rate by serving as a proxy for the electron column density once coupled with assumptions about the magnetic field.  Equipartition between kinetic, magnetic and gravitational energy is often assumed to relate the electron density and magnetic field (e.g. Bower et al. 1999a; Melia \\& Falcke 2001; Marrone et al. 2006).  The discovery of variations in both the polarization angle (Bower et al.\\,2005) and fraction (Marrone et al.\\,2006) suggests two potential sources of variability.  The variations may be intrinsic and, therefore, offer evidence on the nature and structure of the emission region on a scale of $\\sim 10 R_s$.  Changes in the RM along the line of sight may also induce external polarization variability. Structure in the accretion region on all scales within the Bondi radius can contribute to RM variations.  There is thus strong motivation to accurately characterize the RM, its fluctuations, and the intrinsic variability of the polarized source.  In \\S2 we present the detection of linear polarization of Sgr A* at 3.5\\,mm, showing that it varies on short timescales as well as relative to historical non-detections.  Combining this in \\S3 with other data, our detection yields the best constraints so far on the RM.   We discuss the nature of the accretion flow and on the intrinsic source properties in \\S4. ",
        "conclusions": "\\label{Discussion}    The accretion rate implied by this RM depends on density and magnetic profiles assumed.  Following the prescription outlined by Marrone et al.\\,(2006),\\footnote{Note the typographical error in eq.\\,(9) of Marrone et al.\\,(2006) in which one should have RM$ \\propto r_{\\rm in}^{-7/4}$.} we write $n \\propto r^{-\\beta}$, for  $r_{\\rm in} < r <r_{\\rm out}$ and assume equipartition between magnetic, kinetic and gravitational energy.  The radius $r_{\\rm in}$ designates the point at which the flow becomes relativistic, and so ceases to contribute to the RM because its contribution to Faraday rotation is decreased by a factor $\\gamma^{2}/\\log \\gamma$.  The outer scale is a proxy for the coherence length of the magnetic field fluctuations.  The estimated accretion rate for various ADAF ($\\beta=3/2$) and CDAF ($\\beta=1/2$) models is shown in Figure \\ref{FigMdotlabel} as a function of $r_{\\rm in}$. The accretion rate is in the range $0.2 - 4 \\times 10^{-8}\\,$M$_\\odot$\\,yr$^{-1}$ for $2 \\, R_{\\rm S} < r_{\\rm in} < 10 \\,R_{\\rm S}$, $r_{\\rm out} > 3 \\, r_{\\rm in}$ and $1/2<\\beta < 3/2$.   The fact that a single RM accounts for the frequency dependence of all mean polarization p.a.s implies that the Faraday rotation occurs external to the polarized source.  Internal rotation would cause the RM to vary as a function of frequency.  The emission is optically thick at all frequencies at which polarization is detected (Falcke et al.\\,1998; Zhao et al.\\,2003, Yuan et al.\\,2003), so if any internal Faraday rotation did occur, the diminuation of opacity effects with frequency, which increases the depth down to which one observes emission, would cause a corresponding increase in the Faraday rotation path length.    An interpretation of the p.a.~jitter observed at 85\\,GHz in terms of RM variability would imply rms deviations of $1.2 \\times 10^4\\,$rad\\,m$^{-2}$, far smaller than the $\\sim 2 \\times 10^5\\,$rad\\,m$^{-2}$ rms deviations implied by the 340\\,GHz Marrone et al.\\,(2006) fluctuations interpreted similarly. However, the absence of a clear frequency dependence in the inter-day jitter makes it difficult to ascribe to RM fluctuations, and hence variations in the accretion rate.  Table 2 shows the rms p.a. deviations for the multi-epoch observations at 85, 216, 230 and 340\\,GHz.  These are inconsistent with the $\\nu^{-2}$ dependence expected of RM fluctuations from a magnetoionic medium external to the source, suggesting instead that the jitter reflects changes in the intrinsic source polarization p.a..  Nonetheless, the fact that the fluctuations at these frequencies were not observed simultaneously, coupled with the short time span of the 85\\,GHz observations compared to the $>2\\,$month -- albeit sporadic -- sampling of the 230 and 340\\,GHz measurements, still admits the possibility that the p.a.~dispersion observed at 85\\,GHz is unrepresentative of its long-term average.  This appears unlikely. If the $\\approx 20^\\circ$ rms p.a. variations observed at 216--230\\,GHz were associated with RM fluctuations we would expect $\\approx 150^\\circ$ fluctuations at 85\\,GHz and would not expect to find $\\chi \\propto \\lambda^2$ over a set of disparate measurements. Moreover, the presence of intrinsic p.a. changes is unsurprising given that the polarization fraction is also intrinsically variable, as discussed below.  The lack of p.a.~jitter attributable to RM fluctuations can be interpreted in terms of an upper limit in accretion rate fluctuations.  The absence of clearly identifiable inter-day RM fluctuations suggests it is uniform on inter-day timescales. The consistency of the observations over a large number of disjoint epochs and frequencies with a single RM further suggests that the underlying accretion rate is constant on the timescale over which the observed polarization pas were averaged. The uncertainty in our RM fit places an upper bound on the accretion rate variations using ${\\rm RM} \\propto \\dot{M}^{3/2}$, valid under the assumption of equipartition between magnetic, kinetic and gravitational energy (e.g. Marrone et al.\\,2006).  The 7\\% uncertainty in the RM implies $\\dot{M}$ fluctuations less than $5$\\%.  The limit on $\\dot{M}$ variations is largely consistent with the limits imposed by source flux density variations.  In the jet model the flux density scales as $\\dot{M}^{17/12}$ (Falcke et al.\\,1993).  The standard deviation of the 3.5\\,mm fluxes is 10\\%, comparable to the errors on the individual measurements, which imposes an upper limit of 14\\% on $\\dot{M}$ fluctuations. Marrone et al.\\,(2006) detect 10\\% rms intensity variations at 340\\,GHz.  The five intensity measurements from Bower et al.\\,(2005) at 216--230\\,GHz exhibit 29\\% modulations, implying $\\dot{M}$ fluctuations of 40\\%.  Variability in polarization fraction is detected by all multi-epoch observations (Fig.\\,\\ref{FigPolnlabel}), at 85, 230 and 340\\,GHz, and presumably occurs at 112\\,GHz, where it was not detected at the 1.8\\% (1-$\\sigma$) level by a previous search (Bower et al. 2001).  A previous limit of 1\\% linear polarization at 86\\,GHz (Bower et al.\\,1999b) demonstrates that it also varies on long timescales.   Since the Faraday screen is external to the source the polarization amplitude fluctuations must be intrinsic.  Instrumental depolarization effects are too low to explain the variability: bandwidth depolarization is only important for RMs greater than $2.7 \\times 10^7\\,$rad\\,m$^{-2}$ at 85\\,GHz, while beam depolarization is similarly improbable given the sub-mas size of  Sgr A* at mm wavelengths (Bower et al. 2004; Shen et al. 2005).  Spatial variation in the RM across the transverse extent of the source could depolarize the emission, but at even 85\\,GHz this requires RM fluctuations $\\delta {\\rm RM} \\ga 7 \\times 10^5\\,$rad\\,m$^{-2}$ (Quataert \\& Gruzinov 2000).  Both the high variability of the mm and sub-mm emission  (Zhao et al.\\,2004; Wright \\& Backer 1993; Tsuboi et al.\\,1999) and the linearly polarized emission of Sgr A* are possibly associated with its excess sub-mm emission (Serabyn et al.\\,1997; Falcke et al.\\,1998; Melia et al.\\,2001).  ADAF models that fit the cm to-sub-mm spectrum include at least two distinct populations of radiating particles (Yuan et al.\\,2003), with the second important at $\\nu \\ga 100\\,$GHz in order to explain the sub-mm bump.    The model of Yuan et al.\\,(2003), in which the sub-mm emission is dominated by thermal electrons, overpredicts the polarization fraction at 85\\,GHz (cf. Fig.\\,\\ref{FigPolnlabel}).  In this model the degree of linear polarization ranges from 32\\% at 85\\,GHz to 70\\% at 400\\,GHz assuming a uniform magnetic field and that Faraday depolarization intrinsic to the source is unimportant (see their Fig.\\,3b). In the context of this model the ratio of the predicted to observed polarization levels can only be attributed to magnetic field inhomogeneity intrinsic to the source.  This ratio is a factor of 3 higher at 85\\,GHz relative to the fraction in the range $150-400\\,$GHz, over which it is constant within the errors.  It is hard to account for such an increase in magnetic field inhomogeneity at 85\\,GHz, particularly if the source size only scales $\\propto \\nu^{-1}$, as expected if its emission is self-absorbed.  On the other hand, including synchrotron self-absorption effects, Goldston, Quataert \\& Igumenshchev (2005) show that the polarization fraction is expected to increase by a factor of three over the range $85-200\\,$GHz. The quasi-spherical accretion polarization model of Melia, Liu \\& Coker (2000) and the two-component model of Agol (2000) predict a 90$^\\circ$ p.a.~flip at $\\sim 280\\,$GHz which is at variance with the measured p.a.s at 150, 230 and 340\\,GHz.   \\begin{figure}[h] \\vskip 5mm \\includegraphics[angle=0,scale=0.9]{f3.eps} \\caption{The mean polarization fraction of Sgr A*.  The error bars plotted are those of Table \\ref{PolnTable}.  Previous 86 and 112\\,GHz non-detections are marked with arrows.  The error bars in the Aitken et al.\\,(2001) measurements, marked with triangles, reflect uncertainty in the contribution from dust emission rather than variability associated with the source.} \\label{FigPolnlabel} \\end{figure}    We have reported here the detection of linear polarization in Sgr A* at 3.5\\,mm.  This enables us to calculate the rotation measure and set a limit on the accretion rate.  The lack of frequency dependence for position angle fluctuations indicates that they are intrinsic to the source.  Our result favors RIAF/CDAF accretion models, with a shallow density distribution, over ADAF and Bondi-Hoyle accretion flows, which have a steep profile and are more likely to produce rapid RM variations. Future wide bandwidth simultaneous observations with CARMA and the SMA will fully characterize intrinsic and extrinsic changes in the polarization properties of Sgr A* and allow us to investigate the accretion environments of other nearby low luminosity AGN, such as M81* (Brunthaler, Bower \\& Falcke 2006)."
    },
    "0606/astro-ph0606348_arXiv.txt": {
        "abstract": "{} {We compare the performance of multiple codes written by different groups for making polarized maps from Planck-sized, all-sky cosmic microwave background (CMB) data.  Three of the codes are based on a destriping algorithm; the other three are implementations of an optimal maximum-likelihood algorithm.} {Time-ordered data (TOD) were simulated using the Planck Level-S simulation pipeline.  Several cases of temperature-only data were run to test that the codes could handle large datasets, and to explore effects such as the precision of the pointing data.  Based on these preliminary results, TOD were generated for a set of four 217~GHz detectors (the minimum number required to produce I, Q, and U maps) under two different scanning strategies, with and without noise.} {Following correction of various problems revealed by the early simulation, all codes were able to handle the large data volume that Planck will produce.  Differences in maps produced are small but noticeable; differences in computing resources are large.} {} ",
        "introduction": "Cosmic microwave background (CMB) observations have driven a remarkable advance in cosmology over the past decade (Smoot et al.~\\cite{smoot02}; de~Bernardis et al.~\\cite{debernardis00}; Hanany et al.~\\cite{hanany00}; Beno\\^{\\i}t et al.~\\cite{benoit03}; Bennett et al.~\\cite{bennett03} and references therein), and will continue to furnish invaluable data in the years to come.  As the data volume and precision demanded of these observations increases (for example, Bond et al.~\\cite{bond99}), the complexity of the analysis methods required to deal with the data increases also.  Map-making -- the process of turning time-ordered scan data into an image of the sky -- is an example of a crucial step whose technical complexity has grown significantly.  This is particularly so in the case of total power measurements, where one removes signal drifts due to $1/f$-spectrum noise in the map-making step.  If left unchecked, these drifts leave stripes in the final map with amplitudes greater than the cosmic signal, potentially compromising the scientific goals of a precision instrument such as Planck.  For example, in the simulations presented below, the magnitude of the striping signal (estimated from the rms difference between a simple coadded map and the output map from one of our codes) was 336~\\muK, more than three times the output map's residual rms error of $\\sim$100~\\muK\\ due to white detector noise (see Table~5).  (These numbers are for 1\\farcm7 pixels and four polarized detectors.  For 5\\arcmin\\ pixels, corresponding to the resolution of the detectors, and for the full set of twelve detectors, the residual error is $\\sim$15~\\muK\\ for the temperature map.)  In other words, map-making effectively removed striping with three times the target sensitivity.  Proper map-making is thus crucial to mission objectives.  Planck\\footnote{http://www.esa.int/science/planck/}, to be launched in 2008, will be the third-generation satellite dedicated to observations of CMB anisotropies.  Its primary objective is to measure the temperature anisotropies to the cosmic variance limit out to multipoles $l>2000$; other scientific goals include detailed measurements of the polarized power spectrum, the extraction of catalogs of galaxy clusters and extragalactic sources, searches for non-Gaussianity, and in-depth studies of the Galaxy.  To achieve these goals, Planck will image the sky in nine frequency bands, with resolution and sensitivity in the CMB-dominated bands of 5--15\\arcmin\\ and 5--10~\\muK, respectively.  Crucial to the success of the mission is the production of sky maps approaching the instrumental white-noise limit; drifts and artifacts must be removed and the noise properties well-understood. The formidable challenge of doing this for maps containing millions of pixels lies at the heart of the effort of one of the Planck working groups, the CTP Working Group.  It is important to develop and test well before launch efficient algorithms for inclusion in the data analysis pipeline.  Besides preparing the pipeline, this helps to identify potential sources of systematic error and inform mission operations (e.g., scanning strategy).  It also allows us to better quantify the mission's expected scientific output.  In this paper, we evaluate a suite of map-making techniques using simulations of several channels of the Planck High Frequency Instrument (HFI) 217\\,GHz time-ordered data (TOD).  The simulations model non-white noise and primary CMB temperature and polarization anisotropies.  The suite of methods includes both destriping and optimal map-making algorithms.  We gauge the quality of our recovered temperature and polarization maps by looking at rms pixel residuals, and the residual power spectrum.  This gives us an evaluation through second order statistical measures of noise and artifact residuals. The complexity of the problem requires (at this stage) that we impose a number of simplifying assumptions; these are clearly spelled out in the text and are the focus of on-going work.  We emphasize that the ability to produce the maps shown here is a notable achievement requiring intensive computation.  \\subsection{Planck Scanning Strategy}  Planck will make its observations from the 2nd Earth-Sun Lagrange point, approximately $1.5 \\times 10^6~\\mathrm{km}$ from the Earth. The satellite is spin-stabilized, and during science observations it will rotate on its axis once per minute.  The telescope points at angle of 85 degrees to the spin axis, so the detectors follow small circles on the sky.  The satellite will perform a repointing manouvre once per hour to keep the spin axis close to the anti-solar direction. Thus during the one hour periods between manouvres the detectors will make repeated observations of the same ``rings'' on the sky.  Some of the algorithms presented in this paper can take advantage of these repeated observations to reduce the computational burden.  Within the constraints imposed by Planck's design, there is freedom to choose the precise pointings of the spin axis to optimize the scientific returns of the mission.  The choice of spin axis pointings -- the scanning strategy -- is one of the factors we examine in this paper.  \\subsection{Planck Science Goals for Sky Maps}  Planck is designed to image the sky at nine frequencies from 30 to 857\\,GHz, with angular resolution from 33 to 5\\,arcmin.  The raw sensitivity is sufficient that inferences about the underlying distribution of fluctuations on the sky should be limited not by noise, but rather by cosmic variance.  To achieve this state, systematic errors and processing artifacts must be controlled to microkelvin levels.  It is the latter challenge that we are addressing in this paper and its predecessor (Poutanen et al.~\\cite{poutanen06}). ",
        "conclusions": "The main goal of the simulations reported in this paper was to compare various map-making codes, and to demonstrate that they can deal with Planck-size data sets.  The complexity of the simulated data was kept to a minimum in order to isolate software-induced systematic effects. No instrumental systematic effects other than noise correlations were included in the simulated data.  Comparisons based on more realistic simulations will be made in future papers, which will assess the impact of strong gradients in the signal due to foregrounds, and include a more realistic treatment of the instrumental transfer function, e.g. through the inclusion of beam asymmetries. Nevertheless, some useful results can be identified.  \\subsection{Scanning Strategy}  The Planck design allows considerable flexibility in the choice of the scanning strategy, even in orbit.  The refinement of the scanning strategy will therefore be an important aspect of the pre-launch simulation and analysis work.  Our ability to assess the pros and cons of various candidate strategies will improve as our map-making algorithms include the functionality to deal with increasing levels of realism in the simulated data.  Keeping these qualifications in mind, we compare the map-making residuals for the nominal and cycloidal scanning strategies. The results are summarized in Table~\\ref{tab:Pnoise}. There is a slight preference for the nominal strategy. This preference is a result of the smoother distribution of integration time on the sky for the case of the nominal scanning.  Comparison of Figures 2 and 3, on the other hand, shows that the power in the residuals is higher at the lowest multipoles for the nominal scanning, at least for this particular realization of the noise TOD.  An accurate quantification of the effects of scanning strategy requires, however, ensemble averaging, e.g., using Monte Carlo simulations.  This will be the subject of future work.  The issue is important for proper understanding of the mission's ability to measure, for instance, the reionization bump at low multipole.  Furthermore, we anticipate that the nominal scanning will perform more poorly than the cycloidal scanning when our simulations contain a more realistic set of instrumental systematics.  Once these are included, other features of the scanning strategy will become important, such as the ability to revisit pixels on a range of time scales in order to be able to reject systematic effects (such as residuals from quasi-periodic signals such as cooler noise) and the ability to cross through pixels in several different directions in order to allow the reconstruction of the beam transfer functions.  We will return to the assessment of the relative benefits of scanning strategies in future publications.  \\subsection{Resource requirements}  The map-making codes described in this paper require considerable resources for Planck-sized data.  All codes except Springtide keep the entire data stream in memory.  Memory requirements are therefore dominated by the size of the TOD.  Springtide requires significantly less memory, since it first calculates ring-maps, which compresses the data by a factor of 20--30 (for a Planck-type scanning strategy).  The optimal codes (MADmap, MapCUMBA, ROMA) achieve slightly lower noise in the final maps than Polar and Springtide, but require an order-of-magnitude more CPU time.  Polar achieved lower noise than Springtide, because it worked with shorter baselines (1~minute instead of 1~hour).  Using shorter baselines would increase the memory requirement of Springtide.  MADAM, which combines optimal map-making ideas with destriping, achieves practically the same noise levels as the optimal codes, with similar memory requirements, but in a time comparable to the destriping codes.  \\subsection{Future improvements}  The simulated data prepared for this work are the most advanced so far, but are still far from the reality of the Planck experiment. This holds for instrumental systematics as well as for accurate modeling of the sky signal.  No foregrounds have been included, and several aspects of the CMB emission have been simplified.  In particular, no gravitational lensing or tensor signals have been included.  Both processes have their main impact on the B modes of polarization anisotropy.  Lensing is a well-understood and inevitable effect in cosmology, distorting the CMB and producing B modes in a broad peak in the power spectrum centered at a $\\ell\\approx1000$, with an amplitude much smaller than that of E modes.  Tensor signals show up primarily as degree-scale B modes.  Early reionization could make a tail of that component appear on very large angular scales, with an amplitude which might be detectable by Planck.  The pattern of the total intensity anisotropies on scales of about three degrees or more has been taken directly by the WMAP data.  This affects also a component of the E polarization mode pattern.  The angular extension and reliability of the WMAP pattern in our simulation may certainly benefit from the future releases of WMAP data."
    },
    "0606/astro-ph0606454_arXiv.txt": {
        "abstract": "Measurements of cosmic microwave background (CMB) anisotropies by interferometers offer several advantages over single-dish observations. The formalism for analyzing interferometer CMB data is well developed in the flat-sky approximation, valid for small fields of view.  As the area of sky is increased to obtain finer spectral resolution, this approximation needs to be relaxed.  We extend the formalism for CMB interferometry, including both temperature and polarization, to mosaics of observations covering arbitrarily large areas of the sky, with each individual pointing lying within the flat-sky approximation. We present a method for computing the correlation between visibilities with arbitrary pointing centers and baselines and illustrate the effects of sky curvature on the $\\ell$-space resolution that can be obtained from a mosaic. ",
        "introduction": "The study of anisotropies in the cosmic microwave background (CMB) radiation has revolutionized cosmology.  Key to this revolution have been coupled advances in theory, data analysis, and instrumentation.  In particular, the design of experiments with exquisite systematic error control has been crucial for progress in the field. Interferometers offer several advantages in this respect, with simple optics, instantaneous differencing of sky signals without scanning and no differencing of detectors.  The shape of the beam can be well understood and the measurement is done directly in Fourier space where the theory most naturally lives.  Pioneering attempts to detect CMB anisotropy with interferometers were made by \\cite{MarPar} and \\cite{Sub}. Several groups have successfully detected primary CMB anisotropies \\citep{CAT1,CAT2,DASIT,CBIT,VSA} and polarization \\citep{CBIP,DASIP}, using interferometers.  The formalism for analyzing CMB data from interferometers has been developed by \\cite{HobLasJon,HobMag,WCDH,HobMas}; and \\cite{Mye}; as well as in the experimental papers cited above. \\cite{Parketal} and \\cite{ParkNg} examined interferometric polarimetry.  In the Fraunhofer limit an interferometer measures the Fourier transform of the sky, multiplied by the primary beam.  The primary beam determines the instantaneous field of view of the instrument and its Fourier transform is simply the autocorrelation of the Fourier transform of the point response of the receiver to an electric field.  The angular scale probed by any pair of telescopes being correlated is determined by their spacing in units of the observational wavelength\\footnote{We assume throughout monochromatic radiation; the generalization to a specified frequency band is straightforward.}. The range of scales probed by the interferometer is then determined by the spacing of the elements, while the resolution in spatial wavenumber is determined by the area of sky surveyed.  By ``mosaicking'' several smaller patches together, the resolution in spatial wavenumber can be increased, although the range of spatial scales remains fixed by the geometry of the interferometer elements.  In most cases it has been assumed that the field of view is small, so that one can use the ``small-angle'' or ``flat-sky'' approximation.  However, if we want fine resolution in spatial wavenumber -- which future experiments are driving towards -- we need to survey large areas of sky \\citep{HobMag} and thus relax this assumption. The purpose of this paper is to extend the formalism presented in the above papers to the case where each individual pointing of the interferometer is still within the flat-sky approximation but by mosaicking many pointings together a significant area of sky is surveyed.  Our extension allows one to see how large an error is being made in assuming the flat-sky approximation and shows how corrections can be systematically incorporated.  The central idea of this paper is the following.  The key ingredient in analyzing a mosaic of interferometer pointings is the set of two-point visibility correlations.  For each pair of pointings, we can calculate the correlations in a spherical coordinate system that places both pointing centers on the equator.  If each pointing has a small field of view, then we can approximate the sphere by a cylinder in the vicinity of the equator, allowing the use of Fourier analysis rather than a more cumbersome expansion in spherical harmonics.  The outline of this paper is as follows.  We begin in \\S\\ref{sec:flat} by reminding the reader of some basic results in the flat-sky limit.  We then show how this can be extended using a cylindrical projection in \\S\\ref{sec:cylinder} and make contact with the exact spherical harmonic treatment in \\S\\ref{sec:harmonic}.  Section \\ref{sec:poln} extends our results to include polarization, and we conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  \\begin{figure} \\begin{center} \\resizebox{!}{2.5in}{\\includegraphics{f7.eps}} \\end{center} \\caption{Window functions for a single pointing (solid), the sum of all pointings (dashed), and the sum of all pointings neglecting sky curvature (dotted). \\vskip 0.2in} \\label{fig:polmosaic} \\end{figure}   Interferometers have been used to great effect in measuring CMB temperature and polarization anisotropies.  The formalism for analyzing interferometer data, however, has only been fully developed in the small-field-of-view or flat-sky limit.  Future experiments which aim for exquisite $\\ell$-space resolution will need to survey large areas of sky -- outside the realm of validity of the existing formalism.   In this paper we have extended the formalism to the situation where we can approximate the sky as flat for each individual pointing of the instrument, but we relax the assumption that the angle between pointings is also small.  We have connected the full-sky spherical harmonic approach to the flat-sky Fourier approach in two distinct ways and derived approximations for the visibility covariance matrix in each. We find that the cylindrical method of \\S\\ref{sec:cylinder} and \\S\\ref{sec:polcylinder} works in all cases better than the harmonic method of \\S\\ref{sec:harmonic} and provides accurate approximations to the full-sky expressions for individual pointings smaller than $20^\\circ$ FWHM. Mosaicking together many pointings increases the $\\ell$-space resolution, but in the cases considered here the improvement is less than would be predicted from the flat-sky formalism, in large part due to the effects of baseline rotation. If we neglect sky curvature we overestimate the correlation between distant baselines and hence also overestimate the improvement in $\\ell$-space resolution."
    },
    "0606/hep-ph0606205_arXiv.txt": {
        "abstract": "The properties of string networks at scales well below the horizon are poorly understood, but they enter critically into many observables.  We argue that in some regimes, stretching will be the only relevant process governing the evolution. In this case, the string two-point function is determined up to normalization: the fractal dimension approaches one at short distance, but the rate of approach is characterized by an exponent that plays an essential role in network properties. The smoothness at short distance implies, for example, that cosmic string lensing images are little distorted.  We then add in loop production as a perturbation and find that it diverges at small scales.  This need not invalidate the stretching model, since the loop production occurs in localized regions, but it implies a complicated fragmentation process.  Our ability to model this process is limited, but we argue that loop production peaks a few orders of magnitude below the horizon scale, without the inclusion of gravitational radiation.  We find agreement with some features of simulations, and interesting discrepancies that must be resolved by future work. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606440_arXiv.txt": {
        "abstract": "We present images, integrated photometry, surface-brightness and color profiles for a total of 1034 nearby galaxies recently observed by the Galaxy Evolution Explorer (GALEX) satellite in its far-ultraviolet (FUV; $\\lambda_{\\mathrm{eff}}$=1516\\,\\AA) and near-ultraviolet (NUV; $\\lambda_{\\mathrm{eff}}$=2267\\,\\AA) bands. Our catalog of objects is derived primarily from the GALEX Nearby Galaxies Survey (NGS) supplemented by galaxies larger than 1\\,arcmin in diameter serendipitously found in these fields and in other GALEX exposures of similar of greater depth. The sample analyzed here adequately describes the distribution and full range of properties (luminosity, color, Star Formation Rate; SFR) of galaxies in the Local Universe.  From the surface brightness profiles obtained we have computed asymptotic magnitudes, colors, and luminosities, along with the concentration indices C31 and C42. We have also morphologically classified the UV surface brightness profiles according to their shape. This data set has been complemented with archival optical, near-infrared, and far-infrared fluxes and colors.   We find that the integrated (FUV$-K$) color provides robust discrimination between elliptical and spiral/irregular galaxies and also among spiral galaxies of different sub-types. Elliptical galaxies with brighter $K$-band luminosities (i.e$.$ more massive) are redder in (NUV$-K$) color but bluer in (FUV$-$NUV) (a color sensitive to the presence of a strong UV upturn) than less massive ellipticals. In the case of the spiral/irregular galaxies our analysis shows the presence of a relatively tight correlation between the (FUV$-$NUV) color (or, equivalently, the slope of the UV spectrum, $\\beta$) and the total infrared-to-UV ratio. The correlation found between (FUV$-$NUV) color and $K$-band luminosity (with lower luminosity objects being bluer than more luminous ones) can be explained as due to an increase in the dust content with galaxy luminosity.  The images in this Atlas along with the profiles and integrated properties are publicly available through a dedicated web page at {\\tt http://nedwww.ipac.caltech.edu/level5/GALEX\\_Atlas/} ",
        "introduction": "\\label{introduction}  There are several compelling reasons for observing nearby galaxies in the ultraviolet (UV). First of all, massive, young stars emit most of their energy in this part spectrum and at least in star-forming galaxies they outshine the emission from any other stage of the evolution of a composite stellar population (e.g$.$ Bruzual \\& Charlot 2003). Therefore, the flux emitted in the UV in spiral and irregular is an excellent measure of the current Star Formation Rate (SFR; Kennicutt 1998; Donas et al$.$ 1987).  In the case of quiescent elliptical galaxies the analysis of the UV upturn (the rising part of the FUV spectrum of these galaxies) promises to provide fundamental clues in our understanding of the evolution of low-mass stars on the horizontal branch. Due to its remarkable sensitivity to the physical properties of these stars, the UV upturn could be used, once fully understood, as a powerful diagnostic of old stellar populations (Burstein et al$.$ 1988; O'Connell 1999; Yi et al$.$ 1999; Brown 2004; Rich et al$.$ 2005, 2006, in prep$.$; Boselli et al$.$ 2005). The UV has also revealed the presence of residual star formation in a non-negligible fraction of low-redshift elliptical galaxies (Yi et al$.$ 2005).  Second, the light emitted in the UV can be very efficiently absorbed by dust and then re-emitted at far-infrared (FIR) wavelengths. Therefore an analysis of the energy budget using a comparison of the infrared and UV emission is a powerful tool to determine the dust attenuation of light at all wavelengths (see Buat et al$.$ 2005 and references therein). In this sense, it is worth emphasizing that dust attenuation is the most vexing problem that one has to face when analyzing the observational properties of composite stellar populations and galaxies.  Finally, the observation of nearby galaxies in the UV is fundamental if we are to understand the evolution of galaxies from the high-redshift Universe (where their properties are commonly derived from rest-frame UV observations) to the present.  There have been many attempts in the past to address some of these issues. Sullivan et al$.$ (2000, 2001, 2004) studied the star formation histories in a relatively large and complete sample of UV-selected local galaxies, from which Treyer et al$.$ (1998) derived the SFR density of the local Universe. The nature of the UV upturn in elliptical galaxies has been widely studied by several groups, including O'Connell (1999), Brown et al$.$ (2000), Deharveng, Boselli, \\& Donas (2002). Studies on the dust attenuation in galaxies based on either photometric or spectroscopic UV studies are numerous, including Calzetti et al$.$ (1994), Heckmann et al$.$ (1995), Meurer et al$.$ (1995, 1999), Buat \\& Xu (1996), Gordon et al$.$ (2000, 2003), Buat et al$.$ (2002), Roussel et al$.$ (2005). The analysis of the UV morphology of nearby galaxies as a local benchmark for studies in the optical at high redshift have been also carried out by several authors, including Kuchinski et al$.$ (2000, 2001), Marcum et al$.$ (2001), Windhorst et al$.$ (2002), Lauger, Burgarella, \\& Buat (2005).  However, the results of some of these studies were not conclusive mainly due to the small size of the samples used, which were not representative of the overall population of galaxies in the local Universe. This is particularly true for studies on the dust attenuation in star-forming galaxies and on the rest-frame UV morphology in nearby galaxies. In the case of the UV-upturn studies in early-type galaxies this limitation adds to the lack of spatial resolution and depth of previous UV data and, in some cases, to the availability of UV data in only one band, which leads to a loss of sensitivity to the strength of the UV upturn, best traced by the FUV$-$NUV color (see Gil de Paz et al$.$ 2005 and references therein).  The availability of deep UV observations with moderately-good spatial resolution for large numbers of well-known nearby galaxies is now possible thanks to the launch of the Galaxy Evolution Explorer (GALEX) on April 28th 2003. The compilation of GALEX UV data carried out as part of this paper will allow us (and other researchers making future use of this dataset) to provide fundamental clues for solving some of the still many open questions regarding the UV properties of galaxies in the local Universe. In particular, we will show how the strength of the UV upturn is function of the stellar mass of the galaxy, with more massive elliptical galaxies showing stronger UV upturns. We will also demonstrate that in a sample like ours, which adequately represents the bulk of the galaxy population in the local Universe, the slope of the UV continuum is well-correlated (although with a significant dispersion) with the infrared-to-UV ratio and, therefore, with the UV extinction, and that the (FUV$-$$K$) color provides and excellent segregation between early-type (ellipticals and lenticulars) and late-type (spirals and irregulars) galaxies.  In this ``The GALEX Ultraviolet Atlas of Nearby Galaxies'' we present surface photometry in the two GALEX ultraviolet (FUV \\& NUV) bands, providing integrated photometry and structural parameters for a total of 1034 nearby galaxies, including extensively-studied objects like M31, M32, M~33, M~51, M~81, M~82, M~83, M~87, M~101, etc. We compare the UV properties of this sample with corollary data in the optical, NIR, and FIR, available for the majority of the galaxies in the Atlas. This comparison allows us to obtain insight into fundamental correlations such as the `red sequence' found in the color-magnitude diagram of ellipticals and lenticulars, and a better definition of the IRX-$\\beta$ relation in normal star-forming galaxies.  In Section~\\ref{sample} we extensively describe the sample of galaxies. Section~\\ref{observations} provides a summary of the GALEX observations. The analysis and results are given in Sections~\\ref{analysis} \\& \\ref{results}, respectively. The conclusions are summarized in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have presented an imaging Atlas of 1034 galaxies observed in two UV bands by the GALEX satellite. From these we have derived surface brightness and color profiles in the FUV \\& NUV GALEX bands. Asymptotic magnitudes and colors along with concentration indices have also been obtained. A morphological classification of the profiles is also carried out. Despite a small but non-negligible excess of high-luminosity and paucity of low-luminosity spiral galaxies (compared with the luminosity distribution of ellipticals both in our and the NFGS samples) it is shown that this sample adequately matches the distribution and full range of properties of galaxies in the local Universe. We have augmented this data set with corollary data from the optical (RC3), NIR (2MASS), and far-infrared (IRAS). We emphasize here the special caution should be observed when comparing these results with those derived from a volume-limited sample. From a broad-based initial analysis of the UV properties of this sample we conclude:  \\begin{itemize}  \\item The value of the integrated (FUV$-K$) color of galaxies provides an excellent criterion with which to discriminate elliptical/lenticular galaxies from spirals and irregulars. The best discrimination between these two classes of galaxies (quiescent vs$.$ star-forming) is achieved if a cut-off color (FUV$-K$)=8.8\\,mag is adopted. A reasonably good separation is also obtained by using a (FUV$-$NUV) cut-off color at 0.9\\,mag. These colors also allow for a continuous distinction (although with a significant dispersion) of spiral galaxies of different types.  \\item Elliptical/lenticular galaxies with brighter FUV and $K$-band luminosities show bluer (FUV$-$NUV) colors than ellipticals with fainter luminosities but redder (NUV$-K$) colors. This is true for ellipticals galaxies specifically within the range of absolute magnitudes covered by this Atlas (i.e$.$ M$_K$$<$$-$21\\,mag). This behavior is probably a consequence of luminous elliptical galaxies having stronger UV upturns than their intermediate-mass counterparts (see Boselli et al$.$ 2005).  \\item We do not find a large dispersion in the intrinsic (corrected for internal extinction) (FUV$-$NUV) colors of the spiral/irregular galaxies in the Atlas ($\\sigma_{\\mathrm{(FUV-NUV)}_0}$=0.05\\,mag) neither a strong dependence of it with the galaxy luminosity. Consequently, the variations in the observed (FUV$-$NUV) colors with the luminosity or morphological type of the spiral and irregular galaxies in the sample are plausibly due to variations in the dust content (due for example to changes in metallicity) with these magnitudes. In the case of the (FUV$-K$) color the star formation history necessarily contributes to its dependence on luminosity and morphological type.  \\item The change in the observed (FUV$-$NUV) color with the TIR-to-FUV ratio also suggests that the attenuation law in these galaxies differs from a pure Milky-Way extinction law. In particular, attenuation laws with relatively steep FUV rise and no 2175\\,\\AA\\ bump, like those based on a SMC Bar extinction law or the Calzetti law in the case of the most luminous objects, are favored.  \\item A significant fraction (28\\%) of the UV profiles show some degree of flattening in the inner regions. The galaxies showing this kind of profiles belong to a relatively small range of optical morphological types (compared with the pure-exponential profiles), 2$<$T$<$8, i.e$.$ they are all truly spiral galaxies. We interpret this as a consequence of the high past SFR but comparatively low current gas infall rate in the inner disks of spiral galaxies, leading to an efficient consumption of the gas in these regions and, consequently, to a flattening of the UV profiles compared with the outer disks, where the gas supply is still abundant. This is, indeed, expected to be particularly important in intermediate-type spirals.  \\end{itemize}  The GALEX and corollary photometry data along with the profiles and UV images of galaxies in the sample can be accessed through a dedicated web page at\\newline {\\tt http://nedwww.ipac.caltech.edu/level5/GALEX\\_Atlas/}."
    },
    "0606/astro-ph0606395_arXiv.txt": {
        "abstract": "{A high accuracy photometry algorithm is needed to take full advantage of the potential of the transit method for the characterization of exoplanets, especially in deep crowded fields. It has to reduce to the lowest possible level the negative influence of systematic effects on the photometric accuracy.  It should also be able to cope with a high level of crowding and with large scale variations of the spatial resolution from one image to another. A recent deconvolution-based photometry algorithm fulfills all these requirements, and it also increases the resolution of astronomical images, which is an important advantage for the detection of blends and the discrimination of false positives in transit photometry. We made some changes to this algorithm in order to optimize it for transit photometry and used it to reduce NTT/SUSI2 observations of two transits of OGLE-TR-113b.  This reduction has led to two very high precision transit light curves with a low level of systematic residuals, used together with former photometric and spectroscopic measurements to derive new stellar and planetary parameters in excellent agreement with previous ones, but significantly more precise. ",
        "introduction": "Among the $\\sim$200 exoplanets known so far, only the 10 ones transiting their parent star have measured masses and radii, thanks to the complementarity of the radial--velocity and transit methods. Among them, 5 were detected by the OGLE-III planetary transit survey (\\cite{Udalski1}, \\cite{Udalski2}, \\cite{Udalski3},  \\cite{Udalski4}): OGLE-TR-10b (\\cite{Konacki3}), OGLE-TR-56b (\\cite{Konacki1}, \\cite{Bouchy2}), OGLE-TR-111b (\\cite{Pont1}), OGLE-TR-113b (\\cite{Bouchy1}, \\cite{Konacki2}) and OGLE-TR-132b (\\cite{Bouchy1}). Compared to the other transiting exoplanets, they orbit much fainter stars, leading to a lower amount of information available from their observation. Furthermore, obtaining high accuracy photometry for these stars is difficult with a classical reduction method, even with large telescopes, because of the high level of crowding present in most of the deep fields of view in the Galactic plane. Nevertheless, the accurate photometric monitoring of their transits is important to better constrain the mass-radius relationship of close-in giant planets, and thus the processes of planet formation, migration and evaporation. Besides, high accuracy transit observations may allow the detection of other planets, even terrestrial ones in the best cases, by the measurements of the dynamically induced variations of the period of the transit (\\cite{Miralda1}, \\cite{Agol1}, \\cite{Holman1}).  An image deconvolution algorithm (\\cite{Magain1}) has recently been adapted to the photometric analysis of crowded fields (\\cite{Magain2}), even when the level of crowding is so high that no isolated star can be used to obtain the $PSF$ (Point Spread Function). We made some modifications to this algorithm to optimize it for follow-up transit photometry, with a main goal in mind: to obtain the highest possible level of photometric accuracy, even for faint stars located in  deep crowded fields.  This new method was tested on new photometric observations of two OGLE-TR-113b transits obtained with the NTT/SUSI2 instrument.  This planet was the second one confirmed from the list of planetary candidates of the OGLE-III survey. It orbits around a faint K dwarf star ($I$ = 14.42) in the constellation of Carina. Due to the small radius of the parent star ($R \\sim 0.8$ $R_{\\odot}$), the transit dip in the OGLE-III light curves is the largest one among the planets detected by this survey ($\\sim$ 3 \\%). As OGLE-TR-113 lies in a field of view with a high level of crowding, this case is ideal to validate the potential of our new method.  Sect.\\  2 presents the observational data. Sect.\\  3 summarizes the main characteristics of the deconvolution algorithm and describes the improvements we brought to optimize it for follow-up transit photometry. In Sect.\\  4, our results are presented and new parameters are derived for the planet OGLE-TR-113b. Finally, Sect.\\  5 gives our conclusions. ",
        "conclusions": "The results presented here show that our  new photometry algorithm is well suited for follow-up transit photometry, even in very crowded fields. After analysis  of NTT SUSI2 observations of two OGLE-TR-113b transits, we have obtained two very high accuracy transit light curves with a low level of systematic residuals. Combining our new photometric data with OGLE-III ephemeris, spectroscopic data and radial velocity measurements, we have determined planetary and stellar parameters in excellent agreement with the ones presented in Bouchy et al. (2004) and Konacki et al. (2004), but significantly more precise. We notice that the sampling in time, the sub-millimag photometric accuracy and the systematics residuals level of our light curves would be good enough to allow the photometric detection of a transiting Hot Neptune,  in the case of a small star as OGLE-TR-113.  We have obtained a very precise determination of the transit times, and, combining them with OGLE-III ephemeris, we could determine the orbital period with a very high accuracy. The precisions on the epochs and the period would in fact be high enough to allow the detection of a second planet or a satellite, for some ranges of orbital parameters and masses. Even a terrestrial planet could be detected with such a transit timing precision."
    },
    "0606/astro-ph0606676_arXiv.txt": {
        "abstract": "We present ultra-deep {\\it Spitzer} 70$\\mu$m observations of GOODS-North (Great Observatories Origins Deep Survey).  For the first time, the turn-over in the 70$\\mu$m Euclidean-normalized differential source counts is observed.  We derive source counts down to a flux density of 1.2\\,mJy.  From the measured source counts and fluctuation analysis, we estimate a power-law approximation of the faint 70$\\mu$m source counts of $dN/dS \\propto S^{-1.6}$, consistent with that observed for the faint 24$\\mu$m sources.  An extrapolation of the 70$\\mu$m source counts to zero flux density implies a total extragalactic background light (EBL) of $7.4\\pm1.9\\nW$.  The source counts above 1.2\\,mJy account for about 60\\% of the estimated EBL. From fluctuation analysis, we derive a photometric confusion level of $\\sigma_c = 0.30\\pm0.15$\\,mJy ($q=5$) for the {\\it Spitzer} 70$\\mu$m band. ",
        "introduction": "Deep 24$\\mu$m observations (Chary et al. 2004; Papovich et al. 2004; Fadda et al. 2006) have demonstrated the ability of the Multiband Imaging Photometer for {\\it Spitzer} (MIPS, Rieke et al. 2004) to study the mid-infrared (mid-IR) properties of high-redshift galaxies (Yan et al. 2004; Le Floc'h et al. 2004, 2005; P\\'{e}rez-Gonz\\'{a}lez et al. 2005; Daddi et al. 2005; Papovich et al. 2006; Caputi et al. 2006).  The interpretation of the 24$\\mu$m data are complicated by the presence of strong emission and absorption features (e.g., Armus et al. 2004) redshifted into the 24$\\mu$m band.  Observations at longer wavelengths, such as 70$\\mu$m which is closer to the far-infrared peak of the spectral energy distribution (SED) and is away from the strong mid-IR features, are crucial for constraining the infrared luminosities and star-formation rates.  The previous deep Guaranteed Time Observer (GTO) surveys did not achieve sufficient sensitivity at 70$\\mu$m to detect distant luminous infrared galaxies (LIRGs; $10^{11}\\lsun \\la L_{ir} \\la 10^{12}\\lsun$), without stacking 70$\\mu$m data for a large number of 24$\\mu$m-selected sources (Dole et al. 2006).  Much deeper observations are needed at 70$\\mu$m to individually detect the $z\\sim1$ LIRGs that account for the majority of the extragalactic background light (Elbaz et al. 2002; Lagache et al. 2004).  In this letter, we present initial results for the deepest 70$\\mu$m survey taken to date with {\\it Spitzer}. ",
        "conclusions": "Based on ultra-deep 70$\\mu$m observations, we derive source counts down to a flux density of 1.2\\,mJy, directly resolving about 60\\% of the EBL.  The total fraction of the EBL estimated for sources down to the confusion level ($\\sigma_c\\simeq0.3$\\,mJy, $q=5$) is about 75\\%. A power-law extrapolation to zero flux density implies a total EBL of $7.4\\pm1.9\\nW$ at 71.4$\\mu$m.  This is consistent with the value predicted based on EBL measurements at other wavelengths, the value predicted from the Lagache et al. (2004) model, and the value derived from the extrapolation of the 24$\\mu$m counts and 70$\\mu$m stacking analysis (Dole et al. 2006).  However, the uncertainties on the results leave open the possibility of a significant population of sources at low 70$\\mu$m flux densities that are not accounted for in the models, such as highly obscured $z\\sim 1$ AGNs as proposed to account for the hard X-ray background (e.g., Worsley et al. 2005). Studies of the counterparts of the faint 70$\\mu$m population are ongoing and will help to constrain the infrared luminosities and the relative fraction of AGN versus starburst-dominated galaxies in the high-redshift {\\it Spitzer}-selected surveys.  We thank our colleagues associated with the {\\it Spitzer} mission who have made these observations possible.  This work is based on observations made with the {\\it Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407."
    },
    "0606/astro-ph0606506_arXiv.txt": {
        "abstract": "We calculate the strong lensing probability as a function of the image-separation $\\Delta\\theta$ in TeVeS (tensor-vector-scalar) cosmology, which is a relativistic version of MOND (MOdified Newtonian Dynamics). The lens, often an elliptical galaxy, is modeled by the Hernquist profile. We assume a flat cosmology with $\\Omega_b=1-\\Omega_\\Lambda=0.04$ and the simplest interpolating function $\\mu(x)={\\rm min}(1,x)$. For comparison, we recalculated the probabilities for lenses by  Singular Isothermal Sphere (SIS) galaxy halos in LCDM with Schechter-fit velocity function. The amplification bias is calculated based on the magnification of the second bright image rather than the total of the two brighter images.  Our calculations show that the Hernquist model predicts insufficient but acceptable probabilities in flat TeVeS cosmology compared with the results of the well defined combined sample of Cosmic Lens All-Sky Survey (CLASS) and Jodrell Bank/Very Large Array Astrometric Survey (JVAS); at the same time, it predicts higher probabilities than SIS model in LCDM at small image separations. ",
        "introduction": "Since Bekenstein proposed the relativistic, modified Newtonian dynamics (MOND) theory, named  tensor-vector-scalar \\citep[TeVeS;][]{bekenstein04}, it has become possible to investigate the MOND phenomena in the cosmological sense. In particular, after determining the geometry and background evolution of the Universe, and calculating the deflection of light due to a weak gravitational field, one can test TeVeS and thus MOND with gravitational lensing \\citep{chiu,zhao06a,angus}. Before TeVeS, strong gravitational lensing in the MOND regime could only be manipulated by extrapolating non-relativistic dynamics \\citep{qin,mortlock}, in which the deflection angle is only half the value in TeVeS \\citep{zhaoqin06}.  Needless to say, comparing the predicted results of gravitational lensing with observations is of key importance in testing TeVeS. \\citet{zhao06a} first examined the consistency of the strong lensing predictions in the TeVeS regime for galaxy lenses in the CfA-Arizona Space Telescope Lens Survey (CASTLES). In this {\\it Letter}, we investigate the  statistics of strong lensing in the TeVeS regime, and compare the predicted lensing probabilities to the well defined sample of CLASS/JVAS survey. We adopt the mass function of the stellar component of galaxies \\citep{panters}. As a first approximation, we do not consider galaxy cluster lenses; the lenses in the well defined sample in CLASS/JVAS are believed to be produced by galaxies rather than galaxy cllusters, although a cluster lens, SDSSJ1004, was discovered in Sloan Digital Sky Survey (SDSS) \\citep{inada03,oguri04a}. We consider the simplest MOND interpolating function $\\mu(x)$ and use the Hernquist profile \\citep{hernquist} to model the galaxy lenses. It is now established that, in standard cosmology (LCDM), when galaxies are modeled by a Singular Isothermal Sphere (SIS) and galaxy clusters are modeled by a Navarro-Frenk-White (NFW) profile, the predicted strong lensing probabilities can match the results of CLASS/JVAS quite well \\citep[e.g.,][]{chae03,chena,chenb,chenc,chend,li02,mitchell05, oguri02,oguri03b,oguri04,peng06,sarbu01, wj04,zhang2004,zhang05}. For comparison, we recalculate the lensing probabilities predicted by the SIS modeled galaxy lenses in LCDM cosmology with the velocity function. Note that, in LCDM, baryon infall effect \\citep[e.g.,][]{kochanek01,keeton01} has been well described by SIS model for galaxies \\citep{rusin05,koopmans06}, at least statistically; furthermore, the effects of substructures \\citep{oguri06} are also considered since we use the velocity function to account for the number density of lensing galaxies. Throughout this {\\it Letter}, we assume the source QSOs have a redshift of $z_s=1.27$. ",
        "conclusions": "It has been held that, in the MOND regime, the effect of lensing is inefficient, in particular, that strong lensing never occurs \\citep[e.g., ][]{scarpa}. Our calculations shown in Figure \\ref{figprob} indicate, however, that this is not true. Although the Hernquist model predicts insufficient lensing probabilities in a flat TeVeS cosmology compared with the result of CLASS/JVAS, the result is acceptable considering, at least, that the lensing galaxy can be modeled by steeper slopes and more efficient MOND $\\mu$-functions.  Our results argue that TeVeS (and thus MOND) generates lenses with higher efficiency than CDM if the latter is modelled by SIS profile and in both cases the amplification bias  is calculated based on the magnification of the second bright image (for SIS, the fainter image).  Usually, $B$ is calculated based on the total magnification of the two images cosidered. Because we introduced a cutoff $\\beta_{q_r}$ due to the flux ratio $q_r$, the total magnification is $2\\sim q_r+1$ times larger than that of the second bright image (depending on $\\beta$ when $\\beta\\le\\beta_{q_r}$), which results in the corresponding value of $B$ about 4 times larger (both for SIS in LCDM and Hernquist in TeVeS). As shown in Figure \\ref{figprob}, the lensing probabilities for SIS halos in LCDM with $B=3.976$ (total) matches the results of CLASS/JVAS quite well (dashed line). Similarly, if we apply the total magnification to $B$ for Mondian Hernquist model, the final lensing probabilities would be overpredicted compared with the results of CLASS/JVAS. Note that an SIS profile is more concentrated in mass than a Hernquist profile, so if both profiles are applied in the same regime (LCDM or TeVeS), the SIS profile would be more effective at lensing than Hernquist. Therefore, the fact that the probabilities for SIS model in LCDM ($B=1.09$, dash-dotted line) are lower than the Hernquist model in TeVeS shown in Figure \\ref{figprob} implies that MOND demonstrates a higher lensing efficiency than CDM. This phenomena is, in fact, not difficult to understand. It is well known that MOND, as an alternative to dark matter for solving the ``missing mass\" problem, takes effect in the region surrounding the luminous matter with $r>r_0$, where a CDM halo is assumed to have non-zero density and its acceleration dominates over luminous matter in LCDM cosmology \\citep{kaplinghat}. The deflection angle $\\alpha(b)$ with impact parameter $b>r_0$ can be calculated using Newtonian CDM gravitation or Mondian luminous matter gravitation. We know that the acceleration $g(r)$ in the equation for deflection angle for an SIS modeled CDM halo is $g(r)\\propto r^{-1}$, independent of $b$. So, the image separation is independent of the source position angle $\\beta$ (when $\\beta<\\beta_{cr}$), as is well known in the SIS model. However, for a lensing galaxy (with no dark matter) modeled by a Hernquist profile, we have $g(r)\\propto r^{-1}$  only when $r>r_0$ (Mondian regime). So the higher probabilities indicate a higher lensing efficiency between MOND and CDM.  As a first attempt at investigating strong lensing statistics in the TeVeS scenario, we have used the simplest interpolating function $\\mu(x)$. The deflection angle is, of course, sensitive to $\\mu(x)$ \\citep{zhao06b}. The simplest $\\mu(x)$ adopted in this {\\it Letter} corresponds to the lowest physical (or ``true\") acceleration $g(r)$. Any other forms of $\\mu(x)$ all give stronger physical accelerations than the simplest one \\citep{zhaotian06}. Furthermore, strong lensing  is very sensitive to the concentration and the slope near the center of the density profile of lensing galaxies.  The Hernquist model and an NFW model has $\\rho(r)\\sim r^{-1}$ near the center, both are inefficient in lensing. If elliptical galaxies were modeled as pressure- supported Jaffe model, e.g.,  with $\\rho(r)\\sim \\frac{1}{r^{2}(r+a)^2}$ where $a$ is a core scale length, then the lensing probability would be increased.  Another important point is that we have assumed a flat universe with $\\Omega_\\Lambda=0.96$. However, by fitting to high-z SN Ia luminosity modulus, \\citet{zhao06a} showed that an open universe is more likely in TeVeS (with $\\Omega_\\Lambda=0.5$). In summary, it is promising to constrain TeVeS vs CDM through lensing statistics."
    },
    "0606/astro-ph0606730_arXiv.txt": {
        "abstract": "{Detailed studies of the photometric and kinematical properties of compact groups of galaxies are crucial to understand the physics of galaxy interactions and to shed light on some aspects of galaxy formation and evolution.  In this paper we present a kinematical and photometrical study of a member, NGC4778, of the nearest (z=0.0137) compact group: Hickson 62.} {The aim of this work was to investigate the existence of kinematical anomalies in the brightest group member, NGC4778 in order to constrain the dynamical status and the formation history of the group. } {We used long-slit spectra obtained with FORS1 at VLT, to measure line-of-sight velocity distributions by means of the Fourier Correlation Quotient method, and to derive the galaxy rotation curve and velocity dispersion profile.  } {Our analysis reveals that Hickson 62a, also known as NGC4778, is an S0 galaxy with kinematical and morphological peculiarities, both in its central regions (r $< 5''$) and in the outer halo. In the central regions, the rotation curve shows the existence of a kinematically decoupled stellar component, offset with respect to the photometric center. In the outer halo we find an asymmetric rotation curve and a velocity dispersion profile showing a rise on the SW side, in direction of the galaxy NGC4776.} {The nuclear counterrotation, the distorted kinematics in the outer halo and the X-ray properties of the group suggest that NGC4778 may be the product of a recent minor merger, more reliable with a small late-type galaxy.} ",
        "introduction": "Poor groups of galaxies are the most common cosmic structures and contain a large fraction of the galaxies present in the universe (\\cite{Tully87}, \\cite{Eke04a}). At a difference with rich clusters, they span a wide range of densities, from loose groups, having spatial density of baryonic matter slightly above that of the field, to compact ones having densities comparable or higher than those encountered in the cores of the richest clusters.  For this reason, they are the ideal ground where to test all scenarios for galaxy formation and evolution and where to pinpoint the details of the physics controlling galaxy interactions.  Loose groups having masses in the range $10^{13} - 10^{14} \\ M_{\\odot}$, almost certainly are still collapsing and are therefore crucial to uncover the formation processes shaping cosmic structures (\\cite{Zabludoff98}). Many factors converge in identifying compact groups as good candidates to be one of the regions where some of this processes occur. In first place, their high spatial density of luminous matter and small velocity dispersions imply dynamical lifetimes of the order of a fraction of the Hubble time. This leading to the possibility that the groups observed in the present time and in the local universe are second generation objects, just accreting new members from the loose groups of galaxies in which almost always they are embedded (\\cite{Vennik93}).  Second, compact groups are numerous and contain a non negligible fraction of the baryonic matter in the nearby universe (\\cite{Pildis96}). Therefore, whatever is their ultimate fate, they are bound to have an impact on the observable properties of galaxies and cosmic structures.  As it was stressed in \\cite{Mendes03} (hereafter M03), the influence of environmental effects on the internal dynamics and matter distribution of compact group galaxies has not yet been clearly established, mostly due to lack of reliable kinematic data. Extensive kinematical studies of both the stars and gas in galaxies belonging to compact groups (\\cite{Rubin91}; \\cite{Nishiura00}; M03; \\cite{Rampazzo98}; \\cite{Bonfanti99}) all suggest that peculiar kinematical behaviors are much more common ($\\sim 75\\%$) than in the field.  Moreover, M03 showed that velocity fields of the ionized gas component in galaxies belonging to compact groups are often significantly affected by non-circular motions, local asymmetries and misalignments between the kinematic and stellar axes. These peculiarities, however, tend to smooth out if the rotation curve is derived by averaging the velocity fields of the galaxies over large regions.  If these averaged values are used, a large fraction ($\\sim 80\\%$) of the HCG members follow the same Tully-Fisher (TF) relationship of field galaxies (M03). This may indicate that the haloes of compact group galaxies have not been significantly stripped inside their optical size. However, according to M03, the remaining 20\\% of the galaxies, including the lowest-mass systems, present significant anomalies which could be explained by assuming that compact group galaxies have smaller dark halos than their field counterparts, due to tidal truncation. A result which finds support in numerical simulations (cf. \\cite{Governato91}) and has important consequences on the groups dynamical lifetimes.  In spite of the vast literature existing, due both to the limited statistics and to the problems encountered in disentangling true groups from optical ones as well as in deprojecting the measured kinematic and photometric quantities, our understanding of the dynamical and evolutionary status of compact groups still presents quite a few gaps.  Gaps which can be filled only through detailed multitechnique and multiwavelength analysis of individual cases.  In this respect, the dynamical and evolutive status of a group two observables are crucial: the detailed kinematics of the individual galaxies and the structure of the diffuse hot gas halo.  In this paper and in Paper II (\\cite{Sodani06}) we present a study of the compact group of galaxies Hickson 62 based on archival optical, spectroscopic and X-ray data extracted from the ESO and the Chandra archives.  In this first paper we focus mainly on the peculiar kinematics and on the photometry of the dominant galaxy NGC4778 (Hickson 62a), while in Paper II we shall discuss the diffuse X ray halo embedding at least two of the group components. This paper is structured as follows: in Section \\ref{HCG62} we discuss the main characteristics of Hickson 62, in Section \\ref{thedata} we describe the observations and the data reduction procedure.  The photometric properties of NGC4778 are presented in Section \\ref{phot} and the kinematics in \\ref{kin}. Finally, we draw our conclusions in \\ref{discussion}. Throughout this paper we shall adopt a distance of 60.9 Mpc based on $H_{0}= 70$ km $\\mbox{s}^{-1} \\mbox{Mpc}^{-1}$ and an heliocentric radial velocity $V = 4260$ km $\\mbox{s}^{-1}$, this implies 1 arcsec = 0.29 kpc. ",
        "conclusions": "\\label{discussion}  We have analyzed high signal-to-noise spectra along the major axis of the dominant galaxy of the Hickson Compact Group HCG62. On the whole, the observed kinematics and photometry of NGC4778 is consistent with that of an S0 galaxy.  -\\emph{Nuclear regions}- The higher resolution data enabled us to detect the signature of the nucleus of NGC4778: inside $3''$ we observe an inversion in the velocity profile gradient, which also correspond to some anomalous photometric features, such as bluer colors, and twisting in the position angle of the isophotes. These features strongly suggest the existence of a small core ($\\sim 600\\ pc$) kinematically decoupled from the whole galaxy. Such anomalous Kinematically Decoupled Cores (KDC) are common in early-type galaxies and show very similar features to those observed in the nuclear regions of HCG62a: the velocity profile is characterized by a central asymmetry, to which corresponds an unusual central isophotal flattening (see \\cite{Krajnovic04}).  Similar behavior are observed for instance in NGC3623, belonging to the Leo Triplet (\\cite{Afanasiev05}). The central region of this galaxy, in fact, shows the presence of a chemically distinct core, a relic of a star formation burst, due to interactions, that is shaped as a cold stellar disk with a radius of $\\sim 250-350\\ pc$. Like NGC4778, NGC3623 also shows a drop in the stellar velocity dispersion in the nucleus. Numerical simulations (\\cite{Bournaud03}) predict that peculiar nuclear components may be the result of an interaction event between two galaxies. Both major merging and accretion of external material may induce that some gas does flow in to the nuclear regions of the remnant and quickly forms a small concentration of new stars that maintain the original angular momentum of the initial galaxy, counter-rotate with respect to the host galaxy. Furthermore, N-body simulations including gas, stars and star formation, suggest that galaxies can develope a central velocity dispersion drop due to nuclear gas inflow, then subsequent star formation and the appearance of young luminous stars born from dynamically cold gas (\\cite{Wozniak06}).  -\\emph{Formation and evolution}- Our data are in good agreement with those obtained by \\cite{Rampazzo98} along the direction connecting the nuclei of NGC4778 and NGC4776 (P.A.$=128^\\circ$). They also found that the rotation curve is not symmetric with respect to the center of NGC4778, with a rapid increase in the SE direction at about $15''$ from the center. The velocity dispersion increases on both sides, reaching a maximum at about $10''$ from the center on the SW side. Given that the close galaxy NGC4776 is located on the NW, they suggested that {\\it i)} the rapid variation of the rotation curve and the sharp increase of the velocity dispersion in the SE direction is a real effect, reliable due to a dynamical perturbation; {\\it ii)} while the rise towards NGC4776 (on the NW side) could be partly an artifact due to the apparent superposition of the galaxies. The new kinematics along the SW side (presented in this work) further suggest that the South side of NGC4778 is dynamically perturbed.  As we have discussed in the section \\ref{kin}, the rotation curve of NGC4778 is not symmetric with respect to the center, and this is a feature observed in many other compact groups (\\cite{Bonfanti99}). According to the literature, many compact groups mainly composed by early-type galaxies, like HCG62, show several morphological signature of interactions, and for all of them the apparent kinematical interactions are not explainable as a mere optical superposition. This conclusion is strongly supported by the simulations performed by \\cite{Combes95}, which show that the peculiarities observed in many rotation curves of galaxies belonging to compact groups are due to intrinsic effects and not to contamination along the line of sight. Our results are also in agreement with the estimates coming from merger simulations (\\cite{Combes95}), that predict asymmetry in the kinematical profiles and a distinction between the photometric and the dynamical center.  The asymmetry and the shape of the rotation curve and velocity dispersion profile of NGC4778 do not find correlation with the photometric features of the galaxy, except for the bluer colors in the central region. The absence of correlation between the dynamical and the morphological peculiarities suggests that the dynamical properties of the HCG galaxies may be due to a minor merger event. In fact, as showed by \\cite{Nishiura00}, weak galaxy collisions could not perturb the galaxy rotation curves, but morphological deformations could be induced in the outer parts of the galaxy (tidal tails, bridges etc), while minor mergers could perturb the rotation curves in the inner regions, especially for gas-poor early-type galaxies, without causing morphological peculiarities.  We have estimated the mass-to-light (M/L) ratio of NGC4778 in order to derive some constraints on the amount of dark matter in HCG62. Since the kinematical profiles are not symmetric, for the calculation of the M/L ratio, we have used the value of $v_{max}$ and $\\sigma_{max}$ taken from the unperturbed side (NE) of the curves. Moreover, in absence of an accurate photometric calibration, we used as total B magnitude the value provided by NED, $m_{b} = 13.79$. Choosing the values $v_{max}\\simeq 80\\ km/s$, $\\sigma_{max}\\simeq 270\\ km/s$ and $R_{max}=34$ arcsec ($\\simeq 10\\ Kpc\\ \\simeq 2R_{e}$), by using the virial theorem $(M/L)_{vir}\\simeq\\frac{2R\\ (\\sigma^{2}+v^{2})}{(L\\ G)}$, we obtain $M/L\\simeq20.6$. This abnormally high mass-to-light ratio is compatible with a recent merging which has induced a tidal heating in the center of NGC4778, thus leading to a velocity dispersion which is too high with respect to the actual mass of the galaxy. This result however presents conflicting aspects. In fact, while such behaviour is predicted by numerical simulations (\\cite{Combes95}), a detailed study of the x-ray diffuse halo detected in the central regions of the group leads to a very similar virial estimate of M/L. A more detailed discussion of this point will be presented in \\cite{Sodani06}.  The velocity dispersion of HCGs are generally higher than would be expected given their visible mass (even if the discordant galaxies are ignored): this can also be explained if the bulk of the mass is in a non visible form (\\cite{Hickson97}). Moreover, ROSAT observations revealed a massive hydrogen envelope surrounding HCG62, and showed that this group is dominated by dark matter. Both N-body and hydrodynamic simulations indicate that the dark matter halos of individual galaxies merge first, creating a massive envelope within which the visible galaxies move (Barnes 1984, Bode et al 1993). Kinematic studies of loose groups (e.g. Puche \\& Carignan 1991) indicate that the dark matter is concentrated around the individual optical galaxies. In contrast, the X-ray observations indicate that in most compact groups, the gas and dark matter are more extended and are decoupled from the galaxies. This may be consistent with a M/L 30\\% to 50\\% lower in compact groups respect to isolated galaxies (Rubin et al 1991).  The hierarchical mergers of cluster galaxies might power the emission line gas in the center of the group members (\\cite{Valluri96}): according to the merger scenario, in order to power emission-line nebulae, the merger must include a galaxy or a group of galaxies that are late-types and which bring with them cold gas. The observed $H_{\\alpha}$ emission in NGC4778, and also in NGC4776 and NGC4761, further support the idea that this galaxy has recently experienced a merger event.  The overall scenario depicting NGC4778 as the product of a recent merger, as emerges by previous discussion, is consistent with the results obtained from X-ray observations. The presence of an extended X-ray halo is consistent with scenarios describing current compact group as the result of a first generation of mergers, where the dominant galaxy sits at the bottom of a large common gravitational well. The presence of two X-ray cavities in the hot gaseous halo located on symmetrically with respect to NGC4778 also indicate that the AGN residing in the galaxy core, must have undergone a recent (a few $10^{7}$ yr, \\cite{Birzan04}) active fase during which the radio-emitting relativistic plasma has created two low density regions within the hot IGM. It is commonly believed that such activity can be triggered by merging events, which increases the accretion rate onto the central massive black hole (\\cite{Cattaneo05}). These fits a scenario in which NGC4778 underwent a merger sometime in the past which produced the counter-rotating core and triggered the nuclear activity. However the low incidence of strong, type I AGN activity in interacting galaxies suggests that a delay of several $10^8$ years is generally expected until the peak of the AGN fase (\\cite{Canal06, Grogin05}, however see also \\cite{Koulouridis06}). In the case of NGC4778 we can derive a lower limit of $10^{7}$ yr for such delay from the age of the cavities. On the other end, an upper limit is represented by the age of the merger which, given the typical dynamical timescales of Compact Groups, can be estimated in $~10^8$ yr. This result agrees with the estimate that AGNs have duty cycles of the order of $10^{7-8}$ years."
    },
    "0606/astro-ph0606400.txt": {
        "abstract": "We present a panoramic review of several observational and theoretical aspects of the modern astrophysical research about the origin of the Fundamental Plane (FP) relation for Early-Type Galaxies (ETGs). The discussion is focused on the problem of the {\\it tilt} and the {\\it tightness} of the FP, and on the attempts to derive the luminosity evolution of ETGs with redshift. Finally, a number of observed features in the FP are interpreted from the standpoint of a new theoretical approach based on the two-component tensor virial theorem. ",
        "introduction": "The long-debated question of how many physical independent parameters can be used to describe the overall observational manifold of galaxies, arose approximately 30 years ago, in particular with a pioneering paper by \\citet{aa23-259}. Performing a principal component analysis (PCA) on a sample of spirals with known rotation curves, morphological types, absolute luminosities, colors and [HI] masses, a surprising result arose, that only 2 independent parameters dominate most of the variance of the related manifold. \\citet{brosche88} proved the possibility to scale all the main integral galaxian properties with two parameters, \\eg\\ mass and angular momentum. Since then, the multivariate analysis on galaxian parameters has been carried out by many authors \\citep[e.g.,][]{Lentes,mnras206-453,apj278-61,apj280-7}, substantially confirming the earlier conclusion.  An application of PCA to early-type galaxies (ETGs) yielded a similar result \\citep[][]{BRLen}, providing additional support to two previously discovered correlations, involving residuals from the Faber-Jackson (FJ) relation, $L_T \\propto \\sigma_o^{J}$ \\citep[e.g.,][]{apj204-668}. More precisely, the correlations are between (i) residuals in total luminosity, $L_T$, and metallicity index, $Mg_2$ \\citep[][]{aj100-1416}, and (ii) residuals in central projected velocity dispersion, $\\sigma_o$, and mean effective surface brightness, \\muem\\ \\citep[][]{apj230-697}.  At that time, the idea that ETGs belong to a family of stellar systems controlled by a single parameter, where their physical properties scale according to luminosity (mass), was supported by several observational evidences: the large degree of homogeneity of galaxy light profiles (well represented by the \\r1q\\ de Vaucouleurs law), the uniform color-magnitude diagrams, the existence of the FJ relation and of the Kormendy relation \\citep[KR,][] {Korm77}, $\\muem=3\\log(r_e)+const$, between the effective radius, \\re\\, and \\muem.  The above simple scenario started to change when \\citet{apj313-59} and \\citet{apj313-42}, taking advantage of a large sample of ETGs with available photometric and kinematical data, derived simultaneously the observational evidence that three physical parameters, $\\sigma_o$, \\muem, and \\re , are mutually correlated and unified in a 2-dimensional manifold, since then called {\\it Fundamental Plane} (FP). In this framework, both FJ and KR were interpreted as projections of the FP along the coordinate axes, which provides a natural explanation to the second, mysterious, ``hidden parameter''.  Current astrophysical observations and theoretical speculations have widely extended the spectrum of the structural and dynamical parameters that can be measured or calculated within galaxies, including both the visible stellar (baryonic) component, B, and the invisible dark matter (DM) component, D. Despite this large zoo of parameters, it is very surprising that the main dimensionality of this manifold remains 2. In other words, the cloud of points in the parameter space whose axes are size (mass, luminosity, or radius), density (or surface brightness) and temperature (\\ie\\ kinetic energy per unit mass), does not populate uniformly the three dimensional space, but is distributed approximately along a plane where the scatter maintains small.  The current paper aims to briefly review the most relevant attempts to explain the origin of the FP for ETGs. The mere existence of the FP does indeed indicate that structural properties in ETGs span a narrow range, suggesting some self-regulating mechanism must be at work during formation and evolution. The most relevant observational features of the FP are the tilt with respect to what is expected from one-component scalar virial theorem and homology (see section \\ref{sec1}), and a small thickness called tightness (see section \\ref{sec2}). According to observations, the scatter around the FP is very low, and the position of a galaxy above or below the plane is independent of galaxy flattening, isophotal twisting, velocity anisotropy, and details of radial light distribution. The small thickness corresponds to about 12\\% uncertainty in \\re , implying the FP is a good distance indicator \\citep[e.g.,][]{mnras330-443}.  The paper is organized as follows. In section \\ref{sec1} we review the most relevant works that in the last years provided new data for the FP at different wavelengths, redshifts and environments. Special effort is devoted to both the KR and the $\\kappa$-space. The latter makes a very interesting and debated representation of the FP, built up by \\citet{apj399-462} using a different set of orthogonal variables. In section \\ref{sec2} we summarize the most important theoretical attempts devoted to explain the origin of both the tilt and the tightness of the FP. We will address, in particular, the role of anisotropy and rotation in the stellar velocity distribution, the weak deviation of stellar systems from homology, the role of Initial Mass Function (IMF), the role of DM, and the most relevant aspects of a new theoretical approach which uses the two-component tensor virial theorem for interpreting some of the observed features of the FP. The conclusions are drawn in section \\ref{sec4}. ",
        "conclusions": "\\label{sec4} We have reviewed some of the more relevant observational investigations on the FP, and summarized the most important theoretical attempts to explain its tilt and tightness by means of velocity anisotropy, DM fraction, homology, and IMF. None of the above mentioned mechanisms, by itself, is currently able to provide a complete explanation of the observed properties of the FP.  We have shown that starting from the tensor virial theorem extended to a two-component model for ETGs, more insight on the problem can be gained.  If a maximum of Clausius' virial energy is really the key for a dynamic explanation of FP, it appears surprising how the corresponding special configuration links together many of the involved parameters, in such a way that the basic two-dimensionality of the manifold of ETGs is ensured.  The main observables \\re\\,, \\iem\\ , and $\\sigma_o$, appear to depend on cosmology via $\\gamma'$ (a quantity related to the local slope of the {\\it mass variance}) implying the projections of the FP (\\eg\\, FJ, KR) also depend on it. Nevertheless, the FP equation is found to degenerate with respect to $\\gamma'$.  The debate on the origin of the FP is still currently open, but interesting developments are to be expected as soon as larger field surveys will be available in the near future. \\\\  {\\bf Acknowledgments} This work has partially been supported by {\\em Fondazione Cassa di Risparmio di Padova e Rovigo}, Piazza Duomo 15, Padova (Italy). We thank an anonymous referee for critical remarks."
    },
    "0606/astro-ph0606056_arXiv.txt": {
        "abstract": "In this first paper of a series on the structure of boxy and peanut-shaped (B/PS) bulges, \\kn-band observations of a sample of $30$ edge-on spiral galaxies are described and discussed. \\kn-band observations best trace the dominant luminous galactic mass and are minimally affected by dust. Images, unsharp-masked images, as well as major-axis and vertically-summed surface brightness profiles are presented and discussed. Galaxies with a B/PS bulge tend to have a more complex morphology than galaxies with other bulge types, more often showing centered or off-centered X structures, secondary maxima along the major-axis and spiral-like structures. While probably not uniquely related to bars, those features are observed in three-dimensional N-body simulations of barred discs and may trace the main bar orbit families. The surface brightness profiles of galaxies with a B/PS bulge are also more complex, typically containing $3$ or more clearly separated regions, including a shallow or flat intermediate region (Freeman Type~II profiles). The breaks in the profiles offer evidence for bar-driven transfer of angular momentum and radial redistribution of material. The profiles further suggest a rapid variation of the scaleheight of the disc material, contrary to conventional wisdom but again as expected from the vertical resonances and instabilities present in barred discs. Interestingly, the steep inner region of the surface brightness profiles is often shorter than the isophotally thick part of the galaxies, itself always shorter than the flat intermediate region of the profiles. The steep inner region is also much more prominent along the major-axis than in the vertically-summed profiles. Similarly to other recent work but contrary to the standard `bulge + disc' model (where the bulge is both thick and steep), we thus propose that galaxies with a B/PS bulge are composed of a thin concentrated disc (a disc-like bulge) contained within a partially thick bar (the B/PS bulge), itself contained within a thin outer disc. The inner disc likely formed secularly through bar-driven processes and is responsible for the steep inner region of the surface brightness profiles, traditionally associated with a classic bulge, while the bar is responsible for the flat intermediate region of the surface brightness profiles and the thick complex morphological structures observed. Those components are strongly coupled dynamically and are formed mostly of the same (disc) material, shaped by the weak but relentless action of the bar resonances. Any competing formation scenario for galaxies with a B/PS bulge, which represent at least $45$~per cent of the local disc galaxy population, must explain equally well and self-consistently the above morphological and photometric properties, the complex gas and stellar kinematics observed, and the correlations between them. ",
        "introduction": "\\label{sec:intro} Since the work of \\citet{defis83}, bulges have largely been considered as low-luminosity elliptical galaxies, suggesting a rapid formation dominated either by mergers and accretion of external material (e.g.\\ \\citealt*{k96,kcw96}; but see also \\citealt*{bc01,abp01}) or possibly by dissipative gravitational collapse (e.g.\\ \\citealt*{els62,c84a,c84b,sm95}; but see also \\citealt{n98,n99}). Over the last decade, however, there has been much criticism of this idea. In particular, the realization that most bulges have an inner surface brightness profile shallower than the expected $R^{1/4}$ law \\citep*[e.g.][]{apb95,j96,mch03,bgdp03} argues against both mechanisms \\citep[e.g.][]{hm95,abp01}. Alternative models where bulges grow secularly (i.e.\\ over a long timescale and in relative isolation) have also been developed and studied, many of them bar-driven \\citep*[e.g.][]{pn90,fb93,fb95,nsh96}, and much observational data support them \\citep[e.g.][]{wfmmb95,es02}.  Of primary concern here, several pieces of evidence point to the identification of most boxy and peanut-shaped (B/PS) bulges in edge-on spiral galaxies with the bars of barred spirals. In $N$-body simulations, whenever a disc galaxy forms a bar, a B/PS bar/bulge develops soon after. This process was studied first by \\citet{cs81}, later on by \\citet{cdfp90} and \\citet{rsjk91}, and more recently by \\citet{mwhdb95}, \\citet{am02}, \\citet{a02,a03,a05} and \\citet*{msh05}.  The observed incidence of B/PS bulges is consistent with that expected if they are associated with relatively strong bars. Recent work by \\citet*{ldp00a,ldp00b} demonstrates that $45$ per cent of all bulges are B/PS, while amongst those the exact shape of the bulge depends mainly on the viewing angle to the bar. As shown by the numerical simulations, true peanuts are bars seen side-on, i.e.\\ with the major-axis of the bar roughly perpendicular to the line-of-sight. For less favourable viewing angles, the bulge/bar looks boxy, and if the bar is seen end-on it looks almost spherical. Stronger bars also lead to more prominent peanut shapes, as demonstrated observationally \\citep[e.g.][]{ldp00b} and theoretically \\cite[e.g.][]{am02,ba05}.  The kinematics of discs harboring a B/PS bulge, as measured from both ionized-gas emission lines and stellar absorption lines, show the behaviour expected of barred spirals viewed edge-on. This has been demonstrated by the observations of \\citet{km95}, \\citet{mk99}, \\citet{bf97,bf99} and \\citet{cb04}, and by the modeling of \\citet{ab99} and \\citet{ba99,ba05}. Similar tests are now also available for face-on bars \\citep{dcmm05}.  There is thus ample evidence that edge-on galaxies with a B/PS bulge are simply barred disc galaxies, and that the B/PS bulges themselves represent the thickest parts of the bars (see \\citealt{a05} for a review of all arguments). Yet there has been little convincing evidence for this from surface photometry alone, the best work being that of \\citet{ldp00a,ldp00b}. Early studies of non-spheroidal bulges used mostly optical images \\citep[e.g.][]{j87,sg89,s93} and the interpretation was often hampered by the large amount of extinction. As we will show in this paper, it truly takes the combination of N-body simulations and orbit studies with near-infrared (NIR) images to derive direct and convincing photometric evidence relating B/PS bulges and bars.  \\citet*{spa02a,spa02b} and \\citet*{psa02,psa03a} studied the orbital structure of three-dimensional (3D) bars exhaustively \\citep[but see also][]{p84,pf91}. They find families of orbits which can not only provide the backbone of the boxy and peanut shapes, but can also cause local enhancements within the disc itself. Since extinction is far less important in the NIR, $K$-band images are the ideal tool to study the morphology of galaxies with a B/PS bulge, to look for similarities with barred orbital structures.  In this paper, the first of a series, we study a sample of $30$ edge-on spiral galaxies, most of them with a B/PS bulge, for which we have obtained high-quality $K$-band images. Much complementary data exist for this sample \\citep[e.g.][]{bf97,bf99,cb04,bc06}, but the primary goal here is to study the morphology of the B/PS structures and their host discs, similarly to \\citet{ldp00b}. As we shall see, we find several features ressembling closely those expected of 3D bars: bulges with X shapes, secondary disc enhancements, inner rings, etc. A second paper discusses scaleheight variations in the same galaxies (\\citealt*{aab06}, hereafter \\citeauthor{aab06}; but see also \\citealt{aabbdvp03,abad04}).  We present our sample in \\S~\\ref{sec:sample}, discuss the observations and data reduction in \\S~\\ref{sec:obs}, and then describe and discuss the resulting images and surface brightness profiles in \\S~\\ref{sec:images} and \\ref{sec:sbprofs}, respectively. We examine the direct consequences of our results in \\S~\\ref{sec:discussion} and conclude briefly in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have presented \\kn-band imaging observations of a sample of $30$ edge-on spiral galaxies, most of which harbour a boxy or peanut-shaped (B/PS) bulge. Those data are minimally affected by dust and best trace population II stars, where most of the luminous mass resides. Our multi-faceted analysis suggests that B/PS bulges are simply the thick part of bars viewed edge-on (see Figures~\\ref{fig:bps} and \\ref{fig:control}).  Galaxies with a B/PS bulge tend to have a more complex morphology than galaxies with other bulge types, more often showing centered or off-centered X structures, secondary maxima along the major-axis, and spiral-like structures. Best revealed by unsharp-masking, those features are also observed in three-dimensional N-body simulations of barred discs \\citep{a05}, and can be explained by the orbital structure of bars \\citep[see, e.g.,][]{psa02}, although they need not be uniquely related to them. Only minor-axis extrema may be preferentially related to other bulge types. Whether taken along the major-axis or summed vertically (to simulate a flat galaxy), the surface brightness profiles of galaxies with a B/PS bulge are also more complex, more often showing $3$ or more clearly separated regions, including a rather shallow or flat intermediate region (see Figure~\\ref{fig:sbprof_cartoon}). Such Freeman Type~II profiles are expected from barred galaxies \\citep[e.g.][]{ba05}, but they do not have a natural self-consistent explanation in classic bulge formation scenarios.  The radial breaks observed in the vertically-summed profiles of our objects provide further evidence of the transfer of angular momentum and radial redistribution of disc material mediated by the (presumed) bars \\citep[see, e.g.,][]{a02,a03}. Furthermore, the spatial correlations of the radial breaks with the ionized-gas and stellar kinematics \\citep{bf99,cb04} are as expected for fast bars, currently favoured by observations. The differences between the major-axis and vertically-summed profiles provide evidence for abrupt variations of the scaleheight of the disc material. This is, again, as expected from the diverse orbital families and vertical resonances and instabilities present in barred discs, but contrary to conventional wisdom. A quantitative and robust analysis of those scaleheight variations and a comparison with $N$-body simulations will appear in future papers of this series.  Three other facts stand out. First, the steep inner region of the surface brightness profiles is systematically equal to or shorter than the isophotally thick part of the galaxies. Second, the isophotally thick part is itself systematically contained within the flat intermediate region of the surface brightness profiles. Third, the steep inner region of the surface brightness profiles is much more prominent along the major-axis than in the vertically-summed profiles.  We are thus led to radically alter the classic `bulge + disc' model, composed of a thick and steep spheroidal bulge largely decoupled from a thin (possibly barred) exponential disc. Analogously to \\citet{a05}, we propose here that galaxies with a B/PS bulge are composed of a thin concentrated disc (a disc-like bulge), formed secularly by the bar and responsible for the steep inner region of the surface brightness profiles, contained within a (partially) thick bar (the B/PS bulge), responsible for the flat intermediate region of the surface brightness profiles and the complex morphological structures, itself contained within a thin outer exponential disc (see Figure~\\ref{fig:bulges_cartoon}). Those components are closely intertwined dynamically and are largely made of the same (disc) material, shaped over long timescales by the bar.  The challenge to any competing formation scenario for galaxies with a B/PS bulge, which represent at least $45$~per cent of the local galaxy population \\citep{ldp00a}, is thus to simultaneously and self-consistently explain, equally well or better, their numerous morphological, photometric, and kinematic properties, as well as the correlations between them."
    },
    "0606/astro-ph0606260_arXiv.txt": {
        "abstract": "{We explore the mass-to-light ratio in galaxy clusters and its relation to the cluster mass. We study the relations among the optical luminosity ($L_{op}$), the cluster mass ($M_{200}$) and the number of cluster galaxies within $r_{200}$ ($N_{gal}$) in a sample of 217 galaxy clusters with confirmed 3D overdensity. We correct for projection effects, by determining the galaxy surface number density profile in our cluster sample. This is best fitted by a cored King profile in low and intermediate mass systems. The core radius decreases with cluster mass, and, for the highest mass clusters, the profile is better represented by a generalized King profile or a cuspy Navarro, Frenk \\& White profile. We find a very tight proportionality between $L_{op}$ and $N_{gal}$, which, in turn, links the cluster mass-to-light ratio to the Halo Occupation Distribution $N_{gal}$ vs. $M_{200}$. After correcting for projection effects, the slope of the $L_{op}-M_{200}$ and $N_{gal}-M_{200}$ relations is found to be $0.92\\pm0.03$, close, but still significantly less than unity. We show that the non-linearity of these relations cannot be explained by variations of the galaxy luminosity distributions and of the galaxy M/L with the cluster mass. We suggest that the nonlinear relation between number of galaxies and cluster mass reflects an underlying nonlinear relation between number of subhaloes and halo mass.} \\authorrunning{P. Popesso et al.} ",
        "introduction": "Clusters of galaxies are the most massive gravitationally bound systems in the universe. The cluster mass function and its evolution provide constraints on the evolution of large-scale structure and important cosmological parameters such as $\\Omega_m$ and $\\sigma _8$. Cluster mass-to-light ratios ($M/L$ hereafter) provide one of the most robust determination of $\\Omega_m$ in connection with the observed luminosity density in the Universe via the Oort (1958) method. In this method, a fundamental assumption is that the average $M/L$ of clusters is a fair representation of the universal value. For this reason, many works have focussed on the dependence of the cluster $M/L$ on the mass of the systems.  In general, $M/L$ has been found to increase with the cluster mass. Assuming a power-law relation $M/L \\propto M^{\\alpha}$, and adopting the usual scaling relations between mass and X-ray temperature or velocity dispersion, when needed, most authors have found $\\alpha$ in the range 0.2-0.4, in both optical and near-infrared bands, and over a large mass range (Adami et al. 1998a; Bahcall \\& Comerford 2002; Girardi et al. 2002; Lin et al. 2003, 2004; Rines et al. 2004; Ramella et al. 2004; see however Kochanek et al. 2003 for a discordant result). Why does the cluster $M/L$ increase with the mass? Based on the results of numerical simulations, Bahcall \\& Comerford (2002) have proposed that the trend of $M/L$ with mass is caused by the stellar populations of galaxies in more massive systems being older than the stellar populations of galaxies in less massive systems.  In this scenario, the slope of the $M/L-M$ relation should be steeper in the $B$ and $V$ bands, dominated by the young stellar populations, than at longer wavelengths, eventually becoming flat in the infrared $K$ band, dominated by the light of the old stellar population.  Such a scenario is not consistent with the results of the semi-analytical modeling of Kauffmann et al.  (1999), where the $M/L$ is predicted to increase with mass with approximately the same slope in the $B$ and $I$ band. Also observationally, the slope of the $M/L-M$ relation is found to be the same in different bands, the $B$-band (Girardi et al. 2002) the $V$-band (Bahcall \\& Comerford 2002), the $R$-band (Adami et al. 1998a, Popesso et al. 2005b,2005c) and the $K$-band (Lin et al. 2003, 2004; Rines et al. 2004; Ramella et al. 2004).  An alternative interpretation of the increasing $M/L$ with system mass is provided by Springel \\& Hernquist (2003). They analyze the star formation efficiency within halos extracted from cosmological simulations, with masses in the range $10^8-10^{15} M_{\\odot}$, and find that the integrated star formation efficiency decreases with increasing halo mass by a factor 5--10 over the cluster mass range. This scenario is investigated by Lin et al. (2003), who convert the 2MASS $K$-band cluster luminosities into cluster stellar masses. They find that the fraction of mass in stars is a decreasing function of the cluster mass ($M_{star}/M_{tot} \\propto M_{tot}^{-0.26}$).  In this paper we address the above issues by studying $M/L$ for a sample of 217 clusters, which span the entire cluster mass range. In particular, we study the relations among the cluster optical luminosity $L_{op}$, the mass $M_{200}$, and the number of cluster galaxies $N_{gal}$, within the virial radius $r_{200}$. We find a very tight relation between $L_{op}$ and $N_{gal}$, which links the $L_{op}-M_{200}$ relation (and therefore, the cluster $M/L$), to the Halo Occupation Distribution (HOD hereafter) $N_{gal}-M_{200}$.  The HOD is a powerful tool for describing galaxy bias and modelling galaxy clustering (e.g. Ma \\& Fry 2000; Peacock \\& smith 2000; Seljak 2000; Scoccimarro et al. 2001; Berlind \\& Weinberg 2002). It characterizes the bias between galaxies and mass in terms of the probability distribution $\\rm{P(N|M)}$ that a halo of virial mass M contains N galaxies of a given type, together with relative spatial and velocity distributions of galaxies and dark matter within halos. The HOD is a fundamental prediction of galaxy formation theory (e.g. Kauffmann, Nusser \\& Steinmetz 1997, Kauffmann et al. 1999; White, Hernquist \\& Springel 2001; Yoshikawa et al. 2001; Berlind et al. 2003; Kravtsov et al. 2004; Zheng et al. 2005) and it can be extremely useful to compare the observational results with the theoretical models.  This paper is organized as follows. In section 2 we describe our dataset. In section 3 we describe the methods we use to calculate several cluster properties, like the characteristic radius, the virial mass, the optical luminosity, and the number density profile of cluster galaxies. In section 4 we analyze the $L_{op}-M_{200}$ and the $N_{gal}-M_{200}$ relations, and find that the number of galaxies per given halo mass decreases as the halo mass increases.  In section 5 we seek for a physical explanation of this trend, also by comparing our results with theoretical predictions. Finally, in section 6 we provide our conclusions.  Throughout this paper, we use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ in a flat cosmology with $\\Omega_0=0.3$ and $\\Omega_{\\Lambda}=0.7$ (e.g. Tegmark et al. 2004). ",
        "conclusions": "We have studied the $L-M$ and the $N-M$ relations in the 4 SDSS bands g, r, i, z for a sample of 217 galaxy clusters with confirmed 3D overdensity in the SDSS DR3 spectroscopic catalog. All the quantities are measured within the characteristic cluster radius $r_{200}$. We have remarked upon the direct connection between the two relations due to the proportionality of the cluster optical luminosity and the number of cluster galaxies.  We have studied the galaxy surface number density profile in five bins of cluster mass and discovered that the profile has a strong dependence on the cluster mass.  In the low and intermediate mass systems the best fit is provided by a King profile. The core radius of the best fit is decreasing as a function of the cluster mass, while the central galaxy density is increasing. In the highest mass bins a more concentrated generalized King profile or a cuspy NFW profile provide the best fits. Using the best fit profile in each mass bin, we have converted the observed number of cluster galaxies to the value within the virial sphere. Since clusters of different masses exhibit different surface density profiles, the deprojection correction decreases with the cluster mass.  Applying this mass-dependent correction affects the $L-M$ and $N-M$ relations, by increasing the slope of these relations to the value of $0.92\\pm0.03$. Similarly, also the slope of the $M/L-M$ relation is affected and becomes $0.18\\pm0.04$. Hence, neglecting the dependence of the deprojection correction on the cluster mass, leads one to underestimate the slope of the $L_{op}-M_{200}$ and $N_{gal}-M_{200}$ relations.  Despite the deprojection correction, the derived $N-M$ and the $L-M$ relations are still only marginally consistent with unity, at the 2.5$\\sigma$ level, i.e. direct proportionality between cluster mass and number of cluster galaxies is not supported.    We have compared the properties of our clusters with the prediction of the hierarchical models of structure formation. These models naturally predict that $N \\propto M^{\\gamma}$ with $\\gamma < 1$. This result is generally interpreted as the indication that the galaxies in the low mass systems are older and more luminous per unit mass than the galaxies in high mass clusters. As a consequence, variations of the shape of the cluster LF and of the elliptical FP with the cluster mass are also expected. Such predicted variations are however not seen in our data. Not only we found the LF to be the same for clusters of different masses, but we also proved that this universal LF can be used to accurately predict the magnitudes of the three brightest cluster galaxies, given the LF-normalization of the clusters in which they are located. In other words, the BCG magnitudes are consistent with being drawn from the best-fit magnitude distribution of other cluster galaxies. Moreover we have shown that the FP of cluster ellipticals has the same slope in all the clusters and does not depend on the cluster mass.  We conclude this paper with the following considerations. From the observational point of view, the mean cluster luminosity function and the $N-M$ or the $L-M$ relation determine completely the luminosity distribution of cluster galaxies. The mean cluster LF constrains with high accuracy the shape of the luminosity distribution in clusters, while the $N-M$ relation, calculated in a given magnitude range, fixes the normalization of the LF as a function of the cluster mass. Forthcoming cosmological models of galaxy formation should aim at reproducing this characteristic of the cluster galaxy populations, in order to understand the processes of galaxy formation and evolution in the cluster enviroment.  \\vspace{2cm}  We thank the referee, Christophe Adami, for the useful comments which helped in improving the paper. We acknowledge useful discussions with Stefano Borgani and Simon White.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S.  Despartment of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.  The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0606/astro-ph0606110_arXiv.txt": {
        "abstract": "We present panoramic {\\it Spitzer MIPS} 24-$\\mu$m observations covering $\\sim$9$\\times$9\\,Mpc ($25'\\times 25'$) fields around two massive clusters, Cl\\,0024+16 and MS\\,0451$-$03, at $z=0.39$ and $z=0.55$ respectively, reaching a 5-$\\sigma$ flux limit of $\\sim 200\\mu$Jy. Our observations cover a very wide range of environments within these clusters, from high-density regions around the cores out to the turn-around radius.  Cross-correlating the mid-infrared catalogs with deep optical and near-infrared imaging of these fields, we investigate the optical/near-infrared colors of the mid-infrared sources.  We find excesses of mid-infrared sources with optical/near-infrared colors expected of cluster members in the two clusters and test this selection using spectroscopically confirmed 24$\\mu$m members.  The much more significant excess is associated with Cl\\,0024+16, whereas MS\\,0451$-$03 has comparatively few mid-infrared sources.  The mid-infrared galaxy population in Cl\\,0024+16 appears to be associated with dusty star-forming galaxies (typically redder than the general cluster population by up to $A_V\\sim1$--2\\,mags) rather than emission from dusty tori around active galactic nuclei (AGN) in early-type hosts.  We compare the star-formation rates derived from the total infrared (8--1000$\\mu$m) luminosities for the mid-infrared sources in Cl\\,0024+16 with those estimated from a published H$\\alpha$ survey, finding rates $\\gs5\\times$ than those found from H$\\alpha$, indicating significant obscured activity in the cluster population. Compared to previous mid-infrared surveys of clusters from $z\\sim0$--0.5, we find evidence for strong evolution of the level of dust-obscured star-formation in dense environments to $z=0.5$, analogous to the rise in fraction of optically-selected star-forming galaxies seen in clusters and the field out to similar redshifts. However, there are clearly significant cluster-to-cluster variations in the populations of mid-infrared sources, probably reflecting differences in the intracluster media and recent dynamical evolution of these systems. ",
        "introduction": "\\begin{figure*} \\centerline{ \\psfig{file=f1a.ps,width=3.5in,angle=0}~~ \\psfig{file=f1b.ps,width=3.5in,angle=0}}  \\caption{\\small The $(B - R)$--$R$ color-magnitude diagram for (left) Cl\\,0024+16 and (right) MS\\,0451$-$03. We identify 24-$\\mu$m sources brighter than 200$\\mu$Jy and compare these to the distribution for the optically-selected populations in these fields (for clarity, this has been represented by a smoothed density plot).  Also shown are histograms showing slices through the distributions, comparing the number counts and colors for the mid-infrared sources to the optically-selected population.  To emphasise the location of the sequence of early-type galaxies in the clusters, the $(B-R)$ histogram is limited to $R\\leq 22$. The mid-infrared selected population lies between the red and blue galaxy peaks in these fields, most likely because of the influence of dust on intrinsically blue, star-forming galaxies.  For the mid-infrared sources in the color-magnitude diagram we have removed those objects morphologically classified as stars.  All colors are measured in 2$''$ diameter apertures, and we use the {\\sc best\\_mag} estimate of the total galaxy magnitude. } \\end{figure*}   \\begin{figure*}[t] \\centerline{ \\psfig{file=f2a.ps,width=3.5in,angle=0}~~ \\psfig{file=f2b.ps,width=3.5in,angle=0}} \\caption{\\small We show $(R-K)$ versus $(B-R)$ colors for 24$\\mu$m-detected sources compared to the optical distribution in the Cl\\,0024+16 (left) and MS\\,0451$-$03 (right) fields.  The distributions are limited at $R\\leq24$ and $K\\leq20$ and we represent the optical sample's number density as a slightly smoothed grayscale.  The dashed tracks trace the expected colors of galaxies across $z=0$--2, with star formation histories similar to that of present-day Sdm--E galaxies, and the solid line and points show the location in color--color space of the range of spectral types at the cluster redshift.  The arrow indicates the translation in color--color space corresponding to an increase in reddening of $A_V=1$. To perform a rough photometric cut to select 24$\\mu$m sources at the cluster redshift, we define a region around the predicted colors of cluster members indicated by the rectangular selection boxes in both panels. To avoid contamination by infrared emission from AGN, we designate a sub-region which should contain the passive galaxies (dashed-line box) -- 24$\\mu$m sources in this region are not included in our analysis of the obscured star-forming populations. We show the efficacy of this selection by plotting the colors of known spectroscopically-confirmed 24$\\mu$m cluster members, 85\\% of these fall within the selection box.} \\setcounter{figure}{2} \\end{figure*}   The galaxy populations within the virialised regions of rich clusters at $z\\sim 0$ are characterised by passive elliptical and lenticular (S0) galaxies (Oemler 1974; Dressler 1980).  In contrast, 5-Gyrs ago, at $z\\sim 0.5$, the galaxy populations in the most massive clusters had larger fractions of star-forming late-type spirals, and a corresponding deficit of luminous S0 galaxies (e.g.\\ Dressler et al.\\ 1997; Smail et al.\\ 1997; Couch et al.\\ 1998; Fasano et al.\\ 2000; Treu et al.\\ 2003).  Taken together, these two observations imply that a process (or processes) is transforming many of the star-forming, late-type spirals in these regions into the passive early-type population (specifically S0s) found in local clusters (Poggianti et al.\\ 1999; Moran et al.\\ 2006, in prep).  When considering potential pathways to produce this evolutionary change, we need to bear in mind that the typical luminosities of the star-forming spirals appear to be too low for them to transform into typical S0 galaxies found in local clusters, without the addition of significant numbers of new stars (Poggianti et al.\\ 1999; Kodama \\& Smail 2001).  This problem is exacerbated when we include the fading which is likely to take place after the cessation of star formation in these galaxies.  This then leads us to concentrate on mechanisms which are capable of increasing the luminosity of the galaxies -- mergers and starbursts.  There has been a recent upsurge in interest in the potential for so-called ``dry'' mergers (mergers between dissipationless stellar systems which don't result in additional star formation) to influence the evolution of early-type galaxies (van Dokkum et al.\\ 2003; Bell et al.\\ 2003; van Dokkum 2005).  However, the dynamically hot environments in rich clusters which are the subject of our study are deleterious to the formation and survival of cold, bound-pairs of early-type galaxies -- unless these systems arrive in the cluster as existing bound entities.  It is not clear, therefore, that dry mergers can provide an effective route to substantially increase the number of luminous, early-type S0 galaxies within clusters.  Unfortunately, the alternative mechanism for enhancing the luminosity of the bulge component -- a starburst -- also has strong observational evidence stacked against it.  Surveys of star-forming galaxies in clusters using optical or UV star formation indicators have failed to detect galaxies with strongly enhanced star-formation which would have to exist to explain the growth of the bulge components of early-type galaxies in clusters at $z\\ls 0.5$--1 (Balogh et al.\\ 1999; Poggianti et al.\\ 1999; Gerken et al.\\ 2004).  However, there is growing body of evidence that at least some of the galaxies in distant clusters may be undergoing bursts of star-formation, albeit ones which are heavily shrouded in dust.  Smail et al.\\ (1999) used a deep VLA 1.4-GHz radio map to study a small sample of active galaxies within the core of the cluster Cl\\,0939+4713 ($z=0.41$). Combining the radio data with near-infrared and optical morphological information from the {\\it Hubble Space Telescope (HST)} and ground-based spectroscopy, they found that over half the radio-emitting population in the core are dusty late-type galaxies, presumably undergoing vigorous star formation. However, the spectral classification of these spirals placed them in the post-starburst class, and indeed all the post-starburst galaxies in this small region are radio-emitters.  Dust has also been used to explain the unusual spectral properties of another class of galaxies found in distant clusters and the field: e(a) galaxies, which show enhanced Balmer absorption compared to normal star-forming galaxies (Poggianti et al.\\ 1999).  Poggianti \\& Wu (2000) and Poggianti, Bressan \\& Franceschini (2001) discuss models for these galaxies invoking age-dependent dust obscuration of the younger stellar populations -- enabling significant activity to be hidden from view in these systems. If the passive lenticular galaxies found in local clusters, but absent from the equivalent rich environments at higher redshift, are the result of infalling late-type galaxies undergoing dusty-starburst in high-$z$ clusters, then a possible signature would be evolution in the total level of obscured star-formation in clusters out to $z\\sim1$.  In principle, mid- and especially far-infrared/submillimeter observations give us a direct probe of the level of obscured activity in distant clusters.  In particular, mid-infrared observations with {\\it Infrared Space Observatory (ISO)}, and more recently {\\it Spitzer Space Telescope (SST)}, provide sensitive imaging capabilities which can trace dusty star formation in clusters out to $z\\sim 1$ and beyond. Metcalfe, Fadda \\& Biviano (2005) summarise the results from {\\it ISO} surveys of distant clusters, which have yielded a total of just $\\sim 40$ cluster galaxies detected at 15\\,$\\mu$m across seven clusters between $z=0.18$--0.56 (Duc et al.\\ 2000, 2004; Fadda et al.\\ 2000; Metcalfe et al.\\ 2003; Coia et al.\\ 2005a,b; Biviano et al.\\ 2004). All these studies suggest that there is an increased level of mid-infrared activity in distant clusters, at levels above that suggested by UV/optical tracers of star formation.  Submillimeter observations of more distant clusters have also hinted at possible enhanced activity in these environments (Best 2000; Webb et al.\\ 2005). However, the inhomogeneous mix of coverage and depth in the samples coupled with the modest numbers of sources detected in any individual cluster mean that it has proved difficult to use these data to provide quantitative constraints on the origin and evolution of dust-obscured activity in distant clusters.  The {\\it SST}'s sensitive mid-infrared imaging capabilities provide an unique opportunity to undertake complete and representative surveys of the obscured, active populations in distant clusters.  To search for a population of mid-infrared sources in rich clusters environments, we have therefore used the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) to detect 24-$\\mu$m emission from galaxies in two clusters at $z\\sim0.5$ covering a very wide range in environment from $\\sim1$\\,Mpc out to the turn-around radius ($\\sim5$\\,Mpc) where the clusters merge into the surrounding field. These observations will provide measures of the level of obscured star-formation in these clusters, and so allow us to build up a reliable picture of the evolution of dust-obscured activity in clusters over the past 5\\,Gyrs.  This paper presents a statistical analysis of the 24$\\mu$m populations in two $z\\sim 0.5$ clusters. A subsequent paper (Geach et al.\\ in prep) will discuss the properties of these sources in more detail using the available spectroscopic and morphological surveys of the clusters (Moran et al.\\ 2006). The paper is organised as follows: we describe our observations and their reduction in \\S2, analyse these in \\S3 and discuss our results and present our conclusions in \\S4 and \\S5, respectively.  Throughout, we adopt a geometry with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "We have used the MIPS instrument on {\\it Spitzer} to survey the 24$\\mu$m populations of two optically rich clusters at $z\\sim0.5$: Cl\\,0024+16 and MS\\,0451$-$03. The samples are $\\sim80$\\% complete at $200\\mu$Jy, corresponding to total (8--1000$\\mu$m) infrared luminosities of $6\\times10^{10}$\\,L$_\\odot$ and $12\\times10^{10}$\\,L$_\\odot$ at $z=0.39$ and $z=0.55$ respectively, equivalent to  minimum SFRs of $\\sim10$\\,M$_\\odot$\\,yr$^{-1}$ and $\\sim20$\\,M$_\\odot$\\,yr$^{-1}$. We detect a total of 986 and 1018 mid-infrared sources above this flux limit across $\\sim25'\\times25'$ fields in Cl\\,0024+16 and MS\\,0451$-$03. Our observations probe from around the cores out to the turn-around radius at $\\sim5$\\,Mpc where the clusters merge into the field. In MS\\,0451$-$03, we also analyse archival MIPS observations of the central $\\sim5'\\times5'$ of the cluster, which our observations had to avoid. Similarly in Cl\\,0024+16 we make use of an {\\it ISO} 15$\\mu$m survey from Coia et al.\\ (2005) in the central region to build up a picture of the distribution of mid-infrared sources over the complete range of cluster environments.  We exploit optical-near-infrared colors for the mid-infrared sources to reduce the background field contamination. We find a statistical excess of mid-infrared sources (within $\\sim 5$\\,Mpc of the cluster core) at $S_{\\rm 24\\mu m} > 200\\mu$Jy associated with Cl\\,0024+16: 155$\\pm$18. In contrast MS\\,0451$-$03 has a less significant population of mid-infrared sources, $28\\pm17$, although we note that there are a small number of confirmed 24$\\mu$m members in MS\\,0451$-$03 in our on-going spectroscopic survey of this cluster.  Using our deep optical and near-infrared imaging of both clusters we show that the 24$\\mu$m sources in Cl\\,0024+16 are mostly associated with star-forming galaxies, with typically blue $(B-R)$ colors, but which can be dust reddened by up to $A_V\\sim2.5$ mags. We also compare the infrared star-formation rate to that derived from an optical narrow-band H$\\alpha$ survey of this cluster from Kodama et al.\\ (2004). Typically the H$\\alpha$-derived rates underestimate the extinction-free infrared rate by $\\gs5\\times$, suggesting significant obscuration of the activity in this cluster. We find that the level of obscuration for these individual cluster galaxies is comparable to that found for LIRGs in the field at similar epochs. This suggests that starbursts in clusters are similar (at least in terms of extinction) as those triggered in low-density environments. However the variation in the 24$\\mu$m populations of Cl\\,0024+16 and MS\\,0451$-$03 suggests that the range of triggering and suppression mechanisms in clusters is complex.  We estimate that the total star-formation rate (derived from the infrared) in the central region of Cl\\,0024+16 ($\\ls R_{200}$) is $1000\\pm210$\\,M$_\\odot$\\,yr$^{-1}$. MS\\,0451$-$03 is much poorer in mid-infrared sources, and we derive a total star-formation rate estimate by summing over the small excess of objects, giving a total star-formation rate of $200\\pm100$\\,M$_\\odot$\\,yr$^{-1}$ within the same physical radius.  We note however that our mid-infrared survey can miss some star-formation if a substantial number of lower-luminosity galaxies also exist in these clusters (as shown by the H$\\alpha$ survey of Kodama et al.\\ 2004). However, we show that the majority of the activity is dominated by these dusty-starbursts.  Finally, we look at the evolution of the specific star formation rate per cluster with redshift from $z\\sim 0$--0.5, using our new estimates for the total star-formation rates in Cl\\,0024+16 and MS\\,0451$-$03.  We compare these to estimates for lower and comparable redshift clusters studied with {\\it ISO} and {\\it IRAS}. We find that the high redshift clusters tend to have larger total star-formation rates compared to the more quiescent low-redshift ones, with an evolution similar to that of field LIRGs. However there is considerable scatter in this relation, and the evolution may only apply to the most active clusters. Although it is still unclear exactly what processes govern the star-formation histories of rich clusters, this study has shown that rich environments can sustain significant amounts of hidden star-formation, and that this seems to increase at least out to $z\\sim0.5$. This hidden activity may have a profound influence on the life-cycle of galaxies in high-density regions and the formation of the passive galaxy populations, ellipticals and S0s, which inhabit these environments at the present-day. We will investigate these issues in more detail in our next paper where we bring together spectroscopic and morphological information on the 24$\\mu$m population in these fields. We are also extending our survey with new panoramic observations of distant clusters with {\\it Spitzer} Cycle 3 GO time using both MIPS and the IRS."
    },
    "0606/astro-ph0606236_arXiv.txt": {
        "abstract": "We present \\textit{Spitzer Space Telescope} IRAC and MIPS observations of the galactic globular cluster M15 (NGC 7078), one of the most metal-poor clusters with a [Fe/H] = -2.4.  Our \\textit{Spitzer} images reveal a population of dusty red giants near the cluster center, a previously detected planetary nebula (PN) designated K648, and a possible detection of the intra-cluster medium (ICM) arising from mass loss episodes from the evolved stellar population.  Our analysis suggests $(9 \\pm 2) \\times 10^{-4}$~M$_{\\odot}$ \\ of dust is present in the core of M15, and this material has accumulated over a period of $\\approx 10^{6}$~yrs, a timescale ten times shorter than the last galactic plane crossing event. We also present \\textit{Spitzer} IRS follow up observations of K648, including the detection of the [Ne~II] $\\lambda 12.81$~\\micron \\ line, and discuss abundances derived from infrared fine structure lines. ",
        "introduction": "\\label{sec:intro}  Mass loss in evolved stellar populations affects the chemical evolution of the interstellar medium (ISM) and mass loss from individual stars governs post main sequence evolution. The amount and duration of mass loss that occurs in giant stars remains one of the most uncertain parameters in stellar evolution theory and the effect of these processes on the inferences derived from stellar population models can be significant. Given the wide use of these models (e.g., in inferring stellar masses of high-redshift galaxies), understanding the mass losing process is vital for a range of problems in astrophysics. Although dust constitutes a small fraction of the total mass lost, it is frequently used as a marker as it is optically thin and its thermal emission is readily detectable in a variety of environments.  Globular clusters (GCs), believed to have formed during the assemblage of the Galaxy, are coeval samples of stars at common, well-determined distances with nearly uniform initial compositions. GCs enable study of the chemical enrichment of the interstellar medium arising from mass ejection during the post-main sequence evolution of stars. Red Giant stars, especially those ascending the Asymptotic Giant Branch, are expected to develop winds that inject processed material into the intra-cluster medium (ICM) during post-main sequence evolution.  These winds contain gas and solid phase materials, the latter in the form of dust grains that condense from the metals. However, detection of thermal emission arising from intra-cluster medium dust has been elusive, suggesting that the ICM in GCs is 100 to 1000 times less massive than expected from current stellar evolution theory and observations of mass-losing stars in clusters and in the solar neighborhood.  The circumstellar environments of stars in the late stages of evolution, when most mass loss occurs, are most effectively detected and studied in the infrared (IR). IR surveys conducted by {\\it IRAS}, {\\it 2MASS,} and {\\it ISO} revealed populations of dust-enshrouded asymptotic giant branch (AGB) and red supergiant (RSG) stars in the Galactic bulge \\citep{Jacco03}, the Large Magellanic Cloud (LMC) \\citep{Zijlstra96, jacco97, trams99}, and in galactic Globular Clusters \\citep{Ramdani01,origlia02}. GCs are expected to contain dust from episodes of mass loss in red giant branch (RGB) and AGB stars. The IR excess of their circumstellar dust emission is expected to peak between 20~\\micron{} to 30~\\micron{}, and thus photometry at wavelengths larger than 20~\\micron{} is necessary to estimate accurate dust masses lost by such stars. The amount of dust present in the intra-cluster medium (ICM) will vary depending on the cluster escape velocity, the time since last crossing of the Galactic disk where the ICM can be stripped away by the ISM, and the number of mass-losing stars. In general the dust in the ICM of GCs is expected to be $\\simeq 10^{-2}$ to $10^{-3}$ M$_{\\odot}$ for most galactic clusters. Globular clusters have reasonably homogeneous (in age and metallicity), well-understood stellar populations, so observations of ICM dust are reasonably straightforward to interpret to yield mass loss rates and duty cycles.  Previous attempts to detect the ICM in GCs suggest that the ICM density is well below that expected from predictions of the mass loss input from RGB and AGB stars, even considering the low metallicity of GCs. The lowest 3-$\\sigma$ upper limits to the ICM mass for 70~K dust are $\\sim 6\\times10^{-5}$ M$_\\odot$ \\citep{hopwood99}. Detecting thermal emission from the elusive ICM in GCs is observationally challenging. \\citet{origlia02} reported \\textit{ISO} observations of the IR thermal emission from the winds of individual RGB stars in six massive GCs (47 Tuc, NGC 362, Omega Cen, NGC 6388, M15, and M54) showing that, in those systems, stellar winds from these stars are enriching the ICM.  Though thermal emission from the ICM material might be expected to be detectable, many attempts to do so with IR and millimeter observatories have produced only a single secure detection of ICM dust, a tentative (3.5-$\\sigma$) detection of thermal emission in the core of the  metal-poor GC M15 \\citep{evans03}.  Overall this result suggests that the ICM dust in GCs is significantly less massive than expected from current stellar evolution theory and observations of mass losing stars in GCs and the solar neighborhood.  The causes of the paucity of ICM emission have been proposed to be ram-pressure stripping of ICM gas during Galactic plane passage, blowout by nova explosions, fast winds from the stars themselves, radiative ejection by the sheer luminosity of cluster stars, and continuous ram-pressure from hot gas in the galactic halo.  However, the dominant state of the ICM is unclear. Sensitive searches for neutral H in the ICM at radio wavelengths have yielded upper limits  $\\leq$ 0.1 M$_\\odot$ \\citep{Birk83}, with a possible detection of $\\approx$ 200 M$_\\odot$ in NGC 2808 \\citep{Faulk91} and a 5$\\sigma$ detection of 0.3 M$_\\odot$ in M15 \\citep{Jacco06}. Perhaps much of the ICM is ionized, as suggested by the high electron densities measured from pulsar timing in 47 Tuc \\citep{Paolo01}. However, H$\\alpha$ searches have as yet been unsuccessful.  ICM dust grains in radiative thermal equilibrium should attain temperatures of 50~K to 80~K because of the high energy density of starlight within a GC \\citep{Forte02} and thus be detectable as an IR excess (above the photospheric emission) in GCs at mid- to far-IR wavelengths. Here we present observations of the galactic GC M15 with the \\textit{Spitzer Space Telescope} \\citep{Werner04}, whose instrumental sensitivity enables detection of dust masses as low as $4\\times10^{-9}$ M$_\\odot$ (assuming a population of grains radiating in thermal equilibrium, with integration times down to the background confusion limit) and therefore permits an unequalled opportunity to search for and set stringent limits on the ``missing'' ICM.  \\subsection{M15} \\label{sec:m15}  M15 (NGC 7078), with [Fe/H] = -2.4 \\citep{sneden97}, is one of the most metal poor GCs. It is a well-studied cluster, as it is home to the first planetary nebula (PN) discovered in a GC (K648, also designated as Ps-1, \\citet{pease28,howard97,alves00}) and to the first GC Low-Mass X-ray Binary source (X2127+119, \\citet{auriere84,charles86}).  At least eight millisecond pulsars are also associated with the cluster \\citep{Kulkarni96}.  M15 is generally believed to be a core-collapse GC, with a small, dense core containing approximately 4000~M$_\\odot$ \\citep{phinney96}.  Properties of M15, reproduced from \\citet{hopwood99}, are listed in Table~\\ref{tab:m15_params}. Updated values as listed by \\citet{evans03} include the escape velocity $V_{esc}$ from \\citet{webbink85}, and the time $\\tau_c$ since the last plane crossing from \\citet{oden97}.  The reddening and the distance are updated from \\citet{schlegel98} and \\citet{mcnamara04}, respectively. Galactic coordinates for M15 are $l = 65.01^\\circ$ and $b = -27.31^\\circ$, placing the cluster $\\sim$ 4.5~kpc south of the galactic plane.  The total dust mass expected in a GC can be estimated using the following equation:  \\begin{equation} M_{dust} = \\frac{\\tau_c}{\\tau_{HB}} \\ N_{HB} \\ \\delta M \\ \\frac{10^{[Fe/H]}}{100}, \\end{equation}  \\noindent where $\\tau_{HB}$ is the Horizontal Branch (HB) lifetime, $N_{HB}$ is the number of HB stars, [Fe/H] is the cluster metallicity, and $\\delta M$ is the dust mass lost from each star at the tip of the RGB. The factor of 100 is the solar gas-to-dust ratio, which is scaled for the metallicity of M15 by adding the [Fe/H] factor. Using values typical of population II stars and this relationship, the expected dust mass in M15 has been estimated to be 3.7 $\\times$ 10$^{-3}$ M$_{\\odot}$ by \\citet{evans03} and 2.0 $\\times$ 10$^{-3}$ M$_{\\odot}$ by \\citet{hopwood99}.  M15 is also home to the PN K648.  K648 was the first globular cluster PN discovered, and subsequently it has been extensively studied at UV, optical and IR wavelengths to determine the chemical composition of the ejecta nebula and parameters of the central star \\citep{barker83, adams84, howard97, bianchi01, hld03, garnett93}. The study of post-AGB stellar evolution in old metal-poor, low-mass stars in GCs or the galactic halo can be greatly enhanced if the by-products of stellar nucleosynthesis can be measured. Enriched material produced in the RGB and AGB stages of stellar evolution are dispersed into outer layers of the stellar system and subsequent mass loss processes lead to the formation of a PN. Few PNe in the galactic halo population have been identified and only four of these PNe, including K648, BB-1, DdDM~1, and H4-1 are associated with globular clusters \\citep{jacoby97}.  Below, we present findings derived from a 5-15 \\micron{}~ IR spectrum of K648. ",
        "conclusions": "\\label{sec:concl}  Analysis of our \\textit{Spitzer} image data on the core of the globular cluster M15 show strong evidence for the presence of intercluster medium (ICM) dust in the cluster core, with a mass of $(9 \\pm 2) \\times 10^{-4}$~M$_\\odot$ and with an equilibrium temperature of $\\approx 70$~K. This is the first secure, high signal to noise detection of ICM dust in a globular cluster.  Also present surrounding the core are populations of dusty AGB and post-AGB stars, along with the planetary nebula K648. Using IRS spectral data, we have observed both the [\\ion{S}{4}] and [\\ion{Ne}{2}] fine structure lines in K648 and have derived abundance estimates.  The unique capabilities of {\\it Spitzer} have enabled us to identify both the interstellar dust and the dust producers in M15.  This is surprising at such low metallicity ([Fe/H] = -2.4), and may have implications for dust production in the early universe.  The mass of the ICM dust in M15 suggests that it has been accummulating for $\\sim 10^6$ years, which is a factor of ten shorter than the time since the last galactic plane crossing. The dust mass is also approximately 4 times smaller than the mass predicted by \\citet{evans03}. Both of these results imply that such dust does not survive long compared to its production rate, and is thus part of a stochastic process."
    },
    "0606/astro-ph0606146_arXiv.txt": {
        "abstract": "{We have observed the Crab Pulsar in the optical with S-Cam, an instrument based on Superconducting Tunneling Junctions (STJs) with $\\mu$s time resolution.} {Our aim was to study the delay between the radio and optical pulse.} {The Crab Pulsar was observed three times over a time span of almost 7 years, on two different locations, using three different versions of the instrument, and using two different GPS units.} {We consistently find that the optical peak leads the radio peak by 49$\\pm$90, 254$\\pm$170, and 291$\\pm$100~$\\mu$s. On assumption of a constant optical lead, the weighted-average value is $\\sim$170~$\\mu$s, or when rejecting (based on a perhaps questionable radio ephemeris) the first measurement, 273$\\pm$100~$\\mu$s.} {} ",
        "introduction": "Precise timing of pulsar light curves throughout the electromagnetic spectrum is a powerful tool to constrain theories of the spatial distribution of various emission regions. In recent years, it has become clear that the main and secondary pulses of the Crab Pulsar (PSR J0534+2200) are not aligned in time at different wavelengths. X-rays are leading the radio pulse by reported values of 344$\\pm$40~$\\mu$s (Rots et al.\\ \\cite{rots:2004}; RXTE data) and 280$\\pm$40~$\\mu$s (Kuiper et al.\\ \\cite{kuiper:2003}; INTEGRAL data) and $\\gamma$-rays are leading the radio pulse by 241$\\pm$29~$\\mu$s (Kuiper et al.\\ \\cite{kuiper:2003}; EGRET data. The uncertainty in this value does not include the EGRET absolute timing uncertainty of better than 100~$\\mu$s). At optical wavelengths, the observations present a less coherent picture. Sanwal (1999) has reported a time shift of 140~$\\mu$s (optical leading the radio). The uncertainty in this value is 20~$\\mu$s in the determination of the optical peak and 75~$\\mu$s in the radio ephemeris. Shearer et al.\\ (\\cite{shearer:2003}) have reported a lead of 100$\\pm$20~$\\mu$s for simultaneous optical and radio observations of giant radio pulses. Golden et al.\\ (\\cite{golden:2000}) have reported that the optical pulse {\\it trails} the radio pulse by $\\sim$80$\\pm$60~$\\mu$s. Romani et al.\\ (\\cite{romani:2001}) conclude that the radio and optical peaks are coincident to better than 30~$\\mu$s, but their error excludes the uncertainty of the radio ephemeris (150~$\\mu$s). The internal inconsistency of these results -- if we assume that the optical-radio delay is constant -- has prompted us to look into this matter in detail using recent observations in combination with earlier data. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=8cm,angle=0]{5254f2.eps} \\caption[]{Our three results combined with literature results in the optical. For all results, the uncertainties were obtained by linear addition (i.e.\\ not quadratic) of the uncertainty in the peak determination and the uncertainty in the radio ephemeris. Uncertainties on the peak determination alone are also plotted. The open symbol refers to Shearer et al. (\\cite{shearer:2003}), who observed giant radio pulses.} \\label{fig:results} \\end{figure}  We have observed the Crab Pulsar with three different generations of the S-Cam instrument, on two locations, using two GPS units. We consistently find that the optical pulse is leading the radio pulse.  However, the amount by which the optical is leading the radio differs from observation to observation. When comparing the different results as plotted in Fig.\\ \\ref{fig:results} which are obtained with the Jodrell Bank ephemeris, one should take into account that $\\sigma_{\\mathrm{radio}}$ contains a systematical component (of $\\sim$40~$\\mu$s) affecting all these measurements in the same way.If we subtract the 40~$\\mu$s, for the purpose of comparing the results, from $\\sigma_{\\mathrm{radio}}$ and add the remainder, in quadrature, to our measurement uncertainties, we obtain the following values for S-Cam1, S-Cam2, and S-Cam3: 49$\\pm$41, 254$\\pm$120, and 291$\\pm$50~$\\mu$s.    From the last two observations (S-Cam2 and S-Cam3), we then determine an average lead of 273$\\pm$65~$\\mu$s. The uncertainty in this determination is dominated by uncertainties in the radio ephemeris. The result in X-rays for the same two epochs is 370$\\pm$40~$\\mu$s (Rots, private communication).   From our data, we cannot rigorously exclude the possibility that the delay between the radio and optical peak evolves with time. The S-Cam1 result (49$\\pm$90~$\\mu$s; the uncertainty has been determined by a linear addition of the uncertainty in the peak determination and $\\sigma_{\\mathrm{radio}}$ ) in particular is slightly deviant compared to the S-Cam2 and S-Cam3 data points (254$\\pm$170 and 291$\\pm$100~$\\mu$s).  However, the radio observations at the times of the S-Cam1 observations have a high dispersion measure -- a higher value has not been observed after February 1999. As a result, the radio-ephemeris uncertainty (80~$\\mu$s) might have been slightly underestimated: variations in the dispersion measure could introduce a somewhat higher uncertainty than the standard 20~$\\mu$s (Lyne et al.\\ \\cite{lyne:1993}). We note that Rots et al.\\ (\\cite{rots:2004}) have discarded their contemporaneous (X-ray) data point, because the timing ephemeris is somewhat suspect (the GRO version has a high second derivative indicative of a questionable fit). In the following we will therefore only consider the S-Cam2 and S-Cam3 data points. We emphasize here that corrections for arrival times at infinite frequency (which are substantial: $\\sim$0.6~s at 610 MHz) depend on the dispersion measure. The only (literature) data point which is clearly inconsistent with all observations is that of Golden et al.\\ (\\cite{golden:2000}). We have no explanation for this.   Our value of the optical phase difference between the main and secondary pulse of 0.4054$\\pm$0.0004 is not consistent with the X-ray value from Rots et al.\\ (\\cite{rots:2004}) of 0.4001$\\pm$0.0002 periods. This implies that the details of the pulse profile are different in X-rays and at optical wavelengths.  Our time shift of 273$\\pm$65~$\\mu$s is somewhat smaller than, but consistent with, the time shift as obtained from X-ray measurements. A time shift of $\\sim$270~$\\mu$s indicates that possibly (in a simple geometrical model ignoring relativistic effects) the optical radiation is formed $\\sim$90~km higher in the magnetosphere than the radio emission. Alternatively, the difference in phase of $\\sim$0.008 could be interpreted as an angle between the radio and optical beam of $\\sim$3$^{\\circ}$.  Ideally, a simultaneous radio-optical observation at high frequency should be performed. With an observation like this uncertainties resulting from corrections for interstellar scattering are minimized, and the accuracy will effectively be limited by systematic effects ($\\sim$40~$\\mu$s for Jodrell Bank)."
    },
    "0606/astro-ph0606370_arXiv.txt": {
        "abstract": "We discuss new methods of measuring and interpreting the forbidden-to-intercombination line ratios of helium-like triplets in the X-ray spectra of O-type stars, including accounting for the spatial distribution of the X-ray emitting plasma and using the detailed photospheric UV spectrum. Measurements are made for four O stars using archival {\\it Chandra} HETGS data. We assume an X-ray emitting plasma spatially distributed in the wind above some minimum radius \\(R_0\\). We find minimum radii of formation typically in the range of \\(1.25 < R_0 / R_* < 1.67\\), which is consistent with results obtained independently from line profile fits. We find no evidence for anomalously low \\(f/i\\) ratios and we do not require the existence of X-ray emitting plasmas at radii that are too small to generate sufficiently strong shocks. ",
        "introduction": "Since the discovery of X-ray emission from OB stars by {\\it Einstein} \\citep{Hel79, Sel79}, the exact mechanism for X-ray production has been something of a mystery. X-ray emission from OB stars had been predicted by \\citet{CO79}, who proposed that an X-ray emitting corona could explain the observation of superionized \\ion{O}{6} through Auger ionization of \\ion{O}{4}. However, subsequent observations showing less attenuation of soft X-rays than would be expected from a corona lying below a dense stellar wind made a purely coronal origin seem unlikely \\citep{CS83}. \\citet{Mel93} also found that a distributed X-ray source was necessary to explain the observed \\ion{O}{6} UV P Cygni profile in \\zp. Furthermore, with no expectation of a solar-type \\(\\alpha-\\Omega\\) dynamo in OB stars with radiative envelopes, the coronal model fell out of favor. Subsequently, several scenarios in which magnetic field generation and dynamos could exist in OB stars have been proposed \\citep{CM01, MC03, MM05}. Since these models have been proposed, the primary observational evidence invoked by their proponents is anomalously low $f/i$ ratios in the X-ray emission of a few He-like ions in several stars. Re-examining these line ratios and determining whether they require a coronal model to explain them is one of the main goals of this paper.  Shocks arising from instabilities in the star's radiatively driven wind have been considered to provide a more likely origin for the observed X-ray emission, as they are expected to be present, given the line-driven nature of these winds \\citep*{LW80, L82, KR85, OCR88, F95}. However, there have been difficulties in reproducing the observed X-ray properties of O~stars, such as the overall X-ray luminosity and the spectral energy distribution, from stellar wind instability models \\citep{Hel93, F95, FKPPP97, FPP97}. Until recently, the quality of the available spectral data provided little insight into these problems, since the CCD and proportional counter spectra could not resolve individual spectral lines.  Recent high resolution X-ray spectroscopy of OB stars by the XMM-{\\it Newton} Reflection Grating Spectrometer (RGS) \\citep{Kel01, Mel03, RCMMT05} and the {\\it Chandra} High Energy Transmission Grating Spectrometer (HETGS) \\citep{SCHL00, WC01, CMWMC01, MCWMC02, Cel03, KCO03, Gel05, Cel06} have answered some questions while raising new ones. Some stars have X-ray spectra that appear consistent with emission from shocks in the wind, but the detailed comparisons to predicted spectral models are still problematic. Both \\citet{WC01} and \\citet{CMWMC01} have found low forbidden-to-intercombination line ratios in one set of helium-like triplets each in the X-ray spectra of \\zo\\, and \\zp. They infer from this that some of the X-ray emitting plasma is too close to the star to allow shocks of sufficient velocity to develop.  Other stars ($\\theta^1$~Ori~C and $\\tau$~Sco) have X-ray spectra that are unusually hard and have relatively small line widths. While these stars might be considered prime candidates for a coronal model of X-ray emission - especially after having magnetic fields detected via Zeeman splitting \\citep{Del02, Del06} -  their behavior is better understood in terms of the magnetically channeled wind shock model, rather than a model of magnetic heating \\citep{SCHL00, Cel03, SCHT03, Gel05, Del06}.  Finally, we note that for all of the O giants and supergiants observed, the line profiles are less asymmetric than predicted, given the high mass-loss rates measured for these stars using radio free-free emission, H~$\\alpha$ emission, and UV absorption lines \\citep{WC01, Kel01, CMWMC01, MCWMC02, KCO03, Cel06}. This implies either a lower effective opacity to X-rays in their winds (e.g. due to clumping or porosity effects \\citep{FOH03, OFH04, OFH06, OC06}), or lower mass-loss rates \\citep{Cel02, MFSH03, Hel03, BLH05, FMP06}.  One of the key diagnostic measurements available to us in understanding the nature of X-ray emission in OB stars is the forbidden-to-intercombination line ratio in the emission from ions that are isoelectronic with helium. This ratio is sensitive to the UV flux, and thus to the proximity to the stellar surface. This allows us to constrain the location of the X-ray emitting plasma independently of other spectral data, such as emission line profile shapes.  In this paper we discuss methods for using the $f/i$ ratio to constrain the location of X-ray emitting plasma in O star winds. In particular, we explore the effects of a spatially distributed source motivated by the broad line profiles. We discuss the effects of photospheric absorption lines, as well as the $f/i$ ratio expected for a plasma emitted over a range of radii, taking account of detailed line shapes when signal-to-noise allows. We find that accounting in detail for photospheric absorption lines is not important, as long as the X-ray emission originates over a range of radii.  These methods are then applied to He-like triplet emission in a set of archival {\\it Chandra} observations of O stars. Our primary result is that good fits can be acheived for most lines with models having emission distributed over the wind, with minimum radii of about 1.5 stellar radii. We find that none of the data require the X-ray emitting plasma to be formed very close to the photosphere.  This paper is organized as follows: In \\S~\\ref{sec:model} we review the physics of line formation in He-like species (\\S~\\ref{sub:fi}), explore the effects of spectral structure in the photoexciting UV field (\\S~\\ref{sub:photo}), and of spatial distribution of the X-ray emitting plasma (\\S~\\ref{sub:Rbar}), while incorporating the line-ratio modeling into a self-consistent line-profile model (\\S~\\ref{sub:hewind}). In \\S~\\ref{sec:data} we discuss the reduction and analysis of archival O star X-ray spectra. In \\S~\\ref{sec:results} we give the results of this analysis, fitting high signal-to-noise complexes with the self-consistent line-profile model described in \\S~\\ref{sub:hewind} and fitting the lower signal-to-noise complexes with multiple Gaussians and interpreting these results according to the spatially distributed picture described in \\S~\\ref{sub:Rbar}. In \\S~\\ref{sec:discussion} we discuss the implications of these results, and in \\S~\\ref{sec:conclusions} we give our conclusions. ",
        "conclusions": "} We have investigated the effect of a radially distributed plasma on the forbidden-to-intercombination line ratio in helium-like triplets, as well as variations in the exciting photospheric flux as a function of Doppler shift throughout the wind. We find that the fact that the plasma is likely distributed over a range of radii and Doppler shifts allows us to use an averaged value of the photospheric continuum instead of accounting for it in detail. We also find that the value of \\(R_0\\) derived assuming a distribution of radii is substantially smaller than the value of $R$ derived assuming a single radius.  We have used the $f/i$ ratio of helium-like triplets to constrain the radial distribution of X-ray emitting plasma in four O-type stars. We find that the minimum radius of emission is typically \\(0.6 < u_0 < 0.8\\), or \\(1.25 < R_0 / R_* < 1.67\\) with the emission extending beyond this initial radius with either a constant filling factor or one that increases slightly with radius. This is consistent with the results of line profile fits using the model of \\citet{OC01} \\citep{KCO03, Cel06}. However, some of the minimum radii of formation are well inside the radius of optical depth unity calculated using the mass-loss rates in the literature, implying that either the effective opacities are lower \\citep[e.g. due to porosity effects][]{FOH03, OFH04, OFH06, OC06} or the mass-loss rates are lower than the literature values \\citep{MFSH03, Hel03, BLH05, FMP06} or both. We also measure low values of the characteristic optical depth \\(\\tau_*\\) compared to what one would expect based on the literature mass-loss rates, which is consistent with the same conclusions.  We find that there is no evidence for anomalously low $f/i$ ratios in high-$Z$ species. Our measurements do not require X-ray emission orignating from too close to the star to have sufficiently strong shocks, nor do we need to posit the existence of a magnetically confined corona. This conclusion is based partly on different measured values of $f/i$ ratios and partly on higher photospheric UV fluxes on the blue side of the Lyman edge in the more recent TLUSTY model spectra.  We have fit He-like emission line complexes with profile models that simultaneously account for profile shapes and line ratios. These models constrain the radial distribution of plasma both through the line ratio and the profile parameters \\(u_0\\) and $q$. We find that they are capable of producing good fits to the data, showing that the information contained in the line ratios and profile shapes are mutually consistent."
    },
    "0606/astro-ph0606693_arXiv.txt": {
        "abstract": "\\emph{Accepted for publication in the Astrophysical Journal\\\\} We estimate the intrinsic neutral gas density in Damped Lyman $\\alpha$ systems ($\\Omega_{HI}^{(DLA)}$) in the redshift range $ 2.2 \\lesssim z \\lesssim 5$ from the DLA SDSS DR\\_3 sample of optically selected quasars.  We take into account self-consistently the obscuration on background quasars due to the dust present in Damped Lyman $\\alpha$ systems. We model the column density and redshift distribution of these systems by using both a non-parametric and a parametric approach.  Under conservative assumptions on the dust content of Damped Lyman $\\alpha$ systems, we show that selection effects lead to underestimating the \\emph{intrinsic} neutral gas density by at least $15\\%$ with respect to the \\emph{observed} neutral gas density. Over the redshift range $[2.2;5.5]$ we find $\\Omega_{HI}^{(DLA)}=0.97^{+0.08+0.28}_{-0.06-0.15} \\cdot 10^{-3}$, where the first set of error bars gives the $1\\sigma$ random errors and the second set gives the modeling uncertainty dependent on the fraction of metals in dust - from 0\\% to 50\\%. This value compares with $\\Omega_{HI}^{(DLA)}=0.82^{+0.05}_{-0.05}$ ($1\\sigma$ error bars), which is obtained when no correction for dust is introduced. In the model with half of the metals mass in dust we cannot constraint $\\Omega_{HI}^{(DLA)}$ at a confidence level higher than $90\\%$. In this case there is indeed a probability of about $10\\%$ that the intrinsic column density distribution of DLA systems is a power law $f(N_{HI}) \\propto 1/N_{HI}^{~1.95}$. In contrast, with $25 \\%$ of the metals in dust - the most realistic estimate - a power law is ruled out at $99.5\\%$ of confidence level. ",
        "introduction": "Damped Lyman $\\alpha$ systems (hereafter DLA systems) are quasar absorption systems with a column density above $2\\cdot 10^{20} cm^{-2}$ and represent the high end of the distribution of absorption systems starting from the Lyman $\\alpha$ forest at $N_{HI} \\gtrsim 10^{14} cm^{-2}$. DLA systems represent the most significant reservoir of neutral hydrogen in the universe available for star formation. These systems are considered to be either (cold) massive rotating disks, the progenitors of todays disk galaxies \\citep{pro97,wol05} or compact protogalactic clumps \\citep{hae98,nag04}.  In the era of precision cosmology, an accurate measure of the total mass density of neutral gas as a function of the redshift represents an important constraint for galaxy formation models. Ground observations are able to identify DLA absorption features in the spectra of quasars from $z_{abs} \\gtrsim 1.8$, where the absorption lines enter the atmospheric window, up to $z_{abs} \\approx 5.5$. With an all sky survey like the Sloan Digital Sky Survey, spectra of several thousands of quasars with enough resolution for DLA detection have been acquired and the \\emph{observed} density of neutral gas in DLA systems is now measured with errors below $10\\%$ \\citep{pro05}.  This measurement must be interpreted with some caution, as the presence of DLA systems along a line of sight leads to a potential obscuration due to the dust that they host: the \\emph{observed} gas density is a biased estimator of the \\emph{intrinsic} density unless the dust effects are accurately quantified. Several papers, starting from \\citet{ost84} have attempted to model the influence of dust along the line of sight, often with conflicting results.  A detailed analysis framework for the obscuration of quasars has been developed by \\citet{fal93} (see also \\citealt{fal89,fal89b,pei91,pei95b}) and applied to the quasars sample of \\citet{lan91}. Their study highlighted a potentially severe effect of the dust bias that did not allow to put an upper limit to the intrinsic density of DLA systems. In fact, absorbers with high column densities and/or with high dust-to-gas ratio represent essentially ``bricks'' along the line of sight to a quasar and are very likely to be missed in optically selected surveys. An additional evidence for the dust obscuration came in the form of a detected preferential reddening in the spectra of quasars with DLA absorption with respect to a control sample without detection of these systems \\citep{pei91}.  More recent investigations revised these earlier results on the dust content in DLA systems and on their reddening of background objects \\citep{mur04,ell05}, finding in particular no robust evidence for the reddening of quasars at $z \\approx 3$ with DLA features in their spectra: at $3\\sigma$ \\citet{mur04} find $E(B-V) < 0.02mag$, while \\citet{ell05} have $E(B-V)<0.04mag$ . At the same time radio selected quasars surveys \\citep{ell01}, with complete optical follow-up detection have provided the first bias free constraints on the intrinsic distribution of DLA systems.  Taking advantage of these recent measurements for the number density of DLA systems, we have previously characterized \\citep{tre06} the dust absorption along random lines of sight by means of a Monte Carlo code, finding that, on average, the deviations from unit transmission are effectively modest ($\\langle \\exp{(-\\tau)} \\rangle \\gtrsim 0.9$ at an emitted wavelength $\\lambda_e = 0.14 \\mu m$ over all the redshift range) and of limited impact on most observations. However, this result does not exclude the presence of a small fraction of lines of sight (of the order of a few percent) through the most massive and/or the most metal-rich DLA systems and characterized by a large optical depth. Indeed, \\citet{wil05} and \\citet{wil06} find a significant evidence of reddening in DLA systems with CaII absorption lines at moderate redshift ($z_{abs} \\approx 1$; $\\langle E(B-V) \\rangle \\gtrsim 0.1$). Similarly \\citet{yor06} measure $E(B-V)$ up to 0.085 for MgII selected DLA systems at $z_{abs}\\approx 1.0$. As the determination of the gas density in DLA systems is dominated by these most massive absorbers, the potential bias in this measure, correctly stressed by \\citet{fal93}, must not be dismissed by the recent evidence of a very modest \\emph{average} deviation from unity transmission.  In this paper we take advantage of the large sample of DLA systems identified in Sloan quasars \\citep{pro05} and we investigate the relation between the observed and intrinsic density of neutral gas in these systems. The Sloan sample that we consider has three main advantages over the sample used by \\citet{fal93}. (1) It is larger by a factor 30. (2) The quasars have been selected within a luminosity limit in the I band, significantly less sensitive to dust obscuration than the B band used for the older sample. (3) The color selection algorithm has a good sensitivity to red quasars \\citep{ric02}. In fact even if a quasar with a DLA absorption system is not dropped below the I-band flux limit, its color can be changed so that it ends up lying outside the color selection box used to identify quasar candidates for follow-up spectroscopy. Thanks to the precision of the Sloan photometry, that allows a clear separation of the stellar locus in the color space, and to the use of an extended color selection box, the loss of completeness due to this latter effect is only marginal (see \\citealt{ric02} for a detailed discussion of the Sloan color selection algorithm and for completeness tests) and is therefore not considered in this work.  This paper is organized as follows. In Sec.~\\ref{sec:dust} we characterize the obscuration bias in a magnitude limited survey, in Sec.~\\ref{sec:data} we describe the dataset that we are using. In Sec.~\\ref{sec:para_OM} we present our analysis for the parametric estimation of the intrinsic comoving density of neutral gas, whose uncertainties are quantified in Secs.~\\ref{sec:boot}-\\ref{sec:scatter} by means of Monte Carlo simulations of synthetic observations. In Sec.~\\ref{sec:disc} we discuss the accuracy of a non-parametric estimator for the neutral gas density. We summarize our findings in Sec.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper we have improved the analysis of the column density distribution of DLA systems in the SDSS DR\\_3 DLA survey\\citep{pro05} to take into account the bias due to dust obscuration along the line of sight. A first modeling of the bias was constructed by \\citet{ost84} and improved by \\citet{fal93}. These earlier estimates expected a severe effect of obscuration, with the observed gas density of DLA systems being up to several times smaller than the intrinsic one. The best estimate for $\\Omega_{HI}^{(DLA)}$ given by \\citet{fal93} is $4.9 \\cdot 10^{-3}$ (obtained using a cosmology with $\\Omega_M=1$ and $H_0=70 km/s/Mpc$) and their upper limit $\\Omega_{HI}^{(DLA)} \\leq 3 \\cdot 10^{-2}$, obtained from a standard big bang nucleosynthesis abundance model. More recent works \\citep{mur04,ell05} tend to dismiss the issue of obscuration bias based on the absence of systematic reddening in the spectra of quasars with DLA absorption features.  Here we show that the effect of obscuration, while not being as severe as predicted by \\citet{fal93}, does indeed play an important effect on the precise measurement of $\\Omega_{HI}^{(DLA)}$. In the era of precision cosmology, where the \\emph{observed} density ${\\Omega_{HI}^{(DLA)}}$ is constrained with errors below $10~\\%$, the systematic effects are not to be underestimated. With the typical amount of metals present in DLA systems the \\emph{observed} density ${\\Omega_{HI}^{(DLA)}}$ derived from shallow magnitude limited surveys of quasars underestimates the \\emph{intrinsic} density ${\\Omega_{HI}^{(DLA)}}$ by about $15~\\%$ assuming dust-poor DLA systems (i.e. systems with a fraction of metals in dust of 25\\%). Our best estimation for $z \\in [2.2;5.5]$ gives an intrinsic neutral gas density $\\Omega_{HI}^{(DLA)}=0.97^{+0.08}_{-0.06} \\cdot 10^{-3}$ ($1\\sigma$ error bars) to be compared with the observed gas density $\\Omega_{HI}^{(DLA)}=0.817^{+0.050}_{-0.052} \\cdot 10^{-3}$ derived by \\citet{pro05}. If we leave the dust to metal ratio parameter $\\alpha_{\\kappa}$ free to vary over the relevant range (from 0 to 1), we find $\\Omega_{HI}^{(DLA)}=0.97^{+0.08+0.28}_{-0.06-0.15} \\cdot 10^{-3}$, where the first set of error bars gives the $1\\sigma$ random errors and the second set gives the modeling uncertainty dependent on the fraction of metals in dust.  The obscuration bias therefore represents the main source of uncertainties for the determination of the intrinsic neutral gas content in DLA systems.  Our analysis has been carried out assuming that all the DLA absorbers at a given redshift have the same dust-to-gas ratio. By means of monte carlo simulations of synthetic observations we shoow that this assumption does not introduce a significant bias as long as the dust-to-gas ratio is not correlated with the hydrogen column density. A correlation of the form $k_i \\propto N_{HI}^{\\gamma}$ introduces a systematic error of the order $+14\\%$ in $\\Omega_{HI}^{(DLA)}$ for $\\gamma = - 0.8$ and of $-3 \\%$ for $\\gamma = +0.8$ (the bias is reduced to $+8\\%$ and $0\\%$ for $\\gamma = \\mp 0.4$ respectively).  Caution is also needed when taking the result from the maximum likelihood at face value. In our reference scenario we find that a power law function for column density distribution of DLA systems can be ruled out at a confidence level no greater than 99.5\\%. If the dust content in DLA systems is higher, the dust bias becomes more significant. For $50\\%$ of the metals in dust grains the present data do not allow to put an upper limit to the neutral gas content with confidence level greater than $90\\%$. We show in fact that there is a probability of about 10\\% that the SDSS DLA DR\\_3 data are consistent with an intrinsic power law distribution in column density adopting $\\alpha_{\\kappa}=1$. The slope of this power law is $m_2=-1.95$ and an upper cut-off cannot be derived from the data. In optically selected surveys, absorbers with high column densities of neutral gas and with high metallicity are missed, letting systematic uncertainties go easily out of control.  Radio selected quasar samples would represent an elegant, bias free solution to the measurements of $\\Omega_{HI}^{(DLA)}$. Unfortunately, with the present extension, these surveys constraint $\\Omega_{HI}^{(DLA)}$ with a random uncertainty of the order of $40~\\%$ (plus potential systematic effects due to the limited number of observed DLA systems). To reduce the errors within the $10~\\%$ level without the need to assume a modeling for the dust bias it is therefore necessary to increase the path-length of radio surveys by a factor $10$ at least. The UCSD radio survey \\citep{jor06} has recently made promising progresses in this direction."
    },
    "0606/astro-ph0606416_arXiv.txt": {
        "abstract": "We adapt the Jain--Taylor (2003) shear-ratio geometric lensing method to measure the dark energy equation of state, $w=p_v/\\rho_v$ and its time derivative from dark matter haloes in cosmologies with arbitrary spatial curvature. The full shear-ratio covariance matrix is calculated for lensed sources, including the intervening large-scale structure and photometric redshift errors as additional sources of noise, and a maximum likelihood method for applying the test is presented. Decomposing the lensing matter distribution into dark matter haloes we calculate the parameter covariance matrix for an arbitrary experiment. Combining with the expected results from the CMB we design an optimal survey for probing dark energy. This shows that a targeted survey imaging $60$ of the largest clusters in a hemisphere with 5-band optical photometric redshifts to a median galaxy depth of $z_m=0.9$ could measure $w_0\\equiv w(z=0)$ to a marginal $1$-$\\sigma$ error of $\\Delta w_0=0.5$. We marginalize over all other parameters including $w_a$, where the equation of state is parameterized in terms of scale factor $a$ as $w(a)=w_0+w_a(1-a)$. For higher accuracy a large-scale photometric redshift survey is required, where the largest gain in signal arises from the numerous $\\approx 10^{14}\\rm M_\\odot$ haloes corresponding to medium-sized galaxy clusters. Combined with the expected Planck Surveyor results, such a near-future 5-band survey covering 10,000 square degrees to $z_m=0.7$ could measure $w_0$ to $\\Delta w_0=0.075$ and $\\Delta w_a=0.33$. A stronger combined constraint is put on $w$ measured at the pivot redshift $z_p=0.27$ of $\\Delta w(z_p)=0.0298$. We compare and combine the geometric test with the cosmological and dark energy parameters measured from planned Baryon Acoustic Oscillation (BAO) and supernova Type Ia experiments, and find that the geometric test results combine with a significant reduction in errors due to different degeneracies. A combination of geometric lensing, CMB and BAO experiments could achieve $\\Delta w_0=0.047$ and $\\Delta w_a=0.111$ with a pivot redshift constraint of $\\Delta w(z_p)=0.020$ at $z_p=0.62$. Simple relations are presented that show how our lensing results can be scaled to other telescope classes and survey parameters. ",
        "introduction": "Over the last decade, gravitational lensing has emerged as the simplest and most direct way to probe the distribution of matter in the Universe (Bartelmann \\& Schneider 2001, Refregier 2003). More recently it has become apparent that it can also be used as a probe of the mysterious, negative-pressure ``dark energy'' component of the Universe which gives rise to the observed acceleration of the expansion of the Universe (Hu \\& Tegmark,1999; Huterer, 2002; Jain \\& Taylor, 2003; Hu, 2003; Takada \\& Jain, 2003; Song \\& Knox, 2004; Ishak, 2005; Ma, Hu \\& Huterer, 2006; Heavens et al., 2006)  The dark energy exerts its influence by its effect on the expansion history of the Universe. If the current expansion of the Universe is accelerating, the Universe must be older than if it was decelerating, since the expansion was slower in the past. This changes the distance traveled by a photon, $r(z)$, for a given expansion factor of the Universe, $a(z)=1/(1+z)$, as photons have had more time to travel further than in the decelerating case. The accelerated expansion will also slow the rate of growth of matter perturbations. The simplest phenomenological model of the dark energy can be constructed by simply parameterizing the equation of state of the vacuum, \\be p_v= w\\rho_v, \\ee where $p_v$ is the dark-energy/vacuum-pressure and $\\rho_v$  is its energy-density, and $w=w(a)$ may vary with scale factor.  Gravitational lensing depends upon both the geometry of the Universe, via the observer-lens-source distances, and on the growth of structure which will lens distant galaxies, and so lensing probes both effects. Gravitational lensing is an integral effect and so for a given line of sight these effects are degenerate with each other and other parameters. In order to disentangle the effects of the dark energy we require redshift information for the source images. It has already been shown that such information can be used to reconstruct the 3-D distribution of dark matter (Taylor, 2001; Taylor et al., 2004). For large-scale imaging surveys, the most practical way to get redshifts for each image is from multi-band photometric redshift surveys. The COMBO-17 imaging and photometric survey (Wolf et al., 2003) has already shown the power of combining lensing with photometric redshifts (Brown et al., 2003; Taylor et al., 2004; Gray et al., 2004; Bacon et al., 2005; Semboloni et al., 2006).  The parameters of the dark energy can be extracted from weak gravitational shear measurements by taking correlations of galaxy ellipticities at different redshifts (e.g. Bacon et al., 2005; Hu, 2003; Heavens, 2003; Heavens et al., 2006; Semboloni, 2006), where the expansion history enters both the lens geometry and the dark matter evolution rate. Jain \\& Taylor (2003) proposed an alternative approach, taking the ratio of the galaxy-shear correlation functions at different redshifts. In this case the mass of the lens dropped out leaving behind a purely geometric quantity useful for measuring cosmological parameters. This had the advantages of allowing the analysis to extend into the nonlinear clustering r\\'egimes where modeling the nonlinear matter power spectrum can be inaccurate, and where the shear signal will also be stronger, i.e. in the vicinity of galaxy clusters. In addition, as this relies upon the correlation between galaxies and shear many systematic effects will be averaged over, as in galaxy-galaxy lensing. Following this a number of papers have suggested variations on this theme (Bernstein \\& Jain, 2003; Hu \\& Jain, 2004; Zhang, Hui \\& Stebbins, 2005).   Geometric tests of dark energy not only complement other methods based on the clustering of matter, but directly probe the global evolution of the Universe via the redshift-distance relation, $r(z)$. Other methods measure the combined effect of the growth rate of perturbations and the global geometry. Comparison of the two can be used to test the Einstein-Hilbert action, and extensions and modifications of General Relativity such as extra dimensions.  While the main focus of the Jain-Taylor (2003) paper was a statistic given by the ratio of galaxy-shear correlations (or equivalently power spectra), they illustrated their method with the analysis of a single cluster. In this paper we develop this idea further and focus on applying the geometric test behind individual galaxy clusters. The main difference between this and the original Jain-Taylor approach is that we do not need to first generate galaxy-shear cross-correlation functions, or cross-power spectra, which require large data-sets. Rather the ratios used are just of the shears behind a given cluster at fixed redshifts. This allows the test to be applied to noisy data, since we do not need to estimate correlation functions before applying the ratio test. This is similar to the approach of Bernstein \\& Jain (2004), who considered a ``template matching'' approach, cross-correlating a foreground galaxy template with the background shear pattern. Our approach is different in that we use the galaxies to identify the positions of lensing haloes, and then take shear ratios. In doing so we focus on the dark matter haloes generating the signal, allowing a halo decomposition of the matter distribution, and ask how to maximize the signal. The price we pay for this approach is that we become susceptible to a sampling variance due to lensing by other large-scale structure along the line of sight, which we can beat down using multiple lines of sight. In addition we generalize our methods to non-flat cosmological models.   Zhang, Hui \\& Stebbins (2005) have proposed a different geometric method, which allows them to extend the correlation/power spectrum method to galaxy-galaxy and shear-shear correlations as well as galaxy-shear cross-correlations. They also point out some inaccuracies with the analysis of Jain \\& Taylor (2003) and Bernstein \\& Jain (2003), which we correct here.   In the next Section we lay out the basic lensing equations we will need. In Section \\ref{ML} we derive the statistical properties of the shear ratios, and write down a likelihood function for measuring the dependent cosmological parameters, we then estimate the Fisher matrix and parameter covariance matrix for the dark energy. In Section 4 we outline the survey design formalism, using the dark matter halo model for the distribution of galaxy clusters and group haloes, outlining a realistic photometric redshift analysis and discuss bias and intrinsic ellipticity issues. In Section 5 we discuss survey strategies, considering targeted, wide-field and area limited designs. Using the parameter covariance matrix and a model of photometric redshifts we optimize a weak lensing photometric redshift survey for measuring dark energy parameters from cluster lensing in Section 6. We forecast the expected accuracy of cosmological parameters in Section 7 and compare and combine with other methods. In Section 8 we discuss the required control of systematic effects, and we present our conclusions in Section 9. We begin by introducing the necessary cosmological and weak lensing concepts. ",
        "conclusions": "In this paper we have set out a new method for the analysis of the geometric shear ratio test for measuring the dark energy equation of state, based on the measurement of shear ratios around individual galaxy groups and clusters. The shear ratio test is insensitive to the growth of structure, but sensitive to the geometry of the Universe, via the matter and dark energy density and the dark energy equation of state. This approach allows one to apply the method to individual objects, rather than requiring the measurement of some other statistic such as the galaxy-shear cross-correlation function which may be noisy for small data-sets. The down-side is that the method is now contaminated by structure along the line of sight, which can be overcome by using many independent lines of sight.  Of the parameters which govern the geometry of the Universe, or more properly the photon distance-redshift relation, the shear ratio is most sensitive to a constant dark energy equation of state, $w_0$, and very insensitive to evolution, parameterized here by $w_a$. This can be understood as due to the shear ratios being sensitive only to the change in shape of the shear signal as a function of redshift. As $w_a$ parameterizes the high-redshift effect of the dark energy equation of state, its effects are ``renormalized'' away. This behavior is very different to other probes of dark energy, and so helps to break parameter degeneracies when combined with other probes.  It must be emphasised that the Fisher matrix framework used in this paper may result in overly optimistic constraints. Since the errors are calculated by expanding about a fiducial point in parameter space any higher order effects that may change the shape of the likelihood surface cannot be taken into account. The effect of varying the fiducial dark energy model, in Section \\ref{The effect of changing the fiducial dark energy model}, demonstrates that the errors are sensitive to the choice of the fiducial model. A concrete example of higher order likelihood effects can be seen in a 3D cosmic shear analysis by comparing Fisher matrix calculations of the ($\\sigma_8$, $\\Omega_m$) plane (for example in Heavens et al., 2006) (predicting an ellipse) with the measured constraints from data (for example Kitching et al., 2006) which measure an extended curved constraint. These effects can be investigated by large simulations or by a Monte-Carlo type exploration of the likelihood surface, we leave such investigations for future work.  To account for many of the sources of uncertainty in the method, we have developed a halo decomposition analysis of the lensing dark matter distribution to model the signal from dark matter haloes over a range of mass scales and redshifts. We have also included the effects of shot-noise due to the random intrinsic orientation of each galaxy, photometric redshift errors and the contribution of large-scale structure lensing to the error budget. We have also investigated in detail a model for the photometric redshift error, based on studies of the COMBO-17 data-set, as a function of redshift, number of imaging bands and limiting magnitude. The effect of a bias in the calibration and distribution of photometric redshifts with spectroscopic redshifts is also studied, and we find that we require some $10^4$ galaxies with spectroscopic redshifts to control calibration issues. The limitations of observing the shear signal from the ground and space are also discussed, and we argue that without adaptive optics ground-based lensing studies are seeing limited, suggesting that it will be difficult to use galaxies beyond $z=1.5$.  The halo decomposition analysis of the dark matter lenses has allowed us to probe the origin of the shear signal in different types of survey, taking a 4-metre telescope with a 2 square degree field of view as our default survey. These results can be scaled to any other telescope parameters.  For targeted observations, where the time-limitation translates into the number of clusters and groups one can observe to a given depth, we have shown that we only require around 60 of the largest clusters in a celestial hemisphere to constrain $w_0$ to around $\\Delta w_0\\sim 0.50$, marginalizing over all other parameters, including $w_a$, a factor of $3$ improvement on 4-year WMAP given a marginalization over $w_a$. To achieve a higher accuracy requires the imaging of an unfeasible number of haloes, and instead one should turn to a wide-field imaging and photometric redshift survey. We find for a 4-meter class telescope with a 2 degree field of view that with a 10,000 square degree, 5-band photometric redshift survey with median redshift $z_m =0.7$ ($r=23.8$), we can expect to reach an accuracy of $\\Delta w_0\\sim 0.07$, again marginalizing over all other parameters including $w_a$. Our results can be easily rescaled to other telescope types, and survey strategies.  The halo decomposition allows us to deduce where the main signal comes from in both the targeted and surveying modes. In both cases a significant fraction of the signal comes from the largest hundred clusters in each survey, reaching a sensitivity of $\\Delta w_0 \\sim 0.5$, however the majority of the signal comes from the numerous ($\\sim 10^{5-6}$) $M>10^{14}M_\\odot$ haloes which can push the accuracy up to $\\Delta w_0\\sim 0.07$.  Having determined where the majority of the dark energy signal will come from in a geometric shear ratio test, we then investigate the optimization of such a survey, when combined with the expected results from the Planck Surveyor experiment. We find that for our fiducial telescope for a fixed-time survey, going shallower ($z_m<0.7$) over a wider area decreases the accuracy due to the drop in the number of available background sources and corresponding increase in shot-noise. Going deeper ($z_m>0.7$) over a smaller area increases the clustering noise, since we now have fewer clusters to average over.  We have also studied the effect of varying the number of imaging bands to increase or decrease the photometric accuracy. We find that when combined with Planck an increase from 5, 9 or 17 optical bands makes little difference to the optimal survey. The reason for this insensitivity to higher accuracy photometric redshifts is due to the integral nature of the lensing effect, and the weak effect when combined with another data-set. However decreasing the number of bands is expected to have a strong effect on the accuracy of the lensing survey as redshift information is lost. We discuss how our results can be scaled to other telescope classes and survey parameters.  The dark energy parameters $w_0$ and $w_a$ can be combined to give an uncertainty on $w(z)=w_0+w_a z/(1+z)$, at some optimal redshift. This combination helps distinguish where the survey is most sensitive to the dark energy equation of state. In the case of our optimal lensing survey this is at $z=0.27$ with $\\Delta w(z=0.27)=0.0298$. Again, the reason for the low-redshift sensitivity to $w(z)$ is due to the insensitivity of the shear ratio test to $w_a$.  Having optimized the lensing survey for the geometric test in combination with the expected results from the CMB, we have investigated the effect on the full set of cosmological parameters for the CMB and lensing. The geometric test constrains a narrow sheet in the $(\\Omega_m,\\Omega_v,w_0,w_a)$ parameter-space, which is nicely orthogonal to the CMB parameter constraints. Here we show that the CMB mainly constrains the curvature of the model, while the geometric test constrains $w_0$, and the combination constrain $w_a$.  We have also compared and combined the geometric shear ratio test with the expected results from an Baryon Acoustic Oscillation (BAO) experiment, such as proposed for WFMOS, and a supernova Type Ia survey, such as that proposed for SNAP. Here we have put all of the surveys (lensing, CMB, BAO and SNIa) on an equal footing, using the same curved background cosmology and the same dark energy model parameterization. We find that the degeneracies in the geometric test, in particular the insensitivity to $w_a$, are nicely orthogonal to all these other probes. Combining the geometric test with the CMB, BAO or SNIa will yield accuracies of a $\\Delta w_0\\approx 0.10$ and $\\Delta w_a\\approx 0.5$, and can be compared for systematics. An optimal combination is a geometric lensing test, with the Planck CMB and WFMOS BAO experiment, yielding an expected accuracy of $\\Delta w_0=0.047$ and $\\Delta w_a=0.11$.  Finally we discuss some of the potential systematic effects which could affect the predicted accuracy of lensing.  In summary, the prospects of accurately measuring the dark energy equation of state and its evolution to high accuracy over the next decade are very good. The key to this is the gravitational lensing geometric shear ratio test, which, due to its orthogonal degeneracies, can be optimally combined with a large range of other dark energy probes, such as the CMB, BAO or SNIa. In addition, gravitational lensing can also be a probe via two-point analysis, either from correlation functions or power spectra in redshift-space (Heavens et al. 2006 and Castro et al. 2005). Just as with lensing of the CMB, since the shear ratio analysis does not contain any information on structure, we can expect there to be little correlation between the two methods, even for the same survey. However, the shear ratio covariance may be correlated with the shear power. We shall explore combining these methods elsewhere."
    },
    "0606/astro-ph0606620_arXiv.txt": {
        "abstract": "In this work we connect some measurable properties of the CIV$\\lambda$1549 emission line with the quasar accretion rates (Eddington ratios). A tight correlation is found for a sample of more than a hundred nearby objects, suggesting a possible method for a relatively accurate estimate of the Eddington ratio of high-redshift quasars, at least for the radio-quiet ones. This paper further confirms the existing notion that the CIV changes (shifts) are mostly driven by the accretion rate. ",
        "introduction": "\\begin{figure} \\centering{\\epsfig{file=fig1a.eps,width=80mm}} \\caption{The relation between log($L/L_{\\rm edd}$) and log($P_{\\rm CIV}$), where $P_{\\rm CIV}$ is the EW-normalized ratio of the blueward and redward profile widths at half maximum; see the text. 124 objects are included, 75 of which are radio-quiet. The radio-quiet objects are shown as filled rhombs, the radio-loud ones -- as open squares. Note that RQ objects show a much better correlation, as well as that only they show apparent super Eddington accretion rates.} \\end{figure}   Exploring different quantities, we found a parameter ($P_{\\rm CIV}$), related to the EW-normalized shift of CIV, which correlates very well with the Eddington ratio ($L/L_{\\rm edd}$). This parameter is the ratio of the blueward and redward (measured in respect to the systemic velocity) widths of the profile at half maximum, normalized by the EW (in Angstroms) of the CIV line, i.e. $P_{\\rm CIV}=(FW_{\\rm blue}/FW_{\\rm red})/EW$. This parameter should not be difficult to measure, provided a spectrum of good quality is available and the precise redshift of the object is known. We found a very good correlation between $L/L_{\\rm edd}$ and $P_{\\rm CIV}$ (in logarithmic units), with a Pearson correlation coefficient of 0.64 (Fig. 1). The correlation even improves further if only radio-quiet objects are considered, 0.73, and slightly more if $P_{\\rm CIV}$ is measured at zero-intensity instead of half-maximum level. Thus, one can use $P_{\\rm CIV}$ to estimate $L/L_{\\rm edd}$, following the derived empirical relation (bisector linear fit):   $${\\rm log}(L/L_{\\rm edd}) = 0.64~{\\rm log}(P_{\\rm CIV}) - 0.47$$   with a typical error of 0.5 dex for $L/L_{\\rm edd}$ (Fig. 1). One is to note, however, that $P_{\\rm CIV}$ applies to the broad component of CIV (after the subtraction of a NLR contribution on the top, if such is present, Bachev et al. 2004) and also that this parameter is calculated upon the redshift based on [OIII]$\\lambda$5007 or the narrow top of H$\\beta$ lines, not the top of CIV (often blueshifted). It should also be mentioned, that this relation is found mostly for powerful quasars and may not hold for the lower-luminosity objects (i.e. for $L/L_{\\rm edd} < 0.01$; Fig. 1). ",
        "conclusions": "The exact nature of correlations like the one we demonstrated above is not known, but one may suspect it has something to do with the accretion disk winds. As it is often assumed, CIV has a wind origin (Murray et al. 1995; Proga \\& Kallman  2004) and therefore it is not surprising to find a connection between the line profile and the Eddington ratio, taking into account that the later is often thought to drive the wind strength (details in Bachev et al. 2004, and the references within). Why exactly $P_{\\rm CIV}$, as defined, appears however to be the best surrogate for the Eddington ratio, should be a matter of future modeling. We also note that the \"$P_{\\rm CIV}$ -- $L/L_{\\rm edd}$\" relation could be of importance not only for modeling CIV emission region, but also for explaining the radio-quiet/radio-loud differences (e.g. Fig. 1). From a practical point of view, one can use such empirical relations for roughly estimating the Eddington ratio at high redshifts, where no other estimators are available. Thus, the early accretion history of the quasars can be traced, provided CIV properties do not change significantly with cosmic time."
    },
    "0606/astro-ph0606550_arXiv.txt": {
        "abstract": "{}{We present a new continuum 3D radiative transfer code, MCFOST, based on a Monte-Carlo method.  MCFOST can be used to calculate (i) monochromatic images in scattered light and/or thermal emission, (ii) polarisation maps, (iii) interferometric visibilities, (iv) spectral energy distributions and (v) dust temperature distributions of protoplanetary disks.}{Several improvements to the standard Monte Carlo method are implemented in MCFOST to increase efficiency and reduce convergence time, including wavelength distribution adjustments, mean intensity calculations and an adaptive sampling of the radiation field. The reliability and efficiency of the code are tested against a previously defined benchmark, using a 2D disk configuration.  No significant difference (no more than 10\\%, and generally much less) is found between the temperatures and SEDs calculated by MCFOST and by other codes included in the benchmark.}{  A study of the lowest disk mass detectable by \\textit{Spitzer}, around young stars, is presented and the colours of ``representative'' parametric disks are compared to recent IRAC and MIPS \\textit{Spitzer} colours of solar-like young stars located in nearby star forming regions.}{} ",
        "introduction": "Signatures for the presence of dust are found nearly everywhere in astrophysics. In the context of star and planet formation, dust is abundant in molecular clouds and in the circumstellar environments of a large fraction of stellar objects in the early stages of their evolution. In the circumstellar disks encircling these young stars, where planets are thought to form, the interplay between dust and gas is of paramount importance.  At short wavelengths, dust grains efficiently absorb, scatter, and polarise the starlight. How much radiation is scattered and absorbed is a function of both the geometry of the disk and the properties of the dust. In turns, the amount of absorbed radiation sets the temperature of the dust (and gas) and defines the amount of radiation that is re-emitted at longer, thermal, wavelengths.  The last decade has witnessed the improvement of imaging capabilities with the advent of potent instruments in the optical and near-infrared, and large millimeter interferometers, providing detailed views of the disks around young stars. The sensitivity and wavelength range covered by new instruments is increasing steadily and the mid- and far-infrared ranges are now being explored efficiently by the {\\sl Spitzer Space Telescope}, and new facilities like Herschel and ALMA will soon complete the coverage.  With this unprecedented wealth of data, from optical to radio, fine studies of the dust content and evolution of disks become possible and powerful radiative transfer (RT) codes are needed to fully exploit the data.  In this paper we describe such a code, MCFOST.  That code was used extensively to produce synthetic images of the scattered light from disks around young stars. Examples include the circumbinary ring of GG~Tau \\citep{McCabe02, Duchene04}, the large silhouette disk associated with IRAS 04158+2805 \\citep{Menard05}, and an analysis of the circular polarisation in GSS~30 \\citep{Chrysostomo97}.  In \\S2 below, we briefly describe MCFOST. In section \\S3, tests and validation of MCFOST are presented. Two examples of applications are presented in \\S4. Firstly, a study is presented of the minimum mass of disks detectable by {\\sl Spitzer} around young solar-like stars, young low-mass stars, and young brown dwarfs. Secondly, the colours of parametric disks are compared to the {\\sl Spitzer} colours, for both IRAC and MIPS, of samples of T Tauri stars located in nearby star forming regions. ",
        "conclusions": "We have presented a new continuum 3D radiative transfer code, MCFOST. The efficiency and reliability of MCFOST was tested considering the benchmark configuration defined by P04. MCFOST was shown to calculate temperature distributions and SEDs that are in excellent agreement with previous results from other codes.  Sets of models for young solar-like stars and low-stars are presented and compared to the {\\sl Spitzer}'s detection limits for the programme of the {\\sl Cores to Disks} legacy survey.  Minimal disks masses of $\\approx 10^{-9}$ M$_\\odot$ for a T Tauri star and $\\approx 10^{-7}$ M$_\\odot$ for a brown dwarf are needed for the disk to be detectable by {\\sl Spitzer} in the mapping mode.  IRAC and MIPS colours of Taurus Class II and III objects are compared to passively heated models with a ``representative'' disk geometry. The average location of the two classes of objects are well reproduced, as well as the extreme colours of some of the objects, that may correspond to highly tilted disks.  $R_\\mathrm{in} =1$ AU is found to be maximum inner radius required to account for the observed colours of Class II objects. Further modelling of individual sources, combining optical, infrared and millimeter photometry with images and/or IRS spectroscopy, will be needed to better understand individual disks, their dust properties and their evolution."
    },
    "0606/astro-ph0606599_arXiv.txt": {
        "abstract": "Star-formation history is strongly related to environment; the most massive and least star-forming galaxies reside in the highest density environments.  There are now good observational reasons to believe that the progenitors of these red galaxies have undergone starbursts, followed by a post-starburst phase.  Post-starburst (``K+A'' or ``E+A'') galaxies appear in the SDSS visible spectroscopic data by showing an excess of A star light (relative to K giant light) but deficient $\\Halpha$ line emission.  We investigate the environments of these galaxies by measuring \\textsl{(1)}~number densities in $8\\,h^{-1}~\\mathrm{Mpc}$ radius comoving spheres, \\textsl{(2)}~transverse distances to nearest Virgo-like galaxy clusters, and \\textsl{(3)}~transverse distances to nearest luminous-galaxy neighbors.  We compare the post-starburst galaxies to currently star-forming galaxies identified solely by A-star excess or $\\Halpha$ emission.  We find that post-starburst galaxies are in the same kinds of environments as star-forming galaxies; this is our ``null hypothesis''.  More importantly, we find that at each value of the A-star excess, the star-forming and post-starburst galaxies lie in very similar distributions of environment.  Other studies finding similar results have argued that galaxy transformations occur slowly (time scales $>1$~Gyr), but this is at odds with the observational evidence that red galaxies are formed via starbursts.  The only deviations from our null hypothesis are barely significant: a slight deficit of post-starburst galaxies (relative to the star-forming population) in very low-density regions, a small excess inside the virial radii of clusters, and a slight excess with nearby neighbors. None of these effects is strong enough to make the post-starburst galaxies a high-density phenomenon, or to argue that the starburst events are primarily triggered by external tidal impulses (\\eg, from close passages of massive galaxies).  The small excess inside cluster virial radii suggests that some post-starbursts are triggered by interactions with the intracluster medium, but this represents a very small fraction of all post-starburst galaxies. ",
        "introduction": "How do old, dead, early-type galaxies form?  There are two strong arguments that bulge-dominated galaxies (\\eg, ellipticals, lenticulars, and very early-type spirals) have progenitors that went through a starburst phase.  The first is that bulge-dominated galaxies show enhancements in $\\alpha$-type elements over the Solar chemical abundance mix \\citep[\\eg, ][]{worthey98a, eisenstein03b}.  These abundance patterns are naturally produced when star formation---or the last phase of star formation---has occured in a burst too rapid to allow recycling of the elements ejected by Type~1a supernovae.  The second argument is that the post-starburst galaxies identified spectrally in large surveys of the Local Universe have the colors, surface brightnesses, and radial profiles that would be expected if they are to evolve passively into new bulge-dominated galaxies \\citep{quintero04a}; they cannot evolve passively into disk-dominated or other late types.  Starburst origin is also supported indirectly by the uniformity seen in early-type galaxies' stellar populations \\citep[\\eg, ][]{eisenstein03b} and their lack of large resevoirs of cold gas and dust.  Post-starburst galaxies are identified spectroscopically by having a large contribution to their spectral energy distribution from A stars---\\ie, new stars must have formed within the last $\\sim 1~\\mathrm{Gyr}$---but no, or very little, contribution from O and B stars---\\ie, no new stars have formed within the last $\\sim 0.01~\\mathrm{Gyr}$.  In practice, these ``K+A'' or ``E+A'' galaxies\\footnote{The terminology ``K+A'' is to be preferred to ``E+A'' because the identification is spectral, not morphological, and ``K'' and ``A'' name spectral types.  ``E'' names a morphological type \\citep{franx93a, dressler99a}.} \\citep{dressler83a, zabludoff96a, dressler99a} are identified by having strong Balmer absorption or blue continua (A-star indicators) but no $\\Halpha$ or [O\\,II] emission lines (O and B-star indicators).  Importantly, because A stars have known lifetimes, the evolution of the population can be ``timed'' and rates computed; we find that of order 1~percent of the galaxy population is going through this phase each Gyr at $z\\sim 0.1$ \\citep{quintero04a}.  What is not known is what precedes or triggers the (necessarily rapid) truncation of star-formation.  Is it an external event, such as a tidal impulse, accretion event, or major merger?  Or is it a purely internal event, such as an AGN flare or the abrupt exhaustion of star-formation fuel?  Either way, since disk-dominated galaxies are the galaxies that contain young stars and the cold-gas fuel to make more, post-starburst galaxies must lie on some kind of evolutionary pathway between the disk-dominated and bulge-dominated populations.  The distribution of galaxy star-formation rates is a strong function of environment, with much lower star-formation rates in higher density regions \\citep[\\eg, ][]{kennicutt83a, hashimoto98a, balogh01a, martinez02a, lewis02a, gomez03a, blanton03d, hogg03b, hogg04a, kauffmann04a, blanton05b, quintero06a}.  Although it is customary to think of this as being a consequence of the morphology--density relation \\citep{dressler80a, postman84a}, in fact recent studies with large samples have shown that in fact the star-formation--environment relation has more explanatory or informative power than the morphology--density relation, at least with the morphological proxies currently available for large samples \\citep{hashimoto98a, hashimoto99a, kauffmann04a, blanton05b, quintero06a}.  How is the information about environment transmitted to galaxy star-formation activity?  Are there violent events when galaxies fall into high density regions?  Or do the galaxies reduce their star-formation rates gradually as they find themselves in denser and denser environments? Studies of the (strong, observed) evolution of the fraction of galaxies in clusters that are blue (often called the ``Butcher--Oemler effect''; \\citealt{butcher78a}) have generally concluded that this evolution is gradual, at least when compared to the lifetimes of A stars \\citep{poggianti99a, balogh00a, kodama01a}, although there are certainly some galaxies in clusters that appear to be undergoing rapid evolution \\citep[\\eg, ][]{vogt04a, yang04a}.  It is not clear how to reconcile this conclusion of gradual evolution with the conclusion (mentioned above) that a typical early-type galaxy has undergone a massive starburst in its past.  Galaxy star-formation rates are also evolving very rapidly with redshift in the field; this result comes from many different techniques at many different wavelengths \\citep[\\eg, ][]{lilly96, hammer97, rowan-robinson97, hogg98o2lf, tresse98, cowie99, flores99, mobasher99, haarsma00, juneau05a, schiminovich05a}.  This is usually imagined as being related not to infall into dense regions but rather to the supply of cold gas.  On the other hand, since gravitational clustering brings galaxies into more and more dense environments with cosmic time, this might not be unrelated to the star-formation--environment relation and the Butcher--Oemler effect.  At the same time as star formation in the Universe is declining, the total density of stellar mass on the ``red sequence'' of early-type or bulge-dominated galaxies is increasing \\citep{bell04a, blanton06a, faber05a, brown06a}.  If blue galaxies transform into red galaxies via a post-starburst phase, as we believe they must, then this evolution of the red sequence must be quantitatively matched with an evolving population of post-starburst galaxies.  As transition objects between the star-forming, disk-dominated and dead, bulge-dominated populations, the post-starburst galaxies could in principle have the environmental characteristics of either. Originally, K+A galaxies were found in high-density regions \\citep{dressler83a, couch87a}, and thought to be a ``cluster'' population.  Of course the early searches for such galaxies were made in cluster fields.  Once systematic searches for K+A galaxies were made in large redshift surveys, it was found that they are not particularly concentrated in clusters or high density regions, but rather live in a wide range of environments \\citep[][Helmboldt et~al., in preparation]{zabludoff96a, quintero04a, blake04a, goto05a}.  In the large SDSS and 2dFGRS samples, it can be shown that the mean environment \\citep[][Helmboldt et al., in preparation]{quintero04a} and distribution of environments \\citep{blake04a} are both similar to those of spiral or disk-dominated galaxies.  The environments of disk-dominated galaxies---isolation and small groups, where the virialized mass has a similar velocity dispersion to the contained galaxies---are the best environments in the Local Universe to find galaxy--galaxy mergers, which are the top candidates for triggers for the post-starburst galaxies.  After all, the major mergers observed in the Local Universe are all associated with very high star-formation rates, and major mergers are expected to disrupt disks and leave behind the dynamically hot stellar orbits characteristic of the bulge-dominated population.  With a sample of more than $10^3$ K+A galaxies \\citep{quintero04a}, we are in a position to ask much more detailed questions about the range of environments in which they lie, and the relationships between environmental and star-formation properties.  That is the purpose of this \\textsl{Article}, with the goal of constraining the possible triggering mechanisms for this very important galaxy population.  In what follows, a cosmological world model with $(\\Omega_\\mathrm{M},\\Omega_\\mathrm{\\Lambda})=(0.3,0.7)$ is adopted, and the Hubble constant is parameterized $H_0=100\\,h~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$, for the purposes of calculating distances\\citep[\\eg,][]{hogg99cosm}. ",
        "conclusions": "As we have discussed elsewhere \\citep{quintero04a}, post-starburst galaxies plausibly lie on an evolutionary sequence between disk-dominated galaxies, which are forming stars and contain the neutral gas fuel for further star formation, and bulge-dominated galaxies, which have no star-formation fuel and show chemical signatures of past star-formation bursts.  Post-starburst galaxies might even be the remnants of major mergers.  A priori, the environments of these galaxies could be either like those of disk-dominated galaxies or those of bulge-dominated galaxies, or somewhere in-between.  Of course, prior to this study, it was already known that the mean environments of post-starburst galaxies are more similar to those of disk-dominated galaxies than those of bulge-dominated galaxies \\citep[][Helmboldt et al., in preparation]{zabludoff96a, quintero04a, blake04a}.  Here we have not only confirmed this result, we have shown that for each value of the A-star excess, the post-starburst galaxies with that A-star excess find themselves in similar environments to star-forming galaxies with that same A-star excess.  In other words, A-star excess is more closely related to environment than $\\Halpha$~EW, since the post-starburst galaxies have $\\Halpha$~EWs like bulge-dominated galaxies.  This result confirms our ``null hypothesis'' that the processes that connect a galaxy's star-formation history to its environment act on long timescales, longer than A-star lifetimes ($\\sim 1~\\mathrm{Gyr}$). This is not surprising, since at typical cosmological velocities ($\\sim 100~\\mathrm{km\\,s^{-1}}$), a galaxy can only travel $\\sim 1~\\mathrm{Mpc}$ in a Gyr, and there is now pretty good empirical evidence that everything important about galaxy environments happens on scales $< 1~\\mathrm{Mpc}$ \\citep{blanton04c}.  On the other hand, this result is somewhat difficult to reconcile with the observation from chemical abundances \\citep[\\eg, ][]{worthey98a, eisenstein03b} and central stellar densities \\citep{quintero04a} that a typical red (early-type, or bulge-dominated) galaxy has undergone a brief but strong starburst at some point in its past.  Perhaps the starbursts are triggered very randomly, with just a slight change in starburst trigger probability with environment, so we don't see the relation clearly.  Perhaps the observed relationship between star-formation rate and environment was set down at very high redshift and is in fact diluting at the present epoch.  Perhaps the starbursts are primarily triggered prior to infall into dense environments, in which case the galaxies somehow ``know'' the environment in which they will end up!  We have shown that the fraction of the whole SDSS galaxy sample classified as ``K+A'' (post-starburst) is a function of environment, and that its dependence on environment is very similar to that of the fraction of the sample classified as star-forming.  This is also consistent with our null hypothesis.  The only deviations all have the sense that the K+A fraction is a slightly weaker function of environment than the star-forming fraction: We find slightly fewer K+As in extremely low density environments, and slightly more inside the virial radii of massive clusters and close to luminous neighbors than we would expect by naive scaling of the star-forming fraction.  None of these deviations from the predictions of the null hypothesis are large, but they may point to triggering mechanisms for the starburst and post-starburst phases: It is possible that a small fraction of post-starburst galaxies are triggered by close passages of luminous galaxies.  It seems likely, from Figure~\\ref{fig:clusto}, that a small fraction of post-starbursts are triggered by interactions with intracluster medium, because the deviation of the post-starburst fraction relative to the star-forming fraction does occur at the virial radius.  We can therefore make two negative statements about the triggering of the \\emph{majority} of starbursts, both plausible at the outset: The first is that only a small fraction of post-starburst galaxies are triggered by external tidal impulses from close passages of massive galaxies.  This does not rule out the possibility that they are created by mergers or accretion events, but if they are, the truncation of star-formation must occur after the two merging galaxies are no longer identifiable as separate galaxies.  Our punchline may seem to be at odds with the punchline of some previous work \\citep{goto05a}.  Quantitatively, however, there is no disagreement, given the differences in methodology, and in fact both studies do find small excesses of neighbors at small scales.  This suggests that although tidal-impulse triggering (\\ie, triggering by close passages of massive neighbors) no-doubt occurs, it is not the dominant mechanism.  The second negative statement we can make is that only a small fraction of post-starburst galaxies are created by IGM interactions on infall into clusters.  Certainly some are; the identification of the virial radius in the abundance of post-starburst galaxies relative to star-forming galaxies is intriguing.  The galaxy populations inside clusters are very different in morphological mix and star-formation-history mix than galaxy populations elsewhere.  What physical process are involved in enforcing these differences?  Many have hypothesized---quite naturally---that radical transformations must happen on infall \\citep{poggianti99a, balogh00a, kodama01a}; indeed some galaxies have been ``caught in the act'' of a radical transformation \\citep[\\eg, ][]{vogt04a}.  Quantitatively, however, the fraction of galaxies showing clear evidence for intracluster medium interactions is very small.  This is consistent with prior work in this area \\citep{poggianti99a, balogh00a, kodama01a, vogt04a}.  The orbital time for a galaxy falling into a cluster is only a few times longer than the lifetime of A stars, so the excess of post-starburst galaxies in the infall regions is expected to ``smear'' into the cluster center; \\ie, the morphology of Figure~\\ref{fig:clusto} is about right (given the small numbers) for the intra-cluster medium hypothesis.  What remains to be seen is whether, quantitatively, the tiny number of post-starburst galaxies observed within clusters is consistent with the competing facts that clusters are continuously growing by accretion of field galaxies and groups and that the morphological and star-formation mix in clusters is different inside and outside of clusters; the prodigious expected infall means that a lot of galaxies ought to be changing their morphologies and star-formation rates on infall.  We expect such an analysis---beyond the scope of this paper---to conclude that either the transformation process is slow ($>1$~Gyr, which is not necessarily consistent with observations of abundances in bulge-dominated galaxies), or else that the galaxies in clusters somehow ``knew in advance'' (from their pre-infall group or filament environment) that they were destined to end up in a cluster.  A philosophical point arises here: The fact that galaxy populations inside clusters are different from those outside does not \\emph{necessarily} mean that transformations take place as galaxies move from one environment to the other.  The remaining hypotheses for the triggering of the starburst (or, more properly, star-formation truncation) events that precedes the post-starburst phases of these galaxies are: some kinds of random internal catastrophes or some kinds of galaxy--galaxy mergers.  This latter possibility, which is consistent with all of the results here, is directly supported by the discovery of post-merger morphological signatures (\\eg, tidal arms) in many post-starburst galaxies \\citep{yang04a, goto05a}.  It is also exciting, because merging is one of the fundamental processes of cosmogony, and holds great promise for providing precise connections between cosmological observations and theory at small scales."
    },
    "0606/astro-ph0606199.txt": {
        "abstract": "We study the detection of extragalactic point sources in two-dimensional flat simulations for all the frequencies of the forthcoming ESA's Planck mission. In this work we have used the most recent available templates of the microwave sky: as for the diffuse Galactic components and the Sunyaev-Zel'dovich clusters we have used the ``Plank Reference Sky Model''; as for the extragalactic point sources, our simulations - which comprise all the source populations relevant in this frequency interval - are based on up-to-date cosmological evolution models for sources. To consistently compare the capabilities of different filters for the compilation of the - hopefully - most complete blind catalogue of point sources, we have obtained three catalogues by filtering the simulated sky maps with: the Matched Filter (MF), the Mexican Hat Wavelet (MHW1) and the Mexican Hat Wavelet 2 (MHW2), the first two members of the MHW Family. For the nine Planck frequencies we show the number of real and spurious detections and the percentage of spurious detections at different flux detection limits as well as the completeness level of the catalogues and the average errors in the estimation of the flux density of detected sources. Allowing a $5\\%$ os spurious detections, we obtain the following number of detections by filtering with the MHW2 an area equivalent to half of the sky: 580 (30 GHz), 342 (44 GHz), 341 (70 GHz), 730 (100 GHz), 1130 (143 GHz), 1233 (217 GHz), 990 (353 GHz), 1025 (545 GHz) and 3183 (857 GHz). Our current results indicate that the MF and the MHW2 yield similar results, whereas the MHW1 performs worse in some cases and especially at very low fluxes. This is a relevant result, because we are able to obtain comparable results with the well known Matched Filter and with this specific wavelet, the MHW2, which is much easier to implement and use. ",
        "introduction": "\\label{sec:intro} \\footnotetext{E-mail: caniego@ifca.unican.es} In two years from now, the ESA's Planck\\footnote{http://www.rssd.esa.int/index.php?project=Planck} satellite~\\citep{tauber} will inaugurate a new era in the studies of the Cosmic Microwave Background (CMB) radiation. Planck will observe the microwave sky with unprecedented angular resolution and sensitivity in nine frequency channels ranging from 30 to 857 GHz. Besides CMB anisotropies, Planck will also yield all-sky maps of all the major sources of microwave emission, including a large number of extragalactic sources that have not yet been observed at these frequencies. The study of these sources of microwave emission --often referred to as \\emph{``contaminants''} or \\emph{``foregrounds''}-- is twofold: on one hand, it is necessary to remove them in order to have the cleanest possible view of the CMB and, on the other hand, a better knowledge of the different foregrounds is a scientific goal in itself. As the launch date approaches, a big effort is being made to develop and to test state-of-the-art data processing techniques that will optimise the scientific exploitation of the forthcoming Planck data.  The case of extragalactic foregrounds deserves to be considered in detail. Radio and infrared selected galaxies will be seen by Planck as point-like objects due to their very small projected angular size as compared to the experiment resolution (see Table 1). Hence, they are usually referred to as extragalactic point sources or just as \\emph{point sources} (as we do hereafter). Point sources are expected to be a major contaminant for Planck at multipoles $\\ell\\geq 500-1000$ \\citep{tof98,dz99,hob99,max00} and it is necessary to detect and remove as many of them as possible in order to clean CMB maps. As a very important astrophysics by-product, the detection process will yield all-sky catalogues   of extragalactic point sources in a frequency interval in which they are lacking and that will prove very useful to constrain models of galaxy formation and evolution.  Unfortunately, many of the component separation techniques that are generally used to separate diffuse Galactic foregrounds are not well suited to deal with point sources. This is mainly due to the fact that each galaxy is an independent source, different, in principle, to any other source in the sky. Albeit average energy spectra can be defined for different source populations, the spectral emission law is different for each galaxy and it makes impossible to apply separation methods which exploits a single energy spectrum for this foreground component.   If the lack of knowledge on their spectral emission properties can be troublesome, their projected angular shape is basically the same for all of them: a Dirac's $\\delta$-like response convolved with the instrument beam response function.  Thus, techniques that take into account the specific angular shape of point sources, such as wavelets and band-pass filters, are particularly useful for detecting them. In the last few years a number of techniques have been proposed for the specific case of point source detection in CMB maps, including the Mexican Hat Wavelet \\citep[]{cay00,patri01a,patri01b,patri03}, the Matched filter (MF) \\citep[]{max98}, the Scale-Adaptive filter \\citep{sanz01,yo02a,yo02b}, the Adaptive Top Hat filter \\citep{chi02}, the Biparametric Scale-Adaptive filter \\citep{can05} and the recently introduced Mexican Hat Wavelet Family \\citep[]{jgn06}.  All the previously mentioned methods belong to the class of wavelet and linear filter techniques.  Additionally, non-linear techniques such as Bayesian detection \\citep{hob03} have been successfully applied to point source detection in the CMB context, but since the use of those techniques implies a totally different methodological approach we will focus on the above mentioned filters and wavelets. Whereas some works have made an effort to compare the performances of some of them both theoretically and by using almost-ideal simulations \\citep{bar03,vio04,can05}, an attempt to compare the existent techniques under {\\it (almost)} ``real life'' conditions for the future Planck mission has not been done, yet.  In this work we intend to reproduce the conditions of a blind point source survey as it will be carried out by Planck. We will use the Planck Reference Sky and the nominal Planck instrument characteristics and goal performance to simulate realistic sky emission as it should be observed at the nine frequencies covered by the satellite. Using such realistic simulations we will compare the performance of different filters.  The ultimate goal is to decide which is the best tool of choice for the Planck case. The `goodness' of a filter will be evaluated according to the following criteria:  \\renewcommand{\\labelenumi}{(\\roman{enumi})}  \\begin{enumerate} \\item \\label{launo} The filter must be well suited to conduct a blind survey, that is, it must be able to work well with a minimum number of a priori assumptions about the data. \\item \\label{launo_bis} Besides, it must be robust against the effect of possible systematics. \\item \\label{lados} It must yield, after the detection process, a high number of positive detections. \\item \\label{latres} It must yield, after the detection process, a low number of false detections (not higher that, let us say, $5\\%$ of the integral, or total, number of detections above the corresponding detection threshold). \\item \\label{lacuatro} Moreover, additional factors, such as the flux detection limit of the output catalogue, its completeness and the accuracy with which the positions and the flux densities of the sources are estimated, will be taken into account. \\end{enumerate}  We would like to remind that previous criteria are similar to those required for the Early Release Compact Source Catalogue, ERCSC, a blind catalogue of point sources that is one of the objectives of the Planck collaboration. In the compilation of this kind of catalogue, factors such as quickness and accuracy of the estimation of the flux of the sources have the priority over a high absolute number of detections. Keeping this in mind, criterion~\\ref{launo} eliminates from the competition tools such as the Biparametric Scale-Adaptive filter that, even if they are potentially very powerful, require a detailed knowledge of the probability distribution of the foregrounds\\footnote{Such tools could be very useful for exhaustive data mining the sky down to very low flux limits.}. The Adaptive Top Hat filter is known to produce strong ringing artifacts around the sources which can lead to a high number of false detections --against criterion~\\ref{latres}--, in particular in the vicinity of bright sources. The Scale-Adaptive filter seems to perform similarly or slightly worse than the MF \\citep{yo02c}.  Therefore, in this work we will focus on the comparison of two tools: the Mexican Hat Wavelet Family (MHWF), which includes the standard Mexican Hat Wavelet, and the Matched Filter\\footnote{In particular, \\citet{jgn06} have shown that the second member of the Mexican Hat Wavelet Family (to be introduced in section~\\ref{sec:method}) is the one that gives the best results for the case of the Planck Low Frequency Instrument (LFI) channels, and therefore we will limit the discussion to this wavelet, comparing it with the standard Mexican Hat (the first member of the family), that has been widely used in the literature with good results, and the matched filter.}.  At first glance, the MF should be the obvious winner in the comparison. By definition, it is the best linear operator that can be applied to the data in order to maximise the signal to noise ratio of a signal with a known profile embedded in additive noise. But, in practise, the use of MFs does not lack of subtleties that must be considered here.  In Fourier space the MF is proportional to the inverse of the power spectrum of the noise. That means that the power spectrum of the noise must either be known a priori or be estimated from the data in order to construct the filter. If point sources are scarce, the power spectrum of the noise can be reasonably approximated by the power spectrum of the observed data, that is easy to \\emph{estimate} by means of any of the power spectra toolboxes available in scientific softwares. But any estimated power spectrum, as good as it may be, is just a good guess of the real thing. This leads to a number of problems:  \\begin{itemize} \\item Firstly, it is necessary to estimate the value of the power spectrum for all the Fourier modes present in the image. This implies the estimation of several hundreds of numbers for a typical CMB image. For the typical image size of the sky patches we work with, each Fourier mode must be estimated from a small number of data samples. Therefore, the estimated power spectrum is \\emph{noisy}, especially at low Fourier modes. On the contrary, to use the various members of the Mexican Hat Wavelet Family it is only necessary to determine one single number, the optimal scale of the wavelet and, thus, it is much less sensitive to noise estimation. \\item Since the estimated power spectrum is noisy, if it is directly used to construct the MF it very often happens that the resulting filter is so full of discontinuities and `jumps' in Fourier space that strong ringing effects appear in the filtered image. Therefore, it is necessary to smooth the power spectrum before constructing the filter. The different possible choices used in the smoothing procedure introduce some degree of arbitrarity in the use of MFs: one can, for example, apply some binning and interpolation scheme, or use polynomial fitting to the power spectrum, etc. \\item An additional problem appears when it is not possible to properly estimate some Fourier modes. This is the case when the image which has to be filtered is not complete (for example, if a mask is applied to the data in order to cut bad pixels, or in areas of the sky where there is incomplete coverage by the instrument). In that case the missing modes must be somehow guessed in order to build the MF. This problem is much less relevant for the case of wavelets. \\item Moreover, if we make considerations in the sphere we will have to deal with important problems when using a Matched filter as compared with a wavelet. These problems arise from the fact that the foregrounds are very different in different regions of the sky and therefore we need to use the appropriate filter for every region. The approach followed by \\citet{patri03} using the Spherical MHW (SMHW) was to divide the sky in a number N of regions and obtain the optimal scale for all of them. Then they determined that many of these scales gave similar results and that they could be divided in a small number of different groups. This allowed them to construct just a few filters in the sphere with their corresponding scales and filter the maps a small number of times instead of N. This is important because depending on the resolution, the filtering process may need a lot of CPU time. The problem with the MF is that instead of calculating N optimal scales we would need to calculate N filters, calculating the power spectra from the N regions, and filtering the maps N times, once for every filter. This process will require enormous amounts of CPU time, especially when dealing with high resolution images, and therefore will make it unfeasible in practise.   \\end{itemize}  Therefore, and as any experienced practitioner perfectly knows, the use of the MF is not the same as filtering by the inverse of the squared Fourier transform of the data. On the contrary, it requires a non negligible effort of handmade tuning that is not free from arbitrarities. Besides, all the previous effects lead to an unavoidable degradation of the performance of the MF under realistic conditions.  Furthermore, all the theoretical superiority of the MF with respect to other linear filters comes from the fact that it maximises the signal to noise ratio, that is, it minimises the variance of the filtered noise. But if the noise is not Gaussian, as it is the case in CMB maps due to the emission contributed by Galactic foregrounds, \\textit{the fact that the variance is minimum does not guarantee that the number of false detections be minimum}. There may be outliers that are not removed. In that case, it is not by any means clear that the MF should be better than any other filter.  Taking all the previous points into account, the comparison between the MF and the wavelets is still necessary. The performance of wavelets will degrade as well when going from ideal to realistic conditions. Wavelets, however, are much less sensitive than MFs to the problems described in the paragraphs above. Therefore, it is expected that the performance degradation will be less severe. If we can find a wavelet that performs nearly as well as the MF, but without having to resort to handmade tuning, we will have a detection tool that is as good as the MF regarding criteria~\\ref{lados},~\\ref{latres} and~\\ref{lacuatro} but is better regarding criterion~\\ref{launo} and~\\ref{launo_bis}. Such a wavelet would be preferable to MF for the compilation of a Planck blind point source catalogue.  In Section~\\ref{sec:method} we briefly review the tools to be used in this paper: the MF and the wavelets belonging to the Mexican Hat Wavelet Family. In Section~\\ref{sec:sims} we describe the realistic Planck simulations we use. The results are summarised in Section~\\ref{sec:results}.  In Section~\\ref{sec:discussion} we discuss some additional issues regarding point source detection in microwave satellite missions.  Finally, we describe our main conclusions in Section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have compared the performance of three filters when dealing with the detection of point sources in CMB flat sky maps. These filters are the well known MF, the Mexican Hat Wavelet (MHW1) and the recently introduced Mexican Hat Wavelet 2 (MHW2). The latter is the second member of the so-called Mexican Hat Wavelet Family (the MHW1 is the first), a group of wavelets obtained by applying the Laplacian operator to the Gaussian \\citep{jgn06}. This new wavelet is part of an effort to improve the MHW1, a tool already exploited by our group that has proved very useful in the detection of compact sources and of non-Gaussianity in CMB maps \\citep{patri01a,patri03,patri04}.  We have tested these tools in realistic 2D simulations of the microwave sky prepared in the framework of optimising the efficiency in separating the various astrophysical components in the forthcoming ESA's Planck Satellite all-sky maps. As for the Galactic foregrounds and the S-Z effect in clusters of galaxies, we have used the available ``Planck Reference Sky Model''; we have adopted the Standard \"concordance\" Model for simulating CMB anisotropy maps and, as for extragalactic point sources, we have used simulations obtained by the software discussed in \\citet{jgn05} and with the number count models for sources of \\citet[]{dz05,gr01,gr04} and \\citet[]{ngr04}. We then applied the three considered filters to a sufficient number of flat patches to cover half of the sky $(2 \\pi \\,\\mathrm{sr}, b> |30^\\circ|)$ (see Section 3).  We have found that the MHW2 and the MF outperform the MHW1 in some aspects, especially at the lowest Planck frequencies. The three filters yield approximately the same number of real detections, down to the same flux detection limit, although the MHW1 yields a corresponding higher number of spurious detections. Moreover, the MHW2 and the MF give comparable results for almost every one of the analysed indicators. As shown in Figure \\ref{fig:mfs}, it is remarkable that, even the estimated shape of the MF tends to that of the MHW2, at its optimal scale, for most of the frequencies discussed in this work.  In a previous work \\citep{patri03}, a multi-fit approximation to estimate the flux density of the sources detected by the MHW1 was shown to be able to improve the results (as compared with the direct approach used in this work). We have applied this technique to the simulated maps used in this work and we found different results for the LFI and the HFI Planck frequency channels. In the first case, LFI, the improvement is small and the final results are never comparable to the ones obtained with the MHW2. In the second case, this procedure helps to slightly reduce the number of spurious sources, except for the highest frequency of HFI (857 GHz), where the decrease in the number of spurious is significant and the MHW1 approaches the results obtained with MHW2.  These results are very important because both wavelets, the MHW1 and the MHW2, are represented by a known analytical function. The only parameter that needs to be obtained for each simulation is the ``optimal scale'', which is calculated locally in a very easy way.  On the other hand, the correct definition of the MF implies a number of steps. Firstly, it is necessary to estimate the value of the power spectrum for all the Fourier modes present in the image, especially for the low modes where the power spectrum is \\emph{noisy}. Secondly, the use of such a noisy power spectrum to construct the MF often yields a filter with many discontinuities in Fourier space which, in turn, produces ringing effects in the filtered image. Therefore some smoothing in the spectra needs to be done before constructing the filter, which introduces further arbitrariness. Thirdly, sometimes it will not be possible to properly estimate some Fourier modes, for example when using masks with missing data, and these modes will have to be guessed.  Therefore, the most relevant conclusion of this analysis is that the MHW2 can be surely a better choice for the definition of a blind source catalogue -- like the ERCSC foreseen for the future Planck mission -- because it gives numbers of detected and of spurious sources comparable to the ones obtained with the MF but is easier to implement, more robust and has much lesser CPU requirements, especially in all-sphere applications."
    },
    "0606/astro-ph0606285_arXiv.txt": {
        "abstract": "We develop a scaling ansatz for the master equation in Dvali, Gabadadze, Porrati cosmologies, which allows us to solve the equations of motion for perturbations off the brane during periods when the on-brane evolution is scale-free. This allows us to understand the behavior of the gravitational potentials outside the horizon at high redshifts and close to the horizon today. We confirm that the results of Koyama and Maartens are valid at scales relevant for observations such as galaxy-ISW correlation. At larger scales, there is an additional suppression of the potential which reduces the growth rate even further and would strengthen the ISW effect. ",
        "introduction": "That cosmic acceleration is a fact appears indubitable. Instead of an exotic new form of dark energy driving the acceleration, it may be caused by a modification of gravity. Precise measurements for gravity are only available in the range of scales from a millimeter to that of the solar system---we do not have any direct probe of Einstein gravity beyond these boundaries. Cosmic acceleration may originate in a breakdown of Einstein gravity at distances beyond the range above.  Dvali, Gabadadze and Porrati (DGP) \\cite{dvali00} have proposed a braneworld theory in which our universe is a (3+1)-dimensional brane embedded in an infinite Minkowski bulk. Gravity propagates everywhere, but, on the brane, an additional (3+1)-dimensional gravitational interaction is induced. This allows for gravitational potentials on the brane of a (3+1)-dimensional form at small distances to evolve into (4+1)-dimensional form beyond a crossover scale determined by the unknown energy scale for the bulk gravity. The cosmological solution of this theory was shown to exhibit accelerated cosmic expansion without the aid of an exotic energy component like dark energy \\cite{deffayet00}\\cite{deffayet01}.  It has already been shown that the linearized field theory as defined by the DGP model contains ghost degrees of freedom \\cite{Luty:2003vm}\\cite{Nicolis:2004qq}\\cite{Koyama:2005tx}\\cite{Gorbunov:2005zk}\\cite{Charmousis:2006pn}, or even may violate causality in certain limits \\cite{Adams:2006sv}. It is known, for instance, that the de Sitter background is unstable to classical linear perturbations; however, it is claimed in \\cite{Deffayet:2006wp} that strong-coupling effects at small radii around matter sources ensure that the theory remains stable.  The point of view of our work is to assume that linear perturbation theory remains valid on the largest scales. This is motivated by the fact that the late universe is dominated by the gravitational interaction of dark-matter haloes. The internal structure of the haloes is controlled by the strongly-coupled non-linear theory. On the other hand, the radius below which strong-coupling is important for haloes is approximately equivalent to their size and therefore their interactions should be driven by the linear theory analyzed in this work.  We do find that deep into the accelerated era the spacetime becomes unstable on the timescale of the expansion. However, this is an effect that only becomes important far into the future and is negligible as far as the observational impact today is concerned. We therefore assume that during the early universe, when the theory does not exhibit instabilities, the analysis for DGP proceeds in exactly the same way as that for GR. Then, deep during the acceleration era, instabilities develop and the theory may or may not be saved by non-linear effects---an issue on which we remain agnostic. This evolutionary history appears to be the only one which is capable of reproducing the universe as we see it. If strong coupling effects are important straight away and at all scales, the approximation of a homogeneous background cosmology is completely inapplicable and the DGP model would not be able to reproduce observations such as supernova luminosities. We therefore effectively assume a best-case scenario for DGP: should this analysis fail to predict the observations, the model is excluded. If it passes the observational tests, a more careful study of the effects of the strong coupling regime during the acceleration era would be required to understand fully the future evolution of the universe in the DGP model.  Under the above assumptions, the equations of motion for the theory of gravity on the brane, pertinent to the study of cosmology, do not close owing to the interaction with the bulk at first order in perturbation theory. Koyama and Maartens \\cite{koyama05} have used a quasi-static approximation, valid well within the horizon, to investigate structure formation at smaller scales. This solution shows the essential role that the bulk plays in correcting the gravitational potentials, reducing the growth rate. Also, Lue \\emph{et al.} \\cite{Lue:2004rj} have reached a similar conclusion using a different approach, including non-linearities in their calculations.  In the following, we present a new scaling ansatz for the master equation, allowing us to solve the equation and calculate the resulting cosmological evolution at all scales for high redshifts, and close to the horizon today.  In \\S \\ref{s:eom}, we review the linearized equations of motion for DGP.  We present our scaling solution in \\S \\ref{s:scaling} and discuss its cosmological implications in \\S \\ref{s:cosmo}.  We study the robustness of the scaling solution in the Appendix and discuss these results in \\S \\ref{s:discussion}. ",
        "conclusions": "\\label{s:discussion} We have introduced a new scaling ansatz which allows solutions to linear perturbations in the DGP model on all scales less than the cross-over scale $r_\\cs$ up to the present epoch. The equations of motion for linear perturbations on the brane require knowledge of the gradient of the so-called master variable into the bulk. The master variable obeys a master equation in the bulk.  To solve the master equation, it is sufficient to have two boundary conditions, one on the brane and the other in the bulk.  Our scaling solution begins with an ansatz for the brane boundary condition: that the evolution of the master variable is scale free on the brane.   The second boundary condition is that the master variable vanishes at the causal horizon in the bulk.  With these two boundary conditions, we solve the master equation to determine the gradient. With the gradient known, we can then replace the scale-free ansatz with the dynamical solution and iterate the solution until convergence.  We find that the quasi-static (QS) solution of \\cite{koyama05} is rapidly approached once the perturbation crosses the horizon. Before horizon crossing there are strong deviations from the quasistatic solution. For modes that crossed the horizon only recently during the acceleration epoch, we find that the metric perturbation $\\Phi-\\Psi$ decays more rapidly that the QS solution. The QS solution itself has a stronger decay than the $\\Uplambda$CDM model.  The extra decay compared with $\\Uplambda$CDM is extremely robust to changing the gradient of the master variable into the bulk, the one variable that is required to close the equations of motion on the brane. We consider the observational consequences of these results in a companion paper \\cite{Song:2006}."
    },
    "0606/hep-ph0606137_arXiv.txt": {
        "abstract": "\\hspace*{\\parindent} There are many inflationary models in which inflaton field does not satisfy the slow-roll condition. However, in such models, it is always difficult to generate the curvature perturbation during inflation. Thus, to generate the curvature perturbation, one must introduce another component to the theory. To cite a case, curvatons may generate dominant part of the curvature perturbation after inflation. However, we have a question whether it is unrealistic to consider the generation of the curvature perturbation during inflation without slow-roll. Assuming multi-field inflation, we encounter the generation of the curvature perturbation during inflation without slow-roll. The potential along equipotential surface is flat by definition and thus we do not have to worry about symmetry. We also discuss about KKLT models, in which corrections lifting the inflationary direction may not become a serious problem if there is a symmetry enhancement at the tip (not at the moving brane) of the inflationary throat. ",
        "introduction": "Among many benefits from the inflationary expansion that takes place in the early Universe, an important prediction of inflation would be the generation of a spectrum of the primordial perturbations. Such perturbations naturally arise from the zero point vacuum fluctuations in quantum fields, which are stretched during inflation to cover very large scales in our present Universe. In the standard scenario of the inflationary Universe, the observed density perturbation is produced by a light inflaton field that rolls slowly down its potential. At the end of inflation, the inflaton field oscillates about the minimum of its potential and decays to reheat the Universe. Adiabatic density perturbation is generated because the scale-invariant fluctuations of the light inflaton field are different in different patches. Generically, slow-roll models of inflation predict an almost scale-invariant and Gaussian distribution of primordial density perturbations.  We are not going back to historical developments, however it is easy to understand that the idea of hybrid inflation may provide us with a key idea to construct successful inflationary models\\cite{EU-book}. In hybrid models, the end of inflationary expansion is a second-order phase transition triggered by a trapped field (waterfall field). D-term inflation is an important application of this idea, which is found in the paradigm of supersymmetric particle cosmology. Furthermore, an important variant of D-term inflation is found in the paradigm of brane cosmology, which is called brane inflation\\cite{brane-inflation0, angled-inflation, matsuda_braneinflation}. From phenomenological viewpoints, brane models are sometimes categorized as models with large or intermediate extra dimensions (and vice versa)\\cite{Extra_1}. The idea of large extra dimensions is important for higher-dimensional models, because it may solve or weaken the hierarchy problem. In models with large extra dimensions, fields in the standard model(SM) are localized on a wall-like structure (maybe it is a brane), while the graviton propagates in the bulk. The discrepancy of the volume factor between gauge fields and gravitational fields explains the large hierarchy between gravity and gauge interactions.  Of course, it is an important challenge to find signatures of branes in cosmological observations. Historically, it has been discussed that studying the formation and the evolution of cosmological defects would provide us with important information about branes\\cite{brane-defects, matsuda-defects}.\\footnote{Supergravity provides a natural mechanism for removing cosmological domain walls\\cite{matsuda-wall}.} Inflation models with low fundamental scale are discussed in ref.\\cite{low_inflation}. Scenarios of baryogenesis in such low-scale models are discussed in ref.\\cite{low_baryo, Defect-baryo-largeextra, Defect-baryo-4D}, where defects play distinguishable roles. The curvatons\\cite{curvaton_1, curvaton_2, curvaton_3} will play significant roles in these low-scale models\\cite{curvaton_liberate, topologicalcurvaton}. Moreover, in ref.\\cite{topologicalcurvaton}, it has been discussed that topological defects can play the role of the curvatons. Thus, defects in brane models such as monopoles, strings, domain walls and Q-balls are important\\cite{matsuda_necklace, BraneQball, matsuda_monopoles_and_walls, incidental, matsuda_angleddefect, tit-new}. It might be important to explain why defects other than strings can be produced in brane inflationary models, since (historically) it has been discussed by many authors that only strings are produced in brane inflationary models\\cite{previous-onlystrings}. It is not hard to understand that this conjecture is not so generic as it has been anticipated, as one can see from ref.\\cite{matsuda_JGRG}. Therefore, in order to find signatures of branes, it is important to consider other types of defects as well as cosmic strings. Our scenario of elliptic inflation is important, in a sense that it may provide us with another clue as to how one can find signature of branes from cosmological observations, as we will discuss in Sect.3.5.  When it appeared, D-term inflation was believed to circumvent the well-known eta-problem of supergravity models of inflation based on F-terms. Thus brane inflation, which is in a sense a variant of D-term inflation, was believed to circumvent the eta-problem. However, a similar problem arose later again in brane paradigm. The problem was that the very mechanism that lifts the moduli potential was found to lift inflaton candidates and violate slow-roll conditions\\cite{KKLT-eta}. Directions that are protected by exact global symmetry may be expected to remain flat, however it is still very difficult to find actual symmetry that can survive after moduli stabilization and can be used to construct a successful inflationary scenario.\\footnote{There are interesting approaches in ref.\\cite{DBI-inflation} and \\cite{Sinha}. } From the perspective we discussed above, we believe that finding new mechanism for generating the curvature perturbation without slow-roll inflaton is quite important. In general, the superhorizon spectrum of perturbations is thought to be generated by the amplification of the quantum fluctuations of a light inflaton field, whose mass is much smaller than the Hubble constant during inflation. This is because the quantum fluctuations of the field can reach and exit the horizon only if its Compton wavelength is larger than the horizon during inflation. However, this happens only if the inflaton is effectively massless, i.e. only if $m_{I} \\ll H_I$. The above condition seems to conflict with our aim in this paper. Luckily, we know alternatives for the slow-roll inflaton. The curvatons and their variants can generate the curvature perturbation after inflation\\cite{curvaton_1, curvaton_2}, even if there is no slow-roll during inflation. However, this mechanism still requires flat direction that is supposed to have non-trivial properties. Thus, we think it is still important to find models for generating the curvature perturbation without using neither slow-roll inflaton nor the curvatons. The situations that we will consider in this paper are both general and useful. Our mechanism can be utilized in many kinds of phenomenological models in which slow-roll condition is inevitably violated. As we are considering inflation without slow-roll, the e-foldings of our model may be short, which will require compensation by another period of inflationary expansion. ",
        "conclusions": "The most important point is that an equipotential surface appears whenever there is additional inflaton that couples to waterfall field. Fluctuation appears along the equipotential surface and exits horizon during inflation. The equipotential surface is an ellipsoid, along which fields can fluctuate despite the large mass of the fields. Then, it induces fluctuation of the total number of e-folds at the end of inflation. The end line of hybrid inflation may also be ellipsoid, which becomes another source for the curvature perturbations. The spectrum of the curvature perturbation is determined by the masses and the couplings of the fields.  For brane inflationary models, our mechanism provides us with a new inflationary scenario that circumvent the eta-problem. It is possible to generate the curvature perturbation even in a model where all flat directions are lifted at the UV side of the KS throat. It is also notable that our mechanism can be used to obtain a tilted spectrum. The spectrum is determined by the masses of the position moduli that parameterize the position of the inflationary branes. The mechanism may provide us with an important clue as to how one can obtain information about the structure of extra dimensions from cosmological observations."
    },
    "0606/astro-ph0606291_arXiv.txt": {
        "abstract": "The burgeoning field of astrophotonics explores the interface between astronomy and photonics. Important applications include photonic OH suppression at near-infrared wavelengths, and integrated photonic spectroscopy.  These new photonic mechanisms are not well matched to conventional multi-mode fibre bundles, and are best fed with single or few-mode fibres. We envisage the largest gains in astrophotonics will come from instruments that operate with single or few mode fibres in the diffraction limited or near diffraction limited regime. While astronomical instruments have largely solved the problem of coupling light into multi-mode fibres, this is largely unexplored territory for few-mode and single-mode fibres.  Here we describe a project to explore this topic in detail, and present initial results on coupling light into single and few-mode fibres at the diffraction limit.  We find that fibres with as few as $\\sim5$ guided modes have qualitatively different behaviour to single-mode fibres and share a number of the beneficial characteristics of multi-mode fibres. ",
        "introduction": "\\label{sect:intro}  %  Optical fibres have been in use in astronomical instrumentation for almost 30 years. They were first used for fibre-fed multi-object spectroscopy, which began with the Medusa instrument at Steward Observatory in 1979\\cite{Hilletal80}. The efficiency gains from the ability to observe large numbers of objects at once have made many otherwise impractical scientific programmes possible, including several huge spectroscopic surveys of great importance, for example the 2dF galaxy\\cite{2dFGRS01} \\& QSO\\cite{2QZ04} redshift surveys and the Sloan Digital Sky Survey spectroscopy component\\cite{SDSS00}. Another important application is integral field spectroscopy. The first instruments using fibre-based integral field units were constructed in the late 1980s\\cite{MF40,Silfid,DensePak}, and this is now a well established technique employed by a number of major optical and near infrared instruments (e.g.\\ GMOS\\cite{GMOSIFU}, VIMOS\\cite{VIMOS}, IMACS\\cite{IMACS}, CIRPASS\\cite{CIRPASS}). Optical fibres are also in use in astronomical optical interferometers \\cite{ShaklanRoddier88}.  The importance of optical fibres in astronomy is set to increase even further with the continuing development of adaptive optics (AO). AO systems are now in place on all the world's largest astronomical telescopes, and advanced multi-conjugate adaptive optics (MCAO) systems are planned for both the current generation of telescopes and all of the proposed Extremely Large Telescopes (ELTs).  These MCAO systems will provide large AO-corrected fields of view, and the very large number of spatially resolved elements will make efficient sampling of the focal plane essential\\cite{BlandHawthorn06}.  This will require new developments in multi-object spectroscopy systems, deployable integral field units and deployable imaging systems, and optical fibres are well suited to play an important role in all of these. Consequently optical fibres, and in particular optical fibres combined with adaptive optics systems, are extremely important to the future of astronomy.  The majority of fibre-fed instruments so far have been designed for operation under natural seeing conditions, and the relative ease of coupling light into multi-mode fibres (MMFs), especially in the presence of atmospheric aberrations, has led to their exclusive use in instruments to date.  Developments in astrophotonics now provide a strong motivation to move away from MMFs, however.  Astrophotonics is a broad term used for the astronomical application of a wide range of photonic technology, and this burgeoning field has the potential to revolutionise astronomy.  Important examples which is likely to have a significant impact in the near future are integrated photonic spectrographs\\cite{OtherPaper} and OH suppression fibres based on aperiodic fibre Bragg gratings\\cite{OHsuppr} (AFBGs), which promise to give near infrared instruments sky backgrounds as low or possibly even lower than those seen in the optical. These new devices are capable of greatly benefitting astronomy, however a significant obstacle to realising this potential is the fact that many have been conceived as single-mode devices, meaning that they cannot be fed by conventional MMFs.  The most obvious way to integrate a single-mode photonic device into an instrument is to feed it with a single-mode fibre (SMF), but this approach has its own difficulties. Shaklan \\& Roddier\\cite{ShaklanRoddier88} and Coud\\'e du Foresto et al\\cite{CoudeduForesto00} have shown that the theoretical maximum efficiency with which a stellar image can be coupled into a single mode fibre is $\\sim 80\\%$ in the absence of any atmospheric turbulence effects or obstructions in the telescope pupil.  When the effect of a reasonable circular central obstruction of 20\\% of the primary diameter is included this falls to $\\sim 70\\%$, and the presence of atmospheric aberrations further reduces the coupling efficiency in proportion to the Strehl ratio.  As a result direct coupling of large telescopes to SMFs in natural seeing is rendered impractical by low efficiencies, while the use of SMFs with AO places strong constraints on the necessary performance of the AO system, both in terms of the average Strehl achieved and its variability (which impacts on calibration).  While the continuing development of AO may allow highly efficient and stable coupling of telescopes to SMFs in the future there is an intermediate approach which should allow the efficient integration of astrophotonic devices now. Though many important devices are single-mode it is in general possible to extend them to operate with a few propagating modes.  With OH suppressing fibres, for example, the atmospheric emission lines must be blocked separately for each propagating mode, which can be achieved either with a single, more complex AFBG or by using converters to connect to multiple single-mode AFBGs\\cite{Converter}. An integrated spectrograph can also be made to work with a few modes\\cite{OtherPaper}, at least at low and moderate resolutions.  Using modified astrophotonic devices such as these makes it is possible to use few mode fibres (FMFs) instead of SMFs.  The coupling of light into FMFs is relatively unexplored territory, however as the number of modes increases it will become easier to couple light into the fibres, which is expected to reduce the sensitivity to AO system performance at the expense of increasing the required complexity of attached astrophotonic devices.  Indeed for some, such as OH suppression fibres, it will be practical to use sufficiently many modes to allow use under natural seeing.  In this paper we describe an ongoing investigation into the trade off between coupling efficiency and the number of propagating modes. ",
        "conclusions": "Astrophotonic developments such as OH suppression fibres and integrated photonic spectrographs have enormous potential, however efficiently integrating them into an astronomical instrument presents a challenge. The MMFs conventionally used in astronomy, while efficient at accepting starlight, are not suitable for feeding light into these devices as the devices are not able to accept a large number of modes.  SMFs, on the other hand, while ideal for feeding light into astrophotonic devices are difficult to couple starlight into, even with adaptive optics.  We have begun an investigation into the intermediate territory of FMFs, in order to find the best compromise between the two extremes.  Our initial results on diffraction limited fibre coupling have shown that FMFs exhibit many of the desirable properties of MMFs even when there are only of order 10 guided modes.  For example, FMFs offer higher maximum coupling efficiency than SMFs ($>90\\%$), especially for extended sources.  Also FMFs are less sensitive to the effects of obstructions in the telescope pupil than SMFs are.  Unlike SMFs, FMFs can efficiently couple light over a range of focal ratios from $F_{\\rm min} \\approx 1/(2.\\rm{NA})$ up to a maximum value determined by matching the size of the image to the fibre core.  FMFs are also tolerant of displacement of the image centre from the fibre axis, provided the image remains within the fibre core.  Both SMFs and FMFs exhibit little sensitivity to wavelength.  While these results are encouraging the use of perfect diffraction limited images does represent an idealised case.  In any real ground based telescope the adaptive optics correction will be imperfect, and the residual atmospheric wavefront perturbations will effect the coupling efficiency. It is known that SMF coupling is highly sensitive to imperfect correction, with the coupling efficiency declining in proportion to the Strehl ratio\\cite{CoudeduForesto00}, however the corresponding dependency for FMFs has not yet been investigated. At the time of writing simulated partially corrected atmospheric phase screens were being included in the model system to investigate the effects of various levels of aberrations on fibre coupling performance. The aim of this work is to determine the dependence of FMF coupling efficiency on the order of correction, from natural seeing to the diffraction limit, and thereby establish the number of modes required for acceptable throughput levels under a range of realistic usage conditions.  Preliminary results have also been obtained for pupil-plane coupling to FMFs, which show a similar rapid convergence on MMF behaviour above $\\sim10$ modes as the image-plane results discussed here.  This work will be extended to model lenslet arrays of various geometries and both image and pupil-plane coupling with a view to determining the best approach for FMF integral field spectroscopy."
    },
    "0606/astro-ph0606258_arXiv.txt": {
        "abstract": "Several astrophysics and nuclear physics applications require the detection of photons in the energy range of keV up to several MeV with good position and energy resolution. For certain applications Cadmium Zinc Telluride (CZT) detectors might be the detector option of choice. Up to now, CZT detectors have mainly been used in the energy range between a few keV and $\\sim$1~MeV. They operate at room temperature and achieve excellent  position resolution and substantially better energy resolution than scintillation detectors. Furthermore, they can be built more compact and more economically than Ge detectors and do not require cryogenic cooling.  In this paper, we describe the results of 3-D Monte Carlo simulations of a ``CZT calorimeter'' that can be used to detect  photons in the keV to several MeV range. The main objective of these studies is to evaluate the feasibility of CZT calorimeters, to study their performance and detect and understand performance limiting factors. Such a calorimeter consists of many layers of closely packed pixellated CZT detector units.  Our simulations of single detector units reproduce experimental results, indicating that our simulations capture the main factors that limit the performance of a detector unit.  For a full calorimeter the limiting factors within a range from $\\sim$20~keV to $\\sim$10~MeV are: a) the fact, that the incident energy is not totally deposited within the detector area because secondary particles leave the detector against the direction from which the incident radiation enters, b) signal loss when the interaction is near the pixel edges and near the anodes. In this case signals which are induced in neighboring pixels are discarded when their intensities lie below the trigger threshold. c) the steep weighting potential gradient close to the anodes, which affects about 0.25~cm next to the anode and impairs there the correction of the depth of interaction (DOI). This effect dominates in thin detectors (0.5~cm).\\\\  Understanding the limiting factors we come to the conclusion that 1~cm to 1.5~cm thick detector units can be used to build a calorimeter with good performance over the energy range from $\\sim$20~keV to $\\sim$10~MeV . ",
        "introduction": "\\label{Introduction} Cadmium Zinc Telluride (CZT) has emerged as the detector material of choice for the detection of hard X-rays and soft $\\gamma$-rays when excellent position and energy resolution is needed and cryogenic cooling is impractical. Most commonly, CZT detectors are employed to detect photons in the 10 keV to $\\sim$1 MeV energy range where its high density ($\\rho\\,\\approx$ 5.76 g/cm$^3$) and high average atomic number ($\\simeq$50) result in high stopping power and a large cross section for photoelectric interactions. Several astrophysical, nuclear physics, and homeland security applications require the detection of photons with energies of several MeV with good position and energy resolution and with a detection efficiency close to 100\\%.  A detector built from closely packed CZT detectors units may perform substantially better than a scintillation detector and may be more compact and more economic than a cryogenically cooled Ge detector. So far CZT detectors have mainly been studied in the photoelectric regime (10 keV-$\\sim$300 keV). Two different designs have been widely used, pixellated  detectors and single sided strip detectors \\cite{He2005,Matt,Macri,DOE}. Excellent energy resolutions have been achieved with both designs. For a pixellated detector with 11 $\\times$ 11 pixels and a steering grid Zhang et al. \\cite{He2005} reported 0.8\\% (full width half maximum, FWHM) at 662 keV for single pixel events and 2.3\\% at 662 keV (FWHM) for three pixel events. The detector size was 1.5  $\\times$ 1.5 $\\times$ 1 cm$^3$. For a single sided charge-sharing CZT strip detectors FWHM energy resolutions of 19.5 keV at 662 keV  and 23.7 keV ar 1333 keV were reported \\cite{DOE}.  Whereas detectors with N $\\times$ N pixels require N$^2$ readout channels CZT cross-strip detectors require 2 $\\times$ N channels. The penalty for this advantage are ambiguities in matching strips with the right cross-strip counterparts in case of multiple interactions in one detector. In order to achieve good detector efficiency, energy resolution and imaging properties in the  ``Compton'' regime of several hundred keV requires to take multiple pixel or multiple strip events into account. Strategies exist to ameliorate this problem somewhat, for further details \\cite{DOE}. This paper deals with pixellated CZT detectors only. We use  the Geant 4.0 code \\cite{G4}  to study the performance of very thick detectors ($\\gg$ 10 cm) built from closely packed CZT detector units with the objective of evaluating the theoretically achievable performance of such CZT calorimeters given the electronic properties of present-day CZT substrates. As we will describe further below, the main thrust of our study is to evaluate general performance limitations arising from the combination of the location and spatial extent of the charge clouds below individual detector pixels, trapping of electrons and holes and the influence of the weighting potential. We neglect the performance limitations arising from Te-precipitates and other crystal defects. The good agreement of the simulated detector response and the one experimentally measured shows that electron and hole trapping and the detectors weighting potential indeed limits the performance of current CZT detectors. Te-precipitates and other crystal defects seem to be only important for underperforming detector areas.  The calorimeter performance depends on the shape of the 3-D weighting potential and on the number,  energy and distribution of secondary particles produced in the calorimeter. Simulations with a particle interaction code like Geant 4 and a device simulation code is the most efficient and most accurate way to estimate the performance of such calorimeters.  This paper is structured as follows. In Section \\ref{AstroNuclAppl}, we give a brief description of astrophysics and nuclear physics experiments that may use CZT calorimeters. Subsequently, in Section  \\ref{CZTSimulation} we describe the simulations and present their results for single detector units and for full CZT calorimeters. In Sect. \\ref{con} the results are discussed.  In the following all energy resolutions are given as Full-Width-Half-Maximum energy resolutions. ",
        "conclusions": "\\label{con} In our simulations we have considered weighting potential, electron and hole trapping. The results of a single detector unit are in good agreement with experimental data. Therefore, precipitates and crystal inhomogeneities play a major role only for underperforming detector regions.  The studies in this paper show, that CZT has to be seriously considered as a detector material  for the design of future calorimeters which operate in the range from several~keV up to several MeV. We have  simulated three different calorimeter designs varying the thickness of the individual detector and the number of layers, keeping constant the size of the full calorimeter constant (10.33~$\\times$~10.33~$\\times$~45~cm$^3$).  Calorimeters built out of 1.0~cm and 1.5~cm thick units exhibit almost identical performance with energy resolutions of about 1.7\\% at 662~keV and about 1.9\\% at 10 MeV. The energy resolution is mainly limited by the energy loss owing to secondary particles leaving the detector upward against the direction from which the incident radiation enters the detector and by  the effects of the steep change of the weighting potential close to the anode.   Going to higher energies the resolution is also negatively influenced by the increasing number of events which  have several interactions below one pixel because in these cases DOI corrections are not well defined.  The principal result of our study is the fact that a 0.25~cm thick region below the anodes of the detector units exhibits a substantially poorer energy resolution than the rest of the detector volume.  Reducing the portion of this low resolution volume increases the overall performance. A good calorimeter requires that the single detector unit is much thicker than the pixel pitch and of course much thicker than passive material. (mounting, cables and front-end electronics).  Therefore, calorimeters built with 1.0~cm or 1.5~cm thick CZT crystals have significantly better resolution than those with 0.5~cm thin detector units.  The energy resolution of a CZT calorimeter is shown in Figure \\ref{EffizienzFWHM}. The resolution is limited  a) by the steep weighting potential gradient close to the anodes, which dominates in thin detectors and impairs the correction of the depth of interaction (DOI), b) by the fact, that the incident energy is not totally deposited within the detector area because secondary particles leave the detector upward against the direction from which the incident radiation enters the detector c) because signals induced in neighboring pixels when the interaction is near the pixel edges and near the anodes with intensity below the trigger threshold are lost for signal reconstruction leading to systematically reduced signals.   We found that the energy resolution of a CZT calorimeter will lie between that of CsI-detectors and Ge detectors. In terms of spatial resolution solid state Ge and CZT detectors are of course much better than scintillation detectors. Where compactness and weight are issues as in a next generation space-born $\\gamma$-ray telescope, Cadmium Zinc Telluride (CZT) detectors might be the detector option of choice."
    },
    "0606/astro-ph0606402_arXiv.txt": {
        "abstract": "% We present observations of V838~Mon and its close vicinity in the three lowest rotational transitions of CO. The $J$ = 2$\\to$1 and 3$\\to$2 data were obtained using the 3 m KOSMA telescope. They include on-the-fly maps covering a large area ($\\sim$3.4 sq. deg.) around V838~Mon and long integrations on the star position. Complementary observations in the CO $J$ = 1$\\to$0 transition were obtained using the 13.7 m Delingha telescope. The star position as well as 25 other points preselected in the near vicinity of the object have been measured in this transition.  We report on a detection of two narrow emission components in $J$ = 2$\\to$1 and 3$\\to$2 transitions at the position of V838~Mon. Lines were found at radial velocities of $V_{\\rm lsr}=53.3$~km~s$^{-1}$ and $V_{\\rm lsr}=-11.0$~km~s$^{-1}$. Their origin is unclear. We also shortly discuss results of the observations of the vicinity of V838~Mon. ",
        "introduction": "The enigmatic eruption of V838 Mon, followed by its spectacular light echo, triggered research in different fields of astrophysics. Apart from studies of the evolution of the object, investigations of the circumstellar and interstellar neighborhood of the star can also be important for understanding the nature of the event. We present observations undertaken to search for molecular matter in the vicinity of V838 Mon.  Using the results of the CO $J$~=~1$\\to$0 galactic survey of \\citet{dame} \\citet{loon} have suggested that V838 Mon is situated in a bubble of CO emission of a diameter of about 1$^{\\circ}$ (see Fig. 1). According to these authors the structure is of circumstellar origin due to the AGB activity of the V838 Mon progenitor. Several critical points against this finding and interpretation have been raised in \\citet{tyl}. However, the structure in the CO map reported in \\citet{loon} is quite suggestive and we have found important to obtain CO observations of the same ragion with a better sensitivity and angular resolution  There are also other important reasons for observing the star and its vicinity in molecular lines. These are the nature of the echoing matter and a search for matter lost during and after the 2002 outburst. The detection and monitoring of the SiO maser emission from V838~Mon \\citep{deguchi,claus} shows that a molecular activity started close to the star. Complementary observations in molecular thermal transitions would be important to better understand what is going on there.  \\begin{figure} \\includegraphics[scale=0.6, origin=r, trim = 0 0 0 100]{kaminski_fig1.ps} \\caption{A $3^{\\circ}.5\\times3^{\\circ}.5$ map centered at the V838~Mon position from the CO~(1--0) galactic survey \\citep[data taken from {\\it SkyView} -- {\\tt \\small http://skyview.gsfc.nasa.gov/};][]{dame}. The star symbol marks V838 Mon position. The polygon drawn with a solid line defines the region observed in the on-the-fly mode in the CO (2--1) and (3--2) transitions. Triangles mark the positions observed in the CO (1--0) transition. Dashed line indicates the Galactic equator.} \\end{figure} ",
        "conclusions": ""
    },
    "0606/astro-ph0606687_arXiv.txt": {
        "abstract": "We have compiled a complete extragalactic sample based on $\\sim25,000 \\rm \\, deg^2$ to a limiting flux of $3 \\times 10^{-11} \\rm \\, ergs \\, cm^{-2} \\, s^{-1}$ ($\\sim7,000 \\rm \\, deg^2$ to a flux limit of $10^{-11} \\rm \\, ergs \\, cm^{-2} \\, s^{-1}$) in the 20 -- 40 keV band with {\\it INTEGRAL}. We have constructed a detailed exposure map to compensate for effects of non-uniform exposure. The flux-number relation is best described by a power-law with a slope of $\\alpha = 1.66 \\pm 0.11$. The integration of the cumulative flux per unit area leads to $f_{20 - 40 \\rm \\, keV} = 2.6 \\times 10^{-10} \\rm \\, ergs \\, cm^{-2} \\, s^{-1} \\, sr^{-1}$, which is about 1\\% of the known $20 - 40 \\rm \\, keV$ X-ray background. We present the first luminosity function of AGN in the 20--40 keV energy range, based on 38 extragalactic objects detected by the imager IBIS/ISGRI on-board {\\it INTEGRAL}. The luminosity function shows a smoothly connected two power-law form, with an index of $\\gamma_1 = 0.8$ below, and $\\gamma_2 = 2.1$ above the turn-over luminosity of $L_* = 2.4 \\times 10^{43} \\rm \\, ergs \\, s^{-1}$. The emissivity of all {\\it INTEGRAL} AGNs per unit volume is $W_ {20 - 40 \\rm \\, keV}(> 10^{41} \\rm \\, ergs \\, s^{-1}) = 2.8 \\times 10^{38} \\rm \\, ergs \\, s^{-1} \\, h_{70}^3 \\, Mpc^{-3}$. These results are consistent with those derived in the $2 - 20 \\rm \\, keV$ energy band and do not show a significant contribution by Compton-thick objects. Because the sample used in this study is truly local (${\\bar z} = 0.022$), only limited conclusions can be drawn for the evolution of AGNs in this energy band. ",
        "introduction": "The Galactic X-ray sky is dominated by accreting binary systems, while the extragalactic sky shows mainly active galactic nuclei (AGN) and clusters of galaxies. Studying the population of sources in X-ray bands has been a challenge ever since the first observations by rocket borne X-ray detectors (Giacconi et al. 1962). At soft X-rays (0.1 -- 2.4 keV) deep exposures by {\\it ROSAT} have revealed an extragalactic population of mainly broad line AGNs, such as type Seyfert 1 and quasars \\citep{Hasinger98,ROSATdeep}. In the 2 - 10 keV range surveys have been carried out with {\\it ASCA} (e.g. Ueda et al. 2001), {\\it XMM-Newton} (e.g. Hasinger 2004), and {\\it Chandra} (e.g. Brandt et al. 2001) and have shown that the dominant extragalactic sources are more strongly absorbed than those within the {\\it ROSAT} energy band. For a summary on the deep X-ray surveys below 10 keV see Brandt \\& Hasinger (2005). At higher energies the data become more scarce. Between a few keV and $\\sim 1 \\rm \\, MeV$, no all-sky survey using imaging instruments has been performed to date. The {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) sky survey in the $3 - 20 \\rm \\, keV$ energy band revealed about 100 AGNs, showing an even higher fraction of absorbed ($N_H > 10^{22} \\rm \\, cm^{-2}$) sources of about 60\\% \\cite{RXTENGC}. The {\\it International Gamma-Ray Astrophysics Laboratory} ({\\it INTEGRAL}; Winkler et al. 2003) offers an unprecedented $>20 \\rm \\, keV$ collecting area and state-of-the-art detector electronics and background rejection capabilities. Notably, the imager IBIS with an operating range from $20 - 1000 \\rm \\, keV$ and a fully-coded field of view of $10^\\circ \\times 10^\\circ$ enables us now to study a large portion of the sky. A first catalog of AGNs showed a similar fraction of absorbed objects as the {\\it RXTE} survey \\cite{intagn}. The Burst Alert Telescope (BAT) of the {\\it Swift} mission \\cite{Swift} operates in the 15 -- 200 keV band and uses a detector similar to IBIS/ISGRI, but provides a field of view about twice the size. The BAT data of the first three months of the mission provided a high galactic latitute survey, including 44 AGNs \\cite{BATsurvey}. Within this sample a weak anti-correlation of luminosity versus intrinsic absorption was found as previously found in the $2-10 \\rm \\, keV$ band \\citep{Ueda03,LaFranca}, revealing that most of the objects with luminosities $L_X > 3 \\times 10^{43} \\rm \\, ergs \\, s^{-1}$ show no intrinsic absorption. Markwardt et al. (2005) also pointed out that this luminosity corresponds to the break in the luminosity function.  Related to the compilation of AGN surveys in the hard X-rays is the question of what sources form the cosmic X-ray background (CXB). While the CXB below 20 keV has been the focus of many studies, the most reliable measurement in the 10 - 500 keV has been provided by the {\\it High Energy Astronomical Observatory} ({\\it HEAO 1}), launched in 1977 (Marshall et al. 1980). The most precise measurement provided by the UCSD/MIT Hard X-ray and Gamma-Ray instrument ({\\it HEAO 1} A-4) shows that the CXB peaks at an energy of about $30 \\rm \\, keV$ (Marshall et al. 1980, Gruber et al. 1999). The isotropic nature of the X-ray background points to an extragalactic origin, and as the brightest persistent sources are AGNs, it was suggested early on that those objects are the main source of the CXB (e.g. Setti \\& Woltjer 1989). In the soft X-rays this concept has been proven to be correct through the observations of the {\\it ROSAT} deep X-ray surveys, which showed that $90 \\%$ of the $0.5 - 2.0 \\rm \\, keV$ CXB can be resolved into AGNs (Schmidt et al. 1998). At higher energies ($2 - 10 \\rm \\, keV$), {\\it ASCA} and {\\it Chandra} surveys measured the hard X-ray luminosity function (XLF) of AGNs and its cosmological evolution. These studies show that in this energy range the CXB can be explained by AGNs, but with a higher fraction of absorbed ($N_H > 10^{22} \\rm \\, cm^{-2}$) objects than in the soft X-rays (e.g. Ueda et al. 2003). A study based on the {\\it RXTE} survey by Sazonov \\& Revnivtsev (2004) derived the local hard X-ray luminosity function of AGNs in the 3--20 keV band. They showed that the summed emissivity of AGNs in this energy range is smaller than the total X-ray volume emissivity in the local Universe, and suggested that a comparable X-ray flux may be produced together by lower luminosity AGNs, non-active galaxies and clusters of galaxies. Using the {\\it HEAO 1}-A2 AGNs, Shinozaki et al. (2006), however, obtained a local AGN emissivity which is about twice larger than the value of Sazonov \\& Revnivtsev (2004) but consistent with the estimates by Miyaji et al. (1994) which was based on the cross-correlation of the {\\it HEAO 1}-A2 map with {\\it IRAS} galaxies.  With the on-going observations of the sky by {\\it INTEGRAL}, a sufficient amount of data is now available to derive the AGN hard X-ray luminosity function. In this paper we present analysis of recent observations performed by the {\\it INTEGRAL} satellite, and compare the results with previous studies. In Section 2 we describe the AGN sample and in Section 3 the methods to derive the number-flux distribution of {\\it INTEGRAL} AGNs are presented together with the analysis of their distribution. Section 4 shows the local luminosity function of AGNs as derived from our data, followed by a discussion of the results in Section 5. Throughout this work we applied a cosmology with $H_0 = 70 \\rm \\, km \\, s^{-1} \\, Mpc^{-1}$ ($h_{70} = 1$), $k = 0$ (flat Universe), $\\Omega_{matter} = 0.3$, and $\\Lambda_0 = 0.7$, although a $\\Lambda_0 = 0$ and $q_0 = 0.5$ cosmology does not change the results significantly because of the low redshifts in our sample. ",
        "conclusions": "The extragalactic sample derived from the {\\it INTEGRAL} public data archive comprises 63 low redshift Seyfert galaxies ($\\langle z \\rangle = 0.022 \\pm 0.003$) and 8~blazars in the hard X-ray domain. Two galaxy clusters are also detected, but no star-burst galaxy has been as yet. This {\\it INTEGRAL} AGN sample is thus the largest one presented so far. 38 of the Seyfert galaxies form a complete sample with significance limit of $5 \\sigma$.  The number flux distribution is approximated by a power-law with a slope of $\\alpha = 1.66 \\pm 0.11$. Because of the high flux limit of our sample the objects account in total for less than $1\\%$ of the $20 - 40 \\rm \\, keV$ cosmic X-ray background. The emissivity of all AGNs per unit volume $W_ {20 - 40 \\rm \\, keV}(> 10^{41} \\rm \\, ergs \\, s^{-1}) = 2.8 \\times 10^{38} \\rm \\, ergs \\, s^{-1} \\, h_{70}^3 \\, Mpc^{-3}$ appears to be consistent with the background estimates in the 2--10 keV energy band based on the cross-correlation of the {\\it HEAO 1}-A2 map with {\\it IRAS} galaxies \\cite{Miyaji94}.  The luminosity function in the $20 - 40 \\rm \\, keV$ energy range is consistent with that measured in the $2 - 20 \\rm \\, keV$ band. Below the turnover luminosity of $L_* = 2.4 \\times 10^{43} \\rm \\, ergs \\, s^{-1}$ the absorbed AGNs become dominant over the unabsorbed ones. The fraction of Compton thick AGNs with known intrinsic absorption is found to be small ($8\\%$) in our AGN sample. For the sources without reliable absorption information we derived an estimate from the comparison with {\\it ROSAT} All-Sky Survey data and find that the data do not require additional Compton thick objects within the sample presented here. It has to be pointed out though, that the sources without RASS counterpart could be Compton thick which would increase the ratio of this source type to 13\\% in the complete sample. Evolution of the source population can play a major role in the sense that the fraction of absorbed sources among AGNs might be correlated with redshift, as proposed for example by Worsley et al. (2005).  Over the life time of the {\\it INTEGRAL} mission we expect to detect of the order of 200 AGNs. Combining these data with the studies based on {\\it Swift}/BAT, operating in a similar energy band as IBIS/ISGRI, will further constrain the hard X-ray luminosity function of AGNs. But we will still be limited to the relatively high flux end of the distribution. Because of this {\\it INTEGRAL} and {\\it Swift}/BAT will most likely not be able to test evolutionary scenarios of AGNs and thus will be inadequate to explain the cosmic X-ray background at $E > 20 \\rm \\, keV$. Future missions with larger collecting areas and/or focusing optics will be required to answer the question of what dominates the Universe in the hard X-rays."
    },
    "0606/astro-ph0606364_arXiv.txt": {
        "abstract": "First discovered in the Magellanic Clouds and in the Milky Way, the largest pools of luminous supersoft X-ray sources (SSSs) now known lie in M31 and in more distant galaxies. Hundreds of newly-discovered SSSs are helping us to test models for Type~Ia supernovae and to identify SSSs that may represent a wider range of physical systems, including accreting intermediate-mass black holes.  In this short report we list ten intriguing facts about distant SSSs. ",
        "introduction": "\\vspace{-.3 true in} Luminous supersoft X-ray sources (SSSs) were discovered and defined in terms of properties observed in a small number of sources ($\\sim 18$) in the Galaxy and Magellanic Clouds (MCs). Specifically, SSSs are defined in terms of their estimated luminosities ($L > 10^{36}$ erg s$^{-1}$) and their broad band spectra, with little or no emission above $1$ keV.  The physical nature of SSSs is not yet determined. In fact, the studies we summarize below indicate that they are likely to be a diverse group, with several different types of physical systems observable as SSSs, including white dwarfs (WDs), neutron stars (NSs), and black holes (BHs). Perhaps this diversity is to be expected, given the simplicity of the  definition.  \\vspace{-.1 true in} SSS effective radii are comparable to those of WDs. Indeed, roughly half of the SSSs in the MCs and Milky Way with optical IDs have counterparts that are consistent with systems known to contain hot WDs: planetary nebulae, recent novae, and symbiotic binaries (Greiner 2000). The remaining SSSs in the Galaxy and MCs have measured periods ranging from a few hours to a few days. The first and best-known model for these binary  was developed by Ed van den Heuvel and collaborators (1992). The close-binary supersoft (CBSS) model, postulates that the prodigious luminosities are generated through the nuclear burning of material accreted by a WD. In order for nuclear burning to occur, the accretion rates must be high (close to or larger than $10^{-7} M_\\odot$  year$^{-1}$). These high rates can be sustained only if the donor star is somewhat more massive than the WD and/or slightly evolved. An interesting characteristic of these models is that they allow the WD to increase its mass. Some SSS binaries may therefore be progenitors of Type~Ia supernovae (Rappaport, Di~Stefano, \\& Smith 1994; Di~Stefano 1996; Kahabka \\& van den Heuvel 1997 and references therein).  While indirect evidence mounts that the WD model applies to some SSSs in the Galaxy and MCs, definite confirmation has been difficult. (See the contribution of Phil Charles to this volume.) One way to learn more about SSSs is to find them in other galaxies. We expect $\\sim 1000$ SSSs to reside in spiral galaxies such as M31. (See Di~Stefano \\& Rappaport 1994, Rappaport, Di~Stefano, \\& Smith 1994, Yungelson et al. 1996.) Thanks to {\\it Chandra} and {\\it XMM-Newton} we have begun to study SSSs in distant galaxies. Below we list some of the things we have learned. ",
        "conclusions": "\\vspace{-.3 true in}  When Ed and his collaborators carried out the first theoretical investigations of the newly established class of SSSs, it was clear that something exciting was going on. Fifteen years later, the study of SSSs in other galaxies is leading to new and unexpected discoveries and challenges. As we begin to develop a statistical sample of SSSs, we are finding that they are a diverse group, and that there are other soft sources, QSSs, that may be related to them. The realms of astrophysics associated with SSSs and QSSs span supermassive BHs, intermediate-mass BHs, X-ray binaries with stellar-remnant accretors, Type~Ia progenitors, and supernova remnants. If, fifteen years hence, we have succeeded in establishing the nature of the sources we detect, and developing theoretical constructs to describe them and the roles they play in the universe, we will have accomplished a good deal.  \\vspace{-.3 true in}"
    },
    "0606/astro-ph0606152_arXiv.txt": {
        "abstract": "We present an analysis of the effects of environment on the photometric properties of galaxies in the core of the Shapley Supercluster at $z=0.05$, one of the most massive structures in the local universe. The Shapley Optical Survey (SOS) comprises archive WFI optical imaging of a 2.0\\,deg$^2$ region containing the rich clusters A3556, A3558 and A3562 which demonstrate a highly complex dynamical situation including ongoing cluster mergers. The \\mbox{$B-R/R$} colour-magnitude relation has an intrinsic dispersion of 0.045\\,mag and is \\mbox{$0.015\\pm0.005$}\\,mag redder in the highest-density regions, indicative of the red sequence galaxy population being 500\\,Myr older in the cluster cores than towards the virial radius. The $B-R$ colours of galaxies are dependent on their environment, whereas their luminosities are independent of the local density, except for the very brightest galaxies \\mbox{(M$_{R}\\!<\\!-22$)}. The global colours of faint \\mbox{($\\ga\\!{\\rm M}^{*}+2$)} galaxies change from the cluster cores where \\mbox{$\\sim\\!90$\\%} of galaxies lie along the cluster red sequence to the virial radius, where the fraction has dropped to just \\mbox{$\\sim\\!20$\\%.} This suggests that processes related to the supercluster environment are responsible for transforming faint galaxies, rather than galaxy merging, which should be infrequent in any of the regions studied here. The largest concentrations of faint blue galaxies are found {\\em between} the clusters, coincident with regions containing high fractions of \\mbox{$\\sim\\!L^{*}$} galaxies with radio emission indicating starbursts. Their location suggests star-formation triggered by cluster mergers, in particular the merger of A3562 and the poor cluster SC\\,1329-313, although they may also represent recent arrivals in the supercluster core complex. The effect of the A3562-SC\\,1329-313 merger is also apparent as a displacement in the spatial distribution of the faint galaxy population from both the centres of X-ray emission and the brightest cluster galaxies for both systems. The cores of each of the clusters/groups are marked by regions that have the lowest blue galaxy fractions and reddest mean galaxy colours over the whole supercluster region, confirming that star-formation rates are lowest in the cluster cores. In the cases of A3562 and SC\\,1329-313, these regions coincide with the centres of X-ray emission rather than the peaks in the local surface density, indicating that ram-pressure stripping may have an important role in terminating any remnant star-formation in galaxies that encounter the dense ICM of the cluster cores. ",
        "introduction": "\\label{intro}  The cluster galaxy population has evolved rapidly over the last 4\\,Gyr \\citep[e.g.,][]{bo78,bo84,dressler94,dressler97,treu,kodama}. Clusters at \\mbox{$z\\!\\ga0.4$} are dominated, particularly at faint magnitudes, by blue spiral galaxies, predominantly irregular or Sc--Sd spirals. Some of these show signs of disturbed morphologies, and many present spectroscopic evidence that they have undergone multiple star-formation events over the last \\mbox{1--2\\,Gyr} \\citep{dressler94}. Conversely, local clusters are completely dominated by passive early-type galaxies: elliptical and lenticular (S0) galaxies at the brighter end, and dwarf spheroids at fainter magnitudes.  In the standard hierarchical cosmological model, \\mbox{$z\\!\\sim\\!0.4$} represents the peak infall rate of field galaxies onto the cluster \\citep{kauffmann95}, and it is the transformation of these infalling field galaxies from star-forming disk-dominated galaxies into passively-evolving spheroids over a period of \\mbox{4--5\\,Gyr} through their encounter with the cluster environment, that produces the observed changes in cluster galaxy populations. Several physical mechanisms related to the cluster environment have been proposed as producing the observed transformations in galaxies, in which interactions with either other cluster galaxies or the hot intra-cluster medium (ICM) affect both their structural and star-formation properties.  To distinguish between these mechanisms it is necessary to examine both where the transformations occur, and how the star-formation and structural properties of the galaxies are changed \\citep[e.g.][]{treu}. For example, ram pressure from the passage of the galaxy through the dense ICM can effectively remove the cold gas supply and thus rapidly terminate new star-formation, either by stripping the gas directly \\citep*{a99}, or by inducing a starburst in which all of the gas is consumed \\citep{fujita}. The most dramatic ICM-galaxy interactions should occur when two clusters merge, as shock fronts created in the ICM may trigger starbursts in galaxies over large scales \\citep*{roettiger}. Importantly, in terms of their environmental effects, these mechanisms all require a dense ICM, and so their evolutionary effects are limited to the cores of clusters.  In contrast, galaxy mergers, which can strongly affect the morphological evolution of disks, cannot occur when the encounter velocities are much greater than the internal velocity dispersion of galaxies \\citep{aarseth}, and so while frequent in small groups, are rare in rich clusters \\citep{ghigna}. Alternatively galaxy harassment, whereby repeated close \\mbox{($<$50\\,kpc)}, high-velocity \\mbox{($>$1000\\,km\\,s$^{-1}$)} encounters with massive galaxies and the cluster's tidal field cause impulsive gravitational shocks that damage the fragile disk of late-type spirals \\citep{moore}, transforming them over a period of several Gyr. Galaxy harassment is effective throughout a cluster, including beyond the virial radius, but its effects should be greater for those clusters with higher velocity dispersions.  Finally, when a galaxy falls into a more massive halo, the diffuse gas in its halo is lost to the ICM, thus preventing further cooling and replenishment of the cold gas supply, ``suffocating'' the galaxy \\citep[e.g.][]{blanton00,diaferio}. Star-formation in the galaxy then declines slowly as the remaining cold gas is used up \\citep*{larson}.  The large datasets provided by the 2dFGRS and SDSS have allowed the environmental effects on galaxy properties to be followed statistically over the full range of environments, from the sparse field to the dense cluster cores \\citep[e.g.][]{lewis,gomez}, at least for the brightest galaxies (\\mbox{${\\rm M}<{\\rm M}^{*}\\!+1$}). They show that star-formation is most closely dependent on local density, and is systematically suppressed above a critical value of density, that is found 3--4 virial radii from clusters, but also in galaxy groups as poor as \\mbox{$\\sigma\\sim100$\\,km\\,s$^{-1}$}. This suppression is observed to be independent of the richness of the structure to which the galaxy is bound \\citep{tanaka}, indicating that mechanisms such as galaxy harassment or ram-pressure stripping are not important for the evolution of bright galaxies. Instead the strongest candidates for driving their transformation are galaxy suffocation and low-velocity encounters, which are effective in both galaxy groups and cluster infall regions.  However, it is not clear if and how this scenario extends to fainter magnitudes, as there has been observed a strong bimodality in the properties of galaxies about a characteristic stellar mass \\mbox{$\\sim3\\times10^{10}{\\rm M}_{\\odot}$} (corresponding to \\mbox{$\\sim{\\rm M}^{*}\\!+1$}), with more massive galaxies predominately passive red spheroids, and less massive galaxies tending to be blue star-forming disks \\citep{kauffmann03}. This bimodality implies fundamental differences in the formation and evolution of giant and dwarf galaxies, and it has been proposed \\citep[e.g.][]{dekel,keres} that these are driven by thermal processes in the gas inflowing from the halo onto the galaxy, with the characteristic mass scale representing the point at which shocks in the halo become stable, heating up the halo gas, and preventing further cooling. Hence if the formation and evolution of giant and dwarf galaxies are so fundamentally different, then they are likely to be affected differently by mechanisms related to the environment. For example, galaxy harassment should be most efficient at transforming low-luminosity late-type galaxies. \\citet{tanaka} find differences in the environmental dependences of faint \\mbox{(${\\rm M}^{*}\\!+1<{\\rm M}<{\\rm M}^{*}\\!+2$)} and bright \\mbox{(${\\rm M}<{\\rm M}^{*}\\!+1$)} galaxy populations, and suggest that faint galaxies are affected by mechanisms related to the structure in which the galaxy is found.  To understand the mechanisms underlying the transformation of faint galaxies requires datasets reaching much fainter luminosities. Crucial discriminators between the different transformation mechanisms are the time- and distance-scales involved: while ram-pressure stripping should rapidly terminate star-formation in a galaxy within \\mbox{$\\la\\! 100$\\,Myr} but requires the dense ICM of the cluster core; a galaxy undergoing suffocation will have its star-formation rate slowly decline over a period of several Gyr. Hence the nature of the transition from regions where the majority of galaxies are star-forming, and those dominated by passive galaxies, will depend strongly on the dominant mechanism involved.  One approach is to use galaxy colours, which can be readily obtained to much fainter magnitudes than spectroscopic star-formation rates, and which through the use of models can be directly related to star-formation histories with minimal assumptions \\citep[e.g.][]{bruzual}. Recent large datasets have shown that the bimodality of galaxies is also manifested through their broadband photometry, in particular a separation can be made on the basis of colour into red and blue galaxy populations \\citep{strateva,blanton03}, which correspond roughly to the two broad types previously known from their morphological and spectroscopical characteristics: passively-evolving early-type and star-forming late-type galaxies. This bimodality has been further quantified, resulting in colour-magnitude (C-M) relations and for both the red and blue galaxy populations \\citep{baldry}, and its evolution observed to \\mbox{$z\\sim1$} \\citep{bell}. \\citet{balogh04} show that the bimodal galaxy colour distribution is strongly dependent on environment, with the fraction of galaxies in the red distribution at a fixed luminosity increasing from 10--30\\% in the lowest density environments, to \\mbox{$\\sim$70\\%} at the highest densities.  The most dramatic effects of environment on galaxy evolution should occur in superclusters, where the infall and encounter velocities of galaxies are greatest \\mbox{($>$1000\\,km\\,s$^{-1}$)}, groups and clusters are still merging, and significant numbers of galaxies will be encountering the dense ICM of the cluster environment for the first time.  With this in mind we are undertaking the Shapley Optical Survey (SOS), an optical photometric study of the core region of the Shapley supercluster \\citep{shapley}, one the most massive structure in the local universe, containing as many as 25 Abell clusters.  In this paper we examine the effect of the supercluster environment on the star-formation histories of galaxies as measured through their galaxy colours. We present the SOS in Section~\\ref{data}, and then describe how we quantify the local environment and statistically subtract the field galaxy population in Sections~\\ref{den} and~\\ref{sub}. We present our analysis of the C-M relation in Section~\\ref{cmrelation}, which then allows us to separate the red and blue galaxy populations whose spatial distributions are presented in Section~\\ref{redandblue}. We examine the environmental dependencies on galaxy colours in Section~\\ref{colours}, and discuss our findings in Section~\\ref{discussion}, before presenting our conclusions in Section~\\ref{conclusions}. Throughout the paper we adopt a cosmology with \\mbox{$\\Omega_M=0.3$}, \\mbox{$\\Omega_\\Lambda=0.7$} and \\mbox{H$_{0}$=70\\,km\\,s$^{-1}$Mpc$^{-1}$}. According to this cosmology 1 arcmin corresponds to 60\\,kpc at \\mbox{$z=0.048$}. ",
        "conclusions": "\\label{conclusions} The area covered by the SOS dataset lies mostly within one virial radius \\mbox{(1--1.5\\,$h_{70}$\\,Mpc)} of one of the clusters. The results presented here and in Paper I show that the global properties of faint galaxies change significantly from the cluster cores to the virial radius, both in terms of the luminosity function, and the mean galaxy colours, which indicates that these galaxies are being transformed by processes related to the supercluster environment. As galaxy mergers should be infrequent in any of the environments covered by the SOS, the finding of such large changes in the mean galaxy colour or fraction of faint blue galaxies within the SOS, indicates that some process other than merging must be responsible. In paper I, we suggested that galaxy harassment is important for shaping the galaxy luminosity function at magnitudes fainter than \\mbox{$\\sim M^{*}+1$}, and here we find additional evidence in favour of faint galaxies being transformed by ram-pressure stripping and undergoing starbursts triggered by shocks in the ICM produced by cluster mergers.  These results indicate that the effect of environment on faint \\mbox{($\\ga M^{*}+1$)} galaxies is quite different from that observed for bright galaxies \\citep{gomez,lewis}. While bright galaxies appear to be transformed by processes that can take place well outside the virial radius, that is galaxy merging or suffocation, we find here that many faint galaxies are affected by their interaction with the supercluster environment, although we cannot rule out that the faint galaxies are also transformed outside the virial radius, beyond the limits of our survey. It is also possible that the observed environmental trends are partly due to a population of primordial, early-type galaxies that formed preferentially in the high-density regions that later became clusters \\citep[e.g.][]{poggianti05}. The differences in the environmental trends of bright and faint galaxies are likely to be related to the observed global bimodality in galaxy properties \\citep[e.g.][]{kauffmann03} about a stellar characteristic stellar mass of \\mbox{$\\sim3\\times10^{10}\\,{\\rm M}_{\\odot}$}. A possible explanation could lie within the context of the hot and cold flow model of galaxy evolution \\citep{dekel,keres}. At low masses, galaxies are able to merge and still retain their gas supply, and hence mergers would have little effect on star-formation or galaxy colours. In contrast, once galaxies merge to become more massive than a characteristic mass, shocks in the halo become stable, and preventing further cooling of gas from the halo, bringing about a transformation in the galaxy star-forming properties.  In a future article, we plan to compare our observational results with predictions from semi-analytical models of galaxy evolution applied to n-body simulations of a region containing a supercluster constrained to match the dynamical structure of the Shapley supercluster region."
    },
    "0606/astro-ph0606708_arXiv.txt": {
        "abstract": "A hypothetical time-variation of the gravitational constant $G$ would cause neutron star matter to depart from beta equilibrium, due to the changing hydrostatic equilibrium. This forces non-equilibrium beta processes to occur, which release energy that is invested partly in neutrino emission and partly in heating the stellar interior. Eventually, the star arrives at a stationary state in which the temperature remains nearly constant, as the forcing through the change of $G$ is balanced by the ongoing reactions. Comparing the surface temperature of the nearest millisecond pulsar, PSR~J0437-4715, inferred from ultraviolet observations, with our predicted stationary temperature, we estimate two upper limits for this variation: (1) $|\\dot G/G| < 2 \\times 10^{-10}$ yr$^{-1}$, if we allow direct Urca reactions operating in the neutron star core, and (2) $|\\dot G/G| < 4 \\times 10^{-12}$ yr$^{-1}$, considering only modified Urca reactions. Both results are competitive with those obtained by other methods, with (2) being among the most restrictive. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606222_arXiv.txt": {
        "abstract": "{ Most statistical tools used to characterize the complex structures of the interstellar medium can be related to the power spectrum, and therefore to the Fourier amplitudes of the observed fields. To tap into the vast amount of information contained in the Fourier phases, one may consider the probability distribution function (PDF) of phase increments, and the related concepts of phase entropy and phase structure quantity. We use these ideas here with the purpose of assessing the ability of radio-interferometers to detect and recover this information. By comparing current arrays such as the VLA and Plateau de Bure to the future ALMA instrument, we show that the latter is definitely needed to achieve significant detection of phase structure, and that it will do so even in the presence of a fair amount of atmospheric phase fluctuations. We also show that ALMA will be able to recover the actual ``amount'' of phase structure in the noise-free case, if multiple configurations are used. ",
        "introduction": "The physics of the interstellar medium (ISM) stands at the crossroads of many astrophysical problems, from stellar formation to galaxy evolution.  Without a proper understanding of the processes taking place in the ISM, and of their interplay, complete and satisfactory solutions to these problems cannot hope to be met.  Turbulence is one such process \\citep[see e.g.][]{spicker88,odell87,miesch94}, and it is thought to play a major role in the shaping of the fractal structures observed \\citep{falgarone91,vogelaar94,elmegreen96,falgarone98a,stutzki98,elmegreen2001}.  Consequently, a quantitative description of these structures is a necessary first step towards understanding the physics of the ISM, and many statistical tools have been used to this end. Let us mention the power spectrum \\citep[see e.g.][]{gautier92,dickey2001,stanimirovic2001}, the autocorrelation function \\citep{kleiner85,perault86}, the $\\Delta$-variance \\citep{stutzki98,bensch2001}, the fractal dimension \\citep{falgarone91} and the wavelet decomposition \\citep{gill90}.  These various tools are not altogether independent from one another. By definition, the autocorrelation function is the Fourier transform of the power spectrum, to which the $\\Delta$-variance and fractal dimension can also be related, albeit less directly \\citep{stutzki98}. Finally, the $\\Delta$-variance can be written as the variance of wavelet transform coefficients \\citep{zielinsky99}. On the whole, it is then fair to say that all of these tools are connected, in one way or another, to the power spectrum, although some of them are of easier and more reliable use depending on the type of observation \\citep{stutzki98,bensch2001}.  Since the power spectrum is given by the squared amplitudes of Fourier components, it basically ignores any structural information that may be contained in the Fourier phases. Now, each Fourier component corresponds to a plane wave in direct space, with a given wave vector, amplitude and phase. The Fourier transform being linear, the observed structures are the result of the interaction between the various plane waves. Ignoring the phases when characterizing the structures is thus comparable to ignoring the interference phenomenon, and therefore marks a major loss in structural information. This has been confirmed by simple numerical experiments \\citep{juvells91,coles2005}.  In the experiment performed by \\citet{coles2005}, the Fourier phases of a numerical simulation of galaxy clustering, which is a highly-structured field, are randomly reshuffled in Fourier space. The resulting field has lost most of the filamentary structure observed in the original image. This shows that it is in the Fourier-spatial distribution of the phases, and not in their values themselves, that most of the structural information must lie.  The importance of this information may be best estimated in the context of interferometry. Indeed, interferometers essentially measure some Fourier components of the observed structures, and thus theoretically provide direct access to their phases. With the forthcoming ALMA instrument, the capacity of interferometers to detect structure in the Fourier phases \\emph{in real time} may be assessed. This is the purpose of this paper, which is organized as follows: Section \\ref{sec_pfa} offers a summary of the Fourier phase analysis technique, whose numerical implementation is presented in section \\ref{sec_psqip}. The main part of the paper, dealing with the application of these techniques to interferometric observations, is the topic of section \\ref{sec_atio}. Finally, section \\ref{sec_conc} gives a summary and conclusions. ",
        "conclusions": "\\label{sec_conc}  In this paper, we have addressed the ability of radio-interferometers to detect and recover the information contained in the Fourier-spatial distribution of phases, which was previously shown to store a vast amount of information about the structure of images. The PDF of phase increments and the related concept of phase entropy were introduced in this perspective. We ourselves have used the phase structure quantity $\\mathcal{Q}$, which is a minor modification of phase entropy leading to $\\mathcal{Q}=0$ for fields with purely random phases.  Our main conclusion is that the dynamical range of spatial frequencies observed by the instrument is the key parameter allowing detection and measurement of phase structure.  Using a turbulent model brightness distribution and instrumental configurations based on the characteristics of the future ALMA interferometer and of two existing arrays (VLA and Plateau de Bure), we have assessed the minimum integration time required by each configuration to have a significant detection of phase structure in the observed field. In the most conservative assessment, it appears that for a whole-field phase structure quantity $\\mathcal{Q} \\simeq 10^{-2}$, detection is achieved with a twenty minute integration in the ALMA E configuration (baselines going from 34 m to 909 m), or with a six hour integration in the VLA D configuration, but is not achieved by any other instrumental configuration tested\\footnote{It should be reminded that the more extended configurations have not been used in this single-configuration approach.}. With a whole-field phase structure quantity $\\mathcal{Q} \\simeq 10^{-3}$, certain detection can only be achieved using the ALMA E configuration, in which case it takes about 7.5 hours of integration.  However, less conservative criteria allow for early hints at the presence of phase structure in the observed field. Indeed, by drawing random phases for the visibilities in real time, it is possible to compare the evolution of the phase structure quantity for the observed field to that for a random-phase field, and check if they start going apart at some point. This is the case for all ALMA and VLA configurations, with whole-field phase structure quantities $\\mathcal{Q} \\simeq 10^{-2}$ and $\\mathcal{Q} \\simeq 10^{-3}$, but not for any of the Plateau de Bure configurations.  Regarding the possibility to recover the actual values of the phase structure quantity for the complete field, only multi-configuration observations with the ALMA instrument seem to allow for it, and it takes 9 hours in each of the 6 configurations to achieve this.  Finally, we have studied the influence of atmospheric phase noise on the single-configuration observations, using the E configuration of ALMA and whole-field $\\mathcal{Q}\\simeq 10^{-2}$. The maximum rms phase fluctuations that can be allowed without completely washing out the actual phase structure lie well above the typical range of variations for the Chajnantor site. Consequently, the use of water vapor radiometers to correct for the atmospheric phase fluctuations does not appear as a necessary feature of the ALMA array in this respect, although it should allow for a more accurate determination of the actual phase structure quantity in the multiconfiguration scheme.  Possible extensions to this work include the study of phase increments along the baseline tracks, which should considerably reduce the effects of atmospheric phase noise, and the evolution of phase structure with frequency in high spectral-resolution observations of line sources."
    },
    "0606/astro-ph0606014_arXiv.txt": {
        "abstract": "Coherent synchrotron emission by particles moving along semi-infinite tracks is discussed, with a specific application to radio emission from air showers induced by high-energy cosmic rays.  It is shown that in general, radiation from a particle moving along a semi-infinite orbit consists of usual synchrotron emission and modified impulsive bremsstrahlung. The latter component is due to the instantaneous onset of the curved trajectory of the emitting particle at its creation. Inclusion of the bremsstrahlung leads to broadening of the radiation pattern and a slower decay of the spectrum at the cut-off frequency than the conventional synchrotron emission. Possible implications of these features for air shower radio emission are discussed. ",
        "introduction": "There is recently growing attention on radio detection of high energy cosmic rays~\\citep{dz89,zhs92,heo96,avz00,fg03,g-etal04,fgp04,f-etal05}. High energy cosmic rays can initiate an electromagnetic cascade in a medium where relativistic electrons and positrons can be produced in a volume with a longitudinal (along the line of sight) dimension being smaller than the relevant radio wavelength. So, particles form a coherent bunch, acting like a single charged particle that emits a short burst of coherent radio emission. Two radiation processes have been considered: coherent Cerenkov radiation \\citep{a62, a65} and coherent synchrotron radiation~\\citep{fg03,hf03}. (Coherent radio emission from air showers was first considered by \\cite{kl66}, also \\cite{c67}, but their theory was not explicitly based on geosynchrotron emission.) The former requires charge asymmetry, say an excess of electrons, and a relatively dense medium for coherent Cerenkov emission to be at radio frequency. For example, the Moon is a good target for high energy neutrinos that can lead to a cascade in the lunar rocks. Excess electrons can develop leading to coherent Cerenkov emission at radio frequency with wavelength comparable with or larger than the longitudinal size of the cascade~\\citep{a62,a65}. For air showers, coherent synchrotron emission is generally considered more important than Cerenkov radiation (in the radio band)~\\citep{fg03,g-etal04}. The shower produces a bunch of relativistic electrons and positrons emitting coherent  synchrotron radiation in geomagnetic fields. Unlike coherent Cerenkov radiation, coherent synchrotron emission does not need charge asymmetry~\\citep{fg03}.   So far, the relevant spectra of coherent synchrotron emission from air showers were commonly calculated using numerical simulation based on the retarded-potential method \\citep{sgr03,hf05a,hf05b}. In this method, the radiation is calculated from the retarded potential~\\citep{j98}. Although coherent synchrotron emission has been considered analytically \\citep{ab02,fg03,hf03}, their calculation is based on the standard synchrotron radiation formula, which does not include the effect due to the particle's finite track. For air showers, effective coherent emission occurs in the core of the shower where most of radiating particles are created. Thus, it is of interest to consider radiation by a particle moving along a semi-infinite trajectory. Apart from the usual synchrotron emission there is emission due to the onset of the particle's curved trajectory. The latter component is referred to as the modified impulsive bremsstrahlung (MIB) as it is due to the combined effect of the usual impulsive (or prompt) bremsstrahlung due to particle ($e^\\pm$) creation, which is modeled as an abrupt jump in the particle's velocity from zero to $c$, and curvature of the particle's trajectory. It is worth noting that the finite track effect was considered for Cerenkov radiation~\\citep{t39}  and was taken into account explicitly in calculation of cosmic-ray induced showers in a dense medium~\\citep{zhs92}. In the case of Cerenkov radiation, the finite track leads to a reduction in radiation intensity and modification of angular distribution of the emission. Since the main objective of studying radio emission from air showers is to infer the properties of the cosmic rays that induce the showers, one needs to determine the radio spectrum accurately. In this paper, we present a general formalism for coherent synchrotron radiation from a nonstationary many-particle system, which takes account of MIB. The formalism developed here is based on the single particle treatment~\\citep{mm91}, in which radiation is due to a current associated with particle's motion in a medium. Here, the current is regarded as extraneous as it is different from that due to the plasma response to waves. The spectrum of radiation from a many-particle system is derived from the total current that is obtained by adding all the currents due to individual particles. In the relativistic limit as in the usual synchrotron radiation, the spectrum can be expressed in terms of the Airy functions and as a result, the radiation is highly beamed.  In Sec.2, a general formalism for synchrotron emission from a many-particle system is derived by including the effect of MIB due to the effect of a particle's semi-infinite track. Coherent synchrotron emission is considered in Sec. 3 and the application to air showers is discussed in Sec. 4. ",
        "conclusions": "Synchrotron emission by a particle moving along a semi-infinite trajectory is considered. Since effective coherent synchrotron emission by secondary particles in an air shower occurs in the core of the shower where most emitting particles are created, the initial point of the particle's trajectory need be included explicitly. It is shown that radiation from a particle moving along a semi-infinite track can be separated into the usual synchrotron emission and the bremsstrahlung-like emission (MIB). The latter is due to emission as the result of onset of the particle's curved trajectory.  The spectral intensity of the modified synchrotron emission (usual synchrotron plus MIB) is lower than that for normal synchrotron emission, roughly by a factor of 4. It is interesting to note that such reduction is consistent with the recent result from numerical simulation by~\\cite{hf05b}. The radiation pattern has a broader angular distribution than the usual synchrotron emission. This feature is especially pronounced for the perpendicularly polarized component. The spectrum has a much slower decay above the critical frequency $\\omega\\sim\\gamma^2\\Omega_e$, while the usual synchrotron spectrum has an exponential cutoff above the critical frequency. In the application to radio emission from air showers, the reduced intensity can be verified in principle provided accurate observations of the radio spectrum are available. Although there exist early observations of air shower radio emission, there are some uncertainties in determining the calibration factor \\citep{a71,a78}. The current LOFAR Prototype Station (LOPES) and the future LOFAR may provide a better opportunity to test the predicted spectrum. Since the transition frequency (to spontaneous emission) is much lower than the cut-off frequency, change near the cut-off frequency may not be observable. The broadening of the radiation pattern occurs mainly for the perpendicularly polarized component, which may be observable provided that there is significant charge asymmetry in the emitting plasma.  A major advantage of the formalism presented here is that the initial conditions including the time of particle creation, initial velocity and gyrophase all appear in the phase (\\ref{eq:psi2}) and these quantities can be modeled statistically. Although the single-particle formalism was used by~\\cite{wm82} to treat coherent gyromagnetic emission, in their calculation, the finite track effect was not considered. It is shown here that the distribution of the particle injection time is important in determining the coherence. The semi-infinite track approximation adopted here is valid provided that the emitting particles are highly relativistic with $\\gamma\\gg1$. In the relativistic limit, the synchrotron pulse duration ($1/\\Omega_e\\gamma^2 \\sim 10^{-8}\\,\\rm s$) is much shorter than the typical flight time $T\\sim 10^{-6}\\,\\rm s$ for a particle's free path 500 m, and therefore the orbit can be approximately regarded as semi-infinite. If electrons and positrons from an air shower have an energy cutoff extending to MeV energies with $\\gamma\\sim 1$, the synchrotron approximation is no longer valid and one has cyclotron emission instead. In this case, a finite orbit needs to be considered. One may extend the calculation in Sec. 3 and 4 to include the particle distribution in momentum and this requires a numerical approach, which is not considered here."
    },
    "0606/astro-ph0606086_arXiv.txt": {
        "abstract": "Recent observations from the Swift gamma-ray burst mission  indicate that a  fraction of  gamma ray  bursts are characterized by a canonical behaviour of the X-ray afterglows. We present an effective theory which allows us to account for X-ray light curves of both (short - long) gamma ray bursts and X-ray rich flashes. We propose that  gamma ray bursts originate from massive magnetic powered pulsars. ",
        "introduction": "\\label{Introduction} The unique capability of the Swift satellite has yielded the discovery of interesting new properties of short and long gamma ray burst (GRB) X-ray afterglows. Indeed, recent observations have provided new informations on the early behavior ( $t \\, < \\, 10^3 - 10^4 \\, sec$) of the X-ray light curves of  gamma ray bursts. These early time afterglow observations revealed that~\\cite{Chincarini:2005,Nousek:2005,O'Brien:2006} a  fraction of bursts have a generic shape consisting of three distinct segments: an initial very steep decline with time, a subsequent very shallow decay, and a final steepening (for a recent review, see Piran 2005, Meszaros 2006). This canonical behaviour of the X-ray afterglows of gamma ray bursts is challenging the standard relativistic fireball model, leading to several alternative models (for a recent review of some of the current theoretical interpretations, see M\\'esz\\'aros 2006 and references therein). \\\\ \\indent In order to determine the nature of both short and long gamma ray bursts, more detailed theoretical modelling is needed to establish a clearer picture of the mechanism. In particular, it is important to have at disposal an unified, quantitative description of the X-ray afterglow light curves.  \\\\ \\indent The main purpose of this paper is to present an effective theory which allows us to account for X-ray light curves of both gamma ray bursts and X-ray rich flashes (XRF). In a recent paper~\\cite{cea:2006} we set up a quite general approach to cope with light curves from anomalous $X$-ray pulsars (AXP) and soft gamma-ray repeaters (SGR). Indeed, we find that the canonical light curve of the X-ray afterglows is very similar to the light curve after the June 18, 2002 giant burst from AXP 1E 2259+586 (Woods et al. 2004). This suggests that our approach can be extended also to gamma ray bursts. \\\\ \\indent The plan of the paper is as follows. In Sect.~\\ref{light} we briefly review the  general formalism presented in Cea (2006) to cope with light curves. After that, in Sect.~\\ref{050315} through \\ref{050416A} we carefully compare our theory with the several gamma ray burst light curves. In Sect.~\\ref{origin} we propose that gamma ray bursts originate from massive magnetic powered pulsars, namely pulsars with super strong dipolar magnetic field and mass $M \\sim 10 \\, M_{\\bigodot}$. Finally, we draw our conclusions in Sect.~\\ref{conclusion}. ",
        "conclusions": "} \\label{conclusion} Let us conclude by briefly summarizing the main results of the present paper. We have presented an effective theory which allows us to account for X-ray light curves of both gamma ray bursts and X-ray rich flashes. We have shown that the approach developed to describe the light curves from anomalous $X$-ray pulsars and soft gamma-ray repeaters works successfully even for gamma ray bursts. This leads us the conclusion that the same mechanism is responsible for bursts from gamma ray bursts, soft gamma repeaters, and anomalous X-ray pulsars. In fact, we propose that gamma ray bursts originate by the burst activity from massive magnetic powered pulsars."
    },
    "0606/astro-ph0606279_arXiv.txt": {
        "abstract": "We present the first direct distance determination to a detached eclipsing binary in M33, which was found by the DIRECT Project. Located in the OB~66 association at coordinates ($\\alpha, \\delta)=(01\\!\\!:\\!\\!33\\!\\!:\\!\\!46.17, +30\\!\\!:\\!\\!44\\!\\!:\\!\\!39.9$) for J2000.0, it was one of the most suitable detached eclipsing binaries found by DIRECT for distance determination, given its apparent magnitude and orbital period. We obtained follow-up $BV$ time series photometry, $JHK_s$ photometry and optical spectroscopy from which we determined the parameters of the system. It contains two O7 main sequence stars with masses of $33.4\\pm3.5 \\; \\msun$ and $30.0\\pm3.3 \\; \\msun$ and radii of $12.3\\pm0.4\\; \\rsun$ and $8.8\\pm0.3\\; \\rsun$, respectively. We derive temperatures of $37000\\pm1500$ K and $35600\\pm1500$ K. Using $BVRJHK_s$ photometry for the flux calibration, we obtain a distance modulus of $24.92\\pm0.12$ mag ($964\\pm54$ kpc), which is $\\sim$0.3 mag longer than the Key Project distance to M33. We discuss the implications of our result and the importance of establishing M33 as an independent rung on the cosmological distance ladder. ",
        "introduction": "Starting in 1996 we undertook a long term project, DIRECT (i.e. ``direct distances''), to obtain the distances to two important galaxies in the cosmological distance ladder, M31 and M33. These ``direct'' distances are obtained by measuring the absolute distance to detached eclipsing binaries (DEBs).  M31 and M33 are the nearest and most suitable Local Group galaxies for calibrating the extragalactic distance scale. However, they present a much greater observational challenge than the current anchor of the distance scale, the Large Magellanic Cloud (LMC). Their greater distance makes the brightest stars in them appear $\\sim 6$ mag fainter than the brightest LMC stars, thus pushing the limit of current spectroscopic capabilities. In addition, crowding and blending become more significant with increasing distance \\citep{Mochejska00, Mochejska01c}. Unfortunately, distances are now known to no better than 10-15\\%, as there are discrepancies of $0.2-0.3\\;{\\rm mag}$ between various distance indicators \\citep[e.g.][Figure 8]{Benedict02}. These uncertainties limit the calibration of stellar luminosities and population synthesis models for early galaxy formation and evolution.  DEBs have the potential to establish distances to M31 and M33 with an unprecedented accuracy of 5\\% \\citep[for reviews and history of method see][]{Andersen91, Hilditch96, Paczynski97, Kruszewski99}. They offer a single step distance determination to nearby galaxies and may therefore provide an accurate zero point calibration of various distance indicators -- a major step towards very accurate and independent determination of the Hubble constant. In the last few years, eclipsing binaries have been used to obtain accurate distance estimates to the Large Magellanic Cloud \\citep[e.g.][]{Guinan98,Fitzpatrick03}, the Small Magellanic Cloud \\citep{Harries03,Hilditch05} and most recently to a semi-detached system in M31 \\citep{Ribas05}. Distances to individual eclipsing binaries in the Magellanic Clouds are claimed to be accurate to better than $5\\%$.  Detached eclipsing binaries have yet to be used as distance indicators to M31 and M33. The DIRECT project has initiated a search for DEBs and new Cepheids in the M31 and M33 galaxies. We have analyzed five $11\\arcmin\\times11\\arcmin$ fields in M31, A-D and F \\citep[][hereafter Papers I, II, III, IV, V]{Kaluzny98, Stanek98, Stanek99, Kaluzny99, Mochejska99} and one $22\\arcmin\\times22\\arcmin$ field, Y \\citep[][hereafter Paper IX]{Bonanos03}. A total of 674 variables, mostly new, were found: 89 eclipsing binaries, 332 Cepheids and 253 other periodic, possible long-period or non-periodic variables. We have analyzed two fields in M33, A and B \\citep[][hereafter Paper VI]{Macri01b} and found 544 variables: 47 eclipsing binaries, 251 Cepheids and 246 other variables. Follow up observations of fields M33A and M33B produced 280 and 612 new variables, respectively \\citep[][hereafter Papers VII, VIII]{Mochejska01a, Mochejska01b}, including 101 new eclipsing binaries. Variables from two more DIRECT fields, one in M31 and the other in M33, remain to be reported.  Of the 237 eclipsing binaries found by DIRECT, only four are bright enough ($V_{max}<20$ mag) for distance determination with currently available telescopes. An additional criterion for selection is that they contain deep eclipses, which removes degeneracies in the modeling. D33J013346.2+304439.9 is the brightest of these, discovered in field M33A (Paper VI), and this paper presents the distance we obtained to it with subsequent observations. In \\S 2 we describe the observations and the data reduction, in \\S 3 we present the light curve and radial velocity curve analysis, in \\S 4 the distance determination and in \\S 5 the discussion. ",
        "conclusions": "We present the first distance to a detached eclipsing binary in M33, establishing it as an independent rung on the cosmological distance ladder. This distance determination is a significant step towards replacing the current anchor galaxy of the extragalactic distance scale, the LMC, with galaxies more similar to those in the HST Key Project \\citep{Freedman01}, such as M33 and M31. We have chosen a detached eclipsing binary to simplify the modeling and derived a distance modulus of $24.92\\pm0.12$ mag.  D33J013346.2+304439.9 is located in the rich OB~66 association \\citep{Humphreys80}, which contains a relatively high massive star population \\citep{Massey95}. The presence of one of the best candidate detached eclipsing binaries for distance determination in this association is not surprising. In addition to the DEB, OB~66 contains several other eclipsing binaries (see Paper VI), which suggests a high binary star formation rate for massive stars. Our adopted color excess value $E(B-V)=0.09\\pm0.01$ is smaller than estimations from \\citet{Massey95} for OB~66. Using the ``$q$ method'' and $UBV$ photometry of 36 stars, they derive $E(B-V)=0.15\\pm0.02$ and their spectroscopic sample yields $E(B-V)=0.13\\pm0.01$. However, our multi-band photometry combined with the spectroscopy determines the reddening accurately, thus it is unlikely our distance estimation suffers from systematic errors due to reddening.  There are several avenues for improving the distance to M33 and M31 using eclipsing binaries. \\citet{Wyithe02} propose the use of semi-detached eclipsing binaries to be just as good or better distance indicators as detached eclipsing binaries, which have been traditionally considered to be ideal. The use of new improved stellar atmosphere models to derive surface brightnesses versus calibrations based on interferometry removes the restriction to DEBs for distance determination. Additionally, \\citet{Wyithe02} outline other benefits for using semi-detached binaries: their orbits are tidally circularized and their Roche lobe filling configurations provide an extra constraint in the parameter space, especially for complete eclipses ($i\\sim90$ deg). Bright semi-detached binaries in M33 or M31 are not as rare as DEBs, and are easier to follow-up spectroscopically, as demonstrated by \\citet{Ribas05} in M31. Thus, for the determination of the distances to M33 and M31 to better than $5\\%$ we suggest both determining distances to other bright DEBs and to semi-detached systems found by DIRECT and other variability surveys. Additional spectroscopy of the DEB would also improve the current distance determination to M33, since the errors are dominated by the uncertainty in the radius or velocity semi-amplitude.  How does our M33 distance compare to previous determinations? Table~\\ref{distances} presents a compilation of 13 recent distance determinations to M33 ranging from 24.32 to 24.92 mag, including the reddening values used. Our measurement although completely independent yields the largest distance with a small 6\\% error, thus is not consistent with some of the previous determinations. This possibly indicates unaccounted sources of systematic error in the calibration of certain distance indicators.  The implications of our result on the extragalactic distance scale are significant, especially when comparing to the $HST$ Key Project \\citep{Freedman01} distance to M33. They derive a metallicity corrected Cepheid distance of $24.62\\pm0.15$ mag, using a high reddening value of $E(V-I)=0.27$ and an assumed LMC distance modulus of $18.50\\pm0.10$ mag. If we calculate the LMC distance our result would imply, we derive $18.80\\pm0.16$ mag, which is not consistent with the eclipsing binary determinations. The error is obtained by adding in quadrature the individual errors in the two distance measurements. Taking this one step further, our LMC distance would imply a 15\\% decrease in the Hubble constant to $H_{0}=61\\; \\rm km\\;s^{-1}\\; Mpc^{-1}$. This improbable result brings into question the Key Project metallicity corrections and reddening values not only for M33, but also for the other galaxies in the Key Project. We thus demonstrate the importance of accurately calibrating the distance scale and determining $H_{0}$, which are both vital for constraining the dark energy equation of state \\citep{Hu05} and complementing the cosmic microwave background measurements from the Wilkinson Microwave Anisotropy Probe \\citep[WMAP;][]{Spergel06}."
    },
    "0606/astro-ph0606753_arXiv.txt": {
        "abstract": "} \\newcommand{\\eab}{ ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606109_arXiv.txt": {
        "abstract": "{Relativistic outflows represent one of the best-suited tools to probe the physics of AGN.  Numerical modelling of internal structure of the relativistic outflows on parsec scales provides important clues about the conditions and dynamics of the material in the immediate vicinity of the central black holes in AGN.} {We investigate possible causes of the structural patterns and regularities observed in the parsec-scale jet of the well-known quasar 3C 273.} {We present here the results from a 3D relativistic hydrodynamics numerical simulation based on the parameters given for the jet by Lobanov \\& Zensus (2001), and one in which the effects of jet precession and the injection of discrete components have been taken into account. We compare the model with the structures observed in 3C 273 using very long baseline interferometry and constrain the basic properties of the flow.} {We find growing perturbation modes in the simulation with similar wavelengths to those observed, but with a different set of wave speeds and mode identification. If the observed longest helical structure is produced by the precession of the flow, longer precession periods should be expected.} {Our results show that some of the observed structures could be explained by growing Kelvin-Helmholtz instabilities in a slow moving region of the jet. However, we point towards possible errors in the mode identification that show the need of more complete linear analysis in order to interpret the observations. We conclude that, with the given viewing angle, superluminal components and jet precession cannot explain the observed structures.} ",
        "introduction": "\\label{sect1}  The structure and kinematics of parsec-scale outflows is typically explained in terms of shocks (Marscher \\cite{mar80}, Marscher \\& Gear \\cite{mg85}, G\\'omez et al. \\cite{gam93}, \\cite{gam94}, G\\'omez et al. \\cite{gom+94}) and Kelvin-Helmholtz (K-H) instabilities (Hardee \\cite{har82}, \\cite{har84}, \\cite{har87}, Hardee et al. \\cite{hcr97}, Hardee \\cite{har00}, \\cite{har03}, Hardee et al. \\cite{har05}) developing in a relativistic fluid. Relativistic shocks may dominate the jet dynamics and emission at small scales, but are likely to dissipate at distances larger than $\\sim 10$\\,pc (Lobanov \\& Zensus \\cite{lz99}) due to the interaction with the slower flow. On intermediate scales ($\\sim 10$--100\\,pc) shocks and plasma instabilities may play equally important roles in jets (Lobanov \\& Roland \\cite{lr01}). Distributions of the synchrotron turnover frequency obtained for 3C\\,273 (Lobanov et al. 1997) and 3C\\,345 (Lobanov 1998) indicate that both shocks and instabilities are present on these scales, while larger scales are most likely dominated by plasma instabilities alone.  Recent studies by Hardee (\\cite{har00}) have shown that K-H instability may produce complex, three-dimensional ribbon-like and thread-like patterns inside a relativistic jet. In these ribbons and threads, a substantial increase of particle pressure and radio emissivity can be expected. This model has been successfully applied to the jet in 3C 120 (Hardee \\cite{har03}, Hardee et al. \\cite{har05}). The threaded structure forming a double helix has been detected in a space VLBI image of 3C\\,273 made at 5\\,GHz (Lobanov et al. \\cite{lob+00}). It was explained in terms of K-H instability developing in a relativistic flow with a modest Lorentz factor $\\gamma= 2.1$ and a relativistic Mach number $M = 3.5$ (Lobanov \\& Zensus 2001, hereafter \\cite{LZ01}). The analytical approach used in LZ01 is based on linear perturbation analysis of a K-H instability developed by Hardee (\\cite{har87}, \\cite{har00}). A similar approach applied to kiloparsec-scale jet in M87 allowed for accurate determination of physical parameters and modelling of radio emission to be made (Lobanov et al. \\cite{lhe03}).   However, results from numerical simulations of relativistic flows indicate that, after the linear regime of instability growth, the jets can be easily disrupted (Perucho et al. \\cite{pe+04b}). In addition to this, the bulk Lorentz factor $\\gamma=2.1$ derived in LZ01 is below the values required to explain the apparent speeds of $\\beta_\\mathrm{app}\\sim 5-8\\,c$ of enhanced emission features observed in 3C\\,273. LZ01 suggest that the K-H instability is developing in a slower, underlying flow, and the fast components are most likely faster shock waves produced in the jet by the periodic ejections associated with the nuclear flares. The presence of such shocks may disrupt the linear growth of the K-H instability as well, and it is not clear whether the linear stability analysis can be still applied in the presence of these kind of non-linear effects in the flow.  In view of these concerns, it is important to confront the results of LZ01 with numerical simulations, and attempt to address several fundamental issues about the stability and propagation of relativistic flows similar to the one observed in 3C\\,273. Numerical simulations can be used to verify whether the linear theory can be applied for explaining self-consistently the morphology and kinematics of parsec-scale flows, and whether these flows preserve fingerprints of linear modes even when the non-linear regime has developed or non-linear features, such as fast components, appear. Numerical simulations provide a means to address these problems by following the transition from linear to non-linear regimes of instability development (Perucho et al. \\cite{pe+04a,pe+04b}). The ultimate goal of this work is to probe the advantages and limitations of the combination of different approaches (linear theory, numerical simulations and observations) to studies of parsec-scale jets.   3C\\,273 is the second quasar discovered (Hazard et al. \\cite{hms63}), and the first one for which the emission lines were identified with red-shifted hydrogen lines (Schmidt \\cite{sch63}). In the same work, Schmidt (\\cite{sch63}) also pointed out the presence of a jet-like structure in this object. During the last four decades, the active nucleus and the relativistic outflow in 3C\\,273 have been studied in great detail (Courvoisier \\cite{cou98}).  The parsec-scale radio jet in 3C\\,273 has been monitored for almost three decades (Pearson et al. \\cite{pea+81}, Unwin et al. \\cite{unw+85}, \\cite{unw+89}, Zensus et al. \\cite{zen+88}, \\cite{zen+90}, Davis et al. \\cite{dum91}, Abraham et al. \\cite{abr+96}, Krichbaum et al. \\cite{kwz00}, Lobanov et al. \\cite{lob+00}, Asada et al. \\cite{asa+02}). The emission associated with the relativistic outflow on kiloparsec scales has been probed extensively in the radio (Conway et al. \\cite{con+81}, \\cite{con+93}), near infrared (Neumann et al. \\cite{nmr97}, Hutchings et al. \\cite{hut+04}), optical (Thompson et al. \\cite{tmw93}, Jester \\cite{jes01}, Jester et al. \\cite{jes+01}) and X-ray (R\\\"oser et al. \\cite{roe+00}, Marshall et al. \\cite{mar+01}, Sambruna et al. \\cite{sam+01}) wavebands.  The relativistic jet observed in the quasar 3C273 is one-sided, with no signs of emission on the counter-jet side at dynamic ranges of up to 16,000:1 (Unwin et al.  \\cite{unw+85}).  This is evidence for strong relativistic boosting in an intrinsically double-sided outflow powered by an accretion disk around the central black hole (Begelman et al. \\cite{bbr84}). The mass of the central black hole in 3C\\,273 is estimated to be $M_\\mathrm{bh} = 5.5^{+0.9}_{-0.8}\\times 10^8\\mathrm{M}_{\\sun}$ (Kaspi et al. \\cite{kas+00}).  The enhanced emission features (jet components) identified in the jet on scales of up to $\\sim$20 milliarcseconds (mas) are moving at apparent speeds exceeding the speed of light by factors of 5-8 (Abraham et al. \\cite{abr+96}). Plausible ranges of the Lorentz factor $\\gamma\\approx 5$--10 and viewing angles $\\theta_\\mathrm{jet}\\approx 10\\degr$--$15\\degr$ have been inferred from these measurements.  Ejections of new components into the jet occur roughly once every year (Krichbaum et al. \\cite{kwz00}), and they are likely to be related to weak optical flares observed with a similar periodicity (Belokon \\cite{bel81}).  The position angle at which the components are ejected shows regular variations with a likely period of about 13--15 years (Abraham et al. \\cite{abr+96}, Abraham \\& Romero \\cite{ar99}), correlated with the long-term variability observed in 3C\\,273 in the optical (Babadzhanyants \\& Belokon \\cite{bb93}) and radio (Turler et al. \\cite{tcp99}) bands. Abraham \\& Romero (\\cite{ar99}) have suggested that this periodicity may reflect changes of the jet axis induced by the relativistic precession of the inner part of the accretion disk.   Results from the linear analysis and numerical modelling are presented, compared and discussed in Sect.~\\ref{sect3} in connection to explaining the observed properties of parsec-scale outflow in 3C\\,273. Main results of the work are discussed in Sect.~\\ref{sect4}.  Throughout the paper, we adopt the flat $\\Lambda$CDM Cosmology with the Hubble constant $H_{0}=71\\,h\\,$km\\, s$^{-1}$\\,Mpc$^{-1}$, where $h$ is a constant with a likely value of 1, and matter density $\\Omega_{M}=0.27$. The positive definition of spectral index, $S\\propto\\nu^{\\alpha}$ is used. For 3C\\,273 ($z=0.157$, Strauss et al. \\cite{str+92}), the adopted cosmological parameters correspond to the luminosity distance $D_{\\mathrm L}=0.7h^{-1}$\\,Gpc. The respective linear scale is $2.69h^{-1}$\\,pc\\,mas$^{-1}$, and a proper motion of 1\\,mas/yr corresponds to an apparent speed of $10.1 h^{-1}\\,c$. ",
        "conclusions": "\\label{sect4}  We have performed two numerical RHD simulations with different initial setups in order to study the physical processes generating the observed structures in the parsec-scale radio jet in the quasar 3C\\,273. In the simulation 3C273-A, we have included a general set of helical and elliptical perturbations in a long jet with the basic physical parameters adopted from LZ01. In the simulation 3C273-B, we have used a shorter jet with the same physical parameters and have included precession and injection of fast components. The simulation 3C273-A was aimed to generate structures with wavelengths similar to those measured by LZ01 from the growth of Kelvin-Helmholtz perturbations. The simulation 3C273-B was designed to check if by combining the ejection of \\emph{superluminal} components and jet precession, with the periodicities reported in Babadzhanyants and Belokon (\\cite{bb93}) and Abraham et al. (\\cite{abr+96}), the same structures could be generated.  We find that the structures generated in simulation 3C273-A are of the same order in size as those observed, if the relativistic propagation effects of the waves are taken into account. We observe in the solution of the stability problem that the instability modes found in the simulation propagate at mildly relativistic speeds. These wave speeds differ from those derived from the approximations used in LZ01. This can be due to the uncertainties introduced by the approximations to the characteristic wavelengths in the interpretation of the observations in LZ01. However, we show that wavelengths similar to the observed ones are found for the wave speed given by the solution of the linear problem, although the modes fitted in LZ01 and those used here for the same wavelengths are not coincident. The solutions of the stability problem applied to the adopted wave speeds (0.23 $c$) and line of sight ($15^\\circ$) show that any body modes present in the jet should be much shorter than those fitted in LZ01. It should be noted that these differences do not rule out the presence of Kelvin-Helmholtz instability in parsec-scale jets. Despite difficulties in the mode identifications, the structures generated in the simulation are similar to those observed by LZ01.  Regarding the long-term stability of the flow, we note that the jet in the simulation 3C273-A is disrupted at $\\sim 170\\, \\rm{pc}$ from the inlet, contrary to the observations tracing the jet in 3C273 up to $60\\, \\rm{kpc}$ away from the source. The reasons for this difference may be found in the conjunction of several factors. 1)~The numerical simulation does not run long enough to reach a fully steady-state regime. The disruption point moves downstream along the simulation, which could imply that the disruption is a transitory phase. 2)~Magnetic fields have not been taken into account neither in the linear analysis, nor in the numerical simulation - and it should be noted that the magnetic fields may be dynamically important at parsec scales (Rosen et al. \\cite{ro+99}, Frank et al. \\cite{fra+96}, Jones et al. \\cite{jon+97}, Ryu et al. \\cite{ryu+00}, Asada et al. \\cite{asa+02}). 3)~We only simulate the underlying flow, without considering the faster and possibly denser fluid in the superluminal components. 4) Inaccuracies in the linear analysis approximations can lead to large uncertainties in physical parameters derived. 5)~Differential rotation of the jet, shear layer thickness (Birkinshaw \\cite{bir91}, Hardee \\& Hugues \\cite{hah03}, Perucho et al. \\cite{pe+06}), and a decreasing density external medium could also play an important role (implying jet expansion; see Hardee \\cite{har82,har87}, and Hardee et al. \\cite{har05}). 6)~Arbitrary initial amplitudes of perturbations were chosen for the simulation, so we could have included too large perturbations. A combination of these factors could well change the picture of the evolution of the jet in terms of its stability properties. The effects of the rotation and magnetic fields on the stability of jets remain unclear, since no systematic numerical study has been performed up to now.  In the simulation 3C273-B, we studied the effect of precession on the jet evolution and investigated the possibility that the short wavelength structures found in LZ01 were not due to K-H instabilities but due to the periodicities induced in the flow by the ejection of components. We demonstrate that such non-linear features as superluminal components generate linear structures in the form of Kelvin-Helmholtz instabilities which can be analyzed in the framework of linear perturbation analysis. One of the main conclusions that can be derived from this simulation is that helical twists can be excited by periodic injections if there is some induced helicity in the system. This helicity is induced in our simulation by the helical perturbation frequency, but in real jets this helicity could be induced by helical jet magnetic fields and/or by jet rotation. We have also shown that, in order to explain the observed $18\\,\\rm{mas}$ wavelength in terms of precession, either longer driving periodicities than the $15\\,\\rm{yr}$ suggested by Abraham \\& Romero (\\cite{ar99}) would be needed, or this wavelength must be induced by very fast components observed in a jet moving at a viewing angle $\\theta<15^\\circ$. The fast components could only generate the shorter wavelengths given in LZ01 ($\\lambda=2$ mas and $\\lambda=4$ mas) if a proper combination of the velocities and injection periodicities is used. Altogether, we find that inclusion of faster components and precession with the $15\\,\\rm{yr}$ periodicity does not explain well the observed wavelengths and periodicities. This gives more weight to the general conclusion about K-H instability acting prominently in the flow.  In the future, numerical simulations of this kind may be used to constrain the basic parameters of the flow such as the viewing angle and the component speed. Inclusion of magnetic fields, differential rotation and the effects of an atmosphere with a decreasing density could help reconciling better the simulations with the observed structures. In this way, for example, an increase of the jet radius due to decreasing external density could cause a downstream increase of wavelengths of K-H instability modes (Hardee et al. \\cite{har05}). This remains to be seen with future, full-fledged RMHD simulations of the relativistic jet in 3C~273. The scope of the present work could be expanded to other sources, and applied to prominent jets for which the transversal structure may be resolved, such as 3C~120, extending the work done by Hardee et al. (\\cite{har05}) by performing numerical simulations."
    },
    "0606/astro-ph0606423_arXiv.txt": {
        "abstract": "We have developed an analytical model to describe the evolution of anisotropic galactic outflows. With it, we investigate the impact of varying opening angle on galaxy formation and the evolution of the intergalactic medium. We have implemented this model in a Monte Carlo algorithm to simulate galaxy formation and outflows in a cosmological context. Using this algorithm, we have simulated the evolution of a comoving volume of size $(12h^{-1}{\\rm Mpc})^3$ in the $\\Lambda$CDM universe. Starting from a Gaussian density field at redshift $z=24$, we follow the formation of $\\sim20,000$ galaxies and simulate the galactic outflows produced by these galaxies. When these outflows collide with density peaks, ram pressure stripping of the gas inside the peaks may result. This occurs in around half the cases and prevents the formation of galaxies. Anisotropic outflows follow the path of least resistance and thus travel preferentially into low-density regions, away from cosmological structures (filaments and pancakes) where galaxies form. As a result, the number of collisions is reduced, leading to the formation of a larger number of galaxies. Anisotropic outflows can significantly enrich low-density systems with metals. Conversely, the cross-pollution in metals of objects located in a common cosmological structure, like a filament, is significantly reduced. Highly anisotropic outflows can travel across cosmological voids and deposit metals in other, unrelated cosmological structures. ",
        "introduction": "Galactic outflows play an important role in the evolution of galaxies and the intergalactic medium (IGM). Supernova explosions in galaxies create galactic winds, which deposit energy and metal-enriched gas into the IGM. These outflows are necessary to explain many observations and to solve many problems in galaxy formation, such as the high mass-to-light ratio of dwarf galaxies, the observed metallicity of the IGM, the entropy content of the IGM, the abundance of dwarf galaxies in the Local Group, the overcooling problem, and the angular momentum problem.  High-resolution, gasdynamical simulations of explosions in a single object reveal that outflows generated by such explosions tend to be highly anisotropic, with the energy and metal-enriched gas being channeled along the direction of least resistance, where the pressure is the lowest  \\citep{mf99,ms01a,ms01b}. Furthermore, several observations support the existence of anisotropic outflows (e.g. \\citealt{bt88,fhk90,carignanetal98,sbh98,stricklandetal00,vr02}).  Indirect support for the existence of anisotropic outflows comes from the observed enrichment of systems around the mean density of the universe \\citep{2003ApJ...596..768S,2004MNRAS.347..985P} and the enrichment of systems far from known galaxies at $z\\sim3$ \\citep{psa06, s06}. It may be challenging to enrich such regions with isotropic outflows even with the inclusion of enrichment from poorly understood Population III stars. Early indications are that Population III stars are unlikely to pollute the IGM to a large extent \\citep{nop04}. Anisotropic outflows may also provide an explanation for part of the observed scatter in the metallicity in the IGM, which is still unexplained \\citep{2003ApJ...596..768S,psa06}.  Several simulations of galactic outflows in cosmological contexts have been performed. Typically, these simulations use cubic comoving volumes of size $\\sim(10\\,{\\rm Mpc})^3$, containing thousands of galaxies. The methods used can be divided into two groups : analytical methods and numerical methods. The analytical methods describe the expansion of outflows in simulations that are either Gaussian random realizations of the density power spectrum or N-body simulations combined with prescriptions for galaxy formation. Outflows are represented using an analytical solution (e.g. \\citealt{sb01}, hereafter SB01, \\citealt{tms01,bsw05}) that currently assume isotropy.  The numerical simulations use hydrodynamical algorithms such as SPH, and outflows are generated in a variety of ways: by imparting SPH particles with a large velocity component (e.g. \\citealt{std01,sh03,od06}), depositing additional thermal energy into SPH particles (e.g. \\citealt{2002ApJ...578L...5T}), or taking output from completed SPH simulations and calculating, {\\it a posteriori}, the propagation of outflows into the IGM \\citep{aguirreetal01}. Unlike the analytical methods cited above, which assume isotropic outflows, all these numerical approaches have the potential to generate anisotropic outflows. However, we believe that there are some limitations to these numerical approaches, which motivates us to introduce an analytical model for anisotropic outflows.  With any SPH simulation, we need to be concerned with the limited resolution of the algorithm. Consider for example the simulations of \\citet{std01}. These authors identify galaxies of radius $r_N$, and rearrange the SPH particles located between $r_N$ and $r_0=2r_N$ into two uniform, concentric spherical shells located at radii $0.9r_0$ and $r_0$, which are then given an outward velocity. The outflow is therefore initially isotropic but becomes anisotropic as it propagates into a non-uniform external medium. We can see two potential problems with this approach. Firstly, the structure responsible for generating the anisotropy might be the galaxy or its environment, in which case the rearranging of particles into concentric shells would erase that structure entirely. Secondly, 52 particles per shell (the number they used) provides a good covering of the solid angles initially, but as the outflow expands and the particles in the shells move apart they become more like individual pressure points pushing on the external medium, and this could lead to an artificial mixing of the outflow and the external medium by Raleigh-Taylor instability.  The approach of \\citet{aguirreetal01} is radically different. It consists of identifying galaxies in an output from an SPH simulation and calculating the propagation of outflows from these galaxies in $N_a$ different directions. Since the resistance encountered by the outflows will be direction-dependent, outflows will start isotropic but then become anisotropic as the distance travelled by outflows will vary with direction. There are two limitations to this approach. Firstly, since it uses the output of an SPH simulation and introduces the outflows {\\it a posteriori}, the feedback effect of these outflows cannot be simulated (i.e. outflows do not influence the formation of other galaxies). Secondly, in this approach gas elements move radially and encountering high-pressure gas will dissipate their energy and rapidly slow down. In the real universe, that gas element will likely acquire a tangential velocity component that will redirect it towards regions where the resistance of the external medium is weaker.  To overcome these various limitations and study anisotropic outflows in a cosmological context, {\\it including the effect of feedback}, we have designed an analytical model for galactic outflows. This can then be combined with either an analytical method or a semi-analytic method for the description of galaxy formation. In this paper, we use an analytical Monte Carlo method. Calculations performed with an N-body semi-analytic approach will be presented in a forthcoming paper \\citep{mgp07}.  This paper is set out as follows. In \\S2, we describe our analytical model for anisotropic outflows. In \\S3, we describe our Monte Carlo method for cosmological simulations. Results are presented in \\S4, and their implications are discussed in \\S5. Conclusions are presented in \\S6. ",
        "conclusions": "We have designed an analytical model for anisotropic galactic outflows based on the hypothesis that such outflows are bipolar and follow the path of least resistance through the environment of their source. In this analytical model we vary one parameter: the opening angle, $\\alpha$. We combined this model with an analytical Monte Carlo method for simulating galaxy formation, galaxy mergers, and supernova feedback. With this combined algorithm, we study the evolution of the galaxies and the IGM inside a comoving cosmological volume of size $(12h^{-1}{\\rm Mpc})^3$, from redshifts, $z=24-2$, in a $\\Lambda$CDM model. Our main results are the following:  \\begin{itemize}  \\item Galaxy formation starts at redshift, $z\\sim18$. Since we neglect the formation and evolutionary times of massive stars, each newly formed galaxy immediately produces an outflow that lasts for a time, $t_{\\rm burst}\\sim50{\\rm Myr}$. Such outflows can travel hundreds of kiloparsecs, and eventually collide with other objects. We neglect the effect of a collision with a well-formed galaxy (the cross-section is too small). When an outflow collides with a peak still in the process of collapsing, removal of the gas by ram pressure, preventing the formation of the galaxy, occurs about half of the time. When stripping does not occur, the protogalaxy is enriched in metals. This process occurs for all opening angles and the proportion  stripped or metal-enriched is essentially independent of opening angle.  \\item When metal-enrichment of a peak occurs and this peak collapses to form a halo, the cooling time of that halo is reduced, and the galaxy forms earlier. However, this effect is rather small. In particular, we did not find that metal-enrichment could ``bring to life'' low-mass protogalaxies whose cooling time exceeds the age of the universe.  \\item Anisotropic outflows channel matter preferentially into low-density regions, away from the cosmological structures (filaments or pancakes) in which the galaxies producing the outflows reside. Consequently, the number of halos encountered by outflows decreases with decreasing opening angle. This reduction in the number of hits results in a larger number of galaxies forming, since fewer halos are stripped by the ram pressure of outflows.  \\item The volume filling factor of galactic outflows (that is, the volume fraction of the IGM occupied by outflows) holds constant and then decreases with opening angle. For angles $\\alpha=180^\\circ-110^\\circ$ the constant filling factor is a result of the balance between an increase due to larger numbers of galaxies and a decrease due to a fall in volume per outflow. At smaller angles, the volume of individual outflows drops significantly with $\\alpha$, and the total filling factor decreases since this term wins out.  \\item The decrease in filling factor with decreasing angle is not sufficient to explain the decrease in number of hits. The ratio (number of hits)/(filling factor) decreases with decreasing angles down to $\\alpha\\sim50^\\circ$. This indicates that at these angles, the outflows are efficient at avoiding collisions with halos and channel matter preferentially into low-density region. Hence, if several halos reside in a common cosmological structure, an outflow produced by one of them will tend to avoid encountering the others. For angles $\\alpha<50^\\circ$, we observe the opposite trend: outflows become more efficient in finding halos and hitting them. These narrow outflows can travel across cosmological voids and hit halos located in unrelated structures, like the next filament or pancake. This effect is a continuation of a process begun at $\\alpha\\sim 100^\\circ$ where the mean distance travelled by outflows when they strip collapsing peaks of their baryons begin to increase as the first neighboring structures are hit.  \\item The enrichment of the IGM with metals favors high-density systems since the sources of outflows are located in high-density regions. However, as the opening angle decreases, there is a dramatic reduction of the enrichment of such systems, combined with a dramatic increase in enrichment of low-density systems. Anisotropic outflows enrich around 10\\% larger a volume of the underdense Universe (and $40\\%$ more of the Universe below $\\rho/{\\bar \\rho}=0.1$) than isotropic outflows.  \\end{itemize}  This collection of results is a mere consequence of the fact that outflows follow the path of least resistance. This is an assumption in our model and is motivated by observations as well as high-resolution simulations."
    },
    "0606/astro-ph0606615_arXiv.txt": {
        "abstract": "We use a relativistic ray-tracing code to calculate the light curves observed from a global general relativistic magneto-hydrodynamic simulation of an accretion flow onto a Schwarzschild black hole. We apply three basic emission models to sample different properties of the time-dependent accretion disk. With one of these models, which assumes thermal blackbody emission and free-free absorption, we can predict qualitative features of the high-frequency power spectrum from stellar-mass black holes in the ``Thermal Dominant'' state. The simulated power spectrum is characterized by a power law of index $\\Gamma \\approx 3$ and total rms fractional variance of $\\lesssim 2\\%$ above 10 Hz. For each emission model, we find that the variability amplitude should increase with increasing inclination angle. On the basis of a newly-developed formalism for quantifying the significance of quasi-periodic oscillations (QPOs) in simulation data, we find that these simulations are able to identify any such features with (rms/mean) amplitudes $\\gtrsim 1 \\%$ near the orbital frequency at the inner-most stable orbit. Initial results indicate the existence of transient QPO peaks with frequency ratios of nearly 2:3 at a $99.9\\%$ confidence limit, but they are not generic features because at any given time they are seen only from certain observer directions.  Additionally, we present detailed analysis of the azimuthal structure of the accretion disk and the evolution of density perturbations in the inner disk. These ``hot spot'' structures appear to be roughly self-similar over a range of disk radii, with a single characteristic size $\\delta\\phi=25^\\circ$ and $\\delta r/r=0.3$, and typical lifetimes $T_l\\approx 0.3T_{\\rm orb}$. ",
        "introduction": "Our understanding of accretion onto black holes has grown substantially in recent years. Correlated magneto-hydrodynamic (MHD) turbulence, stirred by an underlying magneto-rotational instability, has now been well-established as the fundamental mechanism of angular momentum transfer in accretion disks \\citep{balbu98}. With that achievement in hand, detailed studies of global accretion disk dynamics have been undertaken via large-scale numerical simulations in both the pseudo-Newtonian \\citep{hawle01,hawle02,machi03,armit03} and general relativistic \\citep{devil03a,devil03b,gammi03} frameworks.  At the same time, we have also come into possession of a wealth of observational data.  High signal-to-noise spectra of active galactic nuclei (AGN; see e.g., the SDSS composite: \\citet{vande01}) as well as of black hole binaries in a variety of spectral states are now available \\citep{mccli05}; so, too, are detailed light curves and power-density spectra, particularly of black hole binaries \\citep{mccli05,vande05}, but also for Seyfert 1 galaxies and other AGN \\citep{marko04}. Fluorescent Fe K$\\alpha$ profiles have been used to infer detailed diagnostics of the disk surface in its innermost portions \\citep{fabia95,reyno97,done00,mille02}.  One particularly exciting area of research has been the discovery of high-frequency quasi-periodic oscillations (QPOs) in accreting black hole binaries \\citep{stroh01a,stroh01b}. In a growing number of sources, these QPOs appear in pairs with integer frequency ratios of 2:3 \\citep{mille01,remil02,homan05}. The high frequencies of these QPOs (near the frequency of the inner-most stable circular orbit; ISCO) suggests a strongly relativistic origin, but there exists a very broad range of theoretical models, none of which at this point appear overwhelmingly convincing \\citep{abram01,rezzo03,rebus04,schni05,petri05}. At the same time, there have not yet been any clear identifications of QPOs in the MHD simulations, much less pairs of QPOs with integer frequency ratios, making it difficult to choose one imperfect theoretical model over another.   Without a physically consistent way of associating light output with numerical simulations of accretion dynamics, we cannot directly compare the predictions of our best theoretical disk models with any spectral or timing observations. Therefore it is the goal of this paper to move toward the objective of linking dynamical simulations to observational diagnostics.   We will do so by applying several phenomenological models to detailed simulation data in order to predict the light that would be produced.  From these models, we generate images and light curves for accretion disks as they might be seen by distant observers.  Statistical analysis of this derived data will lead us to several interesting generalizations about the nature of variability in the light output of accreting black holes.  Although the particular models we employ here for the emissivity and opacity of disks are not entirely realistic, the formalism we create can be readily applied to more realistic models in the future. ",
        "conclusions": "\\label{discussion}  \\subsection{Summary of Findings}  We have developed a generalized post-processor analysis tool that couples ray-tracing in the Kerr metric with global GRMHD simulations to produce light curves and power spectra of accreting black holes. Using a variety of emission models, we probe different regions of the accretion disk and are better able to understand the underlying causes of time variability in the observed flux. The optically thick blackbody emission/absorption model is particularly useful for understanding the behavior of stellar-mass black holes in the Thermal Dominant state, which is characterized by a broad peak in the photon energy spectrum and very low levels of variability.  By fixing the black hole mass and accretion rate, we can convert from dimensionless code units to physical units for local fluid density. Then, assuming a radiation-pressure dominated gas in the inner disk, we can derive a physical temperature, and thus emissivity and absorption coefficients. The temperature and scale height of the disk compare well with the Novikov-Thorne predictions for a Schwarzschild black hole accreting at 50\\% $\\dot{M}_{\\rm Edd}$ with a small torque at the inner boundary. Despite the very different assumptions used in the analytic model and the computer simulations, the agreement is close enough to provide reasonable confidence in our conversion factors for density and temperature, and thus the conclusion that the simulations can be used to understand the Thermal Dominant state.  We have also developed new methods for analyzing the azimuthal structure of the accretion flow, determining the characteristic sizes and lifetimes of density perturbations. Over a large range of radii, we found the perturbations have a nearly exponential distribution of lifetimes, with $T_{\\rm life}\\approx 0.3T_{\\rm orb}$. Also, the characteristic shapes of the hot spots appears to be self-similar throughout the Keplerian regions of the disk, with $\\delta\\phi \\approx 25^\\circ$ and $\\delta r/r \\approx 0.3$.  From these short coherence times, it seems clear that the hot spots formed by MHD turbulence cannot survive long enough to produce the QPOs with quality factors of $Q\\sim 5-10$ described in \\citet{schni05}. Thus, if the transient high frequency QPOs with 2:3 frequency ratios identified in Section \\ref{hfqpos} are in fact robust signatures of GRMHD disk dynamics, they are most likely {\\it not} produced by geodesic hot spots.  This conclusion is supported by the fact that the QPOs are seen only from a single azimuthal viewing angle at any one time, perhaps suggesting some form of localized stationary wave in the disk. Finally, as part of our search for QPOs in the simulation data, we have also developed a general formalism for quantifying the significance of such features in simulations. These analytic results have been combined with Monte Carlo calculations to estimate our confidence limits for the QPOs at $\\gtrsim 99.9\\%$.  \\subsection{Comparison with Observations} As argued throughout this paper, we believe the MHD simulations most closely resemble the Thermal state of black hole X-ray binaries. This high-luminosity state is dominated by a broad thermal peak in the energy spectrum around $E \\sim 0.7-1.5$ keV and a relatively small contribution from a steep power-law tail at higher energies \\citep{mccli05}. The timing properties are characterized by a featureless power spectrum with $P(\\nu) \\sim \\nu^{-\\Gamma}$, with typical power-law index $\\Gamma \\approx 1$. The total power in this state is also small, with (rms/mean)$^2$ Hz$^{-1} \\lesssim 10^{-3}$ above 1~Hz.  Our simulated power spectra predict $\\Gamma \\approx 3-4$, which is rather greater than the typical slope.  On the other hand, just this sort of steep power spectrum has been seen in the hard state observations of \\citet{bello06}.   One robust prediction of the MHD/ray-tracing simulations is that, independent of the emission mechanism, the integrated rms power increases with disk inclination. In principle, this should be an observable prediction that could be tested with {\\it RXTE} data from black hole binaries in the Thermal state. Of course, there are complications when comparing two different binary systems, including different relative contributions to the thermal peak or power-law tail, different total luminosities, and different black hole masses. However, with a sufficiently large number of observations, it should be possible to account for these variables and extract a correlation between integrated spectral power and inclination.  Other observational trends that might be explored with future MHD simulations include the observed linear relation between the X-ray flux and rms, which appears to extend to AGN as well \\citep{uttle04,uttle05}. The timing properties of black hole binaries also vary greatly between their different spectral states, a general feature that is still not understood. To explore either of these relationships, we will require more sophisticated emission models and, quite likely, more simulation data with a broader range of disk parameters.  However, even with improved simulations and light curve models, there still exist some inherent difficulties in comparing theory with observation. {\\it RXTE} typically requires thousands of seconds of observation to detect high-frequency QPOs at rms amplitudes of a few percent (limited largely by photon counting statistics and ``dead time'' corrections). To achieve a similar sensitivity, our MHD simulations require hundreds of hours corresponds to generate only 0.2 seconds of real time for a $10 M_\\odot$ black hole (we have perfect photon statistics, but are limited by the inherent variance in the power spectra of short time series). This is more closely analogous to the time scales of AGN observations. If we were to scale the simulation to an AGN of mass $10^8 M_\\odot$, it would correspond to a sampling rate of every 40 minutes for roughly a month, assuming {\\it perfect} photon counting statistics! Thus it should be no surprise that QPOs are still so difficult to detect with high significance from supermassive black holes \\citep{benll01,vaugh05}."
    },
    "0606/astro-ph0606490_arXiv.txt": {
        "abstract": "\\emph{Context} Star-forming nuclear rings in barred galaxies are common in nearby spirals, and their detailed study can lead to important insights into the evolution of galaxies, their bars, and their central regions. We present integral field spectroscopic observations obtained with SAURON of the bar and circumnuclear region of the barred spiral galaxy M100, complemented by new {\\it Spitzer Space Telescope} imaging of the region.\\\\ \\emph{Aims} We use these data to enhance our understanding of the formation, evolution, and current properties of the bar and ring.\\\\ \\emph{Methods} We derive the kinematics of the gas and the stars and quantify circular and non-circular motions using kinemetry. We analyse this in conjunction with the optical and infrared morphology, and our previously published dynamical modelling. By comparing line indices to simple stellar population models we estimate the ages and metallicities of the stellar populations present within the region, especially in and around the ring.\\\\ \\emph{Results} The stellar and gaseous velocity fields are remarkably similar, and we confirm that the velocity fields show strong evidence for non-circular motions due to the bar and the associated density wave. These are strongest just outside the nuclear ring, where our kinemetric analysis indicates inflow across the spiral armlets and into the ring region.  The line strength maps all indicate the presence of a younger population within this ring, but detailed modelling of the line strengths shows that in addition to this young population, old stars are present. These old stars must have been formed in an event of massive star formation which produced the bulk of the mass, and which ended  some 3~Gyr ago, a constraint set by the age of the stars in the bar and the nucleus. Our best-fitting model is one in which the current star formation is but the latest of a series of relatively short bursts of star formation which have occurred for the last 500~Myr or so.  A clear bi-polar azimuthal age gradient is seen within the ring, with the youngest stars occurring near where the bar dust lanes connect with the ring.\\\\ \\emph{Conclusions} Our kinematic and morphological results all confirm the picture in which the nuclear ring in M100, considered typical, is fed by gas flowing in from the disc under the action of the bar, is slowed down near a pair of resonances, and forms significant amounts of massive stars. Detailed stellar population modelling shows how the underlying bulge and disc were put in place a number of Gyr ago, and that the nuclear ring has been forming stars since about 500~Myr ago in a stable succession of bursts. This confirms that nuclear rings of this kind can form under the influence of a resonant structure set up by a bar, and proves that they are stable features of a galaxy rather than one-off starburst events. ",
        "introduction": "The presence of a bar or other asymmetric perturbation of the gravitational potential within a spiral galaxy can strongly influence the dynamics of the gas and the stars. Bars are commonly formed within dynamically unstable galactic discs. During their evolution they can experience vertical buckling instabilities, which thicken the bar and increase the vertical velocity dispersion \\citep{com90,raha91}, and which we now know can be recurrent \\citep{mart06}. Bars are  present in up to 70$\\%$ of spiral galaxies (Knapen, Shlosman \\& Heller 2000a; Eskridge et al. 2000; Grosb$\\o$l, Patsis \\& Pompei 2004), indicating that they must be long lived components of a galaxy. This strong conclusion is confirmed by recent results suggesting that the bar fraction is roughly constant out to $z\\sim1$ (e.g., Jogee et al. 2004).  Dust lanes are often seen within large scale bars, offset from the bar major axis. The main families of periodic orbits making up the bar are elongated parallel to the bar, and perpendicular orbits are present if one or more inner resonances exist (Contopoulos \\& Papayannopoulos 1980), so called $x_1$ and $x_2$ orbits, respectively. At the location of such resonances, between the bar pattern speed and the period of the radial oscillations of the stars and gas, the stellar orbits abruptly change orientation. The gas orbits can only do this gradually, and this leads to the existence of shocks within the bar. The compression of this gas is manifested by increased dust extinction.  The shocks can lead to angular momentum loss in the gas, as a result of which the gas will flow in towards the nucleus until it stalls in a ring-like region in the vicinity of the Inner Lindblad Resonances (ILRs), where there is no net torque acting on the gas (Knapen et al. 1995a, hereafter K95a; Heller and Shlosman 1996). The accumulation of gas in such a nuclear ring can easily lead to enhanced massive star formation (Heller and Shlosman 1994; K95a). Such star forming rings are found within the central 1-2 kiloparsec of 20$\\%$ of spiral galaxies, practically all barred \\citep{k05}, and are traced by strong line emission. Overall, nuclear rings contribute around 3-5\\% to the local star formation rate \\citep{ken05}, and may play a role in the secular evolution of galaxies (see review by Kormendy \\& Kennicutt 2004).  The galaxy M100 (=NGC~4321) was chosen as the focus of this study due to its relative proximity \\citep[16.1~Mpc,][]{Fer96}, giving a scale of 70 parsec per arcsec, and the moderate strength of its bar. The bar strength ($Q_{\\mbox{b}}$) of M100 has been estimated at 0.2 by Laurikainen \\& Salo (2002), who suggest that significant asymmetric forces are present for bar strengths of $Q_{\\mbox{b}}> 0.05$. M100 exhibits a particularly clear resonant circumnuclear structure which has been described in detail by K95a and Knapen et al. (1995b, hereafter K95b, and 2000b, hereafter K00). The morphology changes dramatically between wavebands. In the \\emph{K} band, a small-scale bar is clearly visible, at the same position angle (PA) as the large-scale bar (see also Section 6.3). The picture, therefore, is of a single bar dissected by a nuclear ring. Hotspots of \\emph{K} band emission are located at the ends of the inner part of the bar. In blue optical light, as well as in H$\\alpha$ emission, a bright star-forming two-armed spiral structure is evident, delineated by the offset dust lanes which can be traced inwards through the main bar, through the ring, and towards the centre. We define the contact points as the locations on the ring dissected by a line through the nucleus and perpendicular to the bar major axis.  In this paper, we present stellar and gas kinematics, absorption line indices (H$\\beta$, Mg$\\emph{b}$ and Fe5015), and {\\it Spitzer Space Telescope} ({\\it SST}; Werner et al. 2004) near- and mid-infrared imaging across the nuclear ring and bar region in M100. We aim to study the detailed interplay between the dynamics on the one hand, and the ancient and recent history of the star formation in the bar and ring on the other.  This is the second paper presenting results from our new dataset on this galaxy. In the first paper (Allard, Peletier \\& Knapen 2005; hereafter Paper~I) we presented emission line and gas velocity dispersion maps which showed how massive star formation occurs at the precise locations of relatively cool gas, and related these results to the dynamical origin of the ring in terms of a resonance structure set up by the bar. The current paper is structured as follows: Section 2 describes the observations, data reduction and analysis of the SAURON and {\\it SST} data, Section 3 briefly describes the morphology and line ratios, Section 4 concentrates on the kinematics, and Section 5 on the stellar populations. Our discussions and conclusions are presented in Sections 6 and 7. ",
        "conclusions": "This paper presents SAURON integral field observations of the central region and bar of M100, which yields the gaseous and stellar kinematics, and emission and absorption line strength maps.  We complement this data set with {\\it SST} near- and mid-IR imaging of M100. The basic emission line morphology is that of the well-known nuclear ring which shines brightly in H$\\beta$, and which is linked to H{\\sc ii} regions at the Eastern end of the bar by a thin arc of ionised gas, running alongside but offset from the dust lanes (see also Paper~I). In [O{\\sc iii}] the morphology shows a central peak, while generally tracing the H$\\beta$ emission.  The H$\\beta$ and [O{\\sc iii}] gas kinematics show regular rotation along the bar, equivalent to an inclined rotating disc, but with characteristic kinks in the velocity field along the bar minor axis, near the radius of the nuclear ring.  The stellar velocity field is similar to that of the gas in the bar, and the non-circular motions are only slightly less evident near the centre, and practically as strong as in the gas along the bar minor axis. As discussed in more detail in Paper~I, the H$\\beta$ gas velocity dispersion shows a lower value where the star formation occurs, in the ring and at the ends of the bar, and also alongside the Eastern dust lane. The [O{\\sc iii}] gas velocity dispersion shows a slight decrease in the ring, and has a much higher central value than H$\\beta$.  The Mg{\\emph b} and Fe5015 absorption line indices show lower values within the ring suggesting that a younger stellar population is diluting the deep absorption features of the older bulge population. The H$\\beta$ absorption line index shows a broken ring, with the breaks occurring where the H$\\beta$ emission is the strongest. All three absorption line indices, as well as  the H$\\beta$ EW, show  that the youngest stellar populations are found at the contact points between the ring and the dust lanes.  Detailed modelling of the line strengths shows that old stars are present in addition to the young population which dominates the appearance of the ring, mostly so in emission lines like that of H$\\beta$. These old stars must have been formed in a past star formation event which produced the bulk of the mass, and which stopped some 3~Gyr ago, a constraint set by the age of the stars in the bar and the nucleus. Our best-fitting model is one in which the current star formation is but the latest of a series of relatively short bursts of star formation which have occurred for the last 500~Myr or so.  A clear bi-polar azimuthal age gradient is seen within the ring, with the youngest stars occurring near where the bar dust lanes connect with the ring.  Our kinematic and morphological results thus confirm a picture in which the nuclear ring in M100, considered typical, is fed by gas flowing in from the disc under the action of the bar, is slowed down near a pair of resonances, and forms significant amounts of massive stars. Detailed stellar population modelling shows how the underlying bulge and disc were put in place a number of Gyr ago, and that the nuclear ring has been forming stars since about 500~Myr ago in a stable succession of bursts. This confirms that nuclear rings can form under the influence of a resonant structure set up by a bar, and proves that they are stable features of a galaxy rather than one-off starburst events."
    },
    "0606/astro-ph0606173_arXiv.txt": {
        "abstract": "I review the progress in research on intracluster planetary nebulae over the last five years.  Hundreds more intracluster planetary nebulae have been detected in the nearby Virgo and Fornax galaxy clusters, searches of several galaxy groups have been made, and intracluster planetary candidates have been detected in the distant Coma cluster. The first theoretical studies of intracluster planetaries have also been completed, studying their utility as tracers of the intracluster light as a whole, and also as individual objects.  From the results to date, it appears that intracluster planetaries are common in galaxy clusters (10-20\\% of the total amount of starlight), but thus far, none have been detected in galaxy groups, a result which currently is not well understood.  Limited spectroscopic follow-up of intracluster planetaries in Virgo indicate that they have a complex velocity structure, in agreement with numerical models of intracluster light.  Hydrodynamic simulations of individual intracluster planetaries predict that their morphology is significantly altered by their intracluster environment, but their emission-line properties appear to be unaffected. ",
        "introduction": "Intracluster starlight, the diffuse starlight which permeates many galaxy clusters, is potentially of great interest to studies of galaxy and galaxy cluster evolution.  Since it is currently believed that the bulk of the intracluster stars were originally formed within galaxies and then were tidally removed from them, they are an important way to study the mechanisms of tidal stripping and interactions that are common within galaxy clusters (\\cite[Dressler 1984]{dressler1984}). Modern numerical simulations of galaxy clusters show that the intracluster light is ubiquitous in galaxy clusters, has a complex spatial and kinematic structure, and can be used to gain information on the dynamical evolution of galaxies and galaxy clusters (Napolitano et al. 2003; Willman et al. 2004; Murante et al. 2004; Sommer-Larsen et al. 2005; Stanghellini et al. 2006; Rudick et al. 2006).  However, the obstacles in observing intracluster light in detail are substantial.  Due to its low surface brightness (at the brightest, less than 1\\% of the night sky background in the $V$ band), it is extremely difficult to image directly.  There have been significant detections of intracluster light in galaxy clusters at low redshifts (z $<$ 0.3) (Feldmeier et al. 2004a; Gonzalez et al. 2004; Mihos et al. 2005; Zibetti et al. 2005; Krick et al. 2006), but nearly all of these observations require specialized observing techniques that are extremely time-consuming to carry out.  In addition, although intracluster imaging observations are crucial for obtaining a global view of the phenomenon, they can only give the spatial distribution of intracluster light, and possibly a color, meaning that detailed comparisons with the theoretical models will be difficult.  Finally, despite heroic efforts, direct imaging cannot yet probe the lowest surface brightness features, which have surface brightnesses of $\\mu_{V}$ = 32 mag/sq. arcsecond.  An alternate way to study intracluster light is to detect luminous individual intracluster stars in nearby galaxy clusters, and gain more detailed information on the distribution, metallicity and velocities of intracluster stars than is possible from surface brightness measurements.  This approach has also been quite successful: intracluster red giant stars (Ferguson, Tanvir, \\& von Hippel 1998; Durrell et al. 2002), intracluster H~II regions (Lee et al. 2000; Gerhard et al. 2002; Ryan-Weber et al. 2004) and intracluster novae and supernovae (Gal-Yam et al. 2003; Neil, Shara, \\& Oegerle 2005) have all been detected in galaxy clusters.  Here, we focus on another luminous tracer of the intracluster light: intracluster planetary nebulae (hereafter, IPN).  IPN have a number of unique advantages over other luminous tracers of the intracluster starlight.  Because IPN are emission-line objects, they can be detected efficiently in [O~III] $\\lambda$ 5007 surveys from the ground. Therefore, using wide-field imagers common on 4-meter class telescopes, samples of hundreds of candidates can be found in a single telescope run.  With spectroscopic follow-up using 6-meter and larger telescopes, the radial velocities of IPN can be determined, offering the ability to study the dynamics of intracluster starlight. ",
        "conclusions": ""
    },
    "0606/astro-ph0606345_arXiv.txt": {
        "abstract": "The recently discovered RRAT sources are characterized by very bright radio bursts which, while being periodically related, occur infrequently. We find bursts with the same characteristics for the known pulsar B0656+14. These bursts represent pulses from the bright end of an extended smooth pulse-energy distribution and are shown to be unlike giant pulses, giant micropulses or the pulses of normal pulsars. The extreme peak-fluxes of the brightest of these pulses indicates that PSR B0656+14, were it not so near, could only have been discovered as an RRAT source. Longer observations of the RRATs may reveal that they, like PSR B0656+14, emit weaker emission in addition to the bursts. ",
        "introduction": "\\object{PSR B0656+14} is one of three nearby pulsars in the middle-age range in which pulsed high-energy emission has been detected. These are commonly known as ``The Three Musketeers'' \\citep{bt97}, the other two being Geminga and PSR B1055--52. PSR B0656+14 was included in a recent extensive survey of subpulse modulation in pulsars in the northern sky at the Westerbork Synthesis Radio Telescope (WSRT) by \\citet{wes06}. In the single pulses analysed for this purpose, the unusual nature of this pulsar's emission was very evident, especially the brief, yet exceptionally powerful bursts of radio emission.  These extreme bursts of radio emission of PSR B0656+14 are similar to those detected in the recently discovered population of bursting neutron stars. These Rotating RAdio Transients (RRATs; \\citealt{mll+06}) typically emit detectable radio emission for less than one second per day, causing standard periodicity searches to fail in detecting the rotation period. From the greatest common divisor of the time between bursts, a period has been found for ten out of the eleven sources. The periods (between 0.4 and 7 s) suggest these sources may be related to the radio-quiet X-ray populations of neutron stars, such as magnetars \\citep{wt06} and isolated neutron stars \\citep{hab04}. However, \\citet{ptp06} have shown that the estimated formation rate of magnetars is too low. Furthermore the spectrum of the only RRAT for which an X-ray counterpart has so far been detected \\citep{rbg+06} seems to be too cool, too thermal and too dim for a magnetar, but is consistent with a cooling middle-aged neutron star like PSR B0656+14 \\citep{szk+06}. Also the pulse period and the slowdown-rate of PSR B0656+14, as well as the derived surface magnetic field strength and characteristic age, are within the range of measured values for RRATs. ",
        "conclusions": "We have shown that PSR B0656+14 could have been identified as an RRAT, had it been at the typical distance of the known RRATs. We have no way of telling whether its capacity to produce intense bursts of emission right across its profile is related to its age, period, inclination, or even its immediate galactic environment, since this behavior has been found in no other pulsar. The pulse-energy distribution is not a power-law, but is better fitted by a lognormal distribution and such distributions are thought to be common for pulsars (e.g. \\citealt{cjd04}).  In a study of 32 pulsars, \\citet{rit76} found that PSR B0950+08 shows the highest degree of pulse-to-pulse intensity variation. Nevertheless, the brightest pulse found in an extensive study of its field statistics by \\citet{cjd04} is approximately 5 {\\Eav}. Vela does not show pulses brighter than 10 {\\Eav} and only 0.5\\% of the pulses are brighter than 3 {\\Eav} \\citep{jvkb01}. For PSR B0656+14 4\\% of the pulses are brighter than 3 {\\Eav}. Therefore the emission of PSR B0656+14 appears to be extremely erratic compared with both normal pulsars and pulsars with giant micropulses.  One could wonder if more known pulsars show a similar kind of sporadic bright emission. To answer this question one should analyse the pulse energy distribution of a sample of pulsars. A complication would be that longer than typical observations are required to detect the presence of a tail of strong pulses. In a large survey for subpulse modulation by \\citet{wes06}, this pulsar was the only pulsar that showed clear evidence for this kind of sporadic emission.  Our identification of PSR B0656+14 with RRATs implies that at least some RRATs could be sources which emit pulses continuously, but over an extremely wide range of energies. This is in contrast to a picture (predicted by \\citealt{zgd06}) of infrequent powerful pulses with otherwise no emission. Therefore, if it indeed turns out that PSR B0656+14 (despite its relatively short period) is a true prototype for an RRAT, we can expect future studies to demonstrate that RRATs emit much weaker pulses among their occasional bright bursts. We would also predict that their integrated profiles will be found to be far broader than the widths of the individual bursts, and will need many thousands of bursts to stabilize.  Hopefully, radio observations of RRATs will soon be able to test these predictions. These, together with the detection of more RRATs and potentially their high-energy counterparts, will shed light on their true nature. The transient nature of these sources makes them difficult to detect. However it is likely that the Galactic population exceeds that of the ``normal'' radio pulsars \\citep{mll+06}. Thus surveys with long pointings, such as those planned with LOFAR, or many observations of the same region of sky are required. Surveys at low frequencies will also be more sensitive to nearer RRATs as the greater degree of dispersion will allow them to be more easily distinguished from radio frequency interference."
    },
    "0606/astro-ph0606035_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} Extreme horizontal branch (EHB) stars play an important role in extragalactic astronomy, since they have been individuated as possibly being responsible for the UV upturn in elliptical galaxies and in the bulges of spiral galaxies, that has been proposed as an independent age indicator for this type of galaxies. In recent years the ``binary scenario'', in which EHB stars formation is related to dynamical interactions inside binary systems, has been proposed as the main channel for their formation. In fact \\cite{Maxted} indicated that 69$\\pm$9\\% of field EHB stars should be close binary systems with short periods P$\\leq$10 days. Nevertheless, more recently \\cite{Napiwotzki} found a noticeably lower binary fraction (40-45\\%), and \\cite{MoniBidin} found no evidence of binarity among 18 EHB stars in globular cluster NGC6752. They estimated that within a 95\\% confidence level the close binary fraction in EHB of this cluster should be lower than 20\\%.  Here we present preliminary results of the extension of the previous survey. ",
        "conclusions": ""
    },
    "0606/astro-ph0606729_arXiv.txt": {
        "abstract": "{} {We present the results of global 3-D MHD simulations of stratified and turbulent protoplanetary disc models. The aim of this work is to develop thin disc models capable of sustaining turbulence for long run times, which can be used for on--going studies of planet formation in turbulent discs.} {The results are obtained using two codes written in spherical coordinates: GLOBAL and NIRVANA. Both are time--explicit and use finite differences along with the Constrained Transport algorithm to evolve the equations of MHD.} {In the presence of a weak toroidal magnetic field, a thin protoplanetary disc in hydrostatic equilibrium is destabilised by the magnetorotational instability (MRI). When the resolution is large enough ($\\sim 25$ vertical grid cells per scale height), the entire disc settles into a turbulent quasi steady--state after about $300$ orbits. Angular momentum is transported outward such that the standard $\\alpha$ parameter is roughly $4-6 \\times 10^{-3}$. We find that the initial toroidal flux is expelled from the disc midplane and that the disc behaves essentially as a quasi--zero net flux disc for the remainder of the simulation. As in previous studies, the disc develops a dual structure composed of an MRI--driven turbulent core around its midplane, and a magnetised corona stable to the MRI near its surface. By varying disc parameters and boundary conditions, we show that these basic  properties of the models are robust.} {The high resolution disc models we present in this paper achieve a quasi--steady state and sustain turbulence for hundreds of orbits. As such, they are ideally suited to the study of outstanding problems in planet formation such as disc--planet interactions and dust dynamics.} ",
        "introduction": "Observational surveys of star forming regions in the Galaxy have revealed the ubiquity of rotationally supported discs of gas and dust orbiting young stars \\citep[e.g.][]{beckwith&sargent96,odelletal93,staufferetal94, siciliaaguilaretal06,kessleretal06}. It is commonly believed that these discs are the likely sites of planetary formation \\citep{safronov69,lissauer93}. The discovery of numerous extrasolar planets has increased the need for greater understanding of their properties so that accurate models of planet formation can be developed.  These ``protoplanetary'' discs often show evidence for active accretion with a canonical mass flow rate onto the central star of $\\sim 10^{-8}$ M$_{\\odot}$ yr$^{-1}$ \\citep[e.g.][]{siciliaaguilaretal04}, requiring a source of anomalous viscosity to transport angular momentum outward. It has long been believed that this is provided by turbulence within the disc \\citep[e.g.][]{shakura&sunyaev73}. So far only one mechanism has been shown to work reliably: MHD turbulence generated by the magnetorotational instability (MRI) \\citep{balbus&hawley91,balbus&hawley91b}. Given the cool and dense nature of protoplanetary discs, there are questions about the global applicability of the MRI in such environments as the ionisation fraction is low \\citep{blaes&balbus94}. Models suggest that protoplanetary discs are likely to have both magnetically active zones, where the disc is turbulent, and adjacent magnetically `dead-zones' where the flow is laminar \\citep[e.g.][]{gammie96,fromang02,ilgner&nelson06a}. In this paper we focus on ideal MHD simulations of protoplanetary discs. We will examine the dynamics of `dead--zones' in future work.  Non linear numerical simulations performed using the local shearing box formalism \\citep[e.g.][]{balbus&hawley91b,hawleyetal96,brandenburgetal96} have shown that the saturated non linear outcome of the MRI is MHD turbulence having an effective viscous stress parameter $\\alpha$ between $\\sim 5 \\times 10^{-3}$ and $\\sim 0.1$, depending on the magnetic field configuration. Outward angular momentum transport can thus occur at the rate required to match observed accretion signatures onto T Tauri stars \\citep[see, for example,][who quote $\\alpha \\sim 0.01$ as being suggested by the observations]{hartmannetal98}.  Much of this early simulation work was performed in discs with no vertical stratification, and so was useful in determining the nonlinear outcome of the MRI, but did not provide insights into the global structure of these discs in either the radial or vertical direction. The question of their vertical structure was addressed by \\citet{stoneetal96} and \\citet{miller&stone00} who performed shearing box simulations of vertically stratified discs. A basic result to come out of these studies is that the discs evolve to a structure consisting of a dense region around the miplane where turbulence is driven by the MRI, sandwiched between a tenuous and magnetically dominated corona which is highly dynamic, but stable against the MRI.  Global MHD simulations of turbulent discs \\citep[e.g.][]{armitage98,hawley00,hawley01,steinacker&pap02,pap&nelson03a} confirm the basic picture provided by the local shearing box simulations. These early global simulations considered the radial structure of turbulent disc models, but employed non stratified cylindrical disc models. Recent work has been performed examining the dynamics of global, vertically stratified turbulent disc models, and has focussed on thick accretion tori around black holes, using either the Paczynski--Witta potential \\citep{hawley00,hawley&krolic01,hawleyetal01} or simulating accretion flows in the full Kerr metric \\citep[e.g.][]{devilliers03}. The starting conditions for these models are usually thick, constant angular momentum tori for which the gas is assumed to be non--radiating. These quickly evolve into thick accretion discs for which $H/R>0.1$. To date, there have been no published simulations of vertically stratified thin discs undergoing MHD turbulence (i.e. discs more akin to protoplanetary discs). In part this is because of the increased computational burden associated with simulating thin discs due to the higher resolution requirements.  The aim of this paper is to present and analyse a suite of MHD simulations of global, stratified, turbulent, protoplanetary disc models with $H/R \\le 0.1$. We focus on the global structure of the discs, and analyse how modifications to the boundary conditions and disc thickness affect the results. The longer term goal is to use these models to examine a range of problems in planetary formation in the context of turbulent, vertically stratified discs. As such this paper is the first in a series. Future publications will address problems such as: the evolution of dust in turbulent discs; the orbital evolution of planetesimals and low mass planets; gap formation and gas accretion by giant protoplanets; the dynamical evolution of dead--zones. A number of previous studies have addressed these issues using cylindrical models: \\citet{nelson&pap03b,papaloizouetal04,nelson&pap04b,nelson05} studied gap formation and planet formation in turbulent discs. They showed in particular that type I migration becomes stochastic because of the turbulent density fluctuations in the disc. \\citet{fromang&nelson05} also used cylindrical models to study the radial migration of solid bodies due to gas drag. They found rapid accumulation of meter size bodies in anticyclonic vortices apparently resulting from the turbulence. Another key problem in planet formation, dust settling in turbulent discs, has been studied recently using shearing box models \\citep{johansenetal05,fromang&pap06,turneretal06}. Because of the local approach, these analyses had to ignore radial drift. All of these issues are likely to be affected by the simultaneous treatment of radial and vertical stratification that we consider in this present work.  The plan of the paper is as follows: in section~\\ref{basic-equations}, we present the basic equations, notations, and diagnostics we use in this work. The set-up of our simulations is detailed in section~\\ref{simulations} and we present their results in section~\\ref{results}. We focus in particular on their global structure, and sensitivity to numerical issues such as resolution and boundary conditions. Finally, in section~\\ref{conclusion}, we summarise our results and highlight future improvements that will be added to the models. ",
        "conclusions": "\\label{conclusion} In this paper we have presented the results of 3-D MHD simulations of stratified and turbulent protoplanetary disc models. Our primary motivation is to develop disc models that can be used to examine outstanding issues in planet formation such as the migration of protoplanets, the growth and settling of dust grains, gap formation and gas accretion by giant planets, and the evolution and influence of dead--zones. Given that these phenomena occur on secular time scales, a key requirement is the development of disc models which are able to achieve a statistical steady state and sustain turbulence over long run times.  We examined the issue of numerical resolution, and found that disc models with $\\simeq 15$ vertical zones per scale height in the vertical direction showed a continuing slow decline in their magnetic activity, and gave rise to relatively small values of $\\langle \\alpha \\rangle$. A suite of higher resolution runs with $\\simeq 25$ zones per scale height achieved statistical steady states with values of $\\langle \\alpha \\rangle \\simeq 4 \\times 10^{-3}$, and it was shown that these models resolve the fastest growing modes of the MRI throughout the disc once a turbulent steady state has been achieved. For this reason we focused on simulations performed using this higher resolution. \\begin{figure} \\begin{center} \\includegraphics[scale=0.5]{figure/S5/thetamfluxS5.ps} \\caption{The sum of the mass fluxes through the vertical boundary of the disc located at $\\theta=\\theta_{max}$ and $\\theta_{min}$.} \\label{thetamflux_S5} \\end{center} \\end{figure} \\noindent The key features of the resulting disc models are: \\begin{itemize} \\item{Any toroidal magnetic flux that is initially present within the disc is quickly expelled from the midplane due to magnetic buoyancy. This occurs on a time scale of $\\sim 100$ orbits, which is the time required for the MRI to grow and develop into non linear turbulence throughout the disc. The disc then evolves as if it is threaded by an approximately zero net flux magnetic field, such that high resolution is required to maintain turbulent activity.}  \\item{A quasi--steady state turbulent disc is obtained after run times of between 250 -- 500 orbits, depending on the model. The volume averaged value of the effective viscous stress paramater $\\langle \\alpha \\rangle \\simeq 4 \\times 10^{-3}$, and time averaged radial profiles of $\\alpha$ yield variations of no more than a factor of two within the active domains of the disc models. These results are in basic agreement with previous studies of cylindrical discs \\citep{hawley01,pap&nelson03a}}  \\item{The discs can be described as having a two--phase global structure as a function of height: a dense, magnetically subdominant turbulent core that is unstable to the MRI in regions within $|Z| < 2.5 H$ of the midplane, above and below which exists a highly dynamic and magnetically dominated corona which supports weak shocks and is stable against the MRI. The engine that drives this structure is the MRI which generates and amplifies magnetic field near the midplane, which then buoyantly rises up into the low density corona where it dissipates and flows out through the boundaries at the disc surface. This is in basic agreement with the shearing box simulations presented by \\citet{miller&stone00} and recent studies of thick tori orbiting around black holes \\citep{hawley00,hawley&krolic01,hawleyetal01,devilliers03}}.  \\item{The velocity and density fluctuations generated by the models were found to be smaller than those obtained in cylindrical disc simulations using toroidal net flux magnetic field configurations \\citep{nelson05,fromang&nelson05}. $\\delta \\rho/\\rho_0 \\simeq 0.08$ in the stratified runs whereas in the cylindrical disc runs with net flux $\\delta \\rho/\\rho_0 \\simeq 0.13$. This has implications for the dynamics of dust, planetesimals and low mass protoplanets in turbulent discs as their stochastic evolution is driven by these fluctuating quantities.}  \\item{The vertically, azimuthally and time averaged values of the radial velocity in some of the disc models were compared with the expectations of viscous disc theory, and were found to give very good agreement. We conclude that, subject to a suitable time average, global evolution of these stratified turbulent models is in good accord with standard viscous disc theory \\citep[e.g.][]{shakura&sunyaev73,balbus&pap99,pap&nelson03a}} \\end{itemize}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.5]{figure/S5/vr_v_rS5.ps} \\caption{Comparison between the time averaged radial velocity obtained in the simulations ({\\it solid} line) and the predicted value of $v_R$ obtained from equation~(\\ref{v_R}) shown by the {\\it dotted} line.} \\label{vr_v_r_S5} \\end{center} \\end{figure}  There are a number of outstanding issues raised by our simulation results. \\citet{fromang&nelson05} reported the presence of anticyclonic vortices in turbulent unstratified cylindrical disc models. In the stratified models we present in this paper, however, we did not find any evidence of vortices. This could be due to a number of reasons. First, the vertical stratification may prevent the formation of vortices near the midplane. A study by \\citet{barranco&marcus05} showed that column--vortices in stratified discs are unstable and are quickly destroyed. They showed that vortices could form in the more strongly stratified upper regions of the disc which are more than one scale height above the midplane. In our simulations these regions are typically dominated by magnetic fields, whose associated stresses may prevent the formation of vortices there. Second, we adopted a smaller azimuthal domain than \\citet{fromang&nelson05}: $\\pi/4$ versus $\\pi/2$ and $2 \\pi$ models. This may prevent the formation of vortices, which were found to be quite extended in $\\phi$ by \\citet{fromang&nelson05}. We did not observe any vortices in the cylindrical disc model C1 described in section~\\ref{cyl_setup}, and this may be partly explained by the smaller azimuthal domain. Finally, the different magnetic field topology contained in the disc may play a role: \\citet{fromang&nelson05} used a net flux toroidal magnetic field. In the stratified models, the toroidal flux is quickly expelled from the disc midplane and the evolution is more similar to that of a zero net flux disc. The cylindrical model C1 contained a zero net flux magnetic field. The properties of the field can affect the turbulence and hence the formation of vortices. In particular, stronger spatial and temporal variations in the stresses may cause the surface density variations to differ systematically between the models presented here and those in \\citet{fromang&nelson05}. If the formation of vortices in the models described in \\citet{fromang&nelson05} are related to the `planet modes' described by \\citet{hawley87}, then these differences may explain the lack of vortices seen in the stratified models and model C1. These issues, and their influence on the evolution of solid bodies will be explored in greater detail in a future paper.  Finally, this is the first paper in a series which describes an approach to setting up models of turbulent, stratified protoplanetary discs capable of sustaining turbulence over long run times. Future papers will present a systematic study of outstanding problems in planet formation theory, such as disc--planet interactions, dust and planetesimal dynamics, and effects related to the presence of a dead zone. We also note that these models themselves can be further improved by including a realistic equation of state, heating and cooling of the disc, and a self--consistent treatment of the evolving ionisation fraction and conductivity of the disc material. At the present time inclusion of these physical processes is beyond current computational resources."
    },
    "0606/astro-ph0606217_arXiv.txt": {
        "abstract": "{{\\it Aims:} The hard and soft spectral states of black hole accretion are understood as connected with ADAF accretion (truncated disk) and standard disk accretion, respectively. However, observations indicate the existence of cool gas in the inner region at times when the disk is already truncated outside. We try to shed light on these not yet understood intermediate states.  {\\it Methods:} The disk-corona model allows to understand the spectral state transitions as caused by changes of the mass flow rate in the disk and provides a picture for the accretion geometry when disk truncation starts at the time of the soft/hard transition, the formation of a gap in the disk filled by an advection-dominated flow (ADAF) at the distance where the evaporation is maximal. We study the interaction of such an ADAF with an inner thin disk below.  {\\it Results:} We show that, when the accretion rate is not far below the transition rate, an inner disk could exist below an ADAF, leading to an intermediate state of black hole accretion. ",
        "introduction": "It is now widely accepted that the hard and soft spectral states of black holes correspond to accretion in form of an advection-dominated flow (ADAF) and a thin disk, respectively (Esin et al. 1997). But observations also show an intermediate state which often appears in transitions between these two states.  The hard states are understood as a consequence of interaction of the corona and the underlying disk which yields evaporation of gas from the disk to the hot coronal flow, and leads to a truncation of the accretion disk at some distance (Meyer et al. 2000b). Inside this truncation radius there is only a hot, optically thin flow. The soft state, on the other hand, is explained as the optically thick, standard accretion disk extending down to the last stable orbit, revealing itself in the characteristic multi-temperature blackbody spectrum.  But observations also show an intermediate state which appears often in transitions between these two states (McClintock and Remillard 2006). After the soft/hard transition the spectra are not always solely produced by an optically thin accretion flow, the occurrence of reflection and a Fe K$\\alpha$ line indicate cool matter in the inner region (\\.Zycki et al. 1998). These intermediate states often persist for times significantly longer than the viscous timescale of the disk and can thus not be explained only by the fading or build-up of a disk during transitions between the two states. We here address the accretion geometry in the intermediate state. (Physically different is the so-called very high/intermediate spectral state (discussion by Fender et al. 2004, model of disk fragmentation Meyer 2004.)  After a short description of the disk corona model in Sect. 2 we analyze in Sect.3 the physics of an ADAF above a disk. One of the key questions is whether heat can be drained from the upper ADAF and radiated away. We investigate the processes in the two-temperature regime which extends over most of the vertical height, as well as those in the thin layer above the disk surface where ion and electron temperature couple. In Sect. 4 we derive under which conditions condensation of matter from the ADAF into the disk occurs, the necessary process to allow an inner disk to survive. In Sect.5 we discuss how this picture can be related to the observed intermediate states. ",
        "conclusions": "We introduce a new feature of accretion flows onto black holes, the advection-dominated accretion flow affected by thermal conduction to a cool disk underneath. Our analysis shows that in the inner region condensation of matter from the hot flow into the disk is possible. This allows the existence of cool matter together with an already formed ADAF as indicated by the observed reflection and Fe K$\\alpha$ line in the intermediate spectral state."
    },
    "0606/astro-ph0606021_arXiv.txt": {
        "abstract": "{ With the new generation of high-resolution X-ray spectrometers the understanding of warm absorbers in Active Galactic Nuclei has improved considerably. However, the main question remains the distance and structure of the photoionised wind. } { We study the absorption and emission properties of the photoionised gas near one of the brightest and most variable AGN, the Seyfert galaxy NGC~4051, in order to constrain the geometry, dynamics and ionisation structure of the outflow.} { We analyse two observations taken with the Low Energy Transmission Grating Spectrometer (LETGS) of Chandra. We study the spectra of both observations and investigate the spectral response to a sudden, long-lasting flux decrease of a factor of 5 that occurred during the second observation. } { We confirm the preliminary detection of a highly ionised component with an outflow velocity of $-4500$~km\\,s$^{-1}$, one of the highest velocity outflow components seen in a Seyfert 1 galaxy. The sudden drop in intensity by a factor of five during the second observation causes a drop in ionisation parameter of a similar magnitude in the strongest and main ionisation component ($v = -610$~km\\,s$^{-1}$), allowing us for the first time to determine the recombination time of this component and thereby its distance in a robust way. We find an upper limit to the distance of $10^{15}$~m, ruling out an origin in the narrow emission line region. In addition, an emission component producing strong radiative recombination continua of \\ion{C}{vi} and \\ion{C}{v} appears during the low state. This can be explained by emission from an  ionised skin of the accretion disk at a distance of only $\\sim 4\\times 10^{12}$~m from the black hole. Finally, the spectra contain a broad relativistic \\ion{O}{viii} line with properties similar to what was found before in this source with XMM-Newton; this line has disappeared during the low flux state,  consistent with the disappearance of the inner part of the accretion disk during that low flux state. } { Combining high-resolution spectroscopy with timing information, we have constrained the geometry of the emission and absorption components in NGC~4051.} ",
        "introduction": "The stupendous amount of energy emitted by an Active Galactic Nucleus is released by gas that flows towards a supermassive black hole in the centre of a galaxy, presumably via an accretion disk. The average  energy of the emitted radiation increases towards the centre of the accretion disk, and thus X-rays provide the best probe of the gas flow in the immediate surroundings of the black hole. From the presence of jets, as well as from blue-shifted ultraviolet absorption lines, we learn that in addition to the flow towards the black hole, there is also gas flowing away from it. With the availability of X-ray spectrographs on Chandra and XMM-Newton, we can study these flows close to the black hole, by studying  lines in the X-ray spectra. The presence of multiple absorption line systems, that differ in their level of ionisation or in their outflow velocity, constrains the geometry of the outflow (for example NGC~3783, \\citet{netzer}; NGC~5548, \\citet{steenbrugge}).  In this paper we describe the X-ray spectra of the Active Galactic Nucleus of NGC\\,4051, obtained with the Low-Energy Transmission Grating on board of the Chandra satellite. Emission from the nucleus of NGC\\,4051 was already found by \\citet{hubble}, and the galaxy is one of the 6 observed by \\citet{seyfert}, and discussed in his seminal paper as one of the first 12 known Seyfert galaxies. Being a member of the Ursa Major cluster \\citep{tullyetal} it has a distance of 18.6~Mpc \\citep{tullypierce} and a redshift of 700~km\\,s$^{-1}$ \\citep{verheijen}. NGC\\,4051 is a relatively bright optical object, at $V\\simeq13.5$. In X-rays, the source was not detected with UHURU or Ariel-V, but was studied extensively with EXOSAT and ROSAT; its X-ray variability was recently analysed on the basis of RXTE and XMM-Newton data \\citep{mchardy}. Its X-ray continuum has been described with a power law with photon index $\\Gamma=$1.8-2.0 \\citep{nandrapounds}. A soft excess with respect to this power law is not well described by a multi-temperature disk, and \\citet{ogle} concluded that it is due to broad emission from relativistic \\ion{O}{vii} emission lines and to radiative recombination.  Two absorption line systems, at $-2340\\pm 130$ and $-600\\pm130$~km\\,s$^{-1}$, were found in the X-ray spectrum with the High Energy Transmission Grating of Chandra \\citep{collinge}. Multiple absorption systems were also found in the ultraviolet with the Space Telescope Imaging Spectrograph and in the Far Ultraviolet with the Far Ultraviolet Spectroscopic Explorer; one of these may coincide with the $-600\\pm130\\,\\mathrm{km\\,s}^{-1}$ X-ray absorption system, all the others have smaller absolute velocities \\citep{collinge,kaspi}.  Interestingly, a preliminary analysis of our first LETGS spectrum \\citet{vandermeer} did not show that high velocity component, but indicated the presence of an outflow component at even higher velocity, $-4500$~km\\,s$^{-1}$. For that reason we requested our second observation. This second observation, taken when NGC\\,4051 was in a much higher flux state, allowed us to investigate the ionised outflow both in terms of outflow velocity and ionisation structure. In this paper we report the analysis of both observations with the LETGS. ",
        "conclusions": "In this study we have investigated the spectral properties of NGC~4051, observed by {\\it Chandra}-LETGS on two occasions for a total of 180\\,ks.  The time averaged spectrum of both the Jan 2002 and the Jul 2003 can be fitted by a modifyed black body and a power law, absorbed by ionised gas.  Two of the ionisation components (2 and 3) are well visibible both in Jan 2002 and in the high state (part C) of Jul 2003. These gas components are consistent with being stable over a long time scale ($\\sim19$\\,months) both in ionisation level and outflow velocity.  In particular, we report the detection of the highest outflow velocity-gas ever observed ($v\\sim-4800$\\,km s$^{-1}$), for component 3.  The lower ionisation phase of the gas (component 1) is well detected only in the state C. This lack of detection is most probably due to the lower statistics affecting the other time segments.  In the last part of the Jul 2003 observation, (part D), only the second component is detected, but it shows a significant variation in the ionisation parameter which linearly responded to the continuum flux variation on a time scale $>$3000\\,s. From this we estimated a lower limit for the gas density ($n\\ga 10^{12}$\\,m$^{-3}$) and as a consequence a distance for the absorber $r\\la 10^{15}$\\,m and a thickness of the gas layer $d\\ga 2\\times 10^{11}$\\,m.  \\noindent The emission spectrum is rich in narrow and broad emission lines:  The narrow lines (e.g. the \\ion{O}{vii} forbidden line) are not variable in time and therefore consistent with being produced in regions very distant from the black hole (e.g. the NLR).  The RRCs of \\ion{C}{v} and \\ion{C}{vi} showed instead a very rapid variability, significantly arising above the continuum after the flux drop in spectrum D. The estimated distance ($r\\sim4\\times10^{12}$\\,m) of the recombining gas shows that these RRCs are produced close to (or in) the accretion disc.  We also find that the column density of the gas producing the RRCs is incompatible with that found for the warm absorber, suggesting a different location for the absorbing and emitting media.  Broad lines of \\ion{O}{viii}, \\ion{O}{vii} and \\ion{C}{vi} are also evident in the spectrum. The \\ion{O}{vii} and \\ion{C}{vi} lines show a simple Gaussian profile with a FWHM ranging from 0.3 to 1.5 \\AA, suggesting an origin in the BLR.  On the contrary, the \\ion{O}{viii} Ly$\\alpha$ line is best modeled by a relativistically broadened profile. The tentatively detected variability  on a months time scale is according to expectations for a line produced this close to a black hole. In accordance with previous studies, we find that the observed anticorrelation between the emissivity slope and the central source flux can be explained in terms of a relativistic accretion disc around a spinning black hole."
    },
    "0606/astro-ph0606398_arXiv.txt": {
        "abstract": "We have used the Very Large Array to search for compact milliarcsecond-size radio sources near methanol masers in high-mass star-forming regions. Such sources are required for Very Long Baseline Interferometry phase-referencing observations.  We conducted pointed observations of 234 compact sources found in the NVSS survey and find 92 sources with unresolved components and synchrotron spectral indexes. These sources are likely the cores of AGNs and, thus, good candidates for astrometric calibrators. ",
        "introduction": "Phase-referenced Very Long Baseline Interferometry (VLBI) observations can measure accurately the position of a target source relative to a reference source.  This, for example, permits one to measure the trigonometric parallax and proper motions of sources in our galaxy relative to extragalactic sources. We are now carrying out a large program to do this for methanol masers in regions of high mass star formation throughout the Milky Way. Table 1 contains 38 sources showing maser emission in 12.2 GHz $2_0-3_{-1}E$ line of methanol that we ultimately intend to observe.  VLBI phase-referenced observations involve a phase calibrator and a target source.  However, the distribution of known calibrators is not sufficiently dense enough to find one within $\\approx2^{\\circ}$ of any given maser source, which is necessary for high precision astrometry: Systematic errors are usually the limiting factor in phase-referenced position measurements, and these errors scale with the angular separation between the sources. Therefore, finding calibrators as close as possible to the maser target sources is mandatory.  For this purpose, we have conducted a survey of NRAO VLA Sky Survey (NVSS; Condon et al. 1998) sources near our methanol maser targets. We chose sources whose 1.4 GHz flux densities are greater than 20 mJy and that are unresolved ($<20''$) at the NVSS resolution. The NVSS survey typically yields about 15 such compact sources within $1^{\\circ}$ of any target position. Since our maser sources are in the Galactic plane many of these NVSS sources are compact HII regions, planetary nebulae (PNe), and, perhaps, compact supernova remnants (SNRs), and, thus, unsuitable for VLBI observations. However, some of the unresolved sources are likely extragalactic synchrotron sources.  In addition to the NVSS sources, we augmented our candidate list with sources from the literature (Gregory \\& Condon 1991; Ma et al. 1998) and included known compact sources as a check on our procedures. In total, we observed 234 compact sources within 1$^{\\circ}$ of the maser sources. ",
        "conclusions": "A VLA search has resulted in the identification of 92 compact sources most likely of extragalactic nature. These sources are likely to be suitable calibrators for phase-referencing VLBI observations of 12.2 GHz CH$_3$OH or/and lower frequency masers."
    },
    "0606/astro-ph0606167_arXiv.txt": {
        "abstract": "{} {We present the results of multi epochs imaging observations of the companion to the planetary host Gliese 86. Associated to radial velocity measurements, this study aimed at characterizing dynamically the orbital properties and the mass of this companion (here after Gliese 86 B), but also at investigating the possible history of this particular system.} {We used the adaptive optics instrument NACO at the ESO Very Large Telescope to obtain deep coronographic imaging in order to determine new photometric and astrometric measurements of Gliese\\,86\\,B.} {Part of the orbit is resolved. The photometry of \\gl~B indicates colors compatible with a $\\sim70$ Jupiter mass brown dwarf or a white dwarf. Both types of objects allow to fit the available, still limited astrometric data. Besides, if we attribute the long term radial velocity residual drift observed for \\gl~A to B, then the mass of the latter object is $\\simeq0.5\\,M_\\odot$. We analyse both astrometric and radial velocity data to propose first orbital parameters for \\gl~B. Assuming \\gl~B is a $\\simeq0.5\\,M_\\odot$ white dwarf, we explore the constraints induced by this hypothesis and refine the parameters of the system.} {} ",
        "introduction": "One of the biggest challenges of today astronomy is to detect and characterize extra solar planetary systems, and to understand the way(s) they form and evolve. Over the past decade, the technical improvements have allowed detections of more than 150 extrasolar planets via radial velocity (hereafter RV) measurements down to 7.5 Earth Masses \\citep{riv05} around solar type stars, while direct imaging allows now the detection of giant planets around young stars \\citep{lagrange2004,chauvin2004}. From the theoretical point of view the influence of the multiplicity or companionship with outer bodies (e.g. brown dwarfs; hereafter BD) on the dynamics and orbital stability of the inner planets has been highlighted. This has led to constant efforts trying to identify outer companions to those stars hosting planets plus long term RV drifts. \\begin{table*}[htb] \\caption{Observation log. $\\nds$ is a CONICA neutral density filter with a transmission of 1.4\\%. S13 and S27 are two CONICA cameras corresponding respectively to a platescale of 13.25 and 27.01~mas. WFS corresponds to the wave front sensor of the adaptive optics system.} \\label{table:1} \\begin{tabular*}{\\textwidth}{@{\\excs}llllllll} \\hline \\noalign{\\smallskip} UT Date & Filter & Camera & Observation type & Exp. Time (s) & WFS & Obs-Program & Platescale calibrator\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 12/11/2003 & \\ks & S27 & coronagraphy $(0.7~\\!'')$ & $100\\times0.6$ & VIS & 072.C-0624 & $\\Theta_1$ Ori C\\\\ 12/11/2003 & $2.17 + \\mbox{ND}_{\\mathrm{short}}$ & S27 & direct & $15\\times4.0$ & VIS &072.C-0624 & $\\Theta_1$ Ori C\\\\ \\noalign{\\smallskip} 22/09/2004&H & S13 & coronagraphy $(0.7~\\!'')$ & $48\\times1.0$ & VIS & 073.C-0468 &$\\Theta_1$ Ori C \\\\ 22/09/2004&H + $\\nds$ & S13 & direct & $42\\times0.35$ & VIS & 073.C-0468 &$\\Theta_1$ Ori C \\\\ \\noalign{\\smallskip} 29/07/2005 & \\ks & S27 & coronagraphy $(0.7~\\!'')$ &$400\\times0.8$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ 29/07/2005 & \\ks + $\\nds$ & S27 &direct &$400\\times0.35$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ 29/07/2005 & H & S13 & coronagraphy $(0.7~\\!'')$ &$360\\times1.$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ 29/07/2005 &H + $\\nds$ & S13 &direct &$400\\times0.35$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ 29/07/2005 & J & S13 & coronagraphy $(0.7~\\!'')$ &$165\\times2.$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ 29/07/2005 & J + $\\nds$ & S13 &direct &$240\\times0.5$ & VIS & 075.C-0813 & $\\Theta_1$ Ori C\\\\ \\noalign{\\smallskip} \\hline \\end{tabular*} \\label{logobs} \\end{table*}  \\gl~A is a K0V star with an estimated mass of $0.8\\,\\msun$ \\citep{siess1997,baraffe1998} and is located at $10.9$\\,pc from the Sun \\citep{perryman1997}. Through RV measurements, \\citet{queloz2000} have detected a 4~$\\mj$ (minimum mass) planet \\gl~b, orbiting \\gl~A at $\\sim 0.11\\mbox{AU}$. This star is also surrounded by a more distant companion \\gl\\,B, discovered at $\\sim 20\\,\\mbox{AU}$ using coronagraphy coupled to adaptive optics imaging \\citep{els2001}. The estimated photometry of \\gl~B is compatible with that expected for a $40$--$70\\,\\mj$ brown dwarf companion. Howewer, \\citet{mug05} showed recently that this was also compatible with a cool white dwarf, and that the latter hypothesis was more likely regarding the K band spectrum of the companion. The absence of near-IR molecular and atomic lines as well as the steep K-band continuum are indeed consistent with what is expected for a high gravity object with an effective temperature higher than 4000~K.  Apart for the RV wobble due to the hot Jupiter companion, \\gl~A also exhibits a long term RV drift measured with \\textsc{Coravel} and \\textsc{Coralie} over 20 years. This drift indicates the possible presence of an additional more distant companion, with a substellar mass and a distance to star greater than $\\simeq$ 20 AUs. \\citet{els2001} claimed that \\gl~B cannot account for this RV drift, due to its too low mass. They postulated instead that an additional companion, located in 2000 ``behind'' the star (i.e., under the coronographic mask), could be responsible for the observed drift.  In the course of a deep search for faint outer companions to stars hosting planets with NACO, we were able to make new images of \\gl~A and B in the near IR. We present the observational results in Sect.~2. In Sect.~3, we report new photometric result of \\gl~B and we present an analysis of both astrometric and RV data, assuming that the RV drift is due to \\gl~B. Finally, in Sect.~4 we discuss the nature of \\gl~B, and we confirm that it is very probably a $\\sim 0.5\\,\\msun$ white dwarf. We discuss the implications of this hypothesis. ",
        "conclusions": "The identification of the orbital motion of \\gl~B around \\gl~A, combined to the measured residuals of the radial velocity data, allow to severely constrain the whole \\gl\\ system and its past evolution. Our dynamical study shows that \\gl~B is very probably a white dwarf, in agreement with the conclusions of totally independent spectrophotometric study by \\cite{mug05}. The brown dwarf hypothesis of \\cite{els2001} can therefore be definitively ruled out.  The mass of \\gl~B is severely constrained by the dynamics. We derive $0.48\\,\\msun\\le m\\le 0.62\\,\\msun$. The orbit is eccentric ($e>0.4$) with a semi-major axis of a few tens of AU. The associated orbital period is several hundreds of years at least, and the stars have recently ($5$--$20$ years ago) passed at periastron. The orbit is retrograde with respect to the plane of the sky, but does not exactly lie in that plane.  Actually we can say that $120\\degr\\la i\\la150\\degr$.  Based on new photometric results on \\gl~B and the dynamical mass constrains, we also re-investigated the physical properties of this white dwarf companion. Using model predictions of \\cite{ber01}, we derived the effective temperature, the gravity and the cooling age of \\gl~B for both hydrogen-rich and helium-rich atmospheres models of white dwarfs.  When \\gl~B was a main sequence star, its mass probably ranged between $0.8\\,\\msun$ and $1.5\\,\\msun$, which implies a spectral type between K2V and F7V. Its orbit was closer. The strong post-main sequence mass loss caused the orbit to widen. If it had been a more massive star, the initial semi-major axis would have been too small to allow orbital stability for the  exoplanet orbiting \\gl~A.  However \\cite{saf05} recently used chromospheric index and metallicity measurements to estimate the age of all known stars harbouring exoplanets. For \\gl~A, they derived an age ranging between 2\\,Gyr and 3\\,Gyr. Given the main sequence lifetimes and the white dwarf cooling times (Table~\\ref{tab:model}), assuming this age for \\gl~B would imply that its progenitor had $m_\\mathrm{init}\\ga2\\,\\msun$. This seems to be incompatible with our dynamical constraints. Obviously, in order to solve this discrepancy, the dynamical evolution of the whole system, including the exoplanet needs to be investigated into further details. There are many open questions associated with this issue: the exoplanet must have survived all the late evolution stages of \\gl~B. If the system is not coplanar, the exoplanet could have been subject to the Kozai resonance in the past. Moreover, the planet must have formed in a large enough circumstellar disk, which implies a minimum initial separation of $\\sim 10\\,$AU. All these issues need to be addressed, and this will be the purpose of forthcoming work."
    },
    "0606/hep-ph0606178_arXiv.txt": {
        "abstract": "In this letter we explore the Higgs instability in the gapless superfluid/superconducting phase. This is in addition to the (chromo)magnetic instability that is related to the fluctuations of the Nambu-Goldstone bosonic fields. While the latter may induce a single-plane-wave LOFF state, the Higgs instability favors spatial inhomogeneity and cannot be removed without a long range force. In the case of the g2SC state the Higgs instability can only be partially removed by the electric Coulomb energy. But this does not exclude the possibility that it can be completely removed in other exotic states such as the gCFL state. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606601_arXiv.txt": {
        "abstract": "We have created a website, called \"Black Holes: Gravity's Relentless Pull\", which explains the physics and astronomy of black holes for a general audience. The site emphasizes user participation and is rich in animations and astronomical imagery. It won the top prize of the 2005 Pirelli INTERNETional Awards competition for the best communication of science and technology using the internet. This article provides a brief overview of the site. The site starts with an opening animation that introduces the basic concept of a black hole. The user is then invited to embark on a journey from a backyard view of the night sky to a personal encounter with a singularity. This journey proceeds through three modules, which allow the user to: find black holes in the night sky; travel to a black hole in an animated starship; and explore a black hole from up close. There are also five ``experiments'' that allow the user to: create a black hole; orbit around a black hole; weigh a black hole; drop a clock into a black hole; or fall into a black hole. The modules and experiments offer goal-based scenarios tailored for novices and children. The site also contains an encyclopedia of frequently asked questions and a detailed glossary that are targeted more at experts and adults. The overall result is a website where scientific knowledge, learning theory, and fun converge. Despite its focus on black holes, the site also teaches many other concepts of physics, astronomy and scientific thought. The site aims to instill an appreciation for learning and an interest in science, especially in the younger users. It can be used as an aid in teaching introductory astronomy at the undergraduate level. ",
        "introduction": "\\label{s:intro}  Some of the Hubble Space Telescope's most ground-braking discoveries have been about black holes, which are arguably the most extreme and mysterious objects in the universe. As a result, black holes appeal to a general audience in a way that almost no other scientific subject does. Unfortunately, most people have little idea of what black holes actually are, and they are more likely to associate them with science fiction than with science. These circumstances recommend black holes as a natural topic for an education and public outreach (E/PO) activity. To this end, we have created a website called \"Black Holes: Gravity's Relentless Pull,\" which serves as the E/PO component of several Hubble observing projects. The new website is part of HubbleSite, the Internet home of all Hubble news. The URL is \\underline{http:/$\\!$/hubblesite.org/go/blackholes/} .  Many sites on the internet already explain black holes in one way or other. Most are encyclopedic, offering detailed text and graphics.  By contrast, our site emphasizes user participation, and it is rich in animations and audio features. This approach was facilitated by the availability of powerful software for authoring multimedia content. The result is a website where scientific knowledge, learning theory, advanced technology, and pure fun all converge. ",
        "conclusions": ""
    },
    "0606/astro-ph0606437_arXiv.txt": {
        "abstract": "Previous calculations of the pregalactic chemistry have found that a small amount of H$_2$, $\\x{{\\rm H}_2}\\equiv n[\\rmH_2]/n[\\rmH] \\approx 2.6\\times 10^{-6}$, is produced catalytically through the H$^-$, H$_2^+$, and HeH$^+$ mechanisms.  We revisit this standard calculation taking into account the effects of the nonthermal radiation background produced by cosmic hydrogen recombination, which is particularly effective at destroying H$^-$ via photodetachment. We also take into consideration the non-equilibrium level populations of H$_2^+$, which occur since transitions among the rotational-vibrational levels are slow compared to photodissociation.  The new calculation predicts a final H$_2$ abundance of $\\x{{\\rm H}_2}\\approx 6\\times 10^{-7}$ for the standard cosmology.  This production is due almost entirely to the H$^-$ mechanism, with $\\sim 1$ per cent coming from HeH$^+$ and $\\sim 0.004$ per cent from H$_2^+$. We evaluate the heating of the diffuse pregalactic gas from the chemical reactions that produce H$_2$ and from rotational transitions in H$_2$, and find them to be negligible. ",
        "introduction": "One of the key problems in cosmology is to understand the physical and chemical state of the baryonic matter in the Universe.  At high redshift, the baryonic matter was fully ionized and co-existed with a thermalized radiation field (the cosmic microwave background, or CMB).  By redshift $z\\sim 10^3$, the Universe had expanded and cooled to $\\sim 3000\\,$K, at which point the ionized nuclei and free electrons of the primordial plasma combined to form neutral atoms. This cosmic recombination was first studied theoretically by \\citet{1968ApJ...153....1P} and \\citet{1968ZhETF..55..278Z}.  The observations of the acoustic peaks in the CMB $TT$ and $TE$ power spectra \\citep{2001ApJ...561L...1L, 2002ApJ...571..604N, 2002ApJ...568...38H, 2003ApJS..148..161K, 2006astro.ph..3450P, 2006astro.ph..3451H} provide direct evidence that cosmic recombination happened, and that it occurred over a narrow range in redshift, in accordance with predictions.  As the Universe continued to expand and cool, the formation of molecules became thermodynamically favourable. Since hydrogen is most abundant, one would expect the most abundant molecule to be H$_2$. However, unlike atomic recombination, which occurs shortly after it becomes thermodynamically favourable and proceeds nearly to completion (e.g. \\citealt{2000ApJS..128..407S}), cosmological formation of molecules is slow and freezes out with a final abundance, $\\x{\\rmH_2} \\equiv n[\\rmH_{2}]/n[\\rmH] \\ll 1$.  Despite their small abundance, molecules in the early universe have been investigated for several reasons. The first is that the primordial gas, mainly of hydrogen and helium atoms, lacks the low-lying excitations necessary for cooling and therefore star formation at low temperatures. On the other hand, molecules (which possess low-lying rotational excitations) could provide the cooling necessary to form the first stars \\citep{1967Natur.216..976S}.  However, recent calculations indicate that the primordial H$_2$ abundance is far too small for this, and that the only H$_2$ important for cooling of early haloes is formed in collapsed haloes (e.g. \\citealt{1997ApJ...474....1T}).  A second reason for studying H$_2$ production is that the heating of the gas, either via rotational transitions induced by the CMB or the chemical energy released by formation of the molecules, could affect the temperature of the pregalactic gas.  Here even a small effect could be important for proposals to study the absorption of the CMB by pregalactic gas in the H~{\\sc i} 21-cm line \\citep{2004PhRvL..92u1301L}. Finally, there is the (perhaps academic) motivation to understand the composition of the primordial gas as part of elucidating the standard cosmological model.  The first calculation of the primordial H$_2$ abundance was by \\citet{1967Natur.216..976S}. Noting that the direct radiative association of two H atoms is forbidden, they proposed that H$_2$ molecules could be built up using H$_2^+$ as an intermediate state (the specific reactions will be given in Section~\\ref{sec:rxn}).  \\citet{1968ApJ...154..891P} and \\citet{1969PThPh..41..835H} suggested that the H$^-$ mechanism dominated the production of molecules in primordial gas clouds.  A number of subsequent studies considered in increasing detail the H$_2$ abundance and cooling in primordial clouds \\citep{1969PThPh..42..523H, 1972PASJ...24...87Y, 1976ApJ...205..103H}.  \\citet{1984ApJ...280..465L} performed a calculation of the abundances of H$_2$ as well as HD and LiH, calculating a final abundance $\\x{\\rmH_2}\\sim 10^{-6}$ in the intergalactic gas. This is essentially today's ``standard'' calculation of the primordial H$_2$ abundance, although some of the reaction rates and cosmological parameters have been updated.  These updated analyses, which include substantial revisions to the deuterium and lithium chemistry, can be found in \\citet{1993A&A...267..337P}, \\citet{1995ApJ...451...44P}, \\citet{1998A&A...335..403G}, \\citet{1998ApJ...509....1S}, and references therein.  The effect of H$_2$ on heating of the gas before collapse has also been considered.  \\citet{1993A&A...267..337P} presented the first analysis of the thermal effect of molecules; they found a moderate effect ($\\sim$ 10 per cent) on the gas temperature mainly due to rotational lines in H$_2$, and a smaller effect due to chemical reactions.  The heating in rotational lines was also considered by \\citet{1996ApJ...464..523H} and \\citet{1998A&A...335..403G}, but was revised downward by \\citet{2000MNRAS.316..901F}, who concluded the effect was insignificant. We revisit the chemical heating here, and conclude that it dominates over H$_2$ rotational lines. However, even this effect is probably too small to be detected by 21-cm experiments in the forseeable future as it changes the gas temperature at the $\\sim 10^{-4}$ level.  All of these analyses, however, have been based on several common assumptions.  One is the assumption of a purely thermal radiation field, which is not completely correct because of the line and continuum radiation emitted during hydrogen and helium recombination (e.g. \\citealt{1993ASPC...51..548R, 2005astro.ph.10634W}).  Indeed, \\citet{2005PhRvD..72h3002S} found that this spectral distortion suppresses the lithium abundance by several orders of magnitude as compared with previous calculations \\citep{1996ApJ...458..401S, 1998A&A...335..403G, 2002ApJ...580...29S}. The other assumption is the use of H$_2^+$ photodissociation rates based either on local thermodynamic equilibrium (LTE) populations of the rotational-vibrational levels, or with all H$_2^+$ ions in the ground state. It is however known that H$_2^+$ forms preferentially in excited states and since radiative transitions between levels are slow, LTE may not apply.  The importance of this was recognized by \\citet{1998A&A...335..403G} and \\citet{2002JPhB...35R..57L}, but a full non-LTE analysis of H$_2^+$ level populations has not been done.  Our purpose here is to revisit the calculation of H$_2$ abundance, including spectral distortions to the CMB and with a level-resolved treatment of H$_2^+$.  In Section~\\ref{sec:rxn}, we introduce the chemical reactions important for H$_2$ production.  The H$^-$ mechanism, including the effect of the spectral distortion, is discussed in Section~\\ref{sec:-}.  The H$_2^+$ mechanism and the level-resolved treatment is in Section~\\ref{sec:+}, and HeH$^+$ is discussed in Section~\\ref{sec:he}.  The abundances of H$_2$, H$^-$, H$_2^+$, and HeH$^+$ are calculated for our presently favoured cosmology in Section~\\ref{sec:results}. The heating of the pregalactic gas by H$_2$ and the chemical reactions leading to its formation is considered in Section~\\ref{sec:heating}.  We conclude in Section~\\ref{sec:discussion}. The theory of the energy levels and transitions of the H$_2^+$ ion is recapitulated in Appendix~\\ref{app:h2p}.  In this paper, we have assumed a primordial helium abundance of $Y_P=0.24$, and a flat $\\Lambda$CDM cosmology with parameters from \\citet{2005PhRvD..71j3515S}: $\\Omega_b=0.0462$, $\\Omega_m=0.281$, and $H_0=71.0\\,$km$\\,$s$^{-1}\\,$Mpc$^{-1}$.  The number density $\\n$ will refer to the total proper density of hydrogen nuclei in all forms (ionized, atomic, and molecular), although in the regime of interest here it is mostly atomic.  The notation $x_i$ (or $\\x i$ for H$^-$, H$_2^+$, and HeH$^+$) will denote the number density of species $i$ relative to the total number density of hydrogen nuclei in all chemical forms (e.g. $\\x{\\rmH_2}= n[\\rmH_2]/n=1/2$ if all hydrogen is molecular). ",
        "conclusions": "\\label{sec:discussion}  We have reconsidered the production of H$_2$ molecules in the pregalactic medium.  In contrast to previous studies, we have included the spectral distortion in our analysis, and resolved all 423 rotational-vibrational levels of the H$_2^+$ ion.  We find that in the level-resolved analysis, the H$_2^+$ reaction pathway is greatly suppressed because newly formed H$_2^+$ ions are photodissociated before they can decay to the ground state or undergo charge transfer to become H$_2$ molecules.  We also find that the H$^-$ ion is easily destroyed by spectral distortion photons at $z>70$, so that the production of H$_2$ by this pathway is suppressed relative to the standard calculation.  We obtain a final H$_2$ abundance $\\x{\\rmH_2}=6\\times 10^{-7}$ assuming standard cosmology.  Unfortunately, the primordial H$_2$ molecules will be very difficult to detect.  The main effect of H$_2$ on the thermal history of the gas actually comes from the formation process (via the H$^-$ sequence, Eqs.~\\ref{eq:-1}, \\ref{eq:-2}) rather than rotational lines; however the effect is only of the order of $10^{-4}$.  In principle the proposed 21-cm tomography of the pre-reionization Universe could reach the sensitivity at which primordial H$_2$ becomes important, since it is sensitive to the temperature of the gas and has many more than $(10^{-4})^{-2}\\sim 10^8$ modes.  However when assessing the prospects, it should be remembered that the high-redshift 21-cm signal has not yet been detected, and measurements at the $10^{-4}$ level are clearly very far in the future.  Aside from H$_2$, there are other molecules with rotational lines such as HD and LiH, which could conceivably have been formed in the early Universe and played a role in the thermal balance.  The treatment of these trace molecules is beyond the scope of this paper, but HD in particular may warrant further study as it has been found to be a significant heating source in some past works (e.g. \\citealt{1993A&A...267..337P}).  Since the main route of formation of HD is via the reaction H$_2$(D$^+$,H$^+$)HD \\citep{1998A&A...335..403G}, any analysis of HD must incorporate the revised H$_2$ calculation presented here.  Finally, one could ask whether the H$_2$ suppression mechanisms discussed here -- the spectral distortion and non-equilibrium populations in H$_2^+$ -- have a significant effect on H$_2$ cooling of protogalaxies.  In the case of the spectral distortion, we have seen that at mean density the photodetachment of H$^-$ becomes unimportant at low redshift, since H$^-$ ions undergo a chemical reaction (usually associative detachment) before being destroyed by radiation.  For example, the branching fraction for H$^-$ at mean density to be destroyed by radiation is 0.13 at $z=40$, 0.09 at $z=30$, and 0.05 at $z=20$.  In overdense gas clouds at $z<40$ the collisional reactions are faster and we conclude that the spectral distortion should be negligible.  We have tried running our H$_2^+$ code for overdense conditions at low redshift (e.g. $z=20$, $\\delta_b=10^4$, $T=10^3\\,$K) and find that the H$_2^+$ levels are still far out of equilibrium, with the lowest levels underpopulated relative to the Boltzmann distribution (at $T_{CMB}$) by many orders of magnitude. However in such clouds H$^-$ is likely to be a more important source of molecules than H$_2^+$ (see e.g. \\citealt{1997ApJ...474....1T}). Therefore we do not expect our changes in H$_2^+$ physics to have large consequences for the cooling of the first collapsed objects in the universe."
    },
    "0606/astro-ph0606571_arXiv.txt": {
        "abstract": "{} {We present a cosmic shear analysis and data validation of 15 square degree high-quality $R$-band data of the Garching-Bonn Deep Survey obtained with the Wide Field Imager of the MPG/ESO 2.2m telescope.} {We measure the two-point shear correlation functions to calculate the aperture mass dispersion. Both statistics are used to perform the data quality control. Combining the cosmic shear signal with a photometric redshift distribution of a galaxy sub-sample obtained from two square degree of {\\it UBVRI}-band observations of the Deep Public Survey we determine constraints for the matter density $\\Omega_{\\rm m}$, the mass power spectrum normalisation $\\sigma_8$ and the dark energy density $\\Omega_\\Lambda$ in the magnitude interval $R\\in [21.5,24.5]$. In this magnitude interval the effective number density of source galaxies is $n=12.5\\,{\\rm arcmin}^{-2}$, and their mean redshift is \\mbox{$\\bar z=0.78$}. To estimate the posterior likelihood we employ the Monte Carlo Markov Chain method.} {Using the aperture mass dispersion we obtain for the mass power spectrum normalisation \\mbox{$\\sigma_8=0.80\\pm 0.10$} ($1\\,\\sigma$ statistical error) at a fixed matter density \\mbox{$\\Omega_{\\rm m}=0.30$} assuming a flat universe with negligible baryon content and marginalising over the Hubble parameter and the uncertainties in the fitted redshift distribution.} {} ",
        "introduction": "Measuring the weak gravitational lensing effect induced by the tidal gravitational field of the large-scale structure (LSS) of the Universe, also called cosmic shear, is a powerful way to explore statistical properties of the LSS and, thus, cosmological models. In contrast to other methods cosmic shear directly probes the total matter distribution of the Universe independent of the baryonic content. Due to the need for high quality data, it has only been in recent years that groups have begun to use this method. The first cosmic shear measurements were published in the beginning of 2000 \\citep{bre00,vme00,kwl00,mvm01,wtk00}. Although some constraints on cosmological parameters were obtained from these measurements, their statistical errors, due to the small areas of the surveys (\\mbox{$\\sim 1\\,{\\rm deg}^2$}), and systematic errors were large. More recently, published cosmic shear measurements use larger survey areas (\\mbox{$\\sim 10\\,{\\rm deg}^2$}) and/or deeper fields to obtain smaller statistical errors and additionally discuss the influence of systematic errors on the cosmic shear signal \\citep[e.g.][]{vmr01,vmp02,vmh05,rrg01,rrc04,rrg02,hms03,bmr03,hyg02,btb03,mrb05,jbf03,jjb05,hbb05}.  We are now entering a new phase of wide-field surveys which are wider and/or deeper so that systematic errors may begin to dominate over statistical errors. Many groups have been working on the problem of reducing systematic errors in cosmic shear measurements and have presented its impact on cosmological parameter estimations. \\citet{hoe04} and \\citet{vmh05} discussed the influence of an imperfect PSF-anisotropy correction on the cosmic shear signal. The impact of the redshift distribution sampling error has recently been studied by \\citet{mhh06} and \\citet{vwh06}. A further, more fundamental, source of bias is the correlation between the weak gravitational shear of distant galaxies and the intrinsic shape of foreground galaxies \\citep[e.g. ][]{his04,hwh06}. In addition, the so-called Shear TEsting Programme (STEP), a world-wide collaborative project to improve the accuracy and reliability of all weak lensing measurements has been initiated in mid 2004 \\citep[see first results in ][]{hvb06}.  In the present paper we first present the results from a simple test of our pipeline using synthetic images. Specifically we give details about the changes we have made to it due to the lessons we have learnt so far from STEP and present the influence of the changes on the cosmic shear signal. To test the pipeline we used $3\\,{\\rm deg}^2$ of synthetic images obtained by ray tracing through $N$-body simulations \\citep{het05}. A cosmic shear analysis is performed using the E- and B-mode decomposition of the aperture mass ($M_{\\rm ap}$-) statistics. We show that within the uncertainties we recover the cosmic shear signal and that our pipeline does not create artificial B-modes.  We then report on the results of a cosmic shear analysis of the Garching-Bonn Deep Survey (hereafter: GaBoDS) data set, obtained with the Wide Field Imager (WFI) of the MPG/ESO 2.2m telescope on La Silla, Chile. The GaBoDS data set consists of $18.6\\,{\\rm deg}^2$ high quality $R$-band observations. Most of the data were obtained in the framework of a virtual survey. That is, we compiled the data set from the ESO archive. The rest of the data set was obtained by our GO programmes.  We effectively use $15\\,{\\rm deg}^2$ of the total data set to create catalogues of galaxies for the shear measurements. Furthermore, the data set includes $2\\,{\\rm deg}^2$ of deep {\\it UBVRI}-band observations from the Deep Public Survey (henceforth: DPS). This yields photometric redshift information for $8\\%$ of all GaBoDS lensing objects in the magnitude interval $R \\in [21.5,24.5]$. The DPS lensing galaxies with photometric redshift information are used to estimate the redshift distribution which in turn is considered for the cosmic shear analysis.  To test the reliability of our measurements we perform four basic tests on the systematics in our data. In addition, we calculate the ellipticity cross-correlation between uncorrected stars and corrected galaxies to quantify the residuals of the PSF anisotropy correction. Furthermore, we utilise the E- and B- mode decomposition of the $M_{\\rm ap}$-statistics to check for possible systematics in the data.  The signal of the aperture mass statistics, the shear two-point correlation functions and their covariance matrix are calculated for the cosmic shear analysis. In contrast to other cosmic shear analyses, we estimate the mass power spectrum normalisation, $\\sigma_8$, using the Monte Carlo Markov Chain (MCMC) technique.  The paper is organised as follows. In Sect. \\ref{sect:cosmicshear} we briefly describe the theory of cosmic shear. We also discuss the statistics that we use to constrain cosmological parameters as well as a control of possible systematics in the data. The data and the weak lensing catalogue creation are presented in Sect. \\ref{sect:data}. The analysis of synthetic images is presented in Sect. \\ref{sect:simulations}. In Sect. \\ref{sect:Bmodecontrol} we elucidate in detail the different tests on the systematics in our data. In Sect. \\ref{sect:cosmicshearanalysis} we give the results from our cosmic shear analysis and present the cosmological parameters calculated using the MCMC technique. In addition we discuss several possible sources of systematic errors. In Sect. \\ref{sect:conclusion} we give a summary and outlook. ",
        "conclusions": "\\label{sect:conclusion} We have performed a cosmic shear analysis based on $15\\,{\\rm deg}^2$ of deep high quality $R$-band imaging data from the Garching-Bonn Deep Survey. Our cosmic shear measurement is B-mode free within the statistical errors, hence there are no significant systematic errors resulting from the data treatment (PSF correction, galaxy selection) left in the data. This encourages us to use the GaBoDS data for further cosmological analyses, such as the galaxy bias studies carried out by Simon et al. (in preparation).  The measured redshift distribution is obtained from {\\it lensing} galaxies from $2\\,{\\rm deg}^2$ of {\\it UBVRI}-band data of the Deep Public Survey, a sub-sample of GaBoDS and not from an external redshift sample.  The measured redshift distribution in combination with the cosmic shear signal is used to perform an unbiased estimate of cosmological parameters where, in contrast to many other cosmic shear analyses, we have estimated the covariance matrix {\\it directly from the data} without any further assumptions and have not computed it analytically. Assuming a flat $\\Lambda$CDM universe with negligible baryon content we have derived the mass power spectrum normalisation and the matter density while marginalising over the uncertainties in the Hubble parameter and the source redshift distribution. As a result we obtain \\mbox{$\\sigma_8=0.61^{+0.31}_{-0.20}$} and \\mbox{$\\Omega_{\\rm m}=0.46^{+0.30}_{-0.22}$} from the $M_{\\rm ap}$-statistics using the Peacock \\& Dodds model of the non-linear mass power spectrum. For a fixed matter density of \\mbox{$\\Omega_{\\rm m}=0.3$}, we obtain \\mbox{$\\sigma_8=0.80 \\pm 0.10$} ($1\\,\\sigma$ statistical error). Within the error bars this is consistent with recent cosmic shear measurements, galaxy clusters and the WMAP three year result, see Fig. \\ref{fig:sigma8vergleich} and table \\ref{tab:sigma8vergleich}.  We discussed various systematic errors and roughly estimated their magnitude. With respect to the magnitude of systematic errors, the most uncertain sources are the redshift distribution and the intrinsic shape-shear correlation. Although the intrinsic shape-shear correlation has been analysed with $N$-body simulations using simple galaxy models, the impact on deep cosmic shear measurements is still unclear. With accurate photometric redshifts at hand, measurements of cross-correlation tomography could be carried out to provide a useful diagnostic tool \\citep{his04}. To reach the full potential of weak gravitational lensing measurements and high accuracy in the determination of cosmological parameters, it is therefore essential to have precise redshift estimates to obtain: a) unbiased redshift distributions of lensing galaxies and b) estimates of the intrinsic shape-shear correlation.  Our cosmic shear result and the ongoing improvements in redshift and shear estimates (e.g. STEP) are quite encouraging for the upcoming wide-field multi-colour surveys as the future KIlo Degree Survey (KIDS), starting beginning 2007. The KIDS aims to image a contiguous area of $1500\\,{\\rm deg}^2$ in five colours. In combination with UKIDS the data will yield photometric redshifts accurate to $\\Delta z/(1+z)=0.03$ for typical  $r^\\prime=23.5$ galaxies, rising to $10\\,\\%$ uncertainty at $r^\\prime=25$. In the primary lensing colour the limiting magnitude is predicted to be $r^\\prime_{\\rm Vega}=24.3$ ($10\\,\\sigma$ sky level measured in a circular aperture of $2^{\\prime \\prime}$ radius) and the median seeing is predicted to be $0\\myarcsec 6$. The depth of KIDS will yield about $10^8$ galaxies to a median redshift of $z\\approx 0.8$.  With the KIDS data set at hand we will be able to measure the angular power spectrum on large scales. The statistical errors of cosmological parameter estimations will be reduced at least by a factor of ten, not only because the number of background galaxies is larger but also the number of galaxy pairs will increase dramatically, in particular for large angular scales. With such a large data set one can start to probe the equation of state of dark energy. \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[width=\\textwidth,clip]{sigma8vergleich}} \\caption{ Recent determinations of $\\sigma_8(\\Omega_{\\rm m}=0.3)$ using galaxy clusters (triangles) and cosmic shear (squares) in comparison to our results ($M_{\\rm ap}$: solid hexagon with bold error bars, $\\xi_\\pm$: hexagon with light error bars) and the WMAP three year result (pentagon). The open triangle and open star on the right are the average of all $\\sigma_8$ determinations between 2002 and 2006 (indicated by the vertical lines) using clusters and cosmic shear, respectively (error bars of single measurements are not taken into account). The WMAP result of $\\sigma_8$ would be larger if $\\Omega_{\\rm m}$ is fixed to $0.3$. Within a year the measurement points are not in chronological order. Values taken from table \\ref{tab:sigma8vergleich}. Points from right to the left in the diagram are associated with values from top down in the table. \\label{fig:sigma8vergleich} } \\end{center} \\end{figure} \\begin{table} \\caption{ Recent determinations of $\\sigma_8$ using clusters, cosmic shear and WMAP. The WMAP result of $\\sigma_8$ would be larger if $\\Omega_{\\rm m}$ is fixed to $0.3$. The result from \\citet{crh04} is obtained from a radio survey assuming that the median redshift is $z_{\\rm m}=2.2$. } \\label{tab:sigma8vergleich} \\begin{center} \\begin{tabular}{l|c|l} \\hline\\hline Reference & $\\sigma_8$ & Method \\\\ \\hline our measurements & $0.80 \\pm 0.10$ & WL: $\\langle M_{\\rm ap}^2\\rangle$ \\\\ & $0.93 \\pm 0.14$ & WL: $\\xi_\\pm$ \\\\ \\citet{smv06} & $0.85 \\pm 0.06$  & WL: $\\xi^{\\rm E}$  \\\\ \\citet{hmv05} & $0.85 \\pm 0.05$  & WL: $\\xi^{\\rm E}$\\\\ \\citet{mrb05} & $1.02 \\pm 0.15$ & WL: $\\xi_\\pm$ \\\\ \\citet{jjb05} & $0.72^{+0.08}_{-0.07}$ & WL: \\\\ \\citet{hbb05} & $0.68 \\pm 0.13$  & WL: $\\xi_\\pm$\\\\ \\citet{vmh05} & $0.83 \\pm 0.07$ & WL: $\\xi^{\\rm E}$\\\\ \\citet{hbh04} & $0.67 \\pm 0.10$ & WL: $$\\\\ \\citet{rrc04} & $1.02 \\pm 0.16$ & WL: $\\langle \\gamma^2\\rangle$ \\\\ \\citet{crh04} & $1.0 \\pm 0.2$   & WL: $\\langle M_{\\rm ap}^2\\rangle$\\\\ \\citet{bmr03} & $0.97 \\pm 0.13$ & WL: \\\\ \\citet{btb03} & $0.72 \\pm 0.09$ & WL: \\\\ \\citet{hms03} & $0.78^{+0.55}_{-0.25}$ & WL: $\\langle M_{\\rm ap}^2\\rangle$ \\\\ \\citet{jbf03} & $0.71^{+0.06}_{-0.08}$ & WL: \\\\ \\citet{hyg02} & $0.86^{+0.09}_{-0.13}$ & WL: $\\langle M_{\\rm ap}^2\\rangle$  \\\\ \\citet{hyg02b}& $0.81^{+0.07}_{-0.10}$ & WL: $\\langle \\gamma^2\\rangle$ \\\\ \\citet{vmp02} & $0.98 \\pm 0.06$ & WL: $\\langle M_{\\rm ap}^2\\rangle$ \\\\ \\citet{rrg02} & $0.94 \\pm 0.24$ & WL:  \\\\ \\citet{rrg01} & $0.91 ^{+0.25}_{-0.30}$ & WL: \\\\ \\citet{mvm01} & $1.04 \\pm 0.05$ & WL:   \\\\ \\hline \\citet{hen04} & $0.62 \\pm 0.04$ & X-ray clusters \\\\ \\citet{asf03} & $0.70 \\pm 0.04$ & X-ray clusters \\\\ \\citet{bdb03} & $0.72 \\pm 0.06$ & X-ray clusters  \\\\ \\citet{sbc03} & $0.76 \\pm 0.01$ & X-ray clusters \\\\ \\citet{pbd03} & $0.77 \\pm 0.05$ & X-ray clusters \\\\ \\citet{pbd03} & $0.78 \\pm 0.06$ & X-ray clusters \\\\ \\citet{vkl03} & $0.84^{+0.16}_{-0.03}$ & X-ray clusters  \\\\ \\citet{vnl02} & $0.61 \\pm 0.05$ & X-ray clusters \\\\ \\citet{reb02} & $0.68 \\pm 0.04$ & X-ray clusters  \\\\ \\citet{rbn02} & $0.72 \\pm 0.02$ & X-ray clusters  \\\\ \\citet{sel02} & $0.76 \\pm 0.06$ & X-ray clusters  \\\\ \\citet{wu01}  & $0.91 \\pm 0.11$ & X-ray clusters \\\\ \\citet{pbd01} & $1.02 \\pm 0.07$ & X-ray clusters  \\\\ \\citet{oua01} & $0.91$ & X-ray clusters  \\\\ \\citet{bsb00} & $0.96$ & X-ray clusters  \\\\ \\hline \\citet{sbd06} & $0.74^{+0.05}_{-0.06}$ & WMAP \\\\ \\end{tabular} \\end{center} \\end{table}"
    },
    "0606/astro-ph0606092_arXiv.txt": {
        "abstract": "We use data from the Northern Sky Variability Survey (NSVS), obtained from the first generation Robotic Optical Transient Search Experiment (ROTSE-I), to identify and study RR Lyrae variable stars in the solar neighborhood.  We initially identified 1197 RRab (RR0) candidate stars brighter than the ROTSE median magnitude $V = 14$. Periods, amplitudes, and mean $V$ magnitudes are determined for a subset of 1188 RRab stars with well defined light curves. Metallicities are determined for 589 stars by the Fourier parameter method and by the relationship between period, amplitude, and [Fe/H]. We comment upon the difficulties of clearly classifying RRc (RR1) variables in the NSVS dataset.  Distances to the RRab stars are calculated using an adopted luminosity-metallicity relation with corrections for interstellar extinction.  The 589 RRab stars in our final sample are used to study the properties of the RRab population within 5 kpc of the Sun.  The Bailey diagram of period versus amplitude shows that the largest component of this sample belongs to Oosterhoff type I.  Metal-rich ($[Fe/H] > -1$) RRab stars appear to be associated with the Galactic disk.  Our metal-rich RRab sample may include a thin disk as well as a thick disk population, although the uncertainties are too large to establish this.  There is some evidence among the metal-rich RRab stars for a decline in scale height with increasing [Fe/H], as was found by \\citet{Layden:1995}.  The distribution of RRab stars with $-1 < [Fe/H] < -1.25$ indicates that within this metallicity range the RRab stars are a mixture of stars belonging to halo and disk populations. ",
        "introduction": "RR Lyrae variable stars are versatile objects for astronomical research.  Not only are they excellent distance indicators, but they are useful as probes for understanding Galactic evolution and structure.  RR Lyrae variables have the virtues of being both luminous and ubiquitous in the Galaxy, tracing the old stellar populations of the bulge, disk, and halo components.  In this paper, we present results from a new survey of Bailey ab-type RR Lyrae stars (RRab or RR0) in the solar neighborhood based upon data in the Northern Sky Variability Survey \\citep{Wozniak:2004a}, hereafter NSVS.  Bailey type c (RRc or RR1) variables in the NSVS will be discussed in a future paper.  To get a clearer picture of the Galaxy, we need as complete a database as possible of both field and globular cluster RR Lyrae stars.  The General Catalog of Variable Stars (GCVS) \\citep{Kholopov:1996} summarizes much previous work on the discovery of RR Lyrae stars in the Galactic field, but is significantly incomplete, especially for RR Lyrae stars of low amplitude \\citep{Akerlof:2000}.  More recent CCD surveys of RR Lyrae stars in the Galactic bulge (OGLE: \\citet{OGLE}, MACHO: \\citet{MACHO2}) and halo (SDSS: \\citet{SDSS}, QUEST: \\citet{Vivas:2004}) have extended the coverage of the GCVS, but do not emphasize the discovery of brighter RR Lyrae stars in the solar neighborhood.  The Northern Sky Variability Survey \\citep{Wozniak:2004a} is based on photometric observations obtained with the first generation telescope of the Robotic Optical Transient Search Experiment (ROTSE-I) (see section 2).  Within the NSVS, RR Lyrae stars as faint as 15th magnitude in V can be detected, extending to a distance of about 7-9 kpc from the Sun \\citep{Akerlof:2000}.  This region includes part of the inner halo and thick disk components of the Galaxy. Hence, this survey is an excellent complement to the bulge and outer halo surveys, bridging the portions of the Galaxy that they cover.  The NSVS data complement those of the QUEST survey \\citep{Vivas:2004}, which includes RR Lyrae stars from 4 to 60 from the sun, but principally samples RR Lyrae stars more distant than those in the NSVS.  Another survey that will complement the NSVS is the All Sky Automated Survey (ASAS) \\citep{Pojmanski:1997}.  The ASAS database surveys the southern sky up to $\\delta = +28$ and has a magnitude limit of $V = 14$.  In this paper, we will limit ourselves to a consideration of those NSVS RRab stars that are brighter than 14th magnitude in the ROTSE-I system.  These stars have light curves sufficiently well defined that the periods are secure, and the classification of Bailey type is clear.  Within certain limits, which we shall describe later, we expect this survey to provide a more complete and unbiased catalog of the RRab stars in the solar neighborhood than is provided by the GCVS. Because of this magnitude restriction, our paper mainly concerns RRab stars to a distance of 5 kpc from the Sun.  Section 2 of the paper briefly describes the ROTSE-I telescope and the NSVS database.  In section 3, we discuss the selection criteria used to identify RRab stars, the determination of periods, the photometric calibration, and methods for determining photometric metallicities.  In section 4, the properties of the RRab sample and its distribution in space are discussed. Results are summarized in section 5. ",
        "conclusions": "The source of the RRab sample in this work primarily comes from the 2MASS correlated version of the NSVS database.  Aside from the constraints on coverage mentioned earlier, the NSVS database should provide a complete catalog of field RR Lyrae stars within the solar neighborhood down to $V \\sim 15$.  As a check, we took an early RRab candidate list that contained 1188 stars and cross correlated it with the GCVS \\citep{Kholopov:1996}.  We took into consideration that some of the stars in the GCVS have poor astrometric positions, with uncertainties up to 1 arcminute.  To match the NSVS RRab candidates with the GCVS RRab stars, our only option was to perform a positional match.  A search grid with sizes varying from 2 arcminutes to 3 arcseconds was used. At the match limit of 3 arcseconds, only 133 NSVS RRab stars were found in the GCVS. Thus, out of 1188 NSVS RRab candidates, almost 90\\% were not found in the GCVS catalog.  For a 2 arcminute positional match, 312 matches occurred, and an 1 arcminute search grid yielded 291 matches. Thus, it is safe to conclude that at least three quarters of the NSVS RRab candidates are not in the GCVS.  We also did a reverse search of the northern GCVS RRab stars in the NSVS database that had been correlated with the 2MASS catalog. We expected that all bright GCVS RRab stars with declination greater than $-30^{\\circ}$ should be in the NSVS database.  We also selected GCVS stars that were between 8th and 14th magnitudes in V.  In this reverse search, 99 GCVS stars were not recovered in our working sample of RRab stars.  To investigate these missing RRab stars, we searched further in the 2MASS correlated NSVS database, as well as in the original NSVS database located at Los Alamos.  All but seven of the GCVS stars could be found in the LANL version of the database.  The reasons for the absence of these seven variable stars (HP Aqr, AX Del, SZ Leo, DI Leo, NQ Lyr, AK Pup, and FI Sge) is unknown at this time.  The search of the 2MASS correlated database recovered 63 of the 99 missing variable stars.  Forty-three GCVS stars were excluded because their entries had fewer than 20 good observations.  Nine other missing stars had a Galactic latitude within $12^{\\circ}$ of the plane.  As noted in \\citet{Wozniak:2004a}, some confusion can exist in the identification of individual stars near the Galactic plane because of crowding.  Eleven stars were excluded due to the selection criteria applied to periods, colors, and amplitudes in identifying RRab candidates as described in section 3.1.  \\citet{Wils:2006} have recently completed a catalog of RR Lyrae stars in the NSVS which was done independently of the present survey.  They identified 785 probable RR Lyrae stars, including 188 that were previously unknown.  Of the 785, 712 are of RRab type, about 468 of which are brighter than magnitude 14. \\citet{Wils:2006} also provide a convenient table for cross identifying known RR Lyrae stars with those in the NSVS.  In comparing our RRab candidates with those in \\citet{Wils:2006}, we find that most of the Wils et al. stars are also in our list, but not all.  We also note that, because of the synonym problem in the NSVS, the stars in our Table \\ref{rrab_properties} are occasionally listed under different NSVS identification numbers than the corresponding variables in \\citet{Wils:2006}.  There are some 60 RRab stars in the Wils et al. list that are brighter than 14 magnitude that are not included in Table \\ref{rrab_properties}.  At least 50 of these appear to be genuine RRab stars that were excluded by one of the selection criteria we used.  Most often the period criteria caused the omission of the star from our sample.  For the remaining 10 stars there is still some uncertainty as to whether an RRab classification is correct.  For consistency, we have not included the additional Wils et al. stars in our analysis, but have retained the original sample from our search criteria.   \\subsection{Period Distribution}  The period distribution of 1188 NSVS RRab stars was examined through a period histogram, shown in Figure \\ref{rrabhist}.  In this histogram, the average period of the RRab stars is $0.563 \\pm 0.001$ days.  This mean value is close to the average period of 0.55 days seen for RRab stars of Oosterhoff type I globular clusters, and is much shorter than the 0.64 days typically seen in clusters of Oosterhoff type II \\citep{Smith:95}.  However, it is clear that the distribution of periods among the field RRab stars is wider than that seen within individual Oosterhoff type I clusters.  Instead, the field RRab period histogram is a consequence of the mixture of several RR Lyrae populations, as can be seen from the distribution of the RRab sample with respect to the Galactic plane.  The outer halo star studies of \\citet{Carney:1996} and cluster studies of \\citet{Lee:1999} hypothesized different origins for Oosterhoff type I and Oosterhoff type II clusters.  \\citet{Lee:1999} used a sample of 125 RRab stars within 3 kpc of the plane and 61 RRab stars farther than 5 kpc from the plane to examine the relative distributions of halo RRab stars belonging to Oosterhoff group I and Oosterhoff group II. \\citet{Lee:1999} concluded that there was evidence that field RRab stars associated with an Oosterhoff II population were more likely to be found in the low Z group, being rarer at $|Z| > 5$ kpc.  Because our sample of RRab stars becomes seriously incomplete as we move to a distance beyond 5 kpc, we cannot extend our survey to distances as large as those studied by \\citet{Lee:1999}.  However, we can investigate trends in RRab populations within 5 kpc of the plane. The RRab sample was divided into regions defined by the distance from the Galactic plane, Z.   Figure \\ref{rrabhist_z} shows the period distribution for RRab stars found close to the plane, $|Z| < 2$, kpc and at larger distances, $2 < |Z| < 5$ kpc. The average period of RRab stars found close to the plane shows a mean period indicative of an Oosterhoff I class.  The average RRab period for the region $|Z| < 2$ kpc was $0.557 \\pm 0.001$ days.  For the regions slightly further away, the average period shifted to an intermediate value between the Oosterhoff I and II types.  The average RRab period for this group of stars in these regions was $0.580 \\pm 0.001$ days. It is clear, however, that the mean period in the low $|Z|$ sample is strongly influenced by the presence of a short period disk population of RRab stars.  To winnow these from the sample requires a more detailed consideration of the RRab properties, which we discuss in the next section.  \\subsection{Period-Amplitude Trends}  A Bailey, or period-amplitude, diagram can be used as a diagnostic tool to investigate the Oosterhoff classification of RRab stars.  \\citet{Preston:59} discovered that the location of field RRab stars in the Bailey diagram was a function of metallicity. \\citet{SKS:1981} noted that the same effect among the RRab stars of globular clusters might be explained if metal-poor RRab stars were more luminous than their metal-rich counterparts. \\citet{Clement:2000} furthered the argument that RRab stars of Oosterhoff type I and II occupy separate and distinctive lines in the Bailey diagram.  We will use the linear trends from \\citet{Clement:2000}, which are based on the RRab stars found in M3 (Oosterhoff I type globular cluster) and $\\omega$ Cen (an Oosterhoff II cluster) as representative of the period-amplitude relations of Oosterhoff I and II clusters.  M3, at [Fe/H]$= -1.5$, is one of the most metal-poor of the Oosterhoff I clusters.  Thus, we would expect field RRab stars with Oosterhoff I properties to lie along or to the left of the M3 locus in the Bailey diagram.  We would expect RRab stars of Oosterhoff type II to lie along, or perhaps slightly to the right, of the $\\omega$ Cen locus in the Bailey diagram.  Figure \\ref{padiag_rrab} is our period-amplitude diagram for the 608 stars for which we had determined the period and amplitude. In this Figure we also overplot the linear Oosterhoff relations of \\citet{Clement:2000}.  The RRab stars do not clearly fall along one or the other relation, as is usually seen in the cases for globular clusters.  However, it appears that a majority of the stars fall near the Oosterhoff I relation, which is in agreement with the result from the period histograms.  \\citet{Cacciari:2005} obtained a mean period-amplitude relation for RRab stars in M3 that is somewhat different from that in \\citet{Clement:2000}.  In particular, Cacciari et al. noted that the more numerous regular and the rarer, more evolved M3 RRab stars occupied somewhat different positions in the period-amplitude diagram.  The Cacciari et al. trend line for the regular (less evolved) RRab stars is generally similar to that of Clement and Rowe, but is slightly displaced and includes a quadratic term, so that the line flattens slightly at large amplitudes.  Use of the Cacciari et al. rather than the Clement and Rowe Oosterhoff I line would not, however, change any of the conclusions we arrive at below. Even with the revised trend lines, significant populations exist near the Oosterhoff II trend line and at short periods to the left of the M3 Oosterhoff I line.  \\subsection{Oosterhoff Dichotomy Classification for NSVS RRab Stars}  A more careful classification for the NSVS RRab stars can be performed with the addition of metallicity information to the period and amplitude.  Three groups of stars are identified in our period-amplitude diagram (Figure \\ref{pa_oo}): the Oosterhoff I RRab stars, the Oosterhoff II RRab stars, and short period group of RRab stars. Figure \\ref{pa_oo} is a reproduction of Figure \\ref{padiag_rrab} but with the Oosterhoff groups identified.  The short period stars in the ``metal rich'' box of Figure \\ref{pa_oo} were scrutinized for RRc stars contaminating the RRab sample.  Almost all of these short period stars were found to be genuine ab-type RR Lyrae stars with metallicities of [Fe/H] $> -1$.  The right edge of this box is located at $[Fe/H] = -1$, as defined by the Sandage metallicity relation (Equation \\ref{ampmet}).  All the stars with the filled circles were identified with a metallicity richer than [Fe/H] $= -1$, according to the best estimate metal abundance in Table \\ref{rrab_feh}.  To examine the spatial distributions of the two Oosterhoff groups and the metal-rich group of RRab stars, we plot in Figure \\ref{feh_z} the [Fe/H] distribution versus the distance from the Galactic plane, $|Z|$.  In Figure \\ref{tdmetz}, we see that the metal rich stars with [Fe/H] $> -1$ all lie close to the Galactic plane.  The metal poor stars exhibit an extended distribution more consistent with the existence of a component with halo properties.  However, below we will argue that a fraction of the RRab stars more metal-poor than [Fe/H] $= -1$ belongs to a thick disk population.  Due to the faintness limit we imposed for our sample of RRab stars, our RRab stars probe to about 4.5 kpc away from the Galactic plane.  \\subsubsection{Oosterhoff I and II}  We divide the RRab stars in Figure \\ref{pa_oo}, excluding the metal rich stars confined in the box, along the dotted line.  Stars to the left of this line are denoted as belonging to Oosterhoff I, while those to the right are credited to Oosterhoff II.  Ideally, we might want to draw the dotted line to correspond to a metal abundance of [Fe/H]$ =-1.7$, which is the approximate boundary between globular clusters of Oosterhoff types I and II.  However, if we use Sandage's amplitude-log period-[Fe/H] relation to identify the location of the [Fe/H]$ = -1.7$ boundary, the dividing line intersects with the Oosterhoff I relation of \\citet{Clement:2000}.  This may suggest a small problem either with Sandage's amplitude-period-metallicity calibration, or with the adopted M3 trend line.  \\subsubsection{Metal Weak Thick Disk Population?}  The Oosterhoff I candidates were separated into two subgroups with the division occurring at $[Fe/H] = -1.25$.  Note in Figure \\ref{oo1subgroup} the existence of a relatively large number of stars of [Fe/H] $> -1.25$ within 2 kpc of the Galactic plane. To investigate the implications of this group of stars, we first consider whether the apparent overabundance of disk stars in the more metal-rich Oosterhoff I group might reflect a bias in the identification of RRab stars.  We used the more metal-poor Oosterhoff II sample as a control, assuming that it represents a pure or nearly pure halo population.  The detection probabilities of the Oosterhoff II RRab stars are similar to those of Oosterhoff I RRab of comparable amplitude within our survey.  We compared the number of Oosterhoff I and II type stars at different $|Z|$ distance intervals.  Table \\ref{numratio} summarizes our calculated number ratios of these stars.  We assume that the two $|Z|$ distance intervals listed in Table \\ref{numratio} best span the disk and inner halo population of our sample of stars.  The regions above 5 kpc are excluded because of our adopted magnitude limit.  In the region closer to the plane ($|Z| < 2.0$ kpc), we also have an additional constraint that the RRab stars have a Galactic latitude greater than $12^{\\circ}$, i.e. $|b| > 12^{\\circ}$.  This constraint was imposed to avoid the regions of heavy interstellar extinction near the plane where the \\citet{Schlegel:1998} reddening maps are uncertain and where our survey incompleteness becomes more serious.  Our number ratio calculations suggest that different Oosterhoff groups are dominant in the two $|Z|$ regions.  In the case of Oosterhoff II (Oo II) to Oosterhoff I (Oo I), the ratio shows that the Oo I stars outnumber the Oo II stars 2 to 1 for the $2.5 < |Z| < 3.5$ kpc region.  However, closer to the plane, the same ratio seems to indicate that Oo I RRab stars increase more rapidly as we approach the plane than is the number of Oo II RRab stars. This result is certainly consistent with the presence of a disk component within our Oosterhoff I group.  We repeat this analysis to investigate the disk and halo populations, but divide the Oosterhoff I sample into two subgroups by metallicity.  The metal richer Oosterhoff I subgroup ([Fe/H]$ > -1.25$) and the metal poorer Oosterhoff I subgroup ([Fe/H]$ < -1.25$) were each compared with the Oosterhoff II group at the same $|Z|$ distance intervals as in the above analysis.  Table \\ref{oo_numbers} provides the numbers of RRab stars assigned to each Oosterhoff group at each $|Z|$ interval.  Table \\ref{ooI_num} lists the number of RRab stars that make up the two Oosterhoff I subgroups.  These number ratios suggest that the disk component of the RRab sample likely extends to metal abundances somewhat lower than our original [Fe/H] $= -1.0$ cutoff.  The numbers of stars in the samples are not so great as to allow this conclusion to be drawn at more than the two sigma level.  Full kinematic data on the sample stars would allow this conclusion to be tested more fully.  Taking the ratios at face value, however, we estimate that approximately 60\\% of our metal-rich Oosterhoff I stars in the region $|Z| <$ 2.0 kpc belong to the thick disk.  There is even some evidence for a metal-weak thick disk component \\citep{Norris:1985, Morrison:1990, Twarog:1994, Beers:1995} but the sample is not large enough to make a completely convincing case.  We note that uncertainties in our derived values of [Fe/H] can scatter metal-rich stars into our metal weak sample and vice versa.  The fact that there are very few stars in our [Fe/H] $> -1$ sample that also have $|Z| > 2$ kpc indicates that this effect is not large and does not account for the overabundance of stars near the disk in our Oosterhoff I metal-rich group.  Such an effect might, however, contribute to the smaller excess near the disk among the Oosterhoff I metal-poor group.  In order to clearly identify whether a star belongs in the thick disk or halo population, the kinematics of the star must be known.  The kinematic information can provide stronger evidence for testing whether there exists a thick disk population in our Oosterhoff I group of stars.  A program to obtain full space motions for these, few of which now have known radial velocities, would certainly be of value.  We note that there is a continuum of properties among these field RRab stars. The stars with higher metal abundances as determined by the Fourier parameter method or (when such are available) by Layden's (1994) spectroscopy, lie to the left in the period-amplitude diagram, whereas the more metal-poor stars lie to the right.  The apparently metal-rich RRab stars in the unusual globular clusters NGC 6388 and NGC 6441 \\citep{Layden:1999, Pritzl:2001,Pritzl:2002, Pritzl:2003, Clementini:2005} break this trend: falling to the right in a period-amplitude diagram although NGC 6388 and NGC 6441 appear to be about as metal-rich as 47 Tucanae.  Thus, there is some similarity in the structure of the RRab stars in the solar neighborhood field, both metal-rich and metal-poor, that may extend to most but not all of their globular cluster counterparts.  \\subsection{Thick Disk Stars}  The metal rich star distribution in Figure \\ref{tdmetz} shows that most of these stars are close to the Galactic plane.  The distribution of these stars with respect to $|Z|$ is broadly consistent with that expected of a thick disk population.  However, the distribution of these stars appears to be a function of metal abundance.  Although there are a few exceptions, there is a trend toward lower $|Z|$ distance as [Fe/H] goes from -1 to 0.  This may indicate the existence of a thin, or at least a less thick, disk component among the most metal-rich of the RRab stars.  To check whether this thick disk trend might be an artifact of our choice of $M_{V}-[Fe/H]$ relation, we rederived the distances using fixed $M_{V}$ values (+0.6 and +0.71, respectively \\citet{Smith:95, Layden:1996}).  Regardless of how we arrived at the distances, the trend was still present, as can be seen in Figure \\ref{td_distfeh}.  Thus, the tendency toward decreasing $|Z|$ with increasing [Fe/H] does not seem to be imposed on the data by our particular choice of $M_{V}-[Fe/H]$ relation.  \\subsubsection{Scale Height of the Thick Disk} Using the 589 RRab stars for which we are able to derive reliable photometric metallicities, we have investigated the scale heights of the thick disk and inner halo for different subgroupings within the data.  To derive scale height, h(Z), of these groups we binned data for each subgroup into an average of either 6 or 10 stars per $|Z|$ bin.  For the thick disk scale height calculation, we considered the full sample as well as a subsample that excluded stars found close to the plane ($|b| < 12^{\\circ}$).  In the latter case, we also implemented a limit to the volume in which we calculated the densities to be used in the determination of the scale heights.  The adopted limits were $|b| > 12^{\\circ}$ and a radius of 2.0 kpc from the Sun.  Therefore, the shapes of the volumes were essentially frustums of a cone until we reached the search radius limit of 2.0 kpc, which then reverted to the volume of a cylinder.  A correction had to be applied to these volumes for areas of the sky that were too far south to be included within the ROTSE-I survey.  In Figure \\ref{aitoff}, this region can be seen as a hole where no RRab stars were detected for our sample.  This region is bounded by Galactic longitude $240^{\\circ} < l < 20^{\\circ}$ and Galactic latitude $-90^{\\circ} < b < +30^{\\circ}$.  Once the densities were calculated, they were binned in $|Z|$, and an exponential function was fitted to the results.  Figures \\ref{logdensity6} and \\ref{logdensity10} show how the log density falls off with the $|Z|$ distance bin for the full sample of 70 metal-rich stars. When the constraints on volume and $|b|$ are applied, 55 of these metal-rich RRab stars remain in the sample. For comparison, we also performed similar scale height calculations for 330 RRab stars within the Oosterhoff I group and for a ``halo'' sample of 428 RRab stars with $[Fe/H] < -1.25$. The stars selected for these groups were subject to the same volume limitations as those in the metal-rich group.  Although it would have been interesting to separately determine scale heights for the Oosterhoff II stars and the more metal-rich and metal-poor Oosterhoff I groups, the numbers of stars within these subgroupings are insufficient to give robust results when the volume constraints are applied.  The halo sample results are, in any case, given merely for completeness. The RRab sample we used does not go deeply enough into the halo to yield a reliable value of scale height.  Results are shown in Table \\ref{sclhgt}.  Figures \\ref{sclhgt6} and \\ref{sclhgt10} show plots of the density of the metal-rich thick disk stars with $|Z|$, where the adopted error bars are indicated.  The weighting of the density points is important to our scale height solution.  If we apply no weighting, the calculated scale heights for the metal-rich sample are much smaller, about 0.3 to 0.4 kpc, as reported in \\citet{Kinemuchi:2005}.  We believe that the scale height calculation including weights depending on the uncertainties in distance and Poisson statistics are the more reliable.  The scale height derived from the 6 stars per $|Z|$ bin case was $0.65 \\pm 0.17$ kpc (Figure \\ref{sclhgt6}).  For the 10 stars per $|Z|$ bin, the scale height was $0.68 \\pm 0.18$ kpc for our thick disk sample of 70 stars (Figure \\ref{sclhgt10}).  These scale heights are smaller than the canonical scale height of a pure thick disk sample ($\\sim 1$ kpc).  Applying the volume constraints gives slightly greater, but also more uncertain results, around 0.8 kpc.  In evaluating these results we must consider whether there exists any selection bias against the identification of RRab variables at large $|Z|$.  However, such a bias seems unlikely since there is no large falloff in the detection efficiency for RRab stars in the NSVS sample until one reaches $V > 14$, which we set as the limiting magnitude for our RRab sample.  Application of the efficiency corrections as a function of apparent magnitude found by \\citet{Amrose:2001} produces only a small change in the calculated metal-rich group scale height, yielding $0.66 \\pm 0.16$ kpc for the 6 stars per bin solution and $0.67 \\pm 0.17$ kpc for the 10 star per bin solution.  The correction from \\citet{Amrose:2001} is an upper limit on the actual correction expected for the stars used in calculating our scale heights since the  number of observed data points in this study is often greater than in the \\citet{Amrose:2001} simulation.  A second source of error can arise from the scattering of the stars across the boundary between the Oosterhoff I and metal-rich groups. However any such scattering would be expected to increase rather than decrease the metal-rich group scale height.  The calculated scale height also depends upon the adopted relationship between luminosity and metallicity.  However, alternative calibrations from the recent literature give similar results.  Because the metal-rich RR Lyrae stars in Figure \\ref{tdmetz} are distributed relatively close to the Galactic plane compared to the more metal-poor RRab stars, we already expected these stars to have a scale height indicative of a disk population.  The fit is, however, strongly influenced by the stars very close to the plane.  If the RRab stars within 400 pc of the plane are removed from the sample, leaving 45 RRab stars, and the calculation is repeated, the resultant scale heights are greater, about 1.1 kpc but with relatively large uncertainties. We can imagine three possible explanations for this: (1) contamination of the RRab sample close to the plane by some other type of variable star; (2) small number statistics; or (3) the presence of a mixture of thin and thick disk stars within the metal-rich RRab sample.  We believe that the first explanation is unlikely.  The light curves of the stars within this sample are characteristically those of RRab stars.  While the light curves of high amplitude $\\delta$ Scuti stars can occasionally resemble those of RRab stars, almost all such stars have periods shorter than $\\sim 0.3$ days.  The second explanation is more difficult to exclude.  The third explanation is the most intriguing of the three -- that the stars comprising our metal-rich sample are a mixture of stars belonging to an old thin disk and the thick disk.  When the scale height of the full metal-rich sample is calculated, the solution in this case would be influenced by the smaller scale height thin disk component.  Values in the literature for the scale height of the old thin disk typically run from about 240 to 330 pc (see \\citet{Chen:2001} and references therein).  In removing stars within 400 pc from the sample, we are removing many of the thin disk stars, leaving a solution dominated by the thick disk component. Values of the scale height of the old thick disk show some scatter, from as low as 0.6 to 0.7 \\citep{Chen:2001} to about 1.5 kpc \\citep{Gilmore:1983}, but are typically in the neighborhood of 1 kpc. Our value for the sample of metal-rich RRab stars more than 400 pc from the plane is comparable to the typical value of the thick disk scale height.  This is not the first suggestion that the metal-rich RRab stars contain a mixture of thin and thick disk stars.  \\citet{Layden:1995} obtained a scale height of 0.7 ($+0.5,-0.3$) kpc for RRab stars of $[Fe/H] > -1$.  However, Layden noted a tendency for the calculated scale height to decrease with increasing metallicity for stars within this group.  Although Layden cautioned that the trend was significant only at the one sigma level, he noted that his scale height was in between the values often quoted for the old thin and thick disk, and that this might indicated that his sample contained a mix of these stars.   One might, however, regard the division of the metal-rich RRab stars into two populations, thick and old thin disk, to be too simple.  There might instead just be a trend of decreasing scale height with increasing metallicity.  More recently, \\citet{Maintz:2005} studied the motions of 217 RR Lyrae stars brighter than $V = 12.5$.  They found evidence for a thick disk component, with a scale height of $1.3 \\pm 0.1$ kpc and a halo component of $4.6 \\pm 0.3$ kpc.  In considering their sample of RR Lyrae stars with $|Z| < 0.5$ kpc, they found some evidence for a thin disk component with a scale height of $0.38 \\pm 0.04$ kpc, though they caution that they cannot draw a firm distinction between possible thin and thick disk components.  Some of the stars in the \\citet{Maintz:2005} sample are common to our own, but our approaches to calculating the scale height are quite different so that the similarity in our results is significant. From a smaller sample of RRab stars \\citet{Amrose:2001} found a larger scale height ($1.8 \\pm 0.5$ kpc) for the thick disk.  Their approach, however, was also different from that adopted here, making no use of metallicity constraints.  A vertical scale height may not be the best method for describing the actual distribution of stars belonging to the Galactic halo.  The scale height of about 4 kpc found for the Oosterhoff type I sample is much larger than the scale height of our sample of the metal-rich RRab stars.  Figure \\ref{fig21} shows the scale height fit for the Oosterhoff I group.  Note that close to the Galactic plane, the data points tend to fall above the fit to the data as a whole.  This is again, suggestive of our mixture of thick disk and halo stars within our Oosterhoff type I sample.  Although our calculated scale height for the ``halo\" group is too uncertain to be very useful here, the \\citet{Maintz:2005} value of 4.6 kpc shows us that our value to be reasonable.  The presence of RR Lyrae stars is often taken as indicating the existence of a stellar population with an age of 10 Gyr or more.  If, in fact, some of the RRab stars in our sample belong to an old thin disk component, then they may have ages slightly lower than the canonical 10 Gyr.  \\citet{Delpeloso:2005} and \\citet{Hansen:2002} found the age of the old thin disk to be $8.3 \\pm 1.8$ Gyr and $7.3 \\pm 1.5$ Gyr, respectively.  Thus, it is possible that some of the RR Lyrae in our sample were formed as recently as 7 or 8 Gyr ago. However, only a small proportion of the relatively metal-rich red giants in the old thin disk appear to lose enough mass to produce RR Lyrae stars \\citep{Layden:1995, Taam:1976}.   \\subsubsection{$\\Delta \\log P$ and Metallicity Gradient}  In \\citet{SKK:1991} figure 8a, a plot of the period shift, $\\Delta \\log P$, with metallicity showed a clear separation of the field RR Lyraes into the Oosterhoff groups.  Since our period-amplitude diagram does not clearly show a sharp distinction between the two Oosterhoff groups, we plotted our sample of RRab stars in the same manner as \\citet{SKK:1991}.  We did not include any of the metal-rich stars ($[Fe/H] > -1$) that we have identified as our thick disk sample.  Figure \\ref{logpfeh} shows our $\\Delta \\log P-[Fe/H]$ plot, with approximately the same axes ranges as in \\citet{SKK:1991}'s figure 8a.  We do not see a gap at $\\Delta\\log P = -0.03$, corresponding to the region of the Oosterhoff gap, but rather a continuous trend from one group to the next Oosterhoff group. Admittedly, our uncertainties in [Fe/H] are larger on a star to star basis than those in Suntzeff et al., which will tend to erase any Oosterhoff gap.  We also looked for a metallicity gradient as a function of Galactocentric distance, R, in kpc.  We assume the Galactocentric distance of the Sun to be 8.5 kpc.  \\citet{SKK:1991} reported a metallicity gradient in their sample of field RR Lyraes, which spanned a region of $R = 4$ to 10 kpc.  Following their steps, we present our result in Figure \\ref{skkgrad}, however, our sample only covers a region of $R \\sim 6$ to 12 kpc.  In the region that overlaps with Suntzeff et al's work, we do not see a significant metallicity gradient.  We note that in Suntzeff et al.'s figure 8a, the greatest change in [Fe/H] occurs at roughly $R < 5$ kpc. Few stars within our sample are that close to the galactic center. Although we have tried to exclude obvious disk stars from Figure \\ref{logpfeh}, the \\citet{SKK:1991} RRab stars are usually fainter than the NSVS sample, and may be a more pure halo sample.  \\subsection{C-Type RR Lyrae Sample}  Although we have been been reporting our results for the NSVS RRab stars in this paper, we have also completed a preliminary search for Bailey type c RR Lyrae (RRc) in the the NSVS database.  This search for RRc stars was done in conjunction with the search for RRab stars, but with different selection criteria than those outlined in section 3.1.  However, the 2MASS correlated database from which we obtained our RRab sample was optimized to find RRab stars and not the lower amplitude, shorter period RRc stars.  We have found that many of the GCVS RRc stars were omitted from our preliminary RRc sample.  A thorough search and analysis of the NSVS RRc stars will be conducted at a later date.  In this section, we will confine ourselves to a description of some of the problems encountered with the identification of RRc stars in the NSVS database.  We initially obtained 2558 RRc candidates from the 2MASS correlated NSVS database.  As with the RRab candidates, we removed duplicate entries but kept the entry with the most observations.  Our selection criteria for the RRc candidate stars are listed in Table \\ref{selection_rrc}.  Periods were obtained using the Supersmoother routine and were compared to the period solutions from the cubic spline method.  After a period was chosen for each star, we visually inspected these candidates and removed those stars with phased light curves of an eclipsing binary or phased light curves of poor quality.  Amplitudes and mean intensity-weighted magnitudes were calculated for the remaining RRc candidates using the spline routine used for RRab stars.  Since our relation for scaling amplitudes was constructed for both RRab and RRc stars, we applied it to scale the amplitudes of the RRc candidates. All of the NSVS RRc parameters will be published in a future paper dealing specifically with these RR Lyrae stars.  We encountered several difficulties in selecting an appropriate RRc star sample that were more severe than in the case of the corresponding RRab stars.  The light curve shape typical of RRc stars is not as distinctive as that of RRab stars.  In particular, when the light curves are noisy, there can be confusion between the RRc stars, W UMa type eclipsing binary stars, and short period $\\delta$ Scuti/SX Phoenicis variable stars.  Moreover, there are fewer RRc stars with excellent NSVS light curves that also have well observed light curves in the standard Johnson $V$ bandpass.  Thus, we have not yet been able to compare $\\phi_{31}$ values on the NSVS system to standard $V$ band $\\phi_{31}$ values in a satisfactory manner for the RRc stars.  Figure \\ref{rrchist} shows the period histogram for a sample of 375 RRc candidates as selected by the criteria in Table \\ref{selection_rrc}.  We also used the selection criteria as described in \\citet{Akerlof:2000} to arrive at this sample.  The sharp increase in the histogram toward shorter periods is, however, suspicious, indicating a possible contamination of the RRc stars by longer period, larger amplitude $\\delta$ Scuti type stars.  This possible contamination is also evident in Figure \\ref{pa_rrc}, in which the RRc candidates are included in the period-amplitude diagram.  The concentration of stars toward the short period and low amplitude corner of the plot may be spurious.  So far, attempts to clearly separate RRc stars from the non-RR Lyrae stars using Fourier decomposition parameters, periods, and amplitudes have not been entirely successful.  \\citet{Wozniak:2004b} have used machine learning techniques to classify long period variable stars, and this method may in the future help with the classification of the short period pulsating variables.  The NSVS RRc variables will therefore be discussed further in a second paper."
    },
    "0606/astro-ph0606098_arXiv.txt": {
        "abstract": "We investigate how well the redshift distributions of galaxies sorted into photometric redshift bins can be determined from the galaxy angular two-point correlation functions.  We find that the uncertainty in the reconstructed redshift distributions depends critically on the number of parameters used in each redshift bin and the range of angular scales used, but not on the number of photometric redshift bins.  Using six parameters for each photometric redshift bin, and restricting ourselves to angular scales over which the galaxy number counts are normally distributed, we find that errors in the reconstructed redshift distributions are large; i.e., they would be the dominant source of uncertainty in cosmological parameters estimated from otherwise ideal weak lensing or baryon acoustic oscillation data.  However, either by reducing the number of free parameters in each redshift bin, or by (unjustifiably) applying our Gaussian analysis into the non-Gaussian regime, we find that the correlation functions can be used to reconstruct the redshift distributions with moderate precision; e.g., with mean redshifts determined to $\\sim 0.01$.  We also find that dividing the galaxies into two spectral types, and thereby doubling the number of redshift distribution parameters, can result in a reduction in the errors in the combined redshift distributions. ",
        "introduction": "There are many different techniques for determining the distance--redshift relation and/or growth--redshift relation motivated by the desire to understand the dark energy.  Those that rely on the distances to a relatively small number of objects, such as the Type Ia supernova method \\citep[e.g.][]{riess98}, can use spectroscopic redshift determinations and thus avoid redshift error as a significant source of uncertainty.  However, when the distance (and/or growth) constraints are derived from measurement of very large numbers of objects spectroscopy can be a practical impossibility.  In such cases one must rely on ``photometric redshifts''; i.e., redshifts estimated from photometry in multiple broad bands \\citep[e.g.][]{loh86,connolly95,sawicki97}.  The relatively low cost per object of imaging surveys compared to spectroscopic surveys is a great advantage and provides significant motivation for pursuing the technique of estimating photometric redshifts.  Imaging surveys can potentially constrain dark energy via a variety of techniques including cluster counting \\citep[e.g.][]{haiman01}, cosmic shear \\citep[e.g.][]{hu02d,huterer02} and baryon acoustic oscillations (BAO)~\\citep[e.g.][]{seo03,blake03,padmanabhan06}.  It may even be possible for imaging surveys to use Type Ia supernovae, without spectroscopic follow-up, to constrain cosmology \\citep{barris04}.  But abandoning spectroscopy has its disadavantages too.  In general, there is some tolerance of redshift error, but less tolerance for uncertainty about the probability distribution of those errors.  The impact of redshift uncertainties on dark energy constraints has been studied for supernovae \\citep{huterer04}, cluster number counts \\citep{huterer04}, weak lensing \\citep{bernstein04,huterer06,ishak05,ma06} and baryon oscillations \\citep{zhan06,zhan06b}.  All of the studies cited in the above paragraph model the error distribution as Gaussian. However, photometric redshift error distributions, due to spectral-type/redshift degeneracies, often have bimodal distributions, with one smaller peak separated from a larger peak by $\\Delta z$ of order unity \\citep[e.g.][]{benitez00,fernandez-soto01,fernandez-soto02}.  Thus a fraction of galaxies have photometric redshifts that are `catastrophically' wrong.  Here we study how well the coarse properties of the true redshift distribution of galaxies in a given photometric redshift bin can be reconstructed from galaxy two-point correlation functions.  The idea is that catastrophic photometric redshift (``\\phz{}'') errors introduce additional correlations between galaxies in different redshift bins.  In general, such errors will alter both the amplitude and shape of the binned angular correlation functions.  Measurements of the correlation functions over a range of angular scales would thus provide valuable information to unravel the effects of large \\phz{} errors.  We emphasize that we are not attempting a forecast of the \\phz{} errors achievable given all possible information.  In particular, we neglect information from spectroscopic calibration of the \\phz{} error distribution.  Spectroscopy, possibly combined with a ``super'' photometric (12 or more bands) photo-z training set will play a critical role\\footnote{ The current plan for \\phz{} calibration is described in http://www.lsst.org/Science/photo-z-plan.pdf.}. In this sense our forecasts here are highly conservative. Further, the catastrophic errors are likely to be avoided by use of luminosity function and surface brightness priors.  For recent results on spectroscopic calibration of \\phz{} measurements for weak lensing, see \\citet{ilbert06}.  The outline of this paper is as follows.  In \\S~\\ref{sc:method} we describe our model for the \\phz{} errors, the Fisher matrix we use to constrain the parameters of this model, and our model for the galaxy angular power spectra.  We present our results in \\S~\\ref{sc:results}, including the details of our fiducial model and its impact on the resulting Fisher matrix constraints.  We discuss some implications of our results in \\S~\\ref{sc:discussion} and draw conclusions on the feasibility of constraining \\phz{} errors with galaxy angular correlation functions. ",
        "conclusions": "}  We have shown that the ability to constrain general (i.e. non-Gaussian) \\phz{} error distributions with galaxy two-point correlation functions depends on the parameterization of the \\phz{} errors, the range of angular scales probed by the correlation function, and prior knowledge of the galaxy bias.  Binning the galaxy sample in \\phz{}, we have presented constraints on the binned redshift distribution and mean redshift in each \\phz{} bin. Parameterizing the redshift distribution by binned values is insensitive to small scatter from \\phz{} errors, but otherwise assumes no {\\it a priori} knowledge of the \\phz{} error distribution.  We find that reducing the number of parameters in each \\phz{} bin can be very helpful, which could be achieved with improved knowledge of the \\phz{} errors from, e.g., spectroscopically calibrated samples or luminosity function priors.  We have limited our use of the galaxy correlation function to angular scales where the galaxy number density is Gaussian distributed.  At low redshifts, this severely limits the amount of data available to constrain the \\phz{} error parameters.  We hypothesize that including information from correlations on non-Gaussian scales could significantly improve the constraints and demonstrate that the constraints on the mean redshift in each \\phz{} bin do improve by two orders of magnitude with a naive extrapolation of our Gaussian calculation to non-Gaussian scales.  If it is possible to separate the galaxies by spectral type, the constraints on the \\phz{} errors may improve further by including information from the cross-correlation of the galaxy sub-samples.  We have demonstrated this in figs.~\\ref{fg:LSSTbias} and \\ref{fg:redblue} by separating our fiducial galaxy sample into ``red'' and ``blue'' spectral types.  We expect this procedure to be particularly helpful if there exists a well-populated spectral class of galaxies whose \\phz{}'s can be estimated unusually well.  Our forecasts are limited to parameters of the \\phz{} error distribution and linear galaxy bias so we cannot make any rigorous conclusions about what kind of dark energy constraints can be achieved in weak lensing and baryon acoustic oscillation surveys with the level of \\phz{} errors forecasted here.  However, we make qualitative comparisons with dark energy forecasts in the literature~\\citep{huterer06,ma06,zhan06,zhan06b} using our constraints on the mean redshift in each \\phz{} bin given in table~\\ref{tb:zbar}.  In the Gaussian regime, the constraints we forecast of $\\sim0.01$ are factors of a few larger than those desired for upcoming dark energy surveys.  However, adding non-Gaussian scales in the correlation function may provide the required constraints.   The galaxy correlation properties are quite likely to provide at least a powerful consistency test for the redshift distributions as determined via spectroscopic and/or ``super'' (12 or more band) calibration subsamples."
    },
    "0606/astro-ph0606051_arXiv.txt": {
        "abstract": "CCD $UBVRI$ photometry is presented for type Ia supernova 2004fu in NGC 6949. The light and colour curves are typical for this class of objects, absolute magnitude at maximum and decline rate are in agreement with the relationship between these parameters, established for SNe Ia. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606267_arXiv.txt": {
        "abstract": "There are various analytical approaches to the mean electromotive force $\\bscE = \\langle \\bu \\x \\bbb \\rangle$ crucial in mean--field electrodynamics, with $\\bu$ and $\\bbb$ being velocity and magnetic field fluctuations. In most cases the traditional approach, restricted to the second--order correlation approximation, has been used. Its validity is only guaranteed for a range of conditions, which is narrow in view of many applications, e.g., in astrophysics. With the intention to have a wider range of applicability other approaches have been proposed which make use of the so--called $\\tau$--approximation, reducing correlations of third order in $\\bu$ and $\\bbb$ to such of second order. After explaining some basic features of the traditional approach a critical analysis of the approaches of that kind is given. It is shown that they lead in some cases to results which are in clear conflict with those of the traditional approach. It is argued that this indicates shortcomings of the $\\tau$--approaches and poses serious restrictions to their applicability. These shortcomings do not result from the basic assumption of the $\\tau$--approximation. Instead, they seem to originate in some simplifications made in order to derive $\\bscE$ without really solving the equations governing $\\bu$ and $\\bbb$. A starting point for a new approach is described which avoids the conflict. ",
        "introduction": "In mean--field electrodynamics the mean electromotive force $\\bscE = \\langle \\bu \\x \\bbb \\rangle$ due to the fluctuations $\\bu$ and $\\bbb$ of the fluid velocity and the magnetic field plays a crucial role \\citep{krauseetal71b,krauseetal80,moffatt78,raedler00b}. A central problem is the determination of $\\bscE$ for a given motion as a functional of the mean magnetic field. Various methods have been used for that.  One approach, which we call ``traditional approach\" or ``approach (i)\" in the following, was established together with mean--field electrodynamics at all \\citep{krauseetal71b,krauseetal80}. Most of the calculations of $\\bscE$ have been done on this basis using the so--called second--order correlation approximation (SOCA) or, what means the same, first--order smoothing approximation (FOSA). This approximation in its original form, that is, applied in the case of a purely hydrodynamic background turbulence, ignores all higher than second--order correlations in the fluctuations $\\bu$ of the velocity field (see section~\\ref{subsec31}). It can be justified only in cases in which these fluctuations are not too large. The usually given simple sufficient conditions for its validity are in view of astrophysical applications rather narrow. Basically it is possible to proceed to higher--order approximations but this requires tremendous efforts and has been done so far only in a few simple cases (Nicklaus and Stix 1988, Carvalho 1992, 1994, R\\\"adler {\\it et al.} 1997, see also section~\\ref{new1}). A slight modification of the second--order correlation approximation applies also to the case of a magnetohydrodynamic background turbulence (see section~\\ref{subsec32}).  In some more recent investigations (which are cited below) other approaches are used, which rely in a sense on the $\\tau$--approximation of turbulence theory \\citep{orszag70} and are called ``$\\tau$-approaches\" or ``approaches (ii)\" in the following. They go in so far beyond the second--order correlation approximation as they consider also higher--order correlations, which are then in the sense of a closure expressed by second--order ones. In approach (i) the relevant equations are simplified by a well--defined approximation and then really solved, and $\\bscE$ is calculated with these solutions. In the approaches (ii) a relation for $\\bscE$ is deduced from the original equations, but without really solving them. Instead, in order to get manageable results, assumptions on the connection of some of the occurring quantities with  $\\bscE$ are introduced. The final result for $\\bscE$ is to a large extent determined by these assumptions. The approaches (ii) cover from the very beginning also the case of a magnetohydrodynamic background turbulence.  The main purpose of this paper is a critical analysis of the approaches (ii). Each step of approach (i) can be justified by the underlying induction equation or, in the case of a magnetohydrodynamic background turbulence, induction equation and momentum balance. Clear, at least sufficient conditions for the applicability of the second--order approximation can be given. There is no doubt in the correctness of its results in the so defined range of applicability. It seems reasonable to assume, and we do so in this paper, that there is at least some overlap of the ranges of applicability of the approaches (i) and (ii). We have then to require that in these overlapping ranges the results of both approaches coincide. Simple versions of approach (ii) as used by \\citet{vainshteinetal83}, \\citet{blackmanetal02b} and \\cite{brandenburgetal05b}, called ``simple $\\tau$--approach\" or ``approach (iia)\" in the following, deliver results which do not in all cases satisfy this requirement. As we will show below the more sophisticated version used in the papers by \\cite{raedleretal03} and by \\citet{rogachevskiietal03,rogachevskiietal04}, called ``spectral $\\tau$--approach\" or ``approach (iib)\", does not satisfy this requirement, too. We have to conclude that these approaches are not in full accordance with the basic equations mentioned. Therefore the results can not be taken for granted. We will propose a starting point for an alternative approach which avoids conflicts with approach (i).  In section~\\ref{model} we define the frame of our considerations and deliver the basic equations. In section~\\ref{seci} we recall the fundamentals of approach (i) and review some of its basic results. In section~\\ref{secii} we explain the approaches (iia) and (iib), derive a few results, restricting attention to the simple case of a non--rotating fluid, and pinpoint shortcomings of these approaches and deviations of the results from those of approach (i). In section~\\ref{new} we explain a proposal for the alternative approach mentioned. Finally in Section~\\ref{sum} we summarize our findings. ",
        "conclusions": "\\label{sum}  The paper provides a critical view on different analytical ways to determine the mean electromotive force $\\bscE$ in mean--field electrodynamics. First some of the findings gained with the traditional approach, reduced to the second--order correlation approximation or approach (i), are summarized (section~\\ref{seci}). In this context also some new results concerning the case of magnetohydrodynamic background turbulence are given (section~\\ref{subsec32}). Then the essentials of two versions of the $\\tau$--approach, or approach (ii), are represented, that is, of the simple $\\tau$--approach, or approach (iia), as used by \\citet{vainshteinetal83}, \\citet{blackmanetal02b} and \\citet{brandenburgetal05b}, and of the spectral $\\tau$--approach, or approach (iib), used by \\cite{raedleretal03} and by \\citet{rogachevskiietal03,rogachevskiietal04} (section~\\ref{secii}).  Approach (i) in its original form, applying to purely hydrodynamic background turbulence, is based on solutions of the induction equation for the fluctuations $\\bbb$ simplified by some approximation. In the case of magnetohydrodynamic background turbulence in addition solutions of the momentum balance for $\\bu$, again in some approximation, are used. In the approaches (ii) these equations are not really solved. Instead, some assumptions on crucial quantities are introduced. There is hardly any doubt in the correctness of the results of approach (i) in the range of its applicability, which is well defined at least by sufficient conditions. It is not surprising that the approaches (ii) deliver the same vectorial structures of the contributions to $\\bscE$ as approach (i), for these are already determined by elementary symmetry principles. The results of the two types of approaches differ however in the dependence of the coefficients of the individual contributions to $\\bscE$ on the correlation properties of the velocity fluctuations $\\bu$. The discrepancies of the two approaches as well as strange aspects or shortcomings of the approaches (ii) are discussed in detail (sections~\\ref{subsubseciia2} and \\ref{seciibc3}). Unless any overlap of the ranges of applicability of the approaches (i) and (ii) can be excluded, what would be very surprising, we have to conclude that the approaches (ii) are in some conflict with the basic equations and not all of their results can be taken for granted. For instance, the magnetic contribution to the $\\alpha$--effect occurs in the approaches (ii) not only in the case of a magnetohydrodynamic background turbulence but for purely hydrodynamic background turbulence, too, where they should not exist. In addition, in the first case it is the same ``correlation time\" which occurs with the kinetic and the magnetic contributions, what seems to be in conflict with numerical simulations (section~\\ref{subsec43}). Further the dependence of $\\alpha$ and $\\beta$ on the magnetic diffusivity $\\eta$ and the kinematic viscosity $\\nu$ which occurs in approach (iib) is not correct. Moreover, it is doubtful whether the approaches (ii) can describe, e.g., the $\\bOmega \\x \\bJ$--effect correctly.  There is no hint that the mentioned shortcomings of the approaches (ii) result from the $\\tau$--approximation, that is, the reduction of third--order correlations of $\\bu$ or $\\bbb$ to second--order ones. An important reason for them seems to be an improper treatment of a term connected with the forcing term in the momentum balance.  A proposal is made for a new approach which avoids any conflict with approach (i) (section~\\ref{new}). In this context also a new formalism for the higher--order correlation approximation is presented (section~\\ref{new1})."
    },
    "0606/astro-ph0606117_arXiv.txt": {
        "abstract": "From an astrobiological point of view, special attention has been paid to the probability of habitable planets in extrasolar systems. The purpose of this study is to constrain a possible range of the mass of a terrestrial planet that can get water. We focus on the process of water production through oxidation of the atmospheric hydrogen---the nebular gas having been attracted gravitationally---by oxide available at the planetary surface. For the water production to work well on a planet, a sufficient amount of hydrogen and enough high temperature to melt the planetary surface are needed. We have simulated the structure of the atmosphere that connects with the protoplanetary nebula for wide ranges of heat flux, opacity, and density of the nebular gas. We have found both requirements are fulfilled for an Earth-mass planet for wide ranges of the parameters. We have also found the surface temperature of planets of $\\leq 0.3 \\ME$ ($\\ME$: Earth's mass) is lower than the melting temperature of silicate ($\\sim 1500$~ K). On the other hand, a planet of more than several~$\\ME$ becomes a gas giant planet through runaway accretion of the nebular gas. ",
        "introduction": "More than 170 extrasolar planets have been detected so far; most of them are Jupiter-like planets (see www.obspm.fr/planets). At present, several projects are in progress to discover terrestrial planets (e.g., TPF, Darwin, etc.). Special attention has thus been paid to habitable planets in extrasolar systems. A prerequisite for a planet being habitable is considered to be the existence of liquid water on it. For a planet to retain liquid water on its surface, it must acquire a sufficient amount of water and be located at a suitable distance from its parent star. This paper focuses on the former issue and constrains the probability of the existence of habitable planets in extrasolar systems.  The latter issue has been discussed using the term ``habitable zone''. The habitable zone (HZ) is defined as a range of orbital distance from a star within which a planet can retain liquid water on its surface. Around a solar-mass main-sequence star, the HZ is located around 1~AU and its width is as narrow as $\\sim$0.4~AU \\citep{KWR93}. The surface temperature of a planet beyond the HZ is below the freezing temperature of $\\rm H_2O$, so that the planet is unable to keep liquid water continuously. Because of high surface temperature of a planet closer to its parent star than the HZ, the concentration of $\\rm H_2O$ in the upper atmosphere is so high that the planet suffers substantial escape of $\\rm H_2O$ due to incident stellar UV radiation, and ends up losing its ocean completely.  Formation of terrestrial planets in that narrow zone, however, seems to be by no means unlikely within the context of the core accretion model for planet formation. \\citet{IL04a,IL04b,IL05}, for example, constructed an integrated model for planet formation including the accretion and dynamical evolution of planets. Their model reproduced the mass-period distribution of detected extrasolar planets and also predicted that of unknown planets. In the predicted mass-period distribution, the HZ is filled with hypothetical terrestrial planets of various masses \\citep[e.g., Figure~12 of][]{IL04a}. This means the existence of planets in the HZ is likely; the remaining issue is thus how likely a planet acquires a sufficient amount of water.  In this paper, we consider the nebular origin of water of terrestrial planets. A planet embedded in a protoplanetary nebula attracts gravitationally the nebular gas to have a hydrogen-rich atmosphere \\citep*{HNM79,MNH78}. The atmospheric hydrogen can be oxidized by some oxide contained in the planet to produce water on the planet, which was first proposed by \\citet{S90}.  The reason why we focus on the process of water production is that it could commonly happen on any extrasolar terrestrial planet. Planets are in general formed in hydrogen-rich protoplanetary nebulae. Oxides are also available on a terrestrial planet if the C/O ratio of the system is less than unity \\citep{L75,LB79}; most main-sequence stars are known to have C/O ratios less than unity \\citep[e.g.,][]{RTLA03,TH05}. The remaining requirements are a sufficient amount of hydrogen and enough high temperature to melt the surface of the planet; the molten planetary surface being called the magma ocean. The second condition is needed because if the planetary surface remains solid, the reaction between the atmospheric hydrogen and the surface oxide would result just in production of membrane covering the surface, not leading to production of a large amount of water.  In this paper, we investigate properties of the atmosphere of nebular origin. We first clarify the conditions for a planet to have a massive hydrogen-rich atmosphere and a magma ocean to constrain the range of the mass of a planet that has a sufficient amount of water. The structure of the nebular-origin atmosphere was investigated by \\citet{HNM79} and \\citet{NMSH85} for wide ranges of planetary accretion rate, grain opacity, and density of the nebular gas. However, they used quite simple forms of opacity and equation of state, both of which have been substantially improved. Furthermore they focused only on the Earth (i.e., Earth-mass planets) and had no discussion on the production of water. Although \\citet{S90} discussed the production of water in order to suggest the deep magma ocean on the early Earth, the ranges of the parameters considered were so restricted that we are unable to get any systematic understanding of the nebular origin of water on terrestrial planets.  We also constrain the mass of a planet that remains a terrestrial one. Habitable planets may be not gas giant but terrestrial ones. If a planet is isolated from planetesimals, it captures a substantial amount of the nebular gas to become a gas giant planet \\citep*{INE00}. We simulate the evolution and accumulation of the planetary atmosphere to obtain the timescale for the substantial gas accretion as a function of planet's mass. The timescale is compared to the lifetime of the nebular gas of $10^6$--$10^7$ years \\citep*{NGM00} to constrain a possible range of the mass of a habitable planet.  Section~2 describes our numerical method. Section~3 presents properties of the atmosphere of nebular origin for various planetary masses. Section~4 shows the timescale for the substantial accretion of the nebular gas. Finally we discuss the probability of water production on terrestrial planets and constrain the masses of the potentially-habitable planets in section~5. ",
        "conclusions": "Based on the numerical results obtained in sections~3 and 4, we discuss production of water from the hydrogen-rich atmosphere on a terrestrial planet, a possible range of the mass of a habitable planet, and the possibility of the nebular origin of water on the Earth.   As described in Introduction, production of water on a terrestrial planet requires a sufficient amount of hydrogen and surface temperature higher than the melting temperature of silicate ($\\sim$ 1500~K). The two conditions are found to be fulfilled on an Earth-mass planet.  In a late stage of terrestrial planet formation, accretion rate of planetesimals (i.e., luminosity) decreases with time because of exhaustion of the planet's feeding zone. The decrease in luminosity increases the amount of hydrogen on an Earth-mass planet (see Fig.~\\ref{f3}a), while surface temperature remains above 2000~K (see Fig.~\\ref{f1}a). For example, if accretion rate of planetesimals is $1 \\times 10^{-9} \\ME \\, {\\rm yr^{-1}}$, corresponding to $L \\sim 1 \\times 10^{23} \\rm erg \\, s^{-1}$ (see eq.~[\\ref{eq: accretion luminosity}]), the mass of atmospheric hydrogen is more than about $1 \\times 10^{25}$~g for $f < 1$, as shown Fig.~\\ref{f3}a: The amount of hydrogen a planet acquires is insensitive to the nebular density (see Fig.~\\ref{f3}b).  Water is produced through reaction between the atmospheric hydrogen and oxides contained in the solid planet. The amount of water depends on what kind of oxide is available. Ion oxides (e.g., w\\\"{u}stite [$\\rm Fe_{0.974}O$], magnetite [$\\rm Fe_3O_4$], etc.) and fayalite ($\\rm Fe_2SiO_4$) react with the atmospheric hydrogen to produce water comparable in mass to hydrogen; the ratios of the partial pressures, $P\\sub{H_2O}/P\\sub{H_2}$, are 0.88, 24.02, and 0.49 at 1500~K for the iron-w\\\"{u}stite ($\\rm 1.894Fe+O_2 \\leftrightarrow 2Fe_{0.947}O$), the w\\\"{u}stite-magnetite ($\\rm 6.696Fe_{0.974}O+O_2 \\leftrightarrow 2.174Fe_3O_4$), and the quartz-iron-fayalite oxygen ($\\rm 2Fe+SiO_2+O_2 \\leftrightarrow Fe_2SiO_4$) buffers, respectively \\citep*{RHF78}. Thus, if a planet acquires hydrogen of $\\sim 1 \\times 10^{25} \\rm g$ and such oxygen buffers are available, the planet obtains water comparable in mass to the current sea water on the Earth ($= 1.4 \\times 10^{24}$~g). However, for the silicon-periclase-forsterite buffer ($\\rm 2MgO+Si+O_2 \\leftrightarrow Mg_2SiO_4$), $P\\sub{H_2O}/P\\sub{H_2}$ is as small as $\\sim 3 \\times 10^{-7}$ \\citep{RHF78}. Fe-bearing minerals might be required to produce sufficient water, although how much water is needed for a planet being habitable is quite uncertain. Whether Fe-bearing minerals commonly exist in extrasolar systems is still a matter of controversy. Equilibrium condensation in a highly-reduced environment like a protoplanetary nebula yields not Fe-bearing minerals but Fe-metal \\citep{WH93}. However, dust grains in a protoplanetary nebula can be considered to have non-equilibrium composition including, at least, fayalite \\citep[][and references therein]{PHBSRF94}.  When the surrounding nebular gas disappears almost completely, the atmosphere and solid planet begin to get cold, and then an ocean forms through the condensation of steam in the atmosphere. However, almost complete dissipation of the nebular gas allows the extremely ultraviolet (EUV) and far-UV radiation from the parent star to penetrate the planetary atmosphere. Such irradiation causes extensive loss of hydrogen (and steam). The timescale for complete loss of a $10^{25}$-g (say) hydrogen-rich atmosphere due to EUV and far-UV can be estimated to be longer than $10^6$~years \\citep*{SNH80,SHN81}, if very strong EUV and far-UV \\citep[up to 100 times higher than the present;][]{GR02} from the young parent star is considered. On the other hand, the timescale for the ocean formation is on the order of $10^3$~years on a planet located in the HZ, according to calculations based on the radiative-convective equilibrium model of an $\\rm H_2O$-$\\rm CO_2$ atmosphere with mass of $\\sim 10^{24}$~g \\citep{A93}. The timescale for the ocean formation for a $\\rm H_2O$-$\\rm H_2$ atmosphere considered here is probably not so different from that for an $\\rm H_2O$-$\\rm CO_2$ atmosphere, because inclusion of $\\rm H_2$ hardly affects the atmospheric structure due to its weak blanketing effect. Therefore, an ocean can forms before the significant loss of the atmosphere due to EUV and far-UV.    We can constrain a possible range of the mass of a habitable planet. As shown in section~3, there is a lower limit to the planetary mass below which the water production proposed in this paper does not work. Figure~\\ref{f5} illustrates that surface temperature of a planet of $\\leq$ 0.3~$\\ME$ is lower than the melting temperature of silicate ($\\sim 1500$~K) for reasonable ranges of the parameters. On the other hand, an upper limit to the mass of a habitable planet can be constrained because a massive planet captures a huge amount of the nebular gas to be a gas giant planet, not a terrestrial planet. The timescale for the gas accretion depends strongly on the planetary mass, as shown Fig.~\\ref{f6}. This timescale should be compared to the lifetime of the nebula that is known to be about $1 \\times 10^7$~years \\citep{NGM00}. The comparison suggests that the upper limit to the planetary mass is $7 \\ME$ for $f = 1$ and $2 \\ME$ for $f = 0.01$.  The water production proposed in this paper may have worked on our Earth. Because of \\textit{N}-body simulations of planetary accretion, details of the terrestrial planet formation in the solar system have been clarified. After the runaway growth of protoplanets \\citep{WS93}, they grow in an oligarchic fashion until they eat almost all of the planetesimals in their feeding zones \\citep{KI98,KI00}. Then several Mars-mass protoplanets form in the terrestrial planet region. The subsequent growth of the protoplanets needs giant impacts between them. The planets formed in the way is likely to have high eccentricities \\citep{CWB96}. Damping of those high eccentricities needs the drag force of the nebular gas \\citep*{KI02,NLT05}. Because the Earth is isolated in the nebular gas, the accretion of the nebular gas inevitably takes place and water is produced on the Earth. Although the amount of the nebular gas required for the damping of the eccentricities are as small as $10^{-4}$ to $10^{-3}$ times that of the minimum-mass solar nebula \\citep{KI02,NLT05}, our numerical results show that this small amount of the nebular gas is sufficient for the Earth to get water comparable in mass with Earth's sea water.  We should be, however, careful when we consider the nebular origin of water on the Earth. This is because the ratio of deuterium to hydrogen (D/H) of the sea water on the present Earth is larger by about a factor of seven than D/H of the solar nebula \\citep[e.g.,][]{DR02}. Moreover, the current Earth's atmosphere includes a tiny amount of noble gas, while the solar nebula was rich in noble gas \\citep[e.g.,][]{P91}. The latter problem may be solved by extensive loss of the noble gases as well as hydrogen from the atmosphere due to EUV and far-UV radiation \\citep{SNH80,SHN81}. Although adequate mixing of water originated from nebular gas and water in comets can produce the present D/H on the Earth's ocean, because D/H in comets is larger (by about a factor of two) than D/H of the present Earth's ocean \\citep[e.g.,][]{DR02}, the former problem is difficult to solve. At present, we have no definite evidences that the sea water of the Earth was originated from the nebular gas.  Our intention in this paper is to claim that water production from the nebular gas on a planet is a possible way for a terrestrial planet in the HZ to acquire water. As shown in this paper, the water production works, if a planet of 0.3 to several $M\\sub{E}$ forms in the protoplanetary nebula. The probability of formation of terrestrial planets in the nebular gas is still open to debate, mainly because the dissipation mechanism of the nebular gas is quite uncertain. However, the existence of many gas giant planets in extrasolar planets has ensured the validity of the core accretion model. That is, accretion of planets generally occurs in a protoplanetary nebula. Also, there is no good reason for terrestrial planet formation to prefer vacuum environment. Therefore, it is rather likely that terrestrial planets also form in the surrounding nebular gas. The production of water from the nebular gas in the HZ thus seems to be a natural consequence of planet formation."
    },
    "0606/astro-ph0606321_arXiv.txt": {
        "abstract": "We present 3.6, 4.5, 5.8, 8.0, 24, and 70 $\\micron$ images of the Crab Nebula obtained with the Spitzer Space Telescope IRAC and MIPS cameras, Low- and High-resolution Spitzer IRS spectra of selected positions within the nebula, and a near-infrared ground-based image made in the light of [Fe II]1.644 $\\micron$. The 8.0 $\\micron$ image, made with a bandpass that includes [Ar II]7.0 $\\micron$, resembles the general morphology of visible H$\\alpha$ and near-IR [Fe II] line emission, while the 3.6 and 4.5 $\\micron$ images are dominated by continuum synchrotron emission. The 24 $\\micron$ and 70 $\\micron$ images show enhanced emission that may be due to line emission or the presence of a small amount of warm dust in the nebula on the order of less than 1\\% of a solar mass. The ratio of the 3.6 and 4.5 $\\micron$ images reveals a spatial variation in the synchrotron power law index ranging from approximately 0.3 to 0.8 across the nebula. Combining this information with optical and X-ray synchrotron images, we derive a broadband spectrum that reflects the superposition of the flatter spectrum jet and torus with the steeper diffuse nebula, and suggestions of the expected pileup of relativistic electrons just before the exponential cutoff in the X-ray. The pulsar, and the associated equatorial toroid and polar jet structures seen in Chandra and HST images \\citep{Hes02} can be identified in all of the IRAC images. We present the IR photometry of the pulsar. The forbidden lines identified in the high resolution IR spectra are all double due to Doppler shifts from the front and back of the expanding nebula and give an expansion velocity of $\\approx 1264$ km s$^{-1}$. ",
        "introduction": "}  In their seminal paper ``Synthesis of the Elements in Stars,'' E. M. Burbidge, G. R. Burbidge, W. A. Fowler, and F. Hoyle (1957) \\nocite{Bur57} described how primordial hydrogen is converted into the other elements by nucleosynthesis in stellar interiors and stellar explosions. Massive stars play a particularly important role in the production of nuclei up to the iron peak during main sequence and post main sequence nucleosynthesis. The Type II Supernova (SN) explosions that result from the collapse of the their iron cores produce the rest of the heavy elements by rapid neutron capture in the expanding ejecta. Thus, SN ejecta enrich the chemical content of the interstellar clouds from which new stars continually form, and it is believed that the contents of the pre-solar nebula were profoundly affected by one or more such events \\citep[see][]{Cla82}.  In particular, the ability of SN explosions to eject copious quantities of carbon and other condensable metals has long led to speculation that SNs are capable of producing vast quantities of stardust as their ejecta cool \\citep{Cla82, Geh87, Dwe88, Geh88}. Clayton has proposed that such grains may survive today in meteorite inclusions. The infrared (IR) is an ideal spectral region in which to test these theories, because dust grains emit strongly in the IR as do forbidden emission lines from condensable metals that remain in the gas phase. On the other hand, previous IR studies that have failed to reveal evidence for any large amount of dust in SN remnants (SNRs) have provided us with somewhat of a mystery \\citep{Geh90, Are99, Dou01, Dwe04, Gre04}.  We have used Guaranteed Observation Time (GTO) on NASA's new Spitzer Space Telescope \\citep{Wer04} to obtain IR images and spectroscopy of the Crab Nebula, formed by a supernova explosion in 1054 AD. It is one of the youngest known SNRs and is one of the most studied objects in the Galaxy. The SN ejecta in the Crab are concentrated in filaments of ionized gas that produce emission line spectra \\citep[see][]{Dav85}. One long-standing puzzle regarding the Crab is the lack of a visible blast wave expected from a massive star explosion \\citep{Sew06}. Previous IR observations have indicated that there is a paucity of dust in the Crab. Infrared Astronomical Satellite (IRAS) observations revealed an IR excess at wavelengths longer than 12 $\\micron$ that was attributed to thermal emission from 0.005-0.03 M$_{\\odot}$ of dust \\citep{Mar84}. Further evidence for dust in the Crab Nebula was found in the form of optical extinction in the filaments \\citep{Hes90, Fes90, Bla97, San98}. In more recent studies using the Infrared Space Observatory (ISO), \\citet{Dou01} report no evidence of spectral features from dust emission in the ISOCAM spectra, while \\citet{Gre04} calculate an upper limit of 0.02 M$_{\\odot}$ of warm dust from the far-IR excess seen by ISOPHOT.  The Crab Nebula is one of the brightest synchrotron sources in the Galaxy. It is a prototype of a class of SNRs called the filled-center SNRs or ``plerions,'' that are powered by a central pulsar. A large fraction of the Crab pulsar's spin down luminosity is converted into the nebular synchrotron luminosity, from radio through gamma-rays. The details of this conversion are quite uncertain. The \\citet{Ken84} steady-state spherical MHD model still provides the best description of the optical and X-ray profiles, but cannot tie them together with the radio synchrotron emission.  The unprecedented sensitivity of the IR imagers and spectrometers of the Spitzer Space Telescope present us with an unparalleled opportunity to search for dust and forbidden line emission in the Crab Nebula and study in detail the spatial variations of synchrotron emission across the remnant. In this paper, we present 3.6, 4.5, 5.8, 8.0, 24, and 70 $\\micron$ images of the Crab Nebula obtained with the Spitzer Space Telescope IRAC and MIPS cameras, IRS spectra of selected positions within the nebula, and a near-IR ground-based image made in the light of [Fe II]1.644 $\\micron$. ",
        "conclusions": "}  We briefly summarize the main conclusions presented by our Spitzer Space Telescope IR imaging and spectroscopic observations of the Crab Nebula:  1.  A comparison of the morphology from the x-ray to the radio shows that the synchrotron component and filaments dominate at different wavelengths.  2.  We have derived a broadband shape to the synchrotron spectrum which shows evidence for multiple components along the line of sight, likely due to the torus and jet which remain flat to large distances from the pulsar, superposed on the broad nebular emission that steepens with increasing distance. At shorter wavelengths, the derived spectral shape is consistent with the expected pileup of relativistic electrons just below an exponential cutoff in X-rays.  3. We have measured the flux density of the nebula and the pulsar. The smooth background of the nebula and the pulsar are dominated by synchrotron emission and a large fraction of the emission from the filaments in the images is due to forbidden line emission from Ar, Ne, O, and Fe.  4. We find a paucity of dust. The small grain component seems to be missing entirely and we see no evidence for silicate emission.  The total emission at long wavelengths from large grains implies a total dust mass in the nebula of less than 1\\% of a solar mass.  5.  In the IRS spectra, we see Doppler shifted emission from both the front and back sides of the expanding shell, and we measure a radial expansion velocity of roughly 1264 km s$^{-1}$."
    },
    "0606/astro-ph0606447_arXiv.txt": {
        "abstract": "s{ Gravitational lensing has now become a popular tool to measure the mass distribution of structures in the Universe on various scales. Here we focus on the study of galaxy's scale dark matter halos with galaxy-galaxy lensing techniques: observing the shapes of distant background galaxies which have been lensed by foreground galaxies allows us to map the mass distribution of the foreground galaxies. The lensing effect is small compared to the intrinsic ellipticity distribution of galaxies, thus a statistical approach is needed to derive some constraints on an average lens population. An advantage of this method is that it provides a probe of the gravitational potential of the halos of galaxies out to very large radii, where few classical methods are viable, since dynamical and hydrodynamical tracers of the potential cannot be found at this radii. We will begin by reviewing the detections of galaxy-galaxy lensing obtained so far. Next we will present a maximum likelihood analysis of simulated data we performed to evaluate the accuracy and robustness of constraints that can be obtained on galaxy halo properties. Then we will apply this method to study the properties of galaxies which stand in massive cluster lenses at $z\\sim0.2$. The main result of this work is to find dark matter halos of cluster galaxies to be significantly more compact compared to dark matter halos around field galaxies of equivalent luminosity, in agreement with early galaxy-galaxy lensing studies and with theoretical expectations, in particular with the tidal stripping scenario. We thus provide a strong confirmation of tidal truncation from a homogeneous sample of galaxy clusters. Moreover, it is the first time that cluster galaxies are probed successfully using galaxy-galaxy lensing techniques from ground based data. }  This contribution is a summary of two galaxy-galaxy lensing papers:\\\\ Limousin et~al., 2005 \\cite{paperI}, hereafter Paper I exposes a theoretical analysis of galaxy-galaxy lensing, and Limousin et~al., 2006 \\cite{paperII}, hereafter Paper II exposes the application of the method tested extensively in Paper I on a sample of homogeneous massive galaxy cluster at $z\\sim0.2$. ",
        "introduction": " ",
        "conclusions": "We have given a revue of galaxy-galaxy lensing detections obtained so far, both on field galaxies as well as on cluster galaxies. We have presented a maximum-likelihood analysis of galaxy-galaxy lensing on simulated data to study the accuracy with which input parameters for mass distributions for galaxies can be extracted.. We have applied this method on a sample of massive galaxy clusters and we have derived some constraints on the dark matter halos of the elliptical cluster galaxies. It is the first time that cluster galaxies are successfully probed from ground based observations. The main result of this work is to find galaxy halos in clusters to be significantly less massive and more compact compared to galaxy halos around field galaxies of equivalent luminosity. This is in good agreement with previous galaxy-galaxy lensing studies. Moreover, this confirmation is based on the analysis of 5 massive clusters lenses whose properties are close one to each other, hence the confirmation we provide is a strong one since it relies on a homogeneous sample.\\\\ This observational result is in good agreement with numerical simulations, in particular with the tidal stripping scenario. The theoretical expectation is that the global tidal field of a massive, dense cluster potential well should be strong enough to truncate the dark matter halos of galaxies that traverse the cluster core (Avila-Reese et~al., 2005 \\cite{avila05}; Ghigna et~al., 2000 \\cite{ghigna00}; Bullock et~al., 2001 \\cite{bullock})."
    },
    "0606/astro-ph0606501_arXiv.txt": {
        "abstract": "The recurrence of a previously documented eclipse of a solar-like pre--main-sequence star in the young cluster IC 348 has been observed. The recurrence interval is 4.7 $\\pm\\ 0.1$ yr and portions of 4 cycles have now been seen. The duration of each eclipse is at least 3.5 years, or  $\\sim 75$\\% of a cycle, verifying that this is not an eclipse by a stellar companion. The light curve is generally symmetric and approximately flat-bottomed. Brightness at maximum and minimum have been rather stable over the years but the light curve is not perfectly repetitive or smooth and small variations exist at all phases. We confirm that the star is redder when fainter. Models are discussed and it is proposed that this could be a system similar to KH 15D in NGC 2264. Specifically, it may be an eccentric binary in which a portion of the orbit of one member is currently occulted during some binary phases by a circumbinary disk. The star deserves sustained observational attention for what it may reveal about the circumstellar environment of low-mass stars of planet-forming age. ",
        "introduction": "Eclipses by non-stellar objects are of great interest for their potential to tell us not only about the stars involved but also about the extended objects -- presumably circumstellar or circumbinary disks. Such events signal their presence by their extended duration. A well known example is $\\epsilon$ Aur, whose eclipse lasts  $\\sim 2$ years  \\citep{b85, l96}. Another is KH 15D \\citep{kh}, an eccentric binary system embedded in an inclined circumbinary disk \\citep{w06}. The longest known eclipse is for the star HMW 15\\footnote{The position listed for this object by \\citet{hmw00} is incorrect; the correct position is given by \\citet{chw04}.}. This object, also known as TJ 108 \\citep{tj}, H 187 \\citep{h98} and LRLL 35 \\citep{lrll}, is a member of  the young cluster IC 348 at a distance of about 300 pc \\citep{c93, lrll, h07} and undergoes an eclipse that lasts $\\sim 3.5$ years \\citep{chw03} (hereafter Paper I). Like KH 15D it was discovered at Wesleyan University's Van Vleck Observatory (VVO) as part of a CCD photometric monitoring program of young clusters that has been going on for 15 years and covers about a thousand stars in several different clusters.  It was not clear from Paper I, which covered only five years of monitoring, whether this lengthy eclipse event would recur and, if so, on what time scale. Now that three additional years have passed, we have obtained enough data to answer these questions rather definitively. We have also obtained new color data that help to constrain models of the system. It is now clear to us that HMW15 is an object worthy of intensive study over the coming years and we hope that the new evidence for periodic recurrence of the eclipse will stimulate additional work on it by other investigators. ",
        "conclusions": "HMW 15 is identified as a normal, PMS member of the young cluster IC 348 by \\citet{lrll} based on its location on the sky and in a color-magnitude diagram. Its spectral type is listed as between G8-K4 in the optical and K3-K6 in the infrared by those authors, an unusually wide range of results although not uniquely so. We may infer that it is a relatively low mass star (0.5 to 1 M$_\\odot$) at an age of 2-4 My and a distance of about 300 pc. It has no detectable infrared excess according to a recent {\\it Spitzer} study by \\citet{lm06} and is not known to be either a spectroscopic or visual binary. Its single distinguishing characteristic is its unusual photometric variability reported in Paper I and in this paper.  Our data demonstrate the importance of three time scales of variation for this star. One is the periodic time of 4.7 years, another is the eclipse duration time of about 3.5 years and a third is the ``wiggle\" time of a few weeks during which coherent departures from smooth variation (sometimes involving reversals in the general trend of brightening or fading) occur. We discuss the plausible physical origins of each of these in turn.  The only plausible explanation for a period of several years, in our view, is that this is an orbital period. The dynamical time scale that would apply to either pulsation or rotation for a star like this is much shorter -- of order hours or days. If HMW 15 is a single star, then we surmise that it must be orbited by an inhomogeneous ring of matter with a period of 4.7 years, corresponding roughly to a semi-major axis of around 3 AU. Presumably this ring would be part of a more extensive structure, perhaps a circumstellar disk. In this model the 3.5 year eclipse time would indicate that for about one quarter of each cycle the material in this ring was either absent or of vanishingly small optical depth.  Alternatively, and perhaps more plausibly, the period may be the orbital period of a binary star system. We imagine a system analogous to KH 15D during the 1960's-80's in which one member of a binary is obscured by a circumbinary disk during a portion of each cycle \\citep{j04, j05}. As in the case of KH 15D this requires an eccentric orbit that is somewhat inclined to the plane of the circumbinary disk. The attractions of this model are the relatively flat bottom to the eclipse (expected if one star is totally obscured during that time) and the amplitude of about 0.75 mag, which is also expected for a total eclipse of stars with nearly the same luminosity. It should be easy to test this model over the next few years since radial velocity variations of the star or stars should easily be detectable. In this model, the 3.5 year extent of the eclipse is set by the fraction of the orbit covered by the circumbinary disk.  Regardless of whether the star is single or a binary, it remains something of a puzzle to explain the shorter time scale variations that superimpose themselves on the overall light curve as small deviations or ``wiggles\". One possibility is that they represent actual variations in the brightness of the star or stars due, perhaps, to starspots, which are ubiquitous on pre-main sequence stars. However, the normal time scale for such variability is much shorter than observed -- typically days, due to the rotation period of the star. Another possibility is that these represent irregularities in the optical depth of the screen -- lumps and clumps in the distribution of matter in the circumstellar or circumbinary disks in the models discussed above. We need to have a longer period of monitoring to see how frequently these events occur and to be able to describe their properties statistically before more can be said about them. Color information during the events would also be interesting to obtain.  The color data do not allow us to distinguish between the single star or binary model for this object. While an interstellar reddening law does not fit the color data very well (solid line on Fig. \\ref{color}) it is also not a terrible fit and it is always possible, perhaps even expected, that grain growth would have occurred in a circucmstellar disk of this age, so that a flatter than normal extinction law might be quite usual. On the other hand, a binary model can fit the data equally well, if not better, with very little effort. As an example, we show (dashed line on Fig. \\ref{color}) the predicted color behavior for a G8 + K6 system in which the G8 star is entirely visible at maximum and entirely obscured at minimum light. In computing this model we have assumed the extinction is gray. If, instead, we allowed the extinction to follow an interstellar reddening line, then the system light would follow the solid line during its brighter part and the dashed line during its fainter part. Obviously this would be an even better fit to the data.  The simplest binary model, with gray extinction and total obscuration of the bluer star, is unlikely to be correct for another reason. The required color difference between the components is quite large for their magnitude difference. If, as one would expect, the components of this putative binary were coeval, then they would be expected to lie along an isochrone in the V, V-I plane. The simple binary model described here does not come close to either a theoretical or observed isochrone. The redder component is too bright by about 1.5 mag in V. This probably means that we need to consider non-gray extinction models. If the binary model is correct and we can obtain the spectral types of both stars, then it should be easily possible to determine the extinction law and the intrinsic properties of the stars. In this case, the system should prove to be a useful test of PMS models since we will have coeval stars of different masses. At present, more refined models are not warranted given the meager available data.  It is interesting that HMW 15 escaped detection as a variable star for so long, in spite of the fact that its amplitude is more than one magnitude in V. Part of the reason is that it varies so slowly. Most variability studies, especially those aimed at young stellar objects, are tuned for detections on time scales of hours to months. While in some seasons this star is detectable as variable over a few weeks, it is not highly variable on that time scale and does not easily emerge from the data until multiple seasons are put together. Perhaps there are more such long time scale, periodic variables in young clusters awaiting detection. If so, we should continue to find them at VVO as the baselines for our variability studies continue to grow.  To conclude, our data do not rule out a single star model for this system, but the binary model is marginally more attractive since it does a better job of explaining the color behavior and perhaps the wide range in reported spectral types \\citep{lrll}. The issue should soon be decided by spectral studies. If the star is a binary with an orbital period of 4.7 y and G or K components, then the radial velocity amplitudes of the stars will be easily detectable (many km/s). Regardless of this, it is now clear that at least one star in the system is periodically eclipsed by extended dust in its vicinity, presumably either a circumstellar or circumbinary disk. As in the case of KH 15D, we have a rare opportunity to learn about possibly planet-forming material by using the starlight as a probe. The relative motions of the star and extended matter provide a time dimension that can be exploited by a determined spectroscopic monitoring program supplementing continued photometry. We recommend this as a promising approach toward improving our knowledge of young stars and disks."
    },
    "0606/astro-ph0606737_arXiv.txt": {
        "abstract": "We study the scattering of low-energy Cosmic Rays (CRs) in a turbulent, compressive MHD fluid. We show that compressible MHD modes -- fast or slow waves with wave lengths smaller than CR mean free paths induce cyclotron instability in CRs.  The instability feeds the new small-scale Alfv\\'enic wave component with wave vectors mostly along magnetic field, which is not a part of the MHD turbulence cascade. This new component gives feedback on the instability through decreasing the CR mean free path. We show that the ambient turbulence fully suppresses the instability at large scales, while wave steepening constrains the amplitude of the waves at small scales. We provide the energy spectrum of the plane-parallel Alfv\\'enic component and calculate mean free paths of CRs as a function of their energy.  We find that for the typical parameters of turbulence in the interstellar medium and in the intercluster medium the new Alfv\\'enic component provides the scattering of the low energy CRs that exceeds the direct resonance scattering by MHD modes. This solves the problem of insufficient scattering of low-energy CRs in the turbulent interstellar or intracluster medium that was reported in the literature. ",
        "introduction": "Cosmic rays (CRs) and magnetic fields are essential components for many astrophysical ecosystems, including galaxies and clusters of galaxies (see Schlickeiser 2002). In many instances, e.g. Milky Way, the pressure of CRs and magnetic fields is larger than the gas pressure. As a rule, astrophysical magnetic fields are frozen in turbulent plasma and move together with it. As a result, CRs interacting with turbulent magnetic fields get scattered and accelerated (see Melrose 1968, Schlickeiser 2002).  The magnetohydrodynamic (MHD) approximation is widely used to describe the actual magnetized plasma turbulence over scales that are much larger than both the mean free path of the particles and their Larmor radius (see Kulsrud 2004). The theory of MHD turbulence has become testable recently due to numerical simulations (see Biskamp 2003) and this provided reliable foundations for describing turbulence-CRs interactions. The simulations (see Cho \\& Lazarian 2005 and ref. therein) confirmed the prediction of magnetized Alfv\\'enic eddies being elongated along magnetic field (see Shebalin, Matthaeus \\& Montgomery 1983, Higdon 1984) and provided results consistent with the quantitative relations for the degree of eddy elongation obtained  in Goldreich \\& Sridhar (1995, henceforth GS95).  Scattering of CRs is an essential part of both CR propagation modes and models of CR acceleration. Efficient scattering is usually postulated (see Schlickeiser 2002), which ensures high degrees of coupling of CRs and magnetized plasma. In addition, efficient scattering provides appreciable second order Fermi acceleration and enables the return of CRs into the shock to ensure the first order Fermi acceleration.  This corner stone of CR physics has been challenged recently when it became clear that Alfv\\'enic eddies are stretched along magnetic field direction.  As the interaction between CRs and such elongated eddies is weak (see discussion in Lerche \\& Schlickeiser 2001), this resulted in the prediction of long mean free paths for Milky Way CRs (Chandran 2000, Yan \\& Lazarian 2002, henceforth YL02). YL02 and Yan \\& Lazarian (2004) attempted to remedy the situation by appealing to CR scattering by isotropic sound-like fast modes. However, plasma-dependent damping of fast modes made the scattering very different in different parts of the interstellar medium. Is such a radical change of the CR scattering picture absolutely necessary?  We note, that the problem of CR scattering goes well beyond the Milky Way physics.  Brunetti (2006) discussed the implications of suppressed CR scattering on the acceleration of CRs in the clusters of galaxies.  Is there any process through which scattering by fast modes can provide high efficiency of CR scattering?  Below we consider such a process that is related to CR feedback on MHD turbulence. We show that compression of CRs induces instability that results in the generation of modes that are parallel to the magnetic field. Such modes that are also frequently referred to as slab modes have been long employed in the models of CR propagation (see, e.g., Jokipii 1966), their origin, however, was somewhat mysterious. This paper provides a physically motivated mechanism for the generation of slab modes and quantify the efficiency of their generation.  In what follows we discuss the properties of compressible MHD turbulence in \\S~2. We describe the kinetic instability that develops in CRs when the magnetic field is compressed on scales less than the CR mean free path in \\S~3. We consider the non-linear saturation of the instability in \\S~4 and its large-scale cut-off that follows from the interaction of the instability waves with the ambient turbulence in \\S~5. The feedback of CRs to compressible turbulence is considered in \\S~6. The implications of our work for various ISM phases and intra-cluster medium (ICM) are considered in \\S~7. The short summary of this work is presented in \\S~8. ",
        "conclusions": "\\subsection{CRs in ISM and galaxy clusters }  Lets consider CRs in galaxy clusters. The magnetic field magnitude and the density of CRs are somewhat uncertain there (see Esslin et al. 2005), so we adopt values similar to our galaxy, namely $B=5\\mu$G, $n_{CR}(E>1~{\\rm GeV})=4\\cdot 10^{-10}~{\\rm cm}^{-3}$ and $\\alpha=2.6$. This corresponds roughly to equipartition between CR and magnetic field energies.  In clusters these energy densities are around 5 per cent of the thermal energy density. We will have then $L_i\\approx 6\\cdot 10^{-7}$~pc. The reference Larmor radius of 1~GeV proton is $r_0\\approx 2\\cdot 10^{-7}$~pc. We take the scale $L=1$~kpc, which, being Alfv\\'enic at this scale, corresponds roughly to driving with the virial velocity at the scale of 30 kpc.  For these numerical values and $\\mu=1/3$ we will have, from Eq.~(\\ref{amplitude}), $\\delta B/B=0.04(r_p/r_0)^{-0.1}$, almost independent on scale, $r_{p,{\\rm crit}}\\approx 10^{3} r_0 \\approx 2\\cdot 10^{-4}$~pc, $\\lambda\\approx 10^{-4}(r_p/r_0)^{1.2}$~pc, and the mean free path corresponding to the turbulent damping ($r_p=r_{p,{\\rm crit}}$) is 0.4~pc which is much smaller than the outer scale. We estimate collisionless cutoff as $l_{\\rm cut}\\sim 10^{-2}$~pc. The feedback mean free path will be around 0.3 pc, so the spectrum of Alfv\\'enic slab motions will be mostly steeper, $k^{-1.2}$ and the mean free paths will be modified according to (10) and (5). The efficient CR scattering entails efficient second-order Fermi acceleration, see eq. (13), the process that may be important for clusters of galaxies (Cassano, Brunetti, 2005).  In our galaxy one can assume same values for $B$, $\\alpha$ and $n$ and value of $L$ around 50~pc.  We assume an acoustic turbulence spectrum for fast waves, taking $\\mu=1/4$. We generally get a smaller range of Alfv\\'enic slab motions, from scales of $r_0$ to about $600r_0$ with $\\delta B/B=0.093(r_p/r_0)^{-0.14}$. The resulting mean free paths $\\lambda$ vary from $2.3\\cdot10^{-5}$~pc to $8\\cdot10^{-2}$~pc.  In the Galactic Corona, fast waves will be damped by collisionless damping (see, e.g., Ginzburg 1961) with a cutoff of around $1.6\\cdot10^{-3}$~pc, which is within the range of $\\lambda$ that we deal with. In the warm ionized medium (WIM) the collisional damping cutoff will be around $10^{-4}$~pc. The feedback mean free path will be around $7\\cdot10^{-2}$. Again the spectrum of slab motions becomes steeper and the mean free paths of CRs are modified accordingly.  In \\S~2 we assumed that the compression factor $A$ is larger than $v_A/c$. This assumption is satisfied in galaxy clusters, as, from the previously adopted values and $n\\approx 10^{-3}$~cm$^{-3}$, $v_A/c \\approx 10^{-3}$, while compression factors for scales between 2~pc and $2\\cdot 10^{-4}$~pc are between $0.12$ and $5\\cdot 10^{-3}$. For the Milky Way ISM this condition is satisfied much better, as for $n\\approx 1$, $v_A/c\\approx 4\\cdot 10^{-5}$, and compression factors are generally larger, due to the fact that the minimum $\\lambda/L$ is smaller.  The slab Alfv\\'en modes had been a part of the CR paradigm from the very start of the research in the field (see Jokipii 1966). Together with anisotropic components they are part of some of the modern models of CR propagation (see Zank \\& Matthaeus 1992, Bieber et al. 1994, Shalchi et al. 2006). In our model the slab plane-parallel modes emerge naturally as the result of the interaction of compressible turbulence with CR. Although this mechanism is different from the earlier considered processes, it may justify some of the earlier calculations invoking slab modes. Unlike earlier theories we predict the dependence on the amount of the slab mode energy on the relative pressure of the CRs.  As we see, in both clusters of galaxies and ionized gas in Milky Way the instability within CR fluid limits the CRs mean free path. Like in scenario discussed in YL04, where the compressible fast modes were identified as the major CRs scattering agent, compressive modes are essential for scattering. However in this treatment, unlike YL04, we show that compressions at scales much larger than the resonance scale are important. This difference is crucial for scattering of low-energy CRs, as the fast mode have collisional or collisionless cut-offs which, depending on the media, may be larger than the low-energy CR gyroradius. In this case YL04 appealed to Transient Time Damping (TTD) processes, which are less efficient for scattering than gyroresonance\\footnote{In fact the gyroresonance instability can be the major source of isotropization during the TTD acceleration.}. Our present work shows that the slab Alfv\\'en mode discussed in the present paper can be responsible for efficient scattering. Another important difference from YL04 is that our new mechanism require relatively large total pressure of CRs (see (\\ref{lili})). In the case when the pressure of CRs is negligible the fast modes could stay the major scattering agent (see Petrosian et al, 2006).  Our model predicts rather small mean free paths, but this does not contradict the estimates on the average lifetime of the CR in the Galaxy. These lifetimes are estimated to be around Galaxy thickness divided by the Alfv\\'en velocity, which is a powerful support for the models with streaming instability. Our model will predict similar lifetimes, because it includes turbulence which advects CRs on outer scale comparable with Galaxy thickness with velocity of around Alfv\\'en velocity. In fact, the turbulence itself could be generated on these scales by Parker instability.  \\subsection{Partially Ionized Gas}  Previous discussion is also applicable to partially ionized gas, if the degree of ionization is larger than $\\sim 90\\%$. Indeed, for such high ionization degrees the Alfv\\'enic turbulence cascades to scales less than the ion-neutral decoupling scale (see Lithwick \\& Goldreich 2001).  If, on the other hand, the degree of ionization is lower, we assume that Alfv\\'enic turbulence is fully damped by ion-neutral collisions at the scale $l_{\\rm damp}$, and it would not be able to provide turbulent damping for $k_\\perp l_{\\rm damp}<1$. As we saw in \\S~5 the damping for slab waves with $k_\\|$ is provided by turbulent eddies with $k_\\perp \\sim k_\\|^{3/4}$, therefore, our slab-type component arising from CR instability will protrude up to scales as large as $l_{\\rm damp}^{4/3}/L^{1/3}$. This scale could be substantially larger than the $r_{p,{\\rm crit}}$ derived in \\S~5.  According to Lazarian, Vishniac \\& Cho (2004) the regime of viscosity damped turbulence emerges for Alfv\\'enic turbulence at scales less than $l_{\\rm damp}$.  This regime is characterized by a shallow $k^{-1}$ spectrum of magnetic perturbations and it persists down to the ion-neutral decoupling scale where it reverts to intermittent Alfv\\'enic turbulence that involves only ions. The detailed treatment of the interactions of CRs with turbulence in partially ionized medium is beyond the scope of this paper, however.   \\subsection{Thermal plasma mean free paths in galaxy clusters}  In the paper above we considered the CR component of the ISM or ICM, which are the high energy particles that interact with the rest of the medium via the magnetic fields. These particles have a power-law distribution that arises from the acceleration terms that are proportional to the CR momenta. The astrophysical plasma, on the other hand, is assumed to have a Maxwellian distribution and provides us with both conductivity and mass density which are required for a MHD treatment.  In a fully ionized plasma particle-to-particle collisions are Coulomb scattering and the rate of the collisions becomes smaller with temperature. With high temperature and small density these mean free paths can be huge. For example, in galaxy clusters it could be as large as 4kpc. This lead to apparent contradiction, as particles with such a huge mean free path will be subjected to acceleration and will not be Maxwellian.  Schekochihin and Cowley (2005) proposed that thermal particles will be scattered by instabilities. They considered hydrodynamic as well as kinetic instabilities and considered the evolution of the cluster from initial state with no magnetic field. Their argument is that the Reynolds number, being initially very low, will increase with increasing magnetic field and the dynamo will self-accelerate. They predict folded magnetic fields due to high-Prandtl number dynamo and their mean free paths are between viscous scale and the reversal scale.  In this subsection we estimate mean free paths of thermal particles in a way similar to the rest of our paper, keeping in mind that there are quite a few other plasma effects and some MHD dynamo effects that might be important, so that these estimates are still rather speculative. This may be excused by the fact that thermal mean free path, viscosity and thermal conductivity are very important for cluster dynamics.  Anisotropic distributions of thermal particles will excite waves with inverse wavevectors of the order of thermal Larmor radius, $r_T\\approx 10^{-9}$~pc as the instability is exponentially slow for smaller wavevectors (see Mikhailovskii 1975, eq. 10.7). All particles will have approximately the same mean free path, and the value of $\\delta B/B$ that provides scattering will now refer to the {\\it total} perturbed magnetic field, in contrast with its definition in \\S~4. Apparently the energy-transfer arguments of \\S 6 will be most important, as the steepening is very fast on thermal Larmor scales. By equating steepening and turbulent energy transfer rates we have $r_T/L\\approx(\\delta B/B)^4$, which gives $\\delta B/B\\approx 10^{-3}$, $\\lambda\\approx 10^{-3}$~pc."
    },
    "0606/astro-ph0606723_arXiv.txt": {
        "abstract": "In this paper we present an analysis of temperature taken at two telescopes located at the Observatorio del Roque de Los Muchachos in the Canary Islands. More than 20 years of measurements at CAMC are included. The analysis of the data from TNG and CAMC are compared in order to check local variations and long term trends. Furthermore, the temperatures at different heights are correlated to the quality of astronomical seeing. We considered the correlation of NAO Index and annual down$-$time with mean annual temperatures. The final aim of this work is to better understand the influence of wide scale parameters on local meteorological data. The analysis is done using a statistical approach. From each long series of data we compute the hourly averages and than the monthly averages in order to reduce the short$-$time fluctuations due to the day/night cycle. A particular care is used to minimize any effect due to biases in case of lacking of data. Finally, we compute the annual average from the monthly ones. The two telescopes show similar trends. There is an increase of temperatures of about 1.0$^{\\circ}$C/10yrs from the annual means and a more rapid increase of the annual minimums then the maximums. We found that positive NAO Index reduces the increase of temperatures, and accelerates the decrease. Moreover, there is no evidence that positive NAO Index corresponds to a lower number of non-observable nights. Finally, seeing deteriorates when the gradient of temperatures between 2 and 10 m above the ground is greater than $-0.6^{\\circ}$C. ",
        "introduction": "Since the year 1970, La Palma Island, located at about 400 km off the Moroccan coast of North$-$West Africa, appeared a favorable geophysical site from the point of view of the sky conditions, due to the proximity of the semi-permanent Azores high pressure system, and it was chosen to host the main astronomical telescopes. It is known that the very good astronomical conditions of the island are mainly due to the stable subsiding maritime airmass that place most of the time the telescopes near the top of the mountain well above the inversion layer occurring in the range between 800 m and 1200 m (McInnes \\& Walker \\cite{mcinnes}).\\\\ All the telescopes are located along the northern edge of the Caldera de Taburiente, at the N$-$W side of La Palma Island, where the irregular shapes produce a complex orography and the crowdedness of the top, due to the presence of all the astronomical observatories, suggests a possible modification of the local microclimate making difficult to foresee in advance the precise local meteorological parameters. Therefore in these last years the ORM has been extensively monitored thanks to the efforts of the several site testing groups belonging to the hosted astronomical observatories.\\\\ Since several years, various astronomical sites are monitored on a continuous basis by automatic weather stations, which provide measurements of a few local meteorological parameters. All these instantaneous and long term records of the meteorological data are important tools for meteorological and climatological studies, as well as for the calibration of satellite remote sensing of the atmospheric and ground conditions (Zitelli et al. \\cite{zitelli}).\\\\ In this paper we present for the first time an analysis of measurements obtained from local meteorological towers and environmental conditions made at three telescopes at ORM. The meteo data of Telescopio Nazionale Galileo (TNG) and Carlsberg Automatic Meridian Circle Telescope (CAMC) are compared in order to check local variations or meteorological conditions. The analysis of these differences in terms of image quality at the telescopes will be discussed in this paper as follows:\\\\ In \\S 2 we discuss the annual temperature means for the two telescopes, in \\S 3 we present differences between day$-$time and night$-$time mean values and their comparison with the down$-$time at CAMC, in \\S 4 it is presented the North Atlantic Oscillation Index and its correlation with temperatures, in \\S 5 we present seeing and temperature analysis for TNG. A preliminary comparison among the three sites \\S 6 it is also presented on the basis of the results of the previous sections. ",
        "conclusions": "We presented for the first time an analysis of long$-$term temperature data directly obtained from local meteorological towers of TNG and CAMC, at a height of about 2300 m above sea level, far from urban concentration and well above the inversion layer.\\\\ Annual mean values show a similar trend between TNG and CAMC. The linear fit of the 20 years long baseline of CAMC data gives an increase of annual mean temperatures of about 1.0$^{\\circ}$C/10yrs.\\\\ It is interesting to note an oscillation of the values with a period of about 3$-$4 years that seems to be slightly smoothed during the last 10 years. Another evidence is the different behaviour of local minimums and local maximums, in fact the minimums increase more rapidly then the maximums.\\\\ A comparison between NAO Index and the annual mean temperatures shows a correlation of 86\\% of significance. In fact, the action of positive NAO Index is like a brake for the increase of temperatures, and like an accelerator for the decrease. Vice versa, negative NAO Index acts in opposite mode. Moreover, no correlation is found between NAO Index and number of non-observing nights.\\\\ We also investigated the influence of the temperature on the astronomical seeing and we have found that the seeing deteriorates when the gradient of temperature measured at 2 and 10 m above the ground is greater than $-0.6^{\\circ}$C. This can be explained as consequence of the lower temperature at 2 m because the higher temperature at 10 m inhibits the thermal convection below the primary mirror height."
    },
    "0606/astro-ph0606209_arXiv.txt": {
        "abstract": "Charged particles (X) decaying after primordial nucleosynthesis are constrained by the requirement that their decay products should not change the light element abundances drastically. If the decaying particle is negatively charged (X$^-$) then it will bind to the nuclei. We consider the effects of the decay of X when bound to Helium-4 and show that this will modify the Lithium abundances. ",
        "introduction": "Many models of physics beyond the standard model have long-lived charged particles  in their spectrum. For example, in the class of supergravity models  where the gravitino is the lightest supersymmetric particle (LSP), the next lightest (NLSP) can be a long-lived charged slepton \\cite{feng03a}. Similar phenomenology arises in Universal Extra Dimension models \\cite{feng03c}. Another example is the class of SUSY models with axino dark matter where again the NLSP could be long-lived \\cite{covi01,covi04,brandenburg05}. A third example is a supersymmetric model with almost degenerate LSP and NLSP such that the dominant decay channel for the NLSP is kinematically suppressed \\cite{sigurdson03,profumo04}.  Here we consider a new aspect of the early universe cosmology of these particles that affects the light element abundances. Other effects of these charged particles were previously considered in \\cite{dimopoulos89,kawasaki94,khlopov94,sedelnikov95,holtmann96,holtmann98,jedamzik99,cyburt03,feng03b,jedamzik04,kawasaki04,jedamzik06a,jedamzik06b,steffen06}  If the particle is negatively charged, it  will eventually form a bound state with the positively charged nuclei. When the particle decays, the decay products will electromagnetically interact with the nearby nucleus. If the nucleus is $^4$He, it may break up, creating $^3$He, T, D, p, n. These $^3$He, T and D nuclei will find other \\hef nuclei and  interact to form \\lisi and \\lise.  In addition to the production mechanism described above, charged particle decays result in a large non-thermal photon background that can destroy the created \\lisi and \\lise nuclei. When both the production and destruction processes are taken into account, we find that the overall \\lisi production could be in the same range as the Lithium isotopic abundance measured in few low metallicity popII stars \\cite{smith93,smith98,nissen99,vangioni-flam99,nissen00,asplund05}.  Other effects can come from the elastic scattering of the charged particle produced in the decay from the \\hef nucleus, and also from the change in the motion of the \\hef  in the bound state. We find that these effects produce insignificant changes to standard BBN abundances.  In this work we only consider singly charged bound state of \\hef nuclei \\cite{cahn81,belotsky05,fargion05,khlopov06}. We ignore the production of neutral atoms of \\hef bound to two X. We also do not consider the bound state of protons. This binding process starts below about 0.5 keV and could be important for lifetimes larger than about a year.  Another interesting process we have not considered here is that of forming stable ${\\rm Be}^8$-X bound states \\cite{cahn81}. A stable ${\\rm Be}^8$-X bound state has important implications for the production of Carbon in the early universe. We hope to return to these issues in future work.  While this work was being completed, two related articles appeared on the arXiv. The three articles all consider different processes and hence complement each other. The first by Pospelov \\cite{pospelov06} pointed out that the nuclear processes with a photon in the final state will be modified by the presence of the strong electric field of the X particle. Pospelov showed that the enhanced D on X-\\hef cross-section could produce orders of magnitude larger \\lisi abundance. However, it is important to note that if the decay lifetime is about $10^6$ seconds or larger, then \\lisi will also be destroyed due to the non-thermal photons from the decay.  The second paper by Kohri and Takayama \\cite{kohri06} worked out the changes in the light element abundances due to the change in coulomb barrier in the presence of the $X^-$, and due to the changes in the kinematics and energy balance for the reactions. They found that the Lithium abundances can be significantly modified from the standard BBN predictions. ",
        "conclusions": "We have concentrated here on the collision of \\het with background \\hef nuclei that would produce \\lisi nuclei. Standard BBN does not produce much \\lisi and thus the expectation is that this could be an important probe. The SBBN yield (without decays) is about $10^{-13}$ compared to the number density of H. This should be compared to the measured lithium isotopic ratios in some low metallicity popII stars which (conservatively) range between 0.01 and 0.1. A reasonable constraint on \\lisi then is that its abundance should not be much larger than that of \\lise nuclei (about $10^{-10}$ relative to the H number density).  It is of course possible that the \\lisi primordial abundance is larger, but then we are left with the question of why \\lisi  is depleted in much larger amounts than \\lise by astrophysical mechanisms.  We found that in some regions of parameter space it is possible to produce \\lisi in quantities that match the observed ``plateau'' abundance. This does not imply that the decaying charged particle scenario can explain the \\lisi abundance. A full calculation including all the light elements would be necessary to ascertain that. We also note that there are astrophysical scenarios wherein the observed abundance of \\lisi finds an explanation (\\eg \\cite{suzuki02,lambert04,rollinde05,reeves05,prantzos06,prodanovic06,nath06}).    The processes we have described here are clearly relevant for determining the light element abundances. Much more work needs to  be done before these processes can be combined with previous analyses to put bounds on the model parameter space. Throughout this work we have neglected the bound states of other light elements. The fraction of bound states of D, T, \\het,  \\lisi and \\lise are small owing to the much smaller abundance of the respective light elements. However, they could still have significant effects because of enhanced nuclear cross-sections. The correct treatment of this problem including all the cross-sections and bound states is beyond the scope of present work.  In summary, we have studied an aspect of the decay of charged particles in the early universe. We looked at the electromagnetic bound states of the charged particle with \\hef nuclei and considered the effects of the decay on the light element abundances. We found that cosmologically relevant quantities of \\lisi could be produced as a result of these decays."
    },
    "0606/astro-ph0606515_arXiv.txt": {
        "abstract": "{} {We present a mode identification based on new high-resolution time-series spectra of the non-radially pulsating $\\delta$ Scuti star FG~Vir (HD 106384, V = 6.57, A5V). From 2002 February to June a global Delta Scuti Network (DSN) campaign, utilizing high-resolution spectroscopy and simultaneous photometry has been conducted for FG~Vir in order to provide a theoretical pulsation model. In this campaign we have acquired 969 Echelle spectra covering 147 hours at six observatories.} {The mode identification was carried out by analyzing line profile variations by means of the Fourier parameter fit method, where the observational Fourier parameters across the line are fitted with theoretical values. This method is especially well suited for determining the azimuthal order $m$ of non-radial pulsation modes and thus complementary with the method of Daszynska-Daszkiewicz (2002) which does best at identifying the degree $\\ell$.} {15 frequencies between 9.2 and 33.5~\\cd~were detected spectroscopically. We determined the azimuthal order $m$ of 12 modes and constrained their harmonic degree $\\ell$. Only modes of low degree ($\\ell \\le 4$) were detected, most of them having axisymmetric character mainly due to the relatively low projected rotational velocity of FG Vir. The detected non-axisymmetric modes have azimuthal orders between -2 and 1. We derived an inclination of 19\\degree, which implies an equatorial rotational rate of 66~\\kms.} {} ",
        "introduction": " ",
        "conclusions": "We carried out a mode identification for the $\\delta$~Sct star FG Vir using high-resolution time-series spectra. 15 frequencies could be detected from the radial velocity and pixel-by-pixel intensity variations of a profile made from the combination of four iron lines.  We have presented the first successful application of the FPF method to line profiles of a $\\delta$ Sct star. The pulsational geometry of 12 independent frequencies was analyzed by means of this method. Seven frequencies were identified with axisymmetric modes having $\\ell$ values between 0 and 2. Five non-axisymmetric modes with $m$ values between -2 and 1 were detected. These modes could be modeled with $\\ell$ values as large as $4$. The observed radial velocity amplitudes range between  about 2 and 0.15~\\kms. The best models fit the profiles with 0.17~\\% relative radius variation for the dominant mode $\\nu_1$ and down to a value as small as 0.02~\\%. No high degree modes were found mainly due to the low rotational broadening of the line.  By simultaneously fitting amplitudes and phases of velocity and light variations we computed $f$-values for eleven frequencies. A comparison of these empirical $f$-values with theoretical values computed for different mixing length values showed that convection seems to be relatively ineffective in FG Vir.  The good identifications of the azimuthal order, combined with the determination of the inclination angle $i \\approx 19^{\\circ}$, will help to improve the treatment of rotational effects by seismic models.  The present fit of the theoretical models to the observations cannot be improved significantly by carrying out additional photometric measurements of FG Vir. Hundreds of unstable modes with $\\ell \\le 5$ are predicted by theory for the frequency range between 7 and 40~\\cd~(Pamyatnykh, priv. comm.). Reducing the observational detection limit from 0.2 to 0.1 mmag would yield many more frequencies but without a knowledge of their pulsational quantum numbers the number of free degrees is increased with every detected frequency. A more promising approach is to start the search for a good stellar pulsation model with the anchor points given by our mode identification."
    },
    "0606/astro-ph0606665_arXiv.txt": {
        "abstract": "Cosmic ray protons interacting with gas at the mean density of the interstellar medium in starburst galaxies lose energy rapidly via inelastic collisions with ambient nuclei.  The resulting pions produce secondary electrons and positrons, high-energy neutrinos, and $\\gamma$-ray photons.  We estimate the cumulative $\\gamma$-ray emission from starburst galaxies.  We find a total integrated background above 100 MeV of $F_\\gamma\\approx10^{-6}$\\,\\,GeV\\,\\,cm$^{-2}$\\,\\,s$^{-1}$ sr$^{-1}$ and a corresponding specific intensity at GeV energies of $\\nu I_\\nu \\approx 10^{-7}$\\,\\,GeV\\,\\,cm$^{-2}$\\,\\,s$^{-1}$\\,\\,sr$^{-1}$. Starbursts may thus account for a significant fraction of the extra-galactic $\\gamma$-ray background.  We show that the FIR-radio correlation provides a strong constraint on the $\\gamma$-ray emission from starburst galaxies because pions decay into both $\\gamma$-rays and radio-emitting electron/positron pairs.  We identify several nearby systems where the potential for observing $\\gamma$-ray emission is the most favorable (M82, NGC 253, \\& IC 342), predict their fluxes, and predict a linear FIR-$\\gamma$-ray correlation for the densest starbursts.  If established, the FIR-$\\gamma$-ray correlation would provide strong evidence for the ``calorimeter'' theory of the FIR-radio correlation and would imply that cosmic rays in starburst galaxies interact with gas at approximately the mean density of the interstellar medium (ISM), thereby providing an important constraint on the physics of the ISM in starbursts. ",
        "introduction": "The magnitude of the extra-galactic $\\gamma$-ray background is uncertain, primarily due to the presence of foreground contaminants (compare Keshet et al.~2004, Strong et al.~2004, Sreekumar et al.~1998). {\\it GLAST\\,}\\footnote{http://glast.gsfc.nasa.gov/; http://www-glast.stanford.edu/} will have an order of magnitude better sensitivity at GeV energies than previous experiments and should provide important constraints on the unresolved extra-galactic $\\gamma$-ray emission.  A number of potential contributors to the GeV $\\gamma$-ray background have been discussed in the literature including blazars (Stecker \\& Salamon 1996), galaxy clusters and groups (Colafrancesco \\& Blasi~1998; Dar \\& Shaviv~1995), intergalactic shocks and structure formation (Loeb \\& Waxman 2000; Keshet et al.~2003; Miniati 2002), dark matter annihilation (e.g., Ullio et al.~2002; Els{\\\"a}sser \\& Mannheim 2005), and normal star-forming galaxies (Pavlidou \\& Fields 2002).  In this paper we calculate the contribution of starburst galaxies to the GeV $\\gamma$-ray background. In particular, we assess the possibility that starbursts are cosmic ray proton calorimeters: the dense ISM of these systems acts as a beam dump for the total energy injected by supernovae into cosmic ray protons, with a portion of the proton energy emerging as $\\gamma$-rays and high-energy neutrinos.  In \\S\\ref{section:calorimeter} we review V\\\"olk's (1989) electron calorimeter model for the observed FIR-radio correlation (see Thompson et al.~2006, hereafter T06).  We then argue that starbursts may also be proton calorimeters.  In \\S\\ref{section:correlations}, we predict the $\\gamma$-ray flux from starburst galaxies and highlight which nearby systems are most likely to have detectable $\\gamma$-ray fluxes. We argue that the observed FIR-radio correlation provides important constraints on the $\\gamma$-ray emission from starbursts because pions from inelastic proton-proton collisions produce both $\\gamma$-rays and secondary electrons and positrons, which then produce radio synchrotron. Section \\ref{section:background} then discusses the cumulative starburst contribution to the extra-galactic $\\gamma$-ray background.  Estimates similar to those presented here for the high-energy neutrino background from starbursts have recently been made by Loeb \\& Waxman (2006).  Although the $\\gamma$-ray flux from individual starbursts has been estimated by several authors (e.g., Torres 2004, Torres et al.~2004, Cillis et al.~2005, Paglione et al.~1996, and Blom et al.~1999), this paper estimates the expected extra-galactic background, defines the requisite conditions for proton calorimetry, and explicitly connects the $\\gamma$-ray predictions with constraints from the FIR-radio correlation. ",
        "conclusions": "\\label{section:discussion}  Figure \\ref{fig:background} shows the numerically calculated $\\gamma$-ray background as a function of $\\gamma$-ray energy ({\\it solid lines}), for $p=2.0-2.4$ (see caption for details).  The inferred background from {\\it EGRET} from Strong et al.~(2004), as well as the conflicting upper limit at 1 GeV derived by Keshet et al.~(2004) are also shown.  Our results indicate that $\\gamma$-ray production from pion decay in starburst galaxies contributes significantly to the observed background above 100 MeV. The considerable uncertainties in the $\\gamma$-ray background determination, primarily because of foreground subtraction (see Keshet et al.~2004), complicate a more direct comparison.  An important point of this paper is that the absolute starburst $\\gamma$-ray flux --- both that from individual galaxies and that of the background --- cannot exceed that predicted in Figure \\ref{fig:background} and Table 1 without over-producing the radio emission from secondary electrons and positrons produced in charged pion decay (see \\S\\ref{section:correlations}, eq.~[\\ref{fir_radio}]). This conclusion is independent of whether starbursts are in fact proton calorimeters (as we have assumed), and can only be circumvented if ionization, inverse Compton, and bremsstrahlung losses significantly exceed synchrotron losses for cosmic ray electrons and positrons in starbursts; such a determination would put strong constraints on the origin of the FIR-radio correlation (T06; \\S\\ref{section:correlations}).  If detected by {\\it GLAST}, the low energy $\\gamma$-ray emission from systems like NGC 253 and M82 (\\S\\ref{section:calorimeter}, Fig.~\\ref{fig:background}, Table \\ref{table:starburst}), may help constrain the contribution of inverse Compton and bremsstrahlung to electron/positron losses.  The primary physical uncertainty in our estimate of the $\\gamma$-ray fluxes from starbursts lies in whether cosmic rays do in fact interact with gas at approximately the mean density of the ISM.  In particular, given that galactic winds efficiently remove mass and metals from galaxies (e.g., Heckman et al.~1990), it is unclear whether the cosmic rays actually interact with the bulk of the ISM, which is required for pion production to be significant (\\S\\ref{section:calorimeter}). Detection of $\\gamma$-rays from starbursts at the level predicted in this paper would thus provide an important constraint on the physics of the ISM in starbursts and on the coupling between supernovae and the dense ISM, which contains most of the mass.  In addition, the ionization and bremsstrahlung energy-loss timescales for electrons/positrons are similar to the proton-proton pion production timescale.  Thus, constraints on pion production in starbursts via $\\gamma$-ray observations directly constrain the importance of ionization and bremsstrahlung losses for shaping the radio emission from starburst galaxies (T06). By contrast, detection of (or upper limits on) the $\\gamma$-ray emission from starbursts well below our predictions would rule out the hypothesis that starbursts are proton calorimeters, although they would, in our opinion, still be electron calorimeters. The alternative possibility suggested by equation (\\ref{gamma_fir}), that $\\eta \\ll 0.05$, is unlikely, given the energetics of Galactic cosmic ray production (e.g., Strong et al.~2004).  We note that the total IR background from the particular model of the star formation history of the universe used to produce the $\\gamma$-ray background in Figure \\ref{fig:background} is $F^{\\rm TIR}\\approx2\\times10^{-5}$ ergs cm$^{-2}$ s$^{-1}$ sr$^{-1}$, or $\\approx20$ nW m$^{-2}$ sr$^{-1}$, with approximately 80\\% ($f_{\\rm cal}\\approx0.8$) coming from starburst galaxies.  A complete calculation, accounting for the contribution from the old stellar population (which contributes significantly to the background at $z \\approx 0$), finds a total background in starlight that is roughly twice as large (e.g., Nagamine et al.~2006).  Equation (\\ref{fir_radio}) and, indeed, the existence of the FIR-radio correlation, implies that the expected radio background is $\\nu I_\\nu({\\rm radio})\\approx7.5\\times10^{-10}\\eta_{0.05}/[2\\ln(\\gamma_{max})]\\approx2.7\\times10^{-11}$ ergs\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$.  A factor of two reduction in $\\nu I_\\nu({\\rm radio})$ has been included in this estimate as per the discussion after equation (\\ref{fir_radio}).  In addition to producing secondary electron/positron pairs and $\\gamma$-rays, pion production is also a significant source of high-energy neutrinos, with $\\approx 5\\%$ of the proton cosmic ray energy going into neutrinos.  Two-thirds of this energy goes to muon-type neutrinos. Therefore, for starburst galaxies that are proton calorimeters, we expect a FIR-neutrino correlation of the form $\\nu L_\\nu(\\mu-{\\rm neutrinos})=(2/3)\\nu L_\\nu({\\rm all\\,\\,neutrinos})\\approx 1.5\\times10^{-4}\\,\\eta_{0.05}L_{\\rm TIR}/\\ln(\\gamma_{max}) \\approx10^{-5}\\,\\eta_{0.05}L_{\\rm TIR}$, identical to the $\\gamma$-ray luminosity (eq.~[\\ref{gamma_fir}]). The corresponding cumulative extra-galactic $\\mu$-neutrino background at GeV energies is then the same as given in equation (\\ref{background}). Using a very similar estimate, Loeb \\& Waxman (2006) have recently argued that starbursts are likely to be an important contributor to the high-energy neutrino background.  Because the ratio of the $\\gamma$-ray to neutrino luminosity from pion decay depends only on the micro-physics of pion production, the $\\gamma$-ray fluxes from starbursts predicted in this paper will provide a crucial calibration of the expected flux of extra-galactic high energy neutrinos."
    },
    "0606/astro-ph0606386_arXiv.txt": {
        "abstract": "{ We describe the number counts and spatial distribution of 239 Distant Red Galaxies (DRGs), selected from the Early Data Release of the UKIDSS Ultra Deep Survey. The DRGs are identified by their very red infrared colours with $(J-K)_{AB}\\!>\\!1.3$, selected over 0.62 deg$^2$ to a $90\\%$ completeness limit of $K_{AB}\\!\\simeq\\! 20.7$.  This is the first time a large sample of bright DRGs has been studied within a contiguous area, and we provide the first measurements of their number counts and clustering.  The population shows strong angular clustering, intermediate between those of $K$-selected field galaxies and optical/infrared-selected Extremely Red Galaxies.  Adopting the redshift distributions determined from other recent studies, we infer a high correlation length of $r_0\\!\\sim\\!11\\,h^{-1}$Mpc. Such strong clustering could imply that our galaxies are hosted by very massive dark matter halos, consistent with the progenitors of present-day $L\\gtrsim L_*$ elliptical galaxies}. ",
        "introduction": "\\label{sec:intro}  A new near-infrared selection technique has been developed in recent years to sample galaxies in the high-redshift Universe.  By relying on purely near-infrared colours, this potentially avoids many biases which are inherent in optical techniques, particularly for detected dusty and/or evolved galaxies.  \\cite{2003ApJ...587L..79F} argue that the simple $(J-K)_{AB}\\!>\\!1.3$ colour selection criteria produces a sample that is mainly populated by galaxies at $z\\!>\\!2$, at least at faint $K$-band magnitudes ($K_{AB}\\!\\gtrsim\\!21$). These are the so-called Distant Red Galaxies (hereafter DRGs).  In the Faint Infrared Extragalactic Survey (FIRES), \\cite{2003ApJ...587L..79F} selected 14 candidate galaxies at $z\\!>\\!2$ to a depth of $K_{AB}\\!<\\!24.4$, of which 6 were spectroscopically confirmed \\citep{2003ApJ...587L..83V}.  \\cite{2005ApJ...624L..81L} found that approximately $70\\%$ of DRGs are dusty star forming galaxies and the remaining $30\\%$ are passively evolved galaxies.   Work by \\cite{2003ApJ...599..847R} suggest that DRGs may be a significant constituent of the $z\\!\\sim\\!2-3$ universe in terms of stellar mass.  \\cite{2004ApJ...616...40F} demonstrated that the average rest-frame optical colours of DRGs fall within the range covered by normal galaxies locally, unlike the Lyman-break galaxies (LBGs -- \\citeauthor{1996AJ....112..352S} \\citeyear{1996AJ....112..352S}) which are typically much bluer. Larger samples of DRGs covering a wide range in stellar mass are now required to fully understand the importance of this population. In particular, studies conducted so far have (by necessity) concentrated on DRGs selected over relatively small areas, and very little is known about the bright end of this population.  In terms of stellar mass, metalicity and star formation rate \\cite{2005ApJ...633..748R} find strong similarities between optically-selected and near-infrared selected galaxy samples. Clustering offers an alternative way to to study these populations. At large scales, the galaxy distribution is dominated by dark mater halo clustering, which is a strong function of halo mass.  Several studies have measured strong clustering strength for high redshift galaxies selected in the near-infrared ($r_0\\!=\\!10-15h^{-1}$Mpc) \\citep{2003ApJ...588...50D,DRG_MUSIC_astroph}, comparable to the most luminous galaxies in the local universe.  In this paper we present a study of the first large sample of DRGs selected at bright infrared magnitudes ($K_{AB}\\!<\\!21$) in a contiguous area.  We analyse the number counts and clustering and draw conclusions about their likely origin.  Throughout this paper, we assume $\\Omega_m\\!=\\!0.3$, $\\Omega_{\\Lambda}\\!=\\!0.7$ and $h\\!=\\!H_0/70$~km~s$^{-1}$~Mpc$^{-1}$.  \\section[]{UKIDSS UDS Early Data Release} \\label{sec:uds}  \\subsection{Survey and Early Data Release} \\label{subsec:uds-edr}  The UKIRT Infrared Deep Sky Survey (UKIDSS -- \\citeauthor{UKIDSS_astroph} \\citeyear{UKIDSS_astroph})\\footnote{\\texttt http://www.ukidss.org} is began observations in Spring 2005, using the Wide-Field Camera (WFCAM -- Casali et al., in prep.) at the 3.8-m United-Kingdom InfraRed Telescope (UKIRT).  Comprising 5 sub-surveys, it will take 7 years to complete and will cover a range of areas and depths.  The deepest of these 5 sub-surveys, the Ultra-Deep Survey (UDS) aims to cover 0.8 deg$^2$ to a depth of $K_{AB}\\!=\\!25.0$, $H_{AB}\\!=\\!25.4$, $J_{AB}\\!=\\!26.0$ ($5\\sigma$, point-source).  It is centred on the Subaru/XMM-Newton Deep Survey field (SXDS - \\citeauthor{sxds} \\citeyear{sxds}) at $02^h18^m00^s$, $-05^{\\circ}00'00''$ (J2000).     Since February 10 2006, the UKIDSS Early Data Release (EDR) has been available to the ESO community\\footnote{\\texttt http://surveys.roe.ac.uk/wsa}. A full description of this data release is given in \\cite{UKIDSS-EDR_astroph}.  \\subsection{Image stacking and mosaics} \\label{subsec:uds-mos}  The stacking of the UDS EDR data was performed by our team using a slightly different recipe to the standard UKIDSS pipeline. Given the (relatively) small size of the field, it is possible to create a full mosaic before extracting catalogues, rather than merge catalogues extracted from individual chips. Given the various jitter and offset sequences, this helps to optimise depth in overlap regions and to produce a more homogeneous final image.  Each observation block consists of a 4-pt mosaic to tile the $0.8$ deg$^2$ field, producing 16 images each of $\\simeq\\!15$ minutes exposure (see \\citeauthor{UKIDSS-EDR_astroph} \\citeyear{UKIDSS-EDR_astroph}). Individual reduced frames for each observation block were extracted from the UKIDSS pipeline as the starting point for our final mosaiced stack (including astrometric and photometric solutions).  We used the variance map produced by the pipeline to weight each frame before stacking and rescaled the pixel flux for each individual image using the pipeline zero-points.  The stacking was carried out in a two-step process by using the \\texttt{SWarp} software, an image resampling tool \\citep{swarp}.  The final mosaiced images in the $J$- and $K$-band have the same pixel scale of $0.1342''$, with identical field centres and image scales to simplify catalogue extraction.  The resulting images were visually inspected, and bad regions were masked out (areas around saturated stars, cosmetic problem areas, and low signal-to-noise borders). After masking, the usable area of this frame with uniform coverage is $0.62$ deg$^2$.  We note that this is smaller than the expected $0.8$ deg$^2$ mosaic since the UDS field centre was moved by $\\sim\\!8$ arcmin shortly after observations began (to allow the use of brighter guide stars).  The image seeing measured from the mosaiced images is $0.69''$ FWHM in $K$ and $0.80''$ in $J$.  The RMS accuracy of our astrometry is $\\simeq\\!0.05''$ (i.e. $<\\!1$ pixel) and for our photometry the RMS accuracy is $\\lesssim\\!2\\%$ in both $J$- and $K$- bands \\citep{UKIDSS-EDR_astroph}. From direct measurements of noise in a $2''$ aperture on the image we estimate $5\\sigma$ limiting magnitudes of $K_{AB}\\!=\\!22.5$ and $J_{AB}\\!=\\!22.5$.  \\subsection{Catalogue extraction} \\label{subsec:uds-cats}  We found that the standard UKIDSS source detection software did not produce optimal catalogues for the UDS. We therefore produced a much improved catalogue for the EDR by using the \\texttt{SExtractor} software \\citep{sextractor}. The $K$-band image was used as the source detection image, since this is measurably deeper for most galaxy colours. All $K_{AB}$ magnitudes quoted below are total magnitudes extracted using the \\texttt{SExtractor} parameter \\texttt{MAG\\_AUTO}, while all colour measurements are obtained from fixed $2''$ aperture magnitudes.   To optimise our catalogue extraction we performed a series of simulations to fine tune the \\texttt{SExtractor} parameters. Artificial point-like sources were added to the real $K$-band image using the observed PSF with FWHM $=\\!0.69''$ (rejecting regions containing bright sources), and distributed with magnitudes in the range $14\\!<\\!K_{AB}\\!<\\!24$.  From the resulting new image a catalogue was extracted using \\texttt{SExtractor} and compared with the list of artificial source positions. This process was repeated 1000 times, and the resulting statistics allow us to estimate the catalogue completeness and the evolution of photometric errors.  Using these simulations, \\texttt{SExtractor} detection parameters were tuned to maximise completeness at the noise-determined $5\\sigma$ depth of $K_{AB}\\!=\\!22.5$, while simultaneously minimising the number of spurious sources.  While formally optimised for point-like sources, we note that these were close to optimal when we generated artificial sources using a substantially more extended PSF (FWHM $\\!=\\!1.2''$).  Using these parameters, we extracted $78709$ sources over $0.62$deg$^2$ from the image, of which $34098$ were determined to be unsaturated, unmasked and from  regions of uniform coverage to  $K_{AB}\\!<\\!22.5$. These form the basis of the analysis outlined below. Assuming the background noise is symmetric about zero, we can estimate the spurious fraction by extracting sources from the inverted image and comparing with the number of sources extracted from the normal image. At our magnitude limit of $K_{AB}\\!=\\!22.5$ the fraction of spurious detections is found to be less than $1\\%$, while the completeness level is above $70 \\%$ (for point sources).   \\section[]{Selection and number counts} \\label{sec:drg}  \\subsection{Selection of DRGs} \\label{subsec:drg-sel}  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\caption{{\\it Lower panel:} $(J-K)_{AB}$ against $K_{AB}$ for sources in the UDS EDR field. Small points represent the full population of K-band selected objects, while larger points are those selected with $(J-K)_{AB}\\!>\\!1.3$ (and visually confirmed).  The  colour criteria, the magnitude limit at $K_{AB}\\!=\\!21.2$ and the magnitude completeness at $K_{AB}\\!=\\!20.7$ are highlighted, as is the crude boundary between galaxies and the stellar locus at $(J-K)_{AB}\\!=\\!0$.  {\\it Upper panel:} Errors as derived by \\texttt{SExtractor} on the $(J-K)_{AB}$ colours (for display purposes only $1/5$th of the points are shown). Mean values are also displayed, as are errors derived from simulations.  } \\label{fig:sel} \\end{center} \\end{figure}  From the catalogue described above we selected objects using the $(J-K)_{AB}\\!>\\!1.3$ criteria. A visual inspection of each source was then required to remove spurious detections, which at these extreme colours was found to be a relatively large fraction ($\\sim\\!20\\%$). The majority are caused by diffraction spikes and cross-talk images \\citep{UKIDSS-EDR_astroph} and are easy to identify and reject. This leaves 369 DRGs at $K_{AB}\\!<\\!21.2$, which represents the largest sample selected over a contiguous area. The surface density derived is $n\\!=\\!0.163\\pm0.009$ arcmin$^{-2}$.  Figure~\\ref{fig:sel} shows the $(J-K)$ colour of these galaxies versus $K$-band magnitude.  The object shown by a star was classified as a point-like source in our global catalogue, and is confirmed to be a star after visual inspection.   \\subsection{Photometric errors and contamination of DRG sample} \\label{subsec:drg-cont}  Since our sample is based on $(J-K)$ colour selection it is vital to carefully consider the effects of photometric errors. Since most galaxies show substantially bluer colours (Figure~\\ref{fig:sel}) we can expect the number density of DRGs to be artificially boosted at fainter magnitudes, as errors push objects above the $(J-K)_{AB}\\!>\\!1.3$ selection boundary.  As a lower limit to this contamination we could use the photometric errors derived from \\texttt{SExtractor}, and these are shown as a function of magnitude in Figure~\\ref{fig:sel}. Our experience suggests that analytically-determined errors from \\texttt{SExtractor} are likely to be underestimates, so we use the mean photometric errors obtained from the simulations described in section~\\ref{subsec:uds-cats}.  We used our simulated errors to estimate the contamination by randomising the real galaxy catalogue using Monte-Carlo simulations. For each object in our full catalogue, we allow the $(J-K)$ colour to vary assuming a Gaussian distribution with a standard deviation corresponding to the chosen photometric error. We then re-select our catalogue using the $(J-K)_{AB}\\!>\\!1.3$ criteria, and repeat this process 1000 times. This should provide an approximate upper limit on contamination, since we are randomising the observed galaxy catalogue (which already suffers from the effects of photometric errors).  Defining the contamination fraction from the number of objects scattered into our selection boundary minus those which are scattered out, our simulated source errors yield contamination fractions of $(46.0\\pm3.8)\\%$ at the limiting magnitude of $K_{AB}\\!<\\!21.2$, falling to $(16.8\\pm3.6)\\%$ at $K_{AB}\\!<\\!20.7$ (the estimated completeness limit; see Section~\\ref{subsec:drg-counts}).  We note that the typical error on the colour is $\\Delta(J-K)_{AB}\\!\\sim\\!0.1$ at $K_{AB}\\!=\\!20.7$.  As shown in \\cite{AEGIS_DRG_astroph}, a slight change of the $(J-K)$ colour selection does not have a major affect on the redshift distribution.   Based on these values we conclude that our number counts and clustering measurements are reasonably robust at $K_{AB}\\!<\\!20.7$, but will become increasingly unreliable at fainter magnitudes. We will therefore adopt a limit of $K_{AB}\\!=\\!20.7$ for further study, which produces a sample of $239$ DRGs.   \\subsection{Number counts of DRGs} \\label{subsec:drg-counts}  Figure~\\ref{fig:counts} shows the $K$-band differential number counts of our sample of DRGs.  The number counts indicate that our sample is complete up to approximately $K_{AB}\\!\\simeq\\!20.7$, after which the counts are clearly dropping. This defines our estimated completeness limit.  Our simulations suggest that the contamination due to photometric errors will be $\\sim\\!16\\%$ at $K_{AB}\\!<\\!20.7$. We conclude the dominant source of error in our number counts will be Poisson counting errors (plotted) and cosmic variance (discussed in section~\\ref{subsec:drg-dxs}).  At bright magnitudes (e.g.  $K_{AB}\\!<\\!20$) our UKIDSS data are entirely unique, and no studies exist in the literature for comparison.  At fainter magnitudes, our counts are in very good agreement with the DRG counts from the \\cite{DRG_MUSIC_astroph} sample. They are also consistent with the number counts from the AEGIS survey (Foucaud et al. in prep.; see also \\citeauthor{AEGIS_DRG_astroph} \\citeyear{AEGIS_DRG_astroph}).  Combining literature data with the present work we can examine the global shape of the DRG number counts over a very wide dynamic range ($18.5\\!<\\!K_{AB}\\!<\\!25.0$).  This strongly suggests a break feature in the slope at $K_{AB}\\!\\sim\\!20.5$ which is an effect already seen in the global K-band number counts (e.g. \\citeauthor{GCW} \\citeyear{GCW}).  We note that the projected density of DRGs is approximately 10 times lower than EROs, and approximately 100 times lower than the global galaxy counts at a given magnitude.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{fig2.eps}} \\caption{$K$-band differential number counts for our sample of DRGs. The errorbars plotted are computed from poissonian small number statistics \\citep{1986ApJ...303..336G}.  For comparison, we have overplotted the number counts for DRGs derived from the DXS EDR survey, with errors representing the field-to-field variance (see section~\\ref{subsec:drg-dxs} -- these points are slightly shifted for display purposes), and from the literature from fainter samples \\citep{2003AJ....125.1107L,2004ApJ...616...40F,2005ApJ...633..748R,DRG_MUSIC_astroph}. EROs number counts from the UDS/DXS SV sample are shown as well, for comparison. The plot inset shows number counts over a larger magnitude range.} \\label{fig:counts} \\end{center} \\end{figure}   \\subsection{Cosmic variance} \\label{subsec:drg-dxs}  As a simple test of cosmic variance, and to investigate whether the UDS is an unusual field, we used the data available from the UKIDSS Deep Extragalactic Survey (DXS) to perform a comparison study.  The DXS is the other deep extragalactic component of UKIDSS \\citep{UKIDSS_astroph}, consisting of 4 fields with a 7-year goal of observing 35 deg$^2$ to depths of $K_{AB}\\!=\\!22.7$ and $J_{AB}\\!=\\!23.2$. We used data from 3 fields observed in both $J$- and $K$-bands in the UKIDSS EDR, covering $\\sim\\!2900$ arcmin$^2$, $\\sim\\!4500$ arcmin$^2$ and $\\sim\\!2900$ arcmin$^2$ respectively. While exposure times are similar to the UDS EDR, the observing conditions are generally poorer. Direct comparison is also complicated by the different source extraction methods used by the DXS.  We applied the same selection method described in section~\\ref{subsec:drg-sel}, except that we did not visually inspect the samples. This selects $1523$ objects in total, of which we estimate approximately $20\\%$ are likely to be spurious (see section~\\ref{subsec:drg-sel}), with a similar fraction likely to be artificially boosted due to photometric errors at faint magnitudes (section~\\ref{subsec:drg-cont}). Since these errors are smaller than the errors in the DXS counts, for simplicity we opt not to make these corrections in our comparison with the UDS. We derived a median surface density of $n\\!=\\!0.176\\pm0.075$ arcmin$^{-2}$ at $K_{AB}\\!=\\!21.2$, in very good agreement with the UDS value.  The resulting median counts from the 3 DXS samples are overplotted in figure~\\ref{fig:counts}, with errors representing the field-to-field RMS variance. The agreement with UDS is very good. Although no corrections were applied, this crude comparison suggests that the density of DRGs is stable and broadly consistent between fields. ",
        "conclusions": "\\label{sec:summ}  We have extracted a large sample of bright Distant Red Galaxies from the UKIDSS UDS EDR. Our catalogue contains $369$ DRGs to a limiting magnitude of $K_{AB}\\!=\\!21.2$, extracted over an area of $0.62$ deg$^2$.  The fainter $K_{AB}\\!>\\!20.0$ number counts are in good agreement with previous estimates, while at brighter magnitudes the sample is unique.  Using simulations we determined that contamination due to photometric errors is below $\\sim\\!16\\%$ at an approximate completeness limit of $K_{AB}\\!<\\!20.7$.  From this sample we extracted a sub-sample of 239 bright DRGs to a limit of $K_{AB}\\!=\\!20.7$. These bright DRGs appear highly clustered, and we determine a correlation length of $r_0\\!\\simeq\\!11 \\,h^{-1}$Mpc and a bias measurement $b\\!\\simeq\\!4.5$, assuming the sample lies at a mean redshift of $\\bar{z}\\!=\\!1.0$ with a standard deviation of $\\sigma\\!=\\!0.25$ (consistent with studies at similar depths -- \\citeauthor{AEGIS_DRG_astroph} \\citeyear{AEGIS_DRG_astroph}). They appear more clustered than fainter samples of DRGs derived at these redshifts, which may be evidence for luminosity segregation, in agreement with biased galaxy formation scenarios."
    },
    "0606/astro-ph0606453_arXiv.txt": {
        "abstract": "\\textbf{{A detailed spectroscopic study covering the blue to near-infrared wavelength range ($\\lambda$3700~\\AA~-1$\\mu$m) was performed for a sample of 34 HII galaxies in order to derive fundamental parameters for their HII regions and ionizing sources, as well as gaseous metal abundances.  All the spectra included the nebular [SIII]$\\lambda$$\\lambda$9069,9532~\\AA~lines, given their importance in the derivation of the S/H abundance and relevant ionization diagnostics.}} {A systematic method was followed to correct the near-IR [SIII] line fluxes for the effects of the atmospheric transmission. A comparative analysis of the predictions of the empirical abundance indicators R$_{23}$ and S$_{23}$ was performed for our sample galaxies.  The relative hardness of their ionizing sources was studied using the $\\eta$' parameter and exploring the role played by metallicity.} {For 22 galaxies of the sample, a value of the electron temperature T$_{e}$[SIII] was derived, along with their ionic and total S/H abundances. Their ionic and total O/H abundances were derived using direct determinations of T$_{e}$[OIII]. For the rest of the objects, the total S/H abundance was derived using the S$_{23}$ calibration. The abundance range covered by our sample goes from 1/20 solar up to solar metallicity. Six galaxies present 12+log (O/H) $<$ 7.8 dex. The mean S/O ratio derived in this work is log (S/O)=-1.68$\\pm$0.20 dex, 1$\\sigma$ below the solar (S/O)$_\\odot$ value. The S/O abundance ratio shows no significant trend with O/H over the range of abundance covered in this work, in agreement with previous findings. There is a trend for HII galaxies with lower gaseous metallicity to present harder ionizing spectra. We compared the distribution of the ionic ratios O$^{+}$/O$^{++}$ vs. S$^{+}$/S$^{++}$ derived for our sample with the predictions of a grid of photoionization models performed for three different stellar effective temperatures. This analysis indicates that a large fraction of galaxies in our sample seem to be ionized by extremely hard spectra, in line with recent suggestions for extra ionizing sources in HII galaxies.} {} ",
        "introduction": "\\label{intro}  HII galaxies are galaxies undergoing violent star formation (Searle \\& Sargent 1972; Terlevich et al. 1991; Cair\\'os et al. 2000). Their optical spectra show strong emission lines (recombination lines of hydrogen and helium, as well as forbidden lines of elements like oxygen, neon, nitrogen, sulfur, among others) that are very similar to the spectra of extragalactic HII regions. Analysis of their spectra shows that they are low-metallicity objects with the metallicity varying from 1/40$Z_{\\odot}$ to 1/2$Z_{\\odot}$ (e.g. Terlevich et al. 1991; Telles 1995 and references therein; V\\'{\\i}lchez \\& Iglesias-P\\'aramo 1998,2003; Thuan \\& Izotov 2005). Among them we can find the least chemically-evolved galaxies in the local Universe.  The study of elemental abundances in emission-line galaxies gives information about their chemical evolution and star formation history. Outside the Local Group, emission lines from ionized gas represent the principal means of deriving abundances, as energy is concentrated in a few conspicuous emission lines. Abundances for the stellar population are derived from absorption features, which are more numerous and require much higher signal-to-noise spectra to be derived meaningfully.  In HII galaxies the metal enrichment of the interstellar medium by supernovae has been operating typically in low-metallicity environments.  Oxygen is the most frequently used element in deriving abundances from emission lines: abundances are easily derived, as the main ionization stages are observable in the optical range. Furthermore, oxygen is particularly suitable for chemical evolution studies, as it traces the overall metallicity very well. It originates quasi exclusively from the nucleosynthesis in type II supernovae progenitors (Meynet \\& Maeder 2002; Pagel 1997; Woosley \\& Weaver 1995). While the sources of oxygen are well-determined and the most important ionization stages can be observed in the optical range, some uncertainties still remain  about the sulfur yields and its sources. In addition, not all the ionization stages can be observed in the optical range and important ionization correction factors (ICFs) must be applied to derive the total sulfur abundance. Hence comparing S and O abundances can give us some clues to sulfur nucleosynthesis and the masses of the stars where the sulfur tends to be formed.   To derive oxygen abundances, one should first derive the electron temperature, which requires the measurement of faint auroral lines, like [OIII]$\\lambda$4363~\\AA, which are often not detected. The alternative is to use strong line-abundance indicators, like R$_{23}$\\footnote{R$_{23}$=([OII]$\\lambda$3727 + [OIII]$\\lambda$$\\lambda$4959,5007)/H$\\beta$}, which calibrated empirically (Pagel et al. 1979, Pilyugin 2001) or through photoionization models (e.g. McGaugh 1991). However, the relation between R$_{23}$ and oxygen abundance presents the noticeable drawback of being double-valued.   V\\'{\\i}lchez \\& Esteban (1996) proposed S$_{23}$\\footnote{S$_{23}$ = ([SII]$\\lambda$$\\lambda$6717,31 + [SIII]$\\lambda$$\\lambda$9069,9532)/H$\\beta$} as an alternative abundance indicator. In contrast to oxygen, S$_{23}$ remains single-valued up to abundances above solar value. Furthermore, sulfur should be as useful as oxygen for tracing metallicity. From an observational point of view, S$_{23}$ has the advantage over R$_{23}$ that the [SII] and [SIII] lines are less affected by reddening (P\\'erez-Montero et al. 2005; hereinafter PM06).   To produce an accurate derivation of S/H abundance, the importance of using the nebular [SIII] lines can not be overlooked (e.g. Dennefeld \\& Stasi\\'nska 1983; V\\'{\\i}lchez et al. 1988, Garnett 1989; Bresolin et al. 2004). Photoionization models indicate that S$^{++}$ is the dominant sulfur ion (Garnett 1989; hereinafter G89), which presents three forbidden transitions at [SIII]$\\lambda$$\\lambda$9069,9532~$\\AA$~ and $\\lambda$6312~$\\AA$~ in the optical to near-IR (NIR) range (analogs to [OIII]$\\lambda$$\\lambda$4959,5007~$\\AA$~ and $\\lambda$4363~$\\AA$). The [SIII]$\\lambda$6312~$\\AA$~line is faint, highly temperature-sensitive, and it can induce several biases in the derived S/H abundance. The NIR [SIII] lines can be quite strong, and a detailed telluric atmosphere correction has to be applied to them. P\\'erez-Montero \\& D\\'{\\i}az (2003) (hereinafter PMD03) and G89 derive the S$^{++}$ ionic abundance for samples of about one dozen emission-line galaxies, both using the nebular [SIII]$\\lambda$9069~$\\AA$~line. Recent work by Izotov et al. (2005) (hereinafter I05) presents S/H abundances for a large number of metal-poor emission-line galaxies from the SDSS-DR3\\footnote{Data Release 3 of Sloan Digital Sky Survey}; however, the auroral line [SIII]$\\lambda$6312~$\\AA$~ was used in this work to calculate the S$^{++}$ ionic abundance.  Here we present long-slit spectrophotometric observations of a sample of 34 HII galaxies to make a detailed analysis of their chemical abundances. The wide coverage of our spectra ($\\lambda$3700~\\AA~ - 1$\\mu$m) for all the galaxies in the sample provides us all the emission lines needed to estimate the oxygen and sulfur abundances directly. All the S/H abundances were estimated using a nebular [SIII] line, so that uncertainties related to the use of the auroral line [SIII]$\\lambda$6312~\\AA~are avoided.  In addition, this wavelength coverage allowed us to study the properties of the ionizing clusters of HII galaxies making use of the $\\eta$' parameter and sequences of photoionization models.  In the next section we describe our sample of galaxies, the observations, and data reduction and present the line intensities. In Sect.3 we perform a comparative study between R$_{23}$ and S$_{23}$ abundance indicators, present an analysis about ionization structure and ionizing sources, and discuss the abundance results for the sample. Finally in Sect.4 we summarize our conclusions. ",
        "conclusions": "\\subsection{Empirical abundance indicators for our sample}  Commonly used strong line empirical abundance indicators are R$_{23}$ (Pagel et al. 1979; Edmunds \\& Pagel 1984; McCall et al.  1985; McGaugh 1991) and S$_{23(4)}$ (V\\'{\\i}lchez \\& Esteban 1996; D\\'{\\i}az \\& P\\'erez-Montero 2000; Oey \\& Shields 2000; PM06).  Though widely used, R$_{23}$ presents the drawback of having a double-valued relation with oxygen abundance, creating an intrinsic uncertainty on the derived O/H abundances. The turnover region of the relation R$_{23}$ vs. O/H takes place for log R$_{23}$ $\\gtrsim$ 0.9, corresponding to 8.0 $\\lesssim$ 12+log (O/H) $\\lesssim$ 8.4. In this region, R$_{23}$ is sensitive to ionization conditions but almost insensitive to O/H. Most of the HII galaxies from our sample show R$_{23}$ values within this ill- defined region, which is what we want to explore.  The S$_{23}$ parameter introduced by V\\'{\\i}lchez \\& Esteban (1996) has been used as an O/H abundance calibrator in D\\'{\\i}az $\\&$ P\\'erez-Montero (2000) and P\\'erez-Montero \\& D\\'{\\i}az 2005 (hereinafter PMD05). It has also been demonstrated that S$_{23}$ is an efficient S/H abundance calibrator in PM06. It presents several advantages over R$_{23}$. First, it has a lower dependence on the ionization parameter and remains single-valued up to metallicities higher than solar, 12+log (O/H)$_{\\odot}$ = 8.69 and 12+log (S/H)$_{\\odot}$ = 7.19 (Lodders 2003). Secondly, the sulfur emission lines are less affected by reddening.  However, the spectral regions around the red [SIII] lines are affected by atmospheric absorption.   The N2\\footnote{N2=log ([NII]$\\lambda$6584/H$\\alpha$)} parameter has also been proposed as an abundance indicator (Denicol\\'o et al. 2002; Van Zee et al. 1998). This parameter offers several advantages, because it involves easily measurable lines that are available for a wide redshift range (up to z $\\sim$ 2.5). The N2 vs. O/H relation seems monotonic and the [NII]/H$\\alpha$ ratio does not depend on reddening correction or flux calibration. The drawbacks are that the [NII] lines can be affected by other excitation sources (see Van Zee et al. 1998). In addition, N2 is sensitive to ionization conditions and relative N/O abundance variations.   \\begin{figure*} \\centering \\includegraphics[bb= 18 510 592 718,width=14cm,clip]{4488fig3.ps} \\caption{The left panel presents the relation between S$_{23}$ and R$_{23}$; the middle and right panels  show the relations between log (1.3{\\it x}[NII]6584/H$\\alpha$) and the empirical indicators of abundances, R$_{23}$ and S$_{23}$ respectively, for all galaxies of our sample.}  \\label{S23_R23} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[bb=20 426 580 692,width=10cm,clip]{4488fig4.ps} \\caption{The left panel shows the relation of log ([SII]6717,31/H$\\alpha$) vs. log (1.3{\\it x}[NII]6584/H$\\alpha$); the right panel presents the relation between log ([SIII]/[SII]) and log S$_{23}$ for all galaxies of our sample.}  \\label{NII_SII} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[bb=20 426 580 692,width=10cm,clip]{4488fig5.ps} \\caption{The left panel shows the relation log $\\eta$' vs. log S$_{23}$ for all galaxies in our sample. Full and empty symbols represent the objects with and without T$_{e}$[OIII], respectively. The right panel presents the relation between log (S$^{+}$/S$^{++}$) and log (O$^{+}$/O$^{++}$) for the galaxies with electron temperature; solid, short-dashed, and long-dashed lines show the loci for three series of photoionization models for T$_{eff}$=50kK, 40kK, and 30kK, respectively. See the text for details.} \\label{cumulos} \\end{figure*}  Figure~\\ref{S23_R23}a shows the relation between S$_{23}$\\footnote{The S$_{23}$ values were derived using the sulfur emission lines quoted in Table~\\ref{fluxes}} and R$_{23}$\\footnote{The R$_{23}$ values were calculated using the oxygen emission lines from our blue spectra (Kehrig et al 2004).}  for the galaxies of our sample. Although log R$_{23}$ values remain approximately constant for most galaxies, log S$_{23}$ values present a variation of approximately 0.8 dex. We can see that for galaxies in the turn-over region of the relation between $R_{23}$ and O/H, R$_{23}$ does not correlate with S$_{23}$. This fact is easily understood since the relationship between $S_{23}$ and O/H is not bivaluate in the metallicity range that we are interested in. Besides, Figs.~\\ref{S23_R23}b and c show that R$_{23}$ does not correlate with [NII]/H$\\alpha$, contrary to the behavior of S$_{23}$. Therefore, for objects located in the ill-defined region of R$_{23}$ vs. O/H, S$_{23}$ can be used to derive chemical abundances, especially the S/H abundance.  \\subsection{Ionization structure and the ionizing sources}   In photoionized regions like the ones we consider here, the physical properties that determine line intensities are the luminosities and temperatures of the ionizing stars, the gas density, the optical thickness to the ionizing photons, and the chemical abundances. Because S$_{23}$ is a combination of strong line intensities, it can be affected by several effects. Taking S$_{23}$ as an abundance indicator, we are not considering, to first order, the detailed effects produced by changes in the physical properties mentioned above. For this reason, it is important to check the sensitivity of S$_{23}$ to some of these properties.  The optical thickness to ionizing photons is the first to assess.  As can be seen in Fig.~\\ref{NII_SII}a, [NII]/H$\\alpha$ and [SII]/H$\\alpha$ present a strong correlation, discarding density boundary effects for the sample galaxies (see e.g. McCall et al. 1985); this correlation implies a statistically significant relation between N$^{+}$/N and S$^{+}$/S, as expected from standard HII region models.   Ratios of line intensities of elements in different ionization stages, such as [OIII]/[OII] or [SIII]/[SII], are sensitive to combinations of the luminosity, the gas density and geometry, and the radiation hardness; but they are insensitive to abundances at first order, as they originate in the same element. Any variation with respect to such line ratios indicates a sensitivity to these physical parameters, though in a combination that might not be straightforward to derive.  Figure~\\ref{NII_SII}b shows the dependence of S$_{23}$ on [SIII]/[SII]\\footnote{[SIII]/[SII] = [SIII]$\\lambda$$\\lambda$9069,9532/[SII]$\\lambda$$\\lambda$6717,6731}. While [NII]/H$\\alpha$ shows a well-known dependence on the excitation degree (e.g. McCall et al. 1985), the dependence is much weaker for S$_{23}$, being mostly marginal. Despite the fact that S$_{23}$ possesses a narrower dynamical range than [NII]/H$\\alpha$, we consider it a better abundance indicator for our sample than [NII]/H$\\alpha$, since S$_{23}$ does not show any strong dependence on the ionization conditions.  Having a wide wavelength coverage has allowed us to study the properties of the ionizing sources in our sample of HII galaxies. This study could help to constrain the range of applicability of photoionization models and stellar atmospheres in order to fit the observations, thus improving our understanding of the mechanisms that heat the HII regions in HII galaxies (Stasi\\'nska \\& Schaerer 1999; Thuan \\& Izotov 2005). A convenient hardness index is the parameter $\\eta$' introduced by V\\'{\\i}lchez \\& Pagel (1988):  \\[\\eta' = \\frac{[OII]\\lambda\\lambda3727,29/[OIII]\\lambda\\lambda4959,5007}{[SII]\\lambda\\lambda6717,31/[SIII]\\lambda\\lambda9069,9532} \\] This parameter has been recommended as a criterion for effective temperature of the ionizing star(s), T$_{eff}$, of HII regions, in the sense that softer ionizing spectra have higher values of $\\eta$' (e.g. V\\'{\\i}lchez \\& Pagel 1988; Kennicutt et al. 2000). In Fig.~\\ref{cumulos}a we present the behavior of the parameter $\\eta$' with respect to S$_{23}$ for our sample. Overall, we can see that $\\eta$' goes up with S$_{23}$, implying that the hardness of the ionizing spectra increases with lower gaseous metallicity for our sample of HII galaxies (e.g. Bresolin \\& Kennicutt 1999; Oey et al. 2000 and references therein; see also Mart\\'{\\i}n-Hernandez et al. 2002). A possible explanation for higher temperatures at lower metallicities has been suggested by Massey et al. (2004). They find that, for a range of stellar atmosphere models, stars of early through mid-O types in a Magellanic Cloud sample are 3000K-4000K hotter than their Galactic (metal-richer) counterparts; and they attribute their higher temperatures to the minor importance of wind emission, wind blanketing, and metal-line blanketing at lower metallicities.    In Fig.~\\ref{cumulos}b we present the relationship between the ionic ratios S$^{+}$/S$^{++}$ and O$^{+}$/O$^{++}$ for the subset of HII galaxies with electron temperature. In this figure we show the loci of the average predictions of three sequences of single-star photoionization models (computed with the photoionization code Cloudy 96; Ferland 2002), performed using CoStar model atmospheres (Schaerer \\& de Koter 1997) T$_{eff}$=50kK, 40kK and 30kK. Along each line, the metallicities vary between $Z_{\\odot}$/20 and $Z_{\\odot}$/2, and the ionization parameter changes from log U = -2 to log U = -3 (a detailed description of the grids of the photo-ionization models used in this work can be found in PMD05). According to these models, a large fraction of the galaxies appear to harbor ionizing sources with spectra harder than the spectrum produced by a 50kK effective temperature CoStar atmosphere (Schaerer \\& de Koter 1997). Kennicutt et al. (2000) have found, for a sample of HII regions (in the Galaxy and Magellanic Cloud), that empirically-based stellar-temperature indices present a decrease in mean stellar temperature with increasing abundance. They show, however, that the typical T$_{eff}$ for their HII regions are below $\\sim$ 55kK (at $Z_{\\odot}$/5), in agreement with the model-based results by Bresolin et al. (1999). Though any calibration of nebular empirical parameters in terms of T$_{eff}$ should be a function of the atmosphere and photoionization models used, it seems that T$_{eff}$ $\\sim$ 55kK  represents a reasonable upper limit for the effective temperature in HII regions in contrast to HII galaxies. These findings suggest the existence of very hard spectral energy distributions as ionizing sources in some HII galaxies. Stasi\\'nska \\& Schaerer (1999), modelling the HII regions in IZw18, argue that extra heating sources might well exist, in addition to ionizing clusters, giving rise to large temperature variations and enhancing the [OIII]$\\lambda$4363 emission. I05 have also invoked extra heating sources (i.e. X-ray ionizing sources) to explain the high-ionization emission lines observed in some metal-poor emission-line galaxies.  More observations, covering a wide range in wavelength, as well as dedicated work using photoionization models for evolving starbursts with a library of different ionizing spectra, are needed to further investigate  the above suggestions.   \\begin{figure} \\centering \\includegraphics[width=9cm,clip]{4488fig6.ps} \\caption{A comparison between the measured line temperatures of [OIII] and [SIII]. The electron temperatures, t$_{e}$[OIII] and t$_{e}$[SIII], are shown in units of 10$^{4}$K. The dashed and solid lines are the photoionization models'relation between these temperatures from PMD05 and I05, respectively.}  \\label{graf_temp} \\end{figure}   \\begin{table*} \\caption{Physical properties and chemical abundances for the galaxies with S/H and O/H derived directly} \\label{6gal} \\centering \\renewcommand{\\footnoterule}{} \\begin{minipage}{\\textwidth} \\begin{tabular}{lcccccccc} \\hline\\hline  Object & \\multicolumn{3}{c}{IIZW40} & Tol0226-390 & Tol1924-416 & Tol0538-416  \\\\ &G89\\footnote{References - G89: Garnett (1989), PMD03: P\\'erez-Montero \\& D\\'{\\i}az (2003)} &PMD03$^{a}$ &This Work  &  &  &  &  & \\\\  \\hline  $n_{e}$([SII])     &100. &290$\\pm$60 &197$\\pm$167   &221 $\\pm$64 &78$\\pm$28  &$\\leq$ 100.\\footnote{Low Density Limit}    \\\\  $T_{e}$([SIII])/10$^{4}$K  &1.35$\\pm$0.12 &1.30$\\pm$0.05  &1.04\\footnote{$T_{e}$([SIII]) derived using [SIII]$\\lambda$6312~\\AA~ line from the red spectra}$\\pm$0.10  &1.40$^{c}$$\\pm$   0.18  &1.43$^{c}$$\\pm$   0.17  &     1.29$^{c}$$\\pm$   0.20  \\\\  $T_{e}$([OIII])/10$^{4}$K  &1.33$\\pm$0.02  &1.34$\\pm$0.03 &1.35$\\pm$   0.06 &    1.19$\\pm$   0.04 &    1.39$\\pm$   0.07 &    1.30$\\pm$   0.06 \\\\ 12+log(S$^{+}$/H$^{+}$) &5.21 $\\pm$0.09 &5.21$\\pm$0.09  &5.23$\\pm$   0.12   &    5.81$\\pm$   0.07   &    5.50$\\pm$   0.06   &    6.18$\\pm$   0.04   \\\\ 12+log(S$^{++}$/H$^{+}$) &6.00 $\\pm$0.11 &5.99$\\pm$0.04  &6.24$\\pm$0.14   &6.17$\\pm$0.16   &5.91$\\pm$   0.13   &    6.20$\\pm$   0.21   \\\\ 12+log[(S$^{+}$ + S$^{++}$)/H$^{+}$] &6.07$\\pm$0.11  &6.07$\\pm$0.05  &    6.28$\\pm$   0.14   &    6.33$\\pm$   0.13   &    6.06$\\pm$   0.11   &    6.49$\\pm$   0.13   \\\\ ICF(S$^{+}$ + S$^{++}$) &1.95 &1.86  &    1.98    &    1.05    &    1.06   &    1.06   \\\\ 12+log (S/H) &6.36$\\pm$0.11 &6.34$\\pm$0.05  &    6.53$\\pm$   0.34   &    6.35$\\pm$   0.14   &    6.08$\\pm$   0.12   &    6.53$\\pm$   0.14   \\\\ 12+log(O$^{+}$/H$^{+}$) &6.95$\\pm$0.07 &7.08$\\pm$0.07  &    6.88$\\pm$   0.26   &    7.37$\\pm$   0.15   &    7.13$\\pm$   0.14   &    6.99$\\pm$   0.07   \\\\ 12+log(O$^{++}$/H$^{+}$) &8.01$\\pm$0.02 &8.03$\\pm$0.08  &    7.96$\\pm$   0.07   &    7.94$\\pm$   0.07   &    7.83$\\pm$   0.08   &    7.79$\\pm$   0.07   \\\\ 12+log (O/H) &8.05$\\pm$0.02 &8.08$\\pm$0.03  & 7.99$\\pm$   0.09   &8.05$\\pm$   0.09   &7.91$\\pm$   0.09   &    7.86$\\pm$   0.07   \\\\  \\hline\\hline  Object &  Cam0840+1201 &  CTS1008 &  UM238 &  Tol0104+388 &  UM306 &  UM307 \\\\ \\hline   $n_{e}$([SII]) &$\\leq$ 100.$^{b}$  &169$\\pm$50 &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$  &111$\\pm$109 &845$\\pm$102      \\\\ $T_{e}$([SIII])/10$^{4}$K  &     1.59$^{c}$$\\pm$   0.38  &     1.22$^{c}$$\\pm$   0.22  &     1.61\\footnote{$T_{e}$([SIII]) derived from the relation $T_{e}$([SIII]) = 10500$T_{e}$([OIII]) - 800, see the text for details.}$\\pm$   0.09  &     0.85\\footnote{$T_{e}$([SIII]) derived using [SIII]$\\lambda$6312~\\AA~ line from the blue spectra}$\\pm$   0.08  &  1.13$^{e}$$\\pm$   0.17  &     1.08$^{e}$$\\pm$   0.07  \\\\ $T_{e}$([OIII])/10$^{4}$K  &    1.36$\\pm$   0.06   &    1.24$\\pm$   0.05   &    1.61$\\pm$   0.09   &    1.55$\\pm$   0.08   &    1.21$\\pm$   0.05   &    1.11$\\pm$   0.07   \\\\ 12+log(S$^{+}$/H$^{+}$) &    5.52$\\pm$   0.05   &    5.50$\\pm$   0.11   &    5.54$\\pm$   0.17   &    6.13$\\pm$   0.03   &    5.83$\\pm$   0.05   &    6.40$\\pm$   0.17   \\\\ 12+log(S$^{++}$/H$^{+}$) &    6.00$\\pm$   0.29   &    6.20$\\pm$   0.22   &    5.34$\\pm$   0.09   &    6.96$\\pm$   0.16   &    6.36$\\pm$   0.24   &    6.31$\\pm$   0.10   \\\\ 12+log[(S$^{+}$ + S$^{++}$)/H$^{+}$] &    6.12$\\pm$   0.24   &    6.28$\\pm$   0.20   &    5.75$\\pm$   0.14   &    7.02$\\pm$   0.14   &    6.48$\\pm$   0.20   &    6.66$\\pm$   0.14   \\\\ ICF(S$^{+}$ + S$^{++}$) &    1.74   &    1.09   &    1.01   &    1.00   &    1.09   &    1.04   \\\\ 12+log (S/H) &    6.36$\\pm$   0.33   &    6.32$\\pm$   0.22   &    5.76$\\pm$   0.14   &    7.02$\\pm$   0.18   &    6.52$\\pm$   0.21   &    6.68$\\pm$   0.14   \\\\ 12+log(O$^{+}$/H$^{+}$) &    6.86$\\pm$   0.14   &    6.90$\\pm$   0.27   &    7.60$\\pm$   0.36   &    8.33$\\pm$   0.09   &    7.75$\\pm$   0.07   &    8.68$\\pm$   0.35   \\\\ 12+log(O$^{++}$/H$^{+}$) &    7.77$\\pm$   0.07   &    8.01$\\pm$   0.07   &    7.65$\\pm$   0.08   &    7.53$\\pm$   0.08   &    7.92$\\pm$   0.07   &    7.57$\\pm$   0.11   \\\\ 12+log (O/H) &    7.82$\\pm$   0.08   &    8.05$\\pm$   0.09   &    7.92$\\pm$   0.23   &    8.40$\\pm$   0.09   &    8.14$\\pm$   0.07   &    8.71$\\pm$   0.34   \\\\  \\hline\\hline  Object &  UM323 &  Tol0140+420 &  UM391 &  UM396 &  UM408 &  Tol0306+405 \\\\ \\hline  $n_{e}$([SII]) &630:$^{f}$ &144:\\footnote{Upper limit} &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$ &220$\\pm$64    \\\\ $T_{e}$([SIII])/10$^{4}$K  &     1.35$^{e}$$\\pm$   0.14  &     2.25$^{e}$$\\pm$   0.65  &     1.09$^{e}$$\\pm$   0.10  &     1.31$^{d}$$\\pm$   0.06  &     1.71$^{d}$$\\pm$   0.11  &     1.94$^{e}$$\\pm$   0.71  \\\\ $T_{e}$([OIII])/10$^{4}$K  &    1.87$\\pm$   0.12   &    1.31$\\pm$   0.06   &    1.12$\\pm$   0.09   &    1.32$\\pm$   0.06   &    1.71$\\pm$   0.10   &    1.29$\\pm$   0.06   \\\\ 12+log(S$^{+}$/H$^{+}$) &    5.88$\\pm$   0.13   &    5.83$\\pm$   0.07   &    6.02$\\pm$   0.05   &    5.30$\\pm$   0.04   &    5.36$\\pm$   0.05   &    5.94$\\pm$   0.04   \\\\ 12+log(S$^{++}$/H$^{+}$) &    5.86$\\pm$   0.13   &    5.83$\\pm$   0.28   &    6.33$\\pm$   0.14   &    6.33$\\pm$   0.12   &    5.47$\\pm$   0.10   &    5.56$\\pm$   0.43   \\\\ 12+log[(S$^{+}$ + S$^{++}$)/H$^{+}$] &    6.17$\\pm$   0.13   &    6.13$\\pm$   0.19   &    6.50$\\pm$   0.11   &    6.37$\\pm$   0.11   &    5.72$\\pm$   0.08   &    6.09$\\pm$   0.19   \\\\ ICF(S$^{+}$ + S$^{++}$) &    1.02   &    1.03   &    1.02   &    1.78   &    1.01   &    1.01   \\\\ 12+log (S/H) &    6.18$\\pm$   0.13   &    6.16$\\pm$   0.18   &    6.51$\\pm$   0.11   &    6.61$\\pm$   0.17   &    5.73$\\pm$   0.08   &    6.13$\\pm$   0.16   \\\\ 12+log(O$^{+}$/H$^{+}$) &    7.99$\\pm$   0.07   &    7.51$\\pm$   0.08   &    7.49$\\pm$   0.12   &    7.03$\\pm$   0.07   &    7.18$\\pm$   0.07   &    7.78$\\pm$   0.07   \\\\ 12+log(O$^{++}$/H$^{+}$) &    7.37$\\pm$   0.08   &    7.78$\\pm$   0.07   &    7.37$\\pm$   0.14   &    7.97$\\pm$   0.07   &    7.56$\\pm$   0.08   &    7.90$\\pm$   0.07   \\\\ 12+log (O/H) &    8.08$\\pm$   0.07   &    7.96$\\pm$   0.08   &    7.74$\\pm$   0.13   &    8.01$\\pm$   0.07   &    7.71$\\pm$   0.08   &    8.15$\\pm$   0.07   \\\\  \\hline\\hline  Object &  Cam0357+3915 &  CTS1006 &  Tol2138+397 &  Tol2146+391 &  Tol2240+384 &  UM151\\footnote{[OII]$\\lambda$3727~\\AA~ line could not be measured for UM151} \\\\ \\hline $n_{e}$([SII]) &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$ &$\\leq$ 100.$^{b}$ &132$\\pm$80 &$\\leq$ 100.$^{b}$      \\\\ $T_{e}$([SIII])/10$^{4}$K  &     1.52$^{d}$$\\pm$   0.08  &     1.42$^{e}$$\\pm$   0.12  &     1.99$^{d}$$\\pm$   0.14  &     1.68$^{d}$$\\pm$   0.10  &     1.60$^{d}$$\\pm$   0.09  &1.88$^{d}$$\\pm$   0.12  \\\\ $T_{e}$([OIII])/10$^{4}$K  &    1.52$\\pm$   0.08   &    1.31$\\pm$   0.06   &    1.97$\\pm$   0.13   &    1.68$\\pm$   0.10   &    1.60$\\pm$   0.09   &    1.86$\\pm$   0.12   \\\\ 12+log(S$^{+}$/H$^{+}$) &    5.07$\\pm$   0.07   &    5.70$\\pm$   0.07   &    5.20$\\pm$   0.04   &    4.93$\\pm$   0.05   &    5.19$\\pm$   0.05   &    5.77$\\pm$   0.06   \\\\ 12+log(S$^{++}$/H$^{+}$) &    5.86$\\pm$   0.11   &    6.03$\\pm$   0.09   &    5.82$\\pm$   0.12   &    6.20$\\pm$   0.08   &    5.68$\\pm$   0.20   &    6.13$\\pm$  0.13   \\\\ 12+log[(S$^{+}$ + S$^{++}$)/H$^{+}$] &    5.93$\\pm$   0.10   &    6.20$\\pm$   0.09   &    5.91$\\pm$   0.10   &    6.22$\\pm$   0.07   &    5.80$\\pm$   0.17   & 6.29$\\pm$0.11   \\\\ ICF(S$^{+}$ + S$^{++}$) &    1.14   &    1.06  &    1.12   &    1.10   &    1.14   & 1.04    \\\\ 12+log (S/H) &    5.98$\\pm$   0.13   &    6.22$\\pm$   0.09   &    5.96$\\pm$   0.12   &    6.26$\\pm$   0.09   &    5.86$\\pm$   0.20   &6.31$\\pm$0.12    \\\\ 12+log(O$^{+}$/H$^{+}$) &6.33$\\pm$0.14   &    7.54$\\pm$   0.15   &    6.69$\\pm$   0.07   &    5.88$\\pm$   0.10   &    6.92$\\pm$   0.08   & ---   \\\\ 12+log(O$^{++}$/H$^{+}$) &7.78$\\pm$   0.08   &7.82$\\pm$   0.07   &    7.31$\\pm$   0.08   &    7.61$\\pm$   0.08   &    7.70$\\pm$   0.08   & 7.05 $\\pm$ 0.08  \\\\ 12+log (O/H) &7.79$\\pm$   0.08   &    8.00$\\pm$   0.10   &    7.41$\\pm$   0.08   &    7.62$\\pm$   0.08   &    7.77$\\pm$   0.08   & ---   \\\\  \\hline  \\end{tabular} \\end{minipage} \\end{table*}  \\subsection{Physical properties and chemical abundances} \\label{total_abund}  The physical properties and chemical abundances of the ionized gas were calculated for these galaxies following the 5-level atom FIVEL program (Shaw \\& Dufour 1994) available in the task IONIC of the STSDAS package. The final quoted errors in the derived quantities were calculated by error propagation including errors in flux measurements, atmospheric corrections, and temperatures. For the [SIII] lines we adopted the most recent atomic coefficients (Tayal \\& Gupta 1999).  Electron densities were obtained from the [SII]$\\lambda$6717/$\\lambda$6731~\\AA~line ratio. We could derive the electron temperature values of T$_{e}$[SIII], T$_{e}$[OIII], T$_{e}$[OII], and T$_{e}$[SII] by combining the data from our blue (Kehrig et al. 2004) and red spectra. Using the [OIII]$\\lambda$4363\\AA/$\\lambda\\lambda$4959,5007~\\AA~line ratio, we derived the T$_{e}$[OIII] for 21 galaxies of the sample. The T$_{e}$[SIII] was calculated from the [SIII]$\\lambda$6312/$\\lambda\\lambda$9069,9532~\\AA~line ratio for 14 galaxies with a measurement of the [SIII]$\\lambda$6312~\\AA~ line. For the 8 galaxies without any measurement of the [SIII]$\\lambda$6312~\\AA~line and with T$_{e}$[OIII], a theoretical relation between [OIII] and [SIII] electron temperatures (PMD05) was used:  \\[T_{e}[SIII] = 10,500T_{e}[OIII]-800\\] In total we have 22 galaxies with a measurement of T$_{e}$[SIII]. In order to derive T$_{e}$[SIII] temperature and S$^{++}$ abundances, whenever possible we used the [SIII]$\\lambda$9069~\\AA~fluxes, given that the flux of [SIII]$\\lambda$9532~\\AA~seems to be affected by partial blending of the P8 line. However, when the [SIII]$\\lambda$9069~\\AA~line falls inside the absorption band head due to the redshift of the galaxy, the flux of the [SIII]$\\lambda$9532~\\AA~line was used if this line was not close to the 1$\\mu$m limit, where the flux calibration becomes highly uncertain.  Regarding [SII] temperatures, for those objects without the [SII] auroral line at $\\lambda$4068~\\AA~we took the approximation T$_{e}$[SII] $\\approx$ T$_{e}$[OII] as valid. We could derive T$_{e}$[OII] using the [OII]$\\lambda$ $\\lambda$3727/7325~\\AA~line ratio for 16 objects of the sample. For the rest of the objects not presenting any auroral line in the low excitation zone, we used the model-predicted relations between T$_{e}$[OII] and T$_{e}$[OIII] found in PMD03, which explicitly take the dependence of T$_{e}$[OII] on electron density into account. In most cases the agreement between our line-intensity measurements in the blue and in the red spectra is good; we thus have adopted the values with the smaller observational errors to derive line temperatures. Otherwise, for T$_{e}$[SIII] and T$_{e}$[SII], we used the line intensities corresponding to the red spectra.   The relationship between both temperatures, T$_{e}$[OIII] and T$_{e}$[SIII], is shown in Fig.~\\ref{graf_temp} for all our galaxies with electron temperature and the sample of HII galaxies and HII regions compiled in PMD03, together with photoionization model relations. We note that there are two behaviors. While most HII regions show lower T$_{e}$[SIII] values than the ones provided by the photoionization models relations, many HII galaxies present higher T$_{e}$[SIII] values than the ones predicted by the models.  The same trend can be noticed for the sample of metal-poor emission-line galaxies in I05 (their Fig.4), for a range of T$_{e}$[OIII] from 1x10$^{4}$K to 2.0x10$^{4}$K. This fact suggests that HII regions and HII galaxies probably present different spatial temperature structures.   In order to compute the total sulfur abundances, we need to evaluate the corresponding ICF. A detailed study of the ICF scheme for sulfur is described in PM06. According to this work, for the objects with log ([SIII]/[SII]) $\\geq$ 0.4, we made use of the formula of Barker with $\\alpha$ =2.5 (Barker 1980). For the rest of the objects in the sample, we used the predictions of photoionization models for CoStar atmospheres (see Fig.3 in PM06). These predictions indicate rather low values of the ICF, independent of the ionizing effective temperature of the models.   With regard to the oxygen ICF, a small fraction of O/H is expected to be in the form of O$^{3+}$ ion in the high-excitation HII regions when the HeII$\\lambda$4686 emission line is detected. We have a measurement of the HeII$\\lambda$4686 emission line in 6 galaxies of our sample. According to the photoionization models from Stasi\\'nska \\& Izotov (2003), the O$^{3+}$/O can be on the order of 1$\\%$ only in the highest-excitation HII regions [O$^{+}$/(O$^{+}$ + O$^{2+}$) $\\leq$ 0.1]; therefore, taking our abundance results into account, we assumed that this correction is negligible in our sample.  Physical conditions, chemical abundances, and ICFs of sulfur for the galaxies with a measurement of the T$_{e}$[SIII] are quoted in Table~\\ref{6gal}. From this Table we can see that there are six galaxies with 12+log (O/H) varying between 7.4 and 7.8. These objects are among the galaxies with very low metallicity.  For the objects without T$_{e}$[SIII], we used the strong line calibration of the total S/H abundance as a function of S$_{23}$, presented by PM06, to derive the total S/H abundance.  \\begin{figure} \\centering \\includegraphics[width=7.0cm,clip]{4488fig7.ps} \\caption{\\footnotesize The distribution of total sulfur abundance for our sample of galaxies. The dashed and empty histograms show the number of galaxies with S/H derived from $T_{e}$[SIII] (see Table~\\ref{6gal}) and obtained from the S$_{23}$ calibration, respectively (see the text for details).}   \\label{hist} \\end{figure}   \\begin{figure} \\centering \\includegraphics[bb=21 260 570 591,width=\\columnwidth,clip]{4488fig8.ps} \\caption{\\footnotesize The observed sulfur-to-oxygen abundance ratio for the subset of galaxies of the sample with $T_{e}$[OIII] plotted as a function of the oxygen abundance. The solar value is shown. The dashed lines are +/- (1$\\sigma$) of the mean as shown by the continuous line.} \\label{S_H_O_H} \\end{figure}   The only galaxy of our sample for which we can compare the S/H abundance derived in this work with previous S/H abundance determinations in the literature is IIZw40. This galaxy has been observed by G89 and PMD03. In Table~\\ref{6gal} we present the results for IIZw40 obtained by the three works. In order to minimize possible reddening corrections effects in the abundance calculation for this galaxy, we referred the flux of each sulfur line we measured to a nearby hydrogen line. In the case of 12+log (S$^{+}$/H$^{+}$) ionic abundance the three values are close to each other; the value of the 12+log (S$^{++}$/H$^{+}$) ionic abundance, derived in this work is higher than the previous ones by up to some 0.2 dex. We believe that this fact could be the result of our systematic absorption correction procedure.   Figure \\ref{hist} shows the distribution of sulfur abundance derived for our sample of HII galaxies. The empty and dashed histograms represent the distribution of S/H derived with S$_{23}$ for all the objects and obtained from T$_{e}$[SIII], respectively. Most of the galaxies present total S/H abundance values that are between 1/20 solar to solar\\footnote{12+log (S/H)$_{\\odot}$ = 7.19 $\\pm$ 0.04 (Lodders 2003)}. This is an expected behavior since our sample is composed mainly of low-luminosity galaxies. Besides, we note that the dashed histogram peak corresponds to total S/H abundance value lower than the S/H maximum of the empty histogram. It suggests that, in order to know the overall metallicity distribution of a sample of galaxies, it would be worth making use of an efficient empirical abundance indicator. Hoyos $\\&$ D{\\'{\\i}}az (2006) found a similar result by studying the O/H abundance for a sample of HII galaxies.   The abundances obtained in this work allow us to study the dependence of S/H as a function of O/H in low metallicity environments.  In Fig.~\\ref{S_H_O_H} we show the relationship between the S/O abundance ratio and total O/H abundance for the subset of galaxies with T$_{e}$[OIII] and T$_{e}$[SIII]. The value of the sigma weighted mean for log(S/O) is -1.68$\\pm$0.20 dex.  The galaxy with nearly solar metallicity (UM307) is classified as an SABd from HYPERLEDA database\\footnote{leda.univ-lyon1.fr (Paturel 2003)}.  Evaluating the contribution of all observational errors to the derivation of these abundances, we can conclude at this level of uncertainty  that there is no statistical evidence of any systematic variation of S/O with O/H for this range of abundances. Therefore, our results agree with a constant S/O ratio and lower (1$\\sigma$) than the solar ratio for this type of emission-line object. This result indicates that sulfur and oxygen appear to be produced by the same massive stars, as expected by current nucleosynthesis prescriptions (see Pagel 1997 and references therein). In recent works, I05 and PM06 indicate that HII galaxy data are consistent with a constant S/O ratio, but somewhat lower than the current solar ratio. Regarding disk HII regions, the dispersion in S/O appears much larger and the assumptions of a constant S/O is questionable there. These results suggest that the assumption that the S/O ratio is constant at all abundances remains controversial (e.g Bresolin et al. 2004) and should be explored further, particularly at the not very well-known metallicity ends: extremely metal deficient HII galaxies (i.e. very low O/H) and HII regions in the inner disk of galaxies (i.e. metal rich central parts with highest O/H)."
    },
    "0606/astro-ph0606060.txt": {
        "abstract": "{ We present the type-1 active galactic nuclei (AGN) sample extracted from the VIMOS VLT Deep Survey first observations of 21000 spectra in 1.75 deg$^2$. This sample, which is purely magnitude limited, free of morphological or color selection biases, contains 130 broad line AGN (BLAGN) spectra with redshift up to 5. Our data are divided into a wide ($I_{AB} \\le 22.5$) and a deep ($I_{AB} \\le 24$) subsample containing 56 and 74 objects respectively. Because of its depth and selection criteria, this sample is uniquely suited to study the population of faint type-1 AGN. Our measured surface density ($\\sim 472 \\pm 48$ BLAGN per square degree with $I_{AB} \\le 24$) is significantly higher than that of any other optically selected sample of BLAGN with spectroscopic confirmation. By applying a morphological and color analysis to our AGN sample we find that: (1) $\\sim 23 \\%$ of the AGN brighter than $I_{AB}=22.5$ are classified as extended; this percentage increases to $\\sim$ 42 \\% for those with $z < 1.6$; (2) a non-negligible fraction of our BLAGN are lying close to the color space area occupied by stars in $u^*-g'$ versus $g'-r'$ color-color diagram. This leads us to the conclusion that classical optical ultraviolet preselection technique, if employed at such deep magnitudes ($I_{AB}=22.5$) in conjuction with a preselection of point-like sources, can miss miss up to $\\sim 35\\%$ of the AGN population. Finally, we present a composite spectrum of our sample of objects. While the continuum shape is very similar to that of the SDSS composite at short wavelengths, it is much redder than it at $\\lambda \\ge 3000$ \\AA . We interpret this as due to significant contamination from emission of the host galaxies, as expected from the faint absolute magnitudes sampled  by our survey. ",
        "introduction": "Recent surveys such as 2QZ and SDSS produce quasi-stellar object (QSO) spectra by the thousands \\citep{Croom2004,Richards2002}, demonstrating  the high efficiency of optical color selection techniques used since the pioneering work by \\citet{Sandage1965}. Spectroscopic targets are preselected based on their location in multidimensional color space (far from main sequence stars) and their morphology (point-like appearance, to avoid galaxies). The drawback  of such prerequisites is that some subsets of the underlying QSO population may be under-represented in these surveys, thereby possibly biasing our current understanding of this population of objects.  The preselection of non-resolved objects  prevents the selection of faint active nuclei in relatively bright galaxies and introduces incompleteness at low redshift. Moreover, the selection of unresolved objects is highly dependent on the quality of the imaging data. The ultraviolet (UV) excess selection is efficient only up to $z\\sim2.3$, and the evolving track of standard QSO crosses over main sequence stars near $z \\sim 2.5 \\pm 0.5$ in most optical broad band color-color planes. While it is difficult to identify AGN candidates in this redshift range, it is also at this epoch that a maximum in the QSO space density seems to be observed \\citep{Wall2005}, although there is now clear indication, at least for the X-ray selected AGN, that this maximum is dependent on the intrinsic luminosity \\citep{Hasinger2005}. % % % It is therefore essential to obtain unbiased samples at these high redshifts to trace and understand the accretion history onto supermassive black holes, which is believed to be the origin of the AGN phenomena.  Other criteria have been used to select AGN samples such as spectral energy distribution of objects in slitless spectroscopic plates \\citep{Hewett1995,Wisotzki2000}, variability and low proper motion \\citep{Brunzendorf2002} or radio emission \\citep{Jackson2002,Ivezic2002}, but they are always biased toward some subsets of the population. X-ray surveys are an efficient alternative to traditional optically selected samples and, indeed, they have allowed to derive, with good statistics, luminosity function and evolution of various classes of X-ray emitting AGN (e.g. absorbed and unabsorbed AGN in the X-ray band, soft and hard X-ray selected, see \\citealp{Miyaji2000,Ueda2003,Hasinger2005}). However, the comparison of the results obtained for X-ray and optically selected AGN is not straightforward both because of the different composition of the samples (X-ray samples contain large fraction of obscured AGN, without broad emission lines in their optical spectra) and of the somewhat different definition of AGN which is usually adopted in the X-ray surveys (e.g. any source with an X-ray luminosity and/or an X-ray to optical luminosity ratio above given thresholds).  The identification of faint type I AGN in a complete magnitude limited optical sample has been attempted in the Canada France Redshift Survey (CFRS) with 6 QSO identified among 943 spectra down to a limiting magnitude of $I_{AB} = 22.5$ \\citep{Schade1996}.  To probe fainter magnitude the COMBO-17 project identifies AGN from their spectral energy distribution (SED)  based on photometry in a set of 17 filters. They selected 192 objects over 0.78 square degree down to a magnitude $R = 24$ with $z >1.2$ \\citep{Wolf2003}. Although the technique developed by this team leads to a good redshift accuracy ($\\Delta z\\sim0.03$), it can be affected by some 'catastrophic errors' in region of the color space where various identifications are possible or in the case of objects with SED deviating from the reference set of templates.  In the present paper we present a spectroscopic catalog of broad emission line objects obtained from the purely flux limited spectroscopic sample of the VIMOS VLT Deep Survey (VVDS). This sample, free of any color or morphological biases, can be considered as a minimally biased sample to explore the population of optically non-obscured QSO at the faint end of the luminosity function. Our completeness is essentially set by the S/N ratio at which a broad emission line can be reliably identified in our spectra.  The data are divided into a wide ($I_{AB} \\le 22.5$) and a deep ($I_{AB} \\le 24$) subsample containing 56 and 74 objects respectively. The limiting magnitude of our deep sample is well beyond the limit of other optical spectroscopic surveys of AGN.  In the following section we describe briefly the VVDS imaging and spectroscopic surveys. In Section \\ref{Catalog} and \\ref{SF} we describe the process of AGN identification in the VVDS and the corresponding selection function. In Sections \\ref{Counts}, \\ref{Zdistrib} and \\ref{Morphcol} we present the surface density of BLAGN, the redshift distribution and the morphological and color properties of our sample. The composite spectrum of our sample of objects is presented and discussed in section \\ref{sec:composite}.  %==================================================================== %   Observations % %==================================================================== ",
        "conclusions": "%\\subsection{selection function} This paper describes a complete sample of 130 BLAGN selected from the VVDS first epoch spectroscopic catalog down to $I_{AB} = 24$. It is the first spectroscopic BLAGN catalog of this size at such faint magnitudes, purely magnitude limited and free of preselection biases.  %\\subsection{counts} The VVDS deep and wide complete samples contain 74 and 56 BLAGN respectively in this first release. A total of about 250 to 300 BLAGN are expected when the VVDS will be completed. We measure cumulative surface densities of $ 472 \\pm 48$ BLAGN per sq. deg. with $I_{AB} \\le 24$.  %\\subsection{z distribution} The redshift distribution ranges from 0 to 5 while the mean redshift in both the WIDE and DEEP samples is 1.8. So, by pushing our magnitude limit from $I_{AB} = 22.5$ to $I_{AB} = 24$ the result is not to increase the mean redshift of the sample, but rather to explore the faint end of the luminosity function at all redshifts.  %\\subsection{Morphological and color analysis} By comparing the $u-g$, $g-r$ color distribution of our AGN population with $z < 2.3$ and $I_{AB} < 22.5$ to the complete VVDS photometric catalog of stellar-like objects,  we find that $\\sim 35\\%$ of the AGN present in our sample would be missed by the usual UV excess and morphological criteria used for the preselection of optical QSO candidates in this redshift range. Most of the extended VVDS BLAGN are below the redshift $z=1.6$, a redshift range where 42\\% of VVDS BLAGN are extended.  The VVDS BLAGN have redder colors at the faint end of the luminosity function. Although we cannot exclude that the intrinsic spectrum of BLAGN is redder at faint luminosity or that the redder color it is due to the presence of dust, there are evidences that this effect is due to contamination by the continuum of the host galaxy at faint magnitudes. Indeed, 35\\% of the VVDS BLAGN have magnitudes fainter than $M_B=-23$.  %\\subsection{spectra analysis} This contamination is also seen in the composite spectrum obtained by co-adding the individual spectra in their rest frame. A comparison of the VVDS and SDSS composite QSO spectra shows that  the VVDS continuum is significantly redder than the SDSS one, especially at long wavelengths.  %\\subsection{prospects} In the context of the study of the VVDS luminosity function and its evolution \\citep{Bongiorno2006}, our BLAGN sample has two interesting properties. First, it is free of biases in the redshift range $2<z<3$, which will help to shed light on the change of QSO evolution suspected to happen in this redshift range. Second, since no bias against low luminosity AGN is present in this sample, this will make easier the comparison of the evolution of low and high luminosity AGN, allowing a better understanding of the QSO-Seyfert connection."
    },
    "0606/astro-ph0606335_arXiv.txt": {
        "abstract": "We have searched the {\\em OGLE-II} archive for candidate counterparts of X-ray sources detected in two low-extinction windows included in our Galactic bulge {\\em Chandra/HST} survey. We find that a significant number---i.e.~in excess of the expected level of random associations---can be matched with probable M-giants.  Their X-ray properties can be understood if these sources are symbiotic binaries where the X-rays are typically, either directly or indirectly, the result of a white dwarf accreting from the wind of a cool giant. Optical and near-infrared properties of selected sources are consistent with a symbiotic nature, although none of the spectra collected for 8 out of 13 candidate counterparts show the high-ionization nebular emission lines observed for many symbiotics. The hard X-ray emission for several sources (power-law photon indices $-1.5\\lesssim\\Gamma\\lesssim1.5$) suggests our sample includes systems similar to the symbiotics recently detected with {\\em INTEGRAL} and {\\em Swift}. ",
        "introduction": "The main goal of the \\Chandra~Multiwavelength Plane Survey ({\\em ChaMPlane}\\footnote{http://hea-www.harvard.edu/ChaMPlane}) is to identify and study the populations of low-luminosity ($L_X\\lesssim10^{33}$\\,erg\\,s$^{-1}$) accretion-powered binaries in the Galaxy \\citep{grinhongea05}. {\\em ChaMPlane} includes a deep survey of the Galactic bulge consisting of (near-)simul\\-taneous \\Chandra~and {\\em Hubble Space Telescope} (\\HST) pointings of three low-extinction regions. With minimal obscuration by dust, for the bulge, and positions progressively closer to the Galactic Center (GC), these ``windows'' are chosen to trace the X-ray point-source population towards the GC. The bulge population of interacting binaries is also of interest for comparison with those in globular clusters, where stellar encounters have enhanced the numbers of close binaries, and with X-ray populations in other galaxies. Here we report initial results of observations of Baade's Window (\\BW, ($l$,$b$)=(1\\fdg06,$-$3\\fdg83)) and Stanek's Window \\citep[\\SW, ($l$,$b$)=(0\\fdg25,$-$2\\fdg15);][]{stan98} where we have found a significant number of sources with candidate M-giant counterparts. In $\\S$\\,\\ref{sec_obs} we describe the optical identification, and X-ray and optical/near-infrared spectra.  As X-ray emission from single M-giants is rarely detected \\citep{huenea98}, we have investigated whether these sources could be symbiotic binaries in which a cool giant typically transfers mass to a white dwarf via a wind ($\\S$\\,\\ref{sec_dis}). More results, including analysis of our ``Limiting Window'' at ($l$,$b$)=(0\\fdg10,$-$1\\fdg43) will be presented in future papers. ",
        "conclusions": "\\label{sec_dis}  We explore possible explanations for the X-ray emission from these giants. Single M-giants are rarely detected in X-rays but \\cite{huenea04} point out a few candidates that are matched securely to X-ray sources ($L_{0.5-8 {\\rm keV}}=(0.1-1.1)\\times10^{31}$ erg s$^{-1}$ for a 10$^{7}$ K coronal plasma, but $L$ is $\\sim$15 times larger for a $\\Gamma=0$ power law, for $n_H=0$ cm$^{-2}$). Their X-rays are yet unexplained, but the absence of optical emission lines is used to argue against a symbiotic nature. In RS\\,CVn binaries, activity of a (sub)giant is increased by tidal coupling of the stellar rotation to the orbit, but no RS\\,CVns with M-giants are known. Tidal forces in a detached binary wide enough to fit an M-giant are too weak to spin up a main-sequence companion.  Moreover, neither single stars nor active binaries are in general as hard as some of our sources.  Of the $\\sim$200 symbiotic binaries currently known, $\\sim$80\\% contain an M3--6 red giant \\citep{muerschm99}. Mass transfer occurs mostly via a wind as opposed to Roche-lobe overflow.  \\cite{muerea97} show in a ROSAT study that symbiotics are typically soft; nuclear burning on the white dwarf produces very soft emission ($\\lesssim0.4$ keV), while somewhat harder X-rays are associated with hot ($kT\\approx1$ keV) shocked gas in the colliding winds of the giant and hot star.  Note that our sample contains many hard sources ($\\gtrsim50$\\%, $\\S$\\ref{sec_xspec}) that may be similar to the hardest sources in M{\\\"u}rset et al.  Emission above a few keV is so far poorly studied for symbiotics; possible explanations include accretion \\citep[CH\\,Cyg;][]{ezukea98} or shocked gas \\citep[Z\\,And;][]{sokoea06}. Recently, symbiotics have been detected also in very hard X-rays by {\\em INTEGRAL} and {\\em Swift} \\citep[e.g.\\,RT\\,Cru,][]{tuelea05}. X-ray luminosities of symbiotics range from 10$^{29}$ to 10$^{34}$ erg s$^{-1}$, up to $\\sim$10$^{37}$ erg s$^{-1}$ for the neutron-star accretor GX\\,1+4.  We estimate the upper limit on the temperature of the shock-heated gas in the colliding-wind model by requiring that the thermal energy per gas particle does not exceed the kinetic energy. A wind velocity $v_w=230$ km s$^{-1}$ gives $kT\\lesssim0.2$ keV, lower than implied for any of our sources by quantile analysis. Thus, if the He\\,I 1.083\\,$\\mu$ lines in BW3 and SW7 (two of the hardest sources) trace the conditions in the shocked winds, this model is an unlikely explanation. To reach $kT\\approx1.5$ keV as in SW3, requires $v_w\\approx660$ km s$^{-1}$. A mass-loss rate $\\dot{M}=10^{-7}$ $M_\\odot$ yr$^{-1}$ gives a total kinetic luminosity 1/2\\,$\\dot{M} v_w^2\\approx2.7 \\times 10^{34}$ erg s$^{-1}$, sufficient to power $L_{0.5-8 {\\rm keV}}$=7.6$\\times$($d$/8kpc)$^2\\times$10$^{32}$ erg s$^{-1}$ for SW3. Looking for such fast outflows would be one test whether shocked winds are indeed a plausible explanation for our soft sources.  To check whether accretion is a viable source of X-rays, we use Eq.~3 from \\cite{liviwarn84} for the total accretion luminosity $L_{acc}$. We choose masses for the giant and white dwarf of 1 and 0.6 $M_\\odot$, a white-dwarf radius of 10$^9$ cm, $\\dot{M}=10^{-7}$ $M_\\odot$ yr$^{-1}$, and $v_w=230$ km s$^{-1}$; the orbital speed of the white dwarf is assumed to be large compared to the local sound speed. The range of known orbital periods $P_b$ for symbiotics (200--5700 d) implies 6$\\times$10$^{34}\\gtrsim L_{acc}\\gtrsim 7\\times10^{33}$ erg s$^{-1}$, scaling linearly with $\\dot{M}$. Thus it is likely that accretion onto a white dwarf can power the X-rays: $L_{0.5-8 {\\rm keV}}$ =(2.0--110)$\\times$($d$/8 kpc)$^2$ 10$^{31}$ erg s$^{-1}$ (Table~\\ref{tab_prop}). Accretion onto a neutron star is less probable: for a 1.4~$M_{\\odot}$, $R=15$ km companion, $L_{acc}\\gtrsim10^{35}$ erg s$^{-1}$ for the assumed $\\dot{M}$.  Assuming we observe the white dwarf through the giant's wind, we estimate $n_H$ for the simplified case of a circular orbit and a spherically symmetric neutral-hydrogen wind of constant velocity: $n_H\\approx\\dot{M}$/(16$m_H$$v_w$$a$)~$(\\cos^2{i}+\\sin^2{\\phi}\\sin^2{i})^{-1/2}$, with $a$ the semi-major axis, $i$ the orbital inclination, and $\\phi$ the orbital phase. For the parameters and range in $P_b$ as above, 9$\\times$10$^{19}\\lesssim n_H \\lesssim$ 2$\\times$10$^{21}$ cm$^{-2}$ for $i=60^{\\circ}$, below or similar to the Galactic $n_H$ for \\BW~and \\SW. \\cite{seaqtayl90} have found a mass-loss rate $\\dot{M} \\approx 10^{-8}$--10$^{-6}$ $M_\\odot$ yr$^{-1}$ for red-giant symbiotics which allows for an enhanced $n_H$, but values as high as for SW3 also require a larger $i$. Quantile analysis suggests 6.4 keV Fe\\,K emission in our hardest sources, which may be another manifestation of the giant's cool wind via interaction with hard X-rays; Fe\\,K lines are observed in e.g.~CH\\,Cyg \\citep{ezukea98} and RT\\,Cru (J.~Sokoloski 2006, private communication).  How do other properties of our sources compare to those of M-giants and symbiotics?  Variability and mass loss, both associated with stellar pulsations, are typical properties of single late M-giants. \\cite{glasschu02} found that many M5 and all M6 and later giants in Baade's Window are variables. Of the 10 matches included in \\cite{woznea02}, the matches to BW1, SW7, 9, and possibly SW3, 4, 5 and 8 show semi-regular variability as commonly observed in late M-giants (amplitudes $\\lesssim2.5$ mag, time scales 20--200 d). \\cite{glasschu02} note that BW3 however shows more regular variability. The period they find in {\\em MACHO} data (135.5d) is the same as what we find in {\\em OGLE-II} data (135.2$\\pm$1.5d). If this is due to ellipsoidal variations then $P_{b}\\approx270$ d. Variability of SW2, 6 and possibly 8 is irregular with $\\Delta I\\lesssim0.1$ in 4 seasons.  From mid-infrared photometry \\citet{alarea01} derive $\\dot{M}\\approx4\\times10^{-7}$ and $4\\times10^{-8}$ $M_{\\odot}$ yr$^{-1}$ for the matches to BW1 and 4, respectively.  Outflow velocities for BW3 and SW7 are high for single red giants ($v_w\\approx10$--30 km s$^{-1}$) but a similar velocity has been derived from He\\,I~1.083$\\mu$ for GX\\,1+4 \\citep{chakea98}.  Optical spectra of our matches lack strong high-ionization nebular emission lines. The presence of such lines is one of the original defining properties of symbiotics.  Examples of ``weakly symbiotic'' (i.e.~without strong nebular emission) systems that are hard in X-rays are 4U\\,1700+24 and 4U\\,1954+319 \\citep[both suspected neutron-star systems that look like ``normal'' M-giants in the optical,][]{maseea06} and the aforementioned {\\em INTEGRAL} and {\\em Swift} sources, for which H$\\alpha$ and H$\\beta$ are the strongest optical emission lines. If these are similar to our sources, the large fraction of such systems in our sample then suggests these properties are characteristic for a bulge sub-population and/or not identified previously due to sensitivity limits or limited astrometric precision. Further study of the \\BW~and \\SW~sources but also those of \\cite{huenea04} is therefore of interest for estimates of the total number of Galactic symbiotics and their progeny, which may include SNIa systems. The evolution of symbiotics depends on the orbital period. Long-period systems become wide double white dwarfs.  For $P_b=270$ d as for BW3, and 1 and 0.6 $M_\\odot$ stars, the Roche-lobe radius is likely $\\sim94R_{\\odot}$. In the models by \\cite{lejescha01}, a 1$M_\\odot$ star reaches $R\\approx110R_{\\odot}$ at the tip of the red-giant branch and expands even further as an asymptotic giant. If BW3 is a binary, it may start unstable mass transfer leaving a close double white dwarf, which could evolve to a SN\\,Ia.  We use the results by \\citet{laycea05} to constrain the number of symbiotics among the initial sample of 1453 still unclassified faint, hard X-ray sources in the 10\\arcmin~$\\times$ 10\\arcmin~region around the GC \\citep{munoea03}. For $A_{K,GC}\\approx3.4$ and $M_K\\lesssim-4$, $K\\lesssim13.9$ for M-giants at the GC. Laycock et al.~find that $\\lesssim7$\\% (90\\% confidence) of the 110 hard sources within 1\\arcmin~from the GC can have counterparts with $K\\leq14$, while this is true for $\\leq4$\\% of the 1343 hard sources at larger offsets. These likely include high-mass X-ray binaries with B0 V primaries but also symbiotics."
    },
    "0606/astro-ph0606429_arXiv.txt": {
        "abstract": "It has been suggested ``that DM particles are strongly interacting composite macroscopically large objects ... made of well known light quarks (or ... antiquarks).\" In doing so it is argued that these compact composite objects (CCOs) are ``natural explanations of many observed data, such as [the] $511$ keV line from the bulge of our galaxy\" observed by INTEGRAL and the excess of diffuse gamma-rays in the 1-20\\,MeV band observed by COMPTEL \\cite{Lawson}. Here we argue that the atmospheres of positrons that surround CCOs composed of di-antiquark pairs in the favoured Colour-Flavour-Locked superconducting state are sufficiently dense as to stringently limit the penetration of interstellar electrons incident upon them, resulting in an extreme suppression of previously estimated rates of positronium, and hence the flux of 511~keV photons resulting from their decays, and also in the rate of direct electron-positron annihilations, which yield the MeV photons proposed to explain the 1-20~MeV excess. We also demonstrate that even if a fraction of positrons somehow penetrated to the surface of the CCOs, the extremely strong electric fields generated from the bulk antiquark matter would result in the destruction of positronium atoms long before they decay. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606759_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0} A distinctive signature of pre-decoupling (adiabatic) initial conditions is that the temperature autocorrelation has the first (Doppler) peak for $\\ell_{\\rm d}\\simeq 220$ \\cite{wmap1,wmap2,wmap3}. For the same set of initial conditions, the cross-correlation power spectrum between temperature and polarization will have the first (anticorrelation) peak for $\\ell_{\\rm c} \\simeq 3 \\,\\ell_{\\rm d}/4 < \\ell_{\\rm d} $ \\cite{wmap1}.  Defining as $k$ the comoving wave-number and as $c_{\\rm sb}$ the sound speed of the baryon-photon system, the large-scale contribution to the temperature autocorrelation oscillates as a $\\cos{(k c_{\\rm sb } \\tau_{\\rm dec} )}$, while the cross-correlation oscillates as $\\sin{( 2 k c_{\\rm sb } \\tau_{\\rm dec}) }$.  The first (compressional) peak of the temperature autocorrelation will then correspond to $ k c_{\\rm sb } \\tau_{\\rm dec}\\sim \\pi$ (i.e. $\\ell_{\\rm d} \\simeq 220$), while  the first peak of the cross-correlation will arise for $k c_{\\rm sb } \\tau_{\\rm dec}\\sim 3\\pi/4$ (i.e., approximately, $\\ell_{\\rm c} \\simeq 3\\, \\ell_{\\rm d}/4$). This result can be obtained analytically by working to first-order in the tight-coupling expansion \\cite{TC1,TC2,TC3}. In this framework, the key assumption is that the initial conditions of the Einstein-Boltzmann hierarchy are characterized, prior to matter-radiation equality, by a solution that is, at least predominantly, adiabatic. A corollary to the previous statements is that the peak in the cross-correlation occurs for typical wavelengths that are larger than the Hubble radius at the the time of photon decoupling.  In a previous paper \\cite{mg1} it has been shown that large-scale magnetic fields, present prior to  equality affect the CMB initial conditions. If only the adiabatic mode is present, the fully inhomogeneous magnetic fields contribute to the Sachs-Wolfe plateau as a subdominant non-adiabatic component that may be either correlated or uncorrelated  with the adiabatic mode. If the pre-decoupling  initial conditions are given by a mixture of adiabatic and non-adiabatic modes, the magnetized contribution may also mix with the isocurvature modes so that the total number of parameters defining the initial conditions increases. It is desirable to complement and extend the analysis of \\cite{mg1} by scrutinizing the impact of pre-decoupling magnetic fields on  intermediate scales. Within the tight coupling expansion, appropriately generalized to include the effects of fully inhomogeneous magnetic fields, the fluctuations in the brightness induced by the (adiabatic and non-adiabatic) scalar modes can be estimated.  The effects of magnetic fields on the scalar brightness perturbations are typically neglected. To give an example, let us consider how the impact of large-scale magnetic fields is usually computed, for instance, in the Faraday effect. This case is interesting, since, the rotation of the CMB polarization plane indeed presupposes that the magnetic fields do not affect the process of formation of polarization. The influence of large-scale magnetic fields on CMB polarization can  be separated, for illustration, into two physically distinct effects. Prior to decoupling, gravitating magnetic fields modify the initial conditions of CMB anisotropies and affect the evolution equations of the photon-baryon system. Therefore, the induced degree of polarization will be sensitive to the presence of magnetic fields whose role will be, at this stage, to modify the features of the adiabatic initial conditions. This is the first effect. Once the polarization is formed, the polarization plane of the CMB may be Faraday rotated \\cite{F1,F2}. From this second effect interesting bounds on the magnetic field intensity are usually derived.  This philosophy motivated various independent studies on the possible Faraday rotation of the CMB polarization degree \\cite{F3,F4,F5,F6}. In some cases the magnetic field has been assumed just uniform \\cite{F3}, but there are also calculations analyzing the situation where the magnetic field is fully inhomogeneous \\cite{F4,F5}.  In spite of the relevant technical differences, the common assumption of the mentioned investigations is that the magnetic field is negligible in the process of formation of CMB polarization. The latter assumption can be rephrased by saying that the initial conditions of the Einstein-Boltzmann hierarchy are set in such a way that the magnetic fields are absent. In the absence of large-scale magnetic fields the only  source of polarization arises to first-order in the tight coupling expansion. The magnitude of the obtained contribution is just a consequence of the adiabaticity of the solution. If large-scale magnetic fields are consistently introduced before matter-radiation equality, the usual adiabatic (as well as non-adiabatic) initial conditions of the Einstein-Boltzmann hierarchy are modified and, consequently, also the degree of polarization of the CMB is affected.  This example motivates a systematic analysis of the impact of magnetized initial conditions on the brightness perturbations of the radiation field. As anticipated above, this calculation will be conducted within the tight-coupling approximation which is known to be rather effective in obtaining the brightness perturbations for large and intermediate scales.  This analysis has never been done before, to the best of our knowledge. A related motivation is that, in the absence of magnetic fields, the tight coupling approximation is employed  in numerical Boltzmann codes at early times to avoid the integration of stiff Euler equations. The results of the present investigation allow to use the same strategy also in the presence of fully inhomogenous magnetic fields.  Several aspects of CMB physics are germane to large-scale magnetization (see \\cite{r1,r2} for two topical review articles). The astrophysical evidence for large-scale magnetization can be summarized as follows. Large-scale magnetic fields are an important element of the physics of the interstellar medium and have been measured, through various techniques in galaxies  \\cite{b1,b2} and  in rich clusters \\cite{l1,l2,l3,l4}. The determination of large-scale magnetic fields associated with loosely gravitationally bound systems (like superclusters) is still debatable \\cite{k1,k2,k3}, but, nonetheless, extremely intriguing. As far as galactic magnetic fields are concerned, there is the hope that in the future all sky survey of Faraday rotation (achievable through the Square Kilometer Array [SKA] \\cite{b2,l1})  will allow a Faraday \"tomography\" of the Milky Way addressing the dark corners (number of reversals, pitch angle) of the morphological structure of the closest large-scale magnetic field we observe through different surveys of pulsars such as the Parkes and the southern galactic plane surveys (see  \\cite{han,parkes,sgps} and references therein). There is the hope, in the present framework, that the present-day coherent field in spiral galaxies may keep the \"memory\" of the structure of the initial seed field (see \\cite{b1} for a review of this long-standing problem). The morphological features of the magnetic fields profiles in clusters are  less understood than in the case of galaxies. However, attempts have been made in constraining the radial profiles of magnetic fields in clusters \\cite{l2,l3}. Determinations of intracluster magnetic fields through statistical samples of rich clusters \\cite{l4} seem to suggest that (ordered?) intra-cluster magnetic fields may be as large as the $\\mu {\\rm G}$ (see, for instance, \\cite{magnetized} for a review on the features of our magnetized Universe).  Compressional amplification (taking place during the gravitational collapse of the protogalaxy) allows to connect the observed field to a protogalactic field, present prior to gravitational collapse, between  $0.1$ nG and $0.01$ nG. It is plausible, within the dynamo hypothesis, that the protogalactic fields could even be much smaller than the nG. In recent years, a lot of progress has been made both in the context of small- and large-scale dynamos i.e. the process by which the kinetic energy of the protogalactic plasma is transferred, by differential rotation, into magnetic energy. This progress \\cite{rev1} (see also \\cite{rev2} for an introduction to astrophysical dynamos) has been driven both by the higher resolution of numerical simulations and by the improvement in the understanding of the largest magnetized system that is close to us, i.e. the sun. In contrast  with the previous lore, it is now clear that the dynamo action demands a change in the topology of the flux lines. As a consequence, large-scale dynamos should also produce small-scale helical fields that quench (i.e. prematurely saturate) the $\\alpha$ effect \\cite{rev1} (see also \\cite{rev3}). The morphology and strength of magnetic fields in clusters may be related to the way the dynamo effect saturates.  The simplest way to understand why large-scale magnetic fields may be relevant for CMB physics is to think of a protogalactic magnetic field in the nG range. If this field is naively blue-shifted at the decoupling epoch, i.e. at a redshift $z_{\\rm dec} \\simeq 1100$ its strength could be as large as the mG, i.e. roughly six orders of magnitude larger. In the past ten years, indeed a lot of work has been done with the purpose of constraining large-scale magnetic fields using CMB physics. Historically the first type of configurations used in this type of exercise have been weakly uniform magnetic fields that would partially break the isotropy of the background geometry \\cite{r2}. It was than realized that this assumption induces a correlation of the $a_{\\ell -1,m}$ and $a_{\\ell +1,m}$ multipole coefficients arising in the angular power spectrum (see \\cite{unif1,unif2,unif3} and references therein). From this observation, uniform magnetic fields can be constrained.  Another (complementary) possibility is that magnetic fields are fully inhomogeneous. In this case the magnetic fields do not break the spatial isotropy of the background but they can affect virtually all observables relevant in the theory of CMB anisotropies. In particular, in a series of papers, Barrow and Subramanian \\cite{bs1,bs2,bs3} pointed out possible effects of magnetic fields on the induced vector CMB anisotropies as well as on the polarization power spectra induced by the same modes. Within a slightly different approach the vector modes (as well as the tensor modes) have been discussed in \\cite{m1}. More recently \\cite{lw1} a full numerical analysis of the vector and of the tensor modes induced by fully inhomogeneous magnetic fields has been presented.  In \\cite{rat1} the signatures of magnetic helicity have been studied always in the perspective of the vector modes. Finally, in a series of papers, the evolution of weakly inhomogeneous magnetic fields has been discussed within the covariant formalism and in the presence of a uniform magnetic background \\cite{ts2}.  In spite of the fact that vector modes are probably one of the characteristic signatures of large-scale magnetic fields, scalar modes are also induced \\cite{mg1,mg2}. This statement can be understood by noticing that the magnetic energy density, the magnetic anisotropic stress and the magnetic pressure contribute, respectively, to the Hamiltonian and momentum constraint and to the dynamical evolution of the curvature perturbations. Furthermore, within the magnetohydrodynamical approximation, the divergence of the Lorentz force contributes to the evolution equation of the divergence-full part of the baryon peculiar velocity.  What is missing, at the moment, is the analysis of the brightness perturbations induced by magnetized initial conditions. As anticipated, one of the purposes of the present paper will exactly be to bridge fill this gap and to allow an approximate evaluation of the brightness perturbations, also at intermediate scales, as a function of the spectral properties of fully inhomogeneous magnetic fields.  The present paper is then organized as follows. In Section 2 the initial condtions of the Einstein-Boltzmann hierarchy will be discussed for typical wavelengths larger than the Hubble radius prior to equality. In Section 3 the tight-coupling equations will be derived and discussed analytically. The numerical treatment of the system is presented in Section 4. Section 5 contains our concluding remarks.  \\renewcommand{\\theequation}{2.\\arabic{equation}} ",
        "conclusions": "\\setcounter{equation}{0} In the recent past, various attempts have been made to constrain large-scale magnetic fields from CMB physics. While interesting results have been obtained for the vector and tensor modes of the geometry, the scalar modes have been, comparatively, less studied. In this study  a systematic strategy to tackle this problem has been proposed and developed. The basic inspiration of the present approach can be summarized  by stating that the effects of fully inhomogeneous magnetic fields on CMB anisotropies is not univocal.  In the standard lore the only relevant parameters to define the effects of large-scale magnetic fields on CMB anisotropies are the ones related to the magnetic field itself. For instance, if the magnetic field is uniform the only parameter will be its intensity; if the magnetic field is fully inhomogeneous the only two parameters will be the spectral amplitude and the spectral slope. The point of view of the present investigation is that this way of thinking is  ambiguous and, to some extent,  not so productive. In fact, the spectral parameters of the magnetic field are certainly a necessary ingredient of the analysis. They are, however, not sufficient to conduct the calculation of the brightness perturbations. The information on the spectral parameters must be complemented by the specific pre-decoupling initial conditions. According to the results of this investigation, a sound procedure to be followed can be summarized as follows: \\begin{itemize} \\item{} include the magnetic fields in the system of scalar modes before matter radiation equality; \\item{} solve the system in the long wave-length approximation prior to decoupling; \\item{} use the obtained solutions for the numerical evaluation of the brightness perturbations, for instance, in the tight coupling approximation. \\end{itemize} The ambiguity mentioned above arises during the first step of the procedure we just outlined. In fact, the inclusion of magnetic fields prior to equality does not lead to a single solution but to different solutions. We will then have quasi-adiabatic and quasi-entropic modes where a magnetized contribution can be accommodated.  In some sense, the present proposal is more conservative than previous attempts: the signature of magnetic fields on CMB physics is not directly accessible but it is mediated by the (supposedly rich) parameter space of the pre-decoupling initial conditions of CMB anisotropies. This approach is indeed common when  mixtures of adiabatic and non-adiabatic modes are studied in the absence of magnetic fields. In the latter case the philosophy is to specify what kind of non-adiabatic mode we want to study and constrain. Here the situation is similar and we do hope, that, in the future, the considerations reported here will serve as an inspiration for analyzing the effects of fully inhomogeneous magnetic fields on the scalar CMB anisotropies.   This analysis reported in this paper allowed to extend the tight-coupling approximation to the case when large-scale magnetic fields participate to the dynamics of the plasma. The evolution equations, in the tight coupling approximation, have then been integrated numerically using, as normalization, the analytical estimate of the Sachs-Wolfe plateau modified by the presence of large-scale magnetic fields. The reported results allow not only to set constraints but also to include a magnetized contribution in the current strategies of parameter extraction.  From the combined analysis of the Sachs-Wolfe plateau and of the phases of Sakharov oscillations at intermediate scales, a fully inhomogeneous magnetic field uncorrelated with the adiabatic mode can be constrained to be smaller than $0.5$ nG for a slightly blue spectral slope $\\varepsilon =0.1$ and for typical comoving wave-number of $1\\,\\, {\\mathrm Mpc}^{-1}$. The conditions implied by the SW plateau ssem to be  necessary but not sufficient since magnetic fields compatible with the SW bound may alter Sakharov oscillations at smaller length-scales. In the most general situation, correlation angles may proliferate: there may be correlations between the magnetized contribution and the adiabatic or even non-adiabatic modes.  In this case the bounds on large-scale magnetic fields become much weaker and sizable magnetic fields may be allowed.  The present analysis also shows that magnetic fields can affect the process of formation of CMB polarization. For instance, the usual approach to the Faraday rotation of CMB polarization is to assume that the polarization itself is only due to the adiabatic mode while the rotation of the polarization plane is only due to the magnetic field. This picture has been complemented by the observation that the magnetic fields affect the pre-decoupling initial conditions and, consequently,  modify the evolution of the degree of polarization.  The tight coupling expansion is rather useful to discuss the integration of brightness perturbations in some simple cases. Moreover, it is an essential ingredient of Boltzmann codes since it is used, at early times, to avoid the integration of the (stiff) Euler equations. It will be interesting, for future studies, to analyze the space of magnetized initial conditions by solving directly the whole Boltzmann hierarchy.   \\newpage"
    },
    "0606/astro-ph0606045_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro} The primary science goal of the Deep Impact mission is the study of the physical and chemical properties of the material at and below the surface of a comet.  This was accomplished by directing a projectile at the nucleus of comet \\tempelfull. The 370\\,kg impactor separated from its fly-by mother spacecraft one day before encounter with the comet and hit the sunlit side of the nucleus with a relative velocity of 10.3\\,\\kms,  releasing a kinetic energy of 19\\,GJ at the impact site. The impact and the ejecta were observed by the fly-by spacecraft of the Deep Impact mission \\citep{AHearn2005} as well as by a large number of Earth-based telescopes and satellite-based observatories \\cite[for an overview, see][]{Meech2005a}.  \\tempelfull, the target comet of the Deep Impact mission, is a low-activity Jupiter-family comet with an orbital period of 5.5\\,years and a slow rotation period of $41.85\\pm0.10$\\,hrs measured before the 2005 apparition \\citep[see][ for a summary]{Meech2005b,Belton2005}. Previous estimates of water vapor production through OH measurements give a typical evaporation rate of $\\qwat\\sim10^{28}$~molecule~s$^{-1}$ near perihelion\\footnote{following the convention commonly used in the literature we write $\\qwat$ in units of \\pers\\ in the remainder of the paper} \\citep[][]{Lisse2005}. The abundance of minor gas species in the coma, such as C$_2$ and CN, is at the low end of the range derived for 'typical' comets \\citep[following the taxonomy by][]{AHearn1995}. The dust probed by observations at optical wavelengths indicates that \\tempelfull\\ is neither particularly dust-poor nor dust-rich.  The choice of \\tempelfull\\ as the target comet for the Deep Impact mission gave rise to a large number of observing campaigns, which have increased our knowledge of the physical properties of this comet, the chemical composition of the dust and volatiles in its coma, and its periodic variations in activity. The observing campaign for \\tempelfull\\ culminated during the weeks leading up to the Deep Impact event and the days following impact.  First results from the impact experiment have been reported by \\cite{AHearn2005}, and an overview on the results from the Earth-based observing campaign is given by \\cite{Meech2005a}. First results constraining the amount and composition of the dust in the ejecta are presented by \\cite{Sugita2005} and \\cite{Harker2005}, while the volatiles released by the impact have been studied by \\cite{Kueppers2005}, \\cite{Mumma2005}, and \\cite{Keller2005}. Most of the observers report that the activity triggered by the impact was relatively short lived, and that the comet (coma) returned to its pre-impact state $\\la 5$ days after the impact. % Observers conducting monitoring observations for \\tempelfull\\ over a longer period reported frequent natural outbursts of the comet \\citep[][and references therein]{Meech2005a}. {\\new With regard to the total brightness of the comet,} these were often much larger than the outburst triggered by the impact, but studied in less detail than the impact itself since most of the observations were conducted around the impact date, July 4, 2005.  Water is the most abundant species among the volatiles in cometary comae. Spectra of water vapor yield information on the amount of water evaporated from the nucleus, and provide reference levels for deriving the relative abundances of minor volatiles while also monitoring overall cometary activity. Water line observations proved useful also for estimating the total amount of water vapor released in response to the Deep Impact event.   Most of the water is rotationally cold throughout the cometary atmosphere \\citep{Crovisier1984,Weaver1984,Bockelee-Morvan1987,Xie1992,Bensch2004b}. The ground-state transition of ortho water is therefore particularly suitable for tracing the water production rate. Present-day heterodyne receivers provide the high velocity resolution ($<1$\\,\\kms) required to resolve cometary emission lines. Observations of low-lying water rotational transitions, however, have to be carried out from space-based platforms because the Earth's atmosphere is opaque at the line frequencies of these transitions in the submillimeter and far-infrared regime.  Three satellites with instruments capable of observing the \\gs\\ ground-state transition of ortho water were available during the 2005 apparition of \\tempelfull\\ and have participated in the Deep Impact observing campaign: the Microwave Instrument for the Rosetta orbiter, MIRO, % the Swedish-led satellite mission Odin {\\new \\citep{Nordh2003},} and the Submillimeter Wave Astronomy Satellite, SWAS {\\new \\citep{Melnick2000}.} In this paper we present the results from the SWAS observing campaign for comet \\tempelfull\\ and Deep Impact. Constrained only by the comet's visibility from Earth orbit and a Moon-avoidance angle, SWAS conducted daily monitoring observations of the comet covering approximately a 3-month period, from 5 June to 1 September 2005. The SWAS observations are presented in Sect.\\ \\ref{swas-obs}. These observations are used in Sect.\\ \\ref{water-prod} to study the variations in water vapor production and to constrain the amounts of water ice vaporized as a result of the Deep Impact event. A discussion of the results is given in Sect.\\ \\ref{discuss}. ",
        "conclusions": "\\label{discuss} No increase in the water rotational-line emission from the impact ejecta was detected by SWAS and we derive a $3$ sigma upper limit to the water vapor generated by the impact, $N_{\\rm tot}<6\\times10^{32}$ water molecules ($1.8\\times10^7$\\,kg). The SWAS non-detection is consistent with results from OH observations at UV wavelengths by \\cite{Kueppers2005} which indicate that only a small amount of water vapor ($\\sim 4.5\\times 10^6$\\,kg) was produced by the impact and suggest that the dust-to-water ice mass ratio of the ejected material exceeded unity \\citep{Keller2005}. However, care has to be taken in arriving at this conclusion.  Ice freshly excavated from the impact crater might sublimate rather slowly on exposure to sunlight {\\new if the ice was ejected in large chunks}, whereas dust ejected into escape orbits by the impact would only linger near the comet for short periods.  Early post-impact observations could thus provide a misleading impression of the longer-term integrated dust to water ratio lost by the comet in response to the impact.  No newly formed active area was detected by SWAS to significantly contribute to the water vapor production rate of the comet for more than 2-3\\,days. The non-detection of water vapor due to the Deep Impact event argues against the impact of meteoritic material as the origin of the natural outbursts observed in \\tempelfull. This is suggested by the small amount of water released by the impact when compared to the natural outbursts, and by the fact that this would require an abundance of meteorites in interplanetary space far higher than all other observations permit in order to explain the number of outbursts observed in \\tempelfull\\ {\\new \\citep[eg.,][and references therein]{Lang1992,Ceplecha1992,Greenberg1994}. } {\\new Another argument against exogenic processes, such as meteorite impacts, as the reason for the outbursts is the correlation of the outbursts with the light curve observed by \\cite{AHearn2005}.}  {\\new The water production rate of \\tempelfull\\ shows large natural variations, and the observed fluctuations are somewhat surprising and difficult to explain.  While the surface material of the comet is largely desiccated, periodic spurts of water vapor emission, nevertheless, are observed. Strictly, a model similar to the outburst model presented for the observations made during the Deep Impact event has to be employed for the water monitoring observations, particularly if one considers that the natural outbursts and the overall variability of the water outgassing rate of the comet is much larger than the amount of water vaporized by the impact. While the present paper for the first time presents a model for radio-line emission in outbursting comets, such a modeling is currently computationally prohibitive for the entire data set of the \\tempelfull\\ observations unless the timing of the outbursts are known (the outburst duration and start time). Several of the strong outbursts recorded by Earth-based instruments are summarized by \\cite{Meech2005a}, but it is not always possible to exactly constrain the timing of these outbursts because of the limited time coverage of these observations and because tracers other than water were observed in most cases. In fact, the opportunity to correlate observed outburst signals with a well-defined impact date was one of the reasons to conduct the Deep Impact experiment.  For the SWAS observations of a natural outburst, because we average over 2-3 days, it is not possible to decide whether the water molecules were released in a relatively short time ($\\tau\\ll 3$\\,days) or at a lower rate over a longer period ($\\tau\\la$3\\,days). Presenting an average $\\qwat$ assuming a constant water production rate during the period where the SWAS-detected signal was averaged is therefore acceptable to trace the variation of the cometary activity on time-scales of $\\ga$days. Since SWAS is sensitive to the total number of water molecules in the extended coma where line excitation is governed by IR-fluorescence, a sudden increase in the water production rate is not immediately visible in the SWAS signal but increases gradually over 1-2 days (Fig.\\ \\ref{double-water-prod}). Assuming a two-fold increase of the water production rate, the SWAS-detected signal has reached 90\\% of the line intensity of the corresponding steady-state model with twice the water production rate within $\\sim30$\\,hours. Thus, one has to note that the SWAS-derived water production rate for \\tempelfull\\ presented with Table \\ref{swas-monitor-tempel1} and Fig.\\ \\ref{tempel1-qwater}) might lag the true water production rate by $\\sim 1.3$\\,days during phases of strong comet variability. A similar conclusion applies to all studies of the water production rate using water-line observations covering a significant fraction of the gaseous coma using a Haser model. }  The main indication of desiccation comes from the high surface temperatures and the associated low thermal conductivities and inertia on the sunlit side of the comet.  \\cite{AHearn2005} find that \\tempelfull\\ has a largely homogeneous color and  uniform albedo as low as $0.04 \\pm 0.02$ over its entire surface. At perihelion, maps of the nucleus show temperature variations across the surface between $260 \\pm 6$ K to $329 \\pm 8$ K, with an absolute calibration uncertainty of order 20\\%. Even at the lowest cited temperature, however, the sublimation rate of ice close to the comet surface would be {\\new significantly} higher than observed. As \\cite{AHearn2005} also point out, the high surface temperature argues for a thermal inertia probably lower than $10^5$ erg K$^{-1}$ cm$^{-2}$ s$^{-1/2}$.  This implies that the 40.83\\,hr periodic heat pulse on the sunlit side corresponding to the observed 40.83\\,hr rotation period observed close to perihelion does not penetrate deep into the comet's surface, and is consistent with the low intra-pulse evaporation rate of water.  {\\new Most of} the frozen water apparently resides deeper beneath the surface than the heat pulse penetrates, and remains at a temperature sufficiently low to limit sublimation. {\\new \\cite{Sunshine2006} detect water ice on the surface of \\tempelfull, but these ice deposits cover too small a fraction of the surface to account for the total water vaporization rate observed for the comet.}  A less compelling indication for dehydration of the comet's surface layers comes from the debris released by the July 4 impact. \\cite{AHearn2005} conclude that ejecta consisted largely of very fine dust in the 1 to 100 micrometer size range, while \\cite{Meech2005a} find the ratio of dust to gas in the impact ejecta to have been significantly higher than in the pre-impact emission of material from the comet. \\cite{Meech2005a} show that the light curve mainly associated with dust emission rose to a maximum within roughly half an hour after impact and then declined after two hours, \\cite{Mumma2005} estimate a doubling of the post-impact water vapor production rate integrated over a 24 hour period. A contribution to the water vapor production over this longer period could have come from newly excavated material ejected from the crater and strewn over a sizeable surrounding area, consistent with the observation that an appreciable fraction of the excavated material fell back onto the comet surface.  The ice in the ejecta would subsequently have been sublimated by sunlight.  As \\cite{Kueppers2005} point out, the heat required for the evaporation of this much water could not have been provided by the energy of impact and would instead have had to come from the Sun.  \\cite{AHearn2005} point out that a number of pre-impact cometary outbursts appear to be correlated with the comet's rotation rate. More specifically, they identify {\\new emission from a topographic feature, possibly a scarp or depression,} that sees the sun rising around the time of each outburst. The sudden heating of the surface of the {\\new topographic feature} by direct sunlight could cause both rapid sublimation and the release of dust grains no longer held together by ice. Smaller grains could be carried off into space by the evaporating water. The gravitational attraction of the comet is sufficient to have retained a significant fraction of  the solid impact ejecta, as shown by the plume of material raining back down onto the comet's surface after impact.  The SWAS water vapor detection history shown in Fig.\\ \\ref{tempel1-qwater} indicates a rough correlation between the times of a natural outburst identified by other observers, and a rise in the 557\\,GHz water vapor emission integrated over a day or two following the outburst. {\\new This time delay is consistent with the expected response of the SWAS-detected signal to an instant increase of the water vaporization rate of the nucleus shown in Fig.\\ \\ref{double-water-prod}.} {\\new In this context it should be noted that the outbursts indicated in Fig. \\ref{tempel1-qwater} are ``big'' natural outburst which were also detected by Earth-based telescopes, while the weaker outbursts reported by \\cite{AHearn2005} were too faint to be detected by Earth-based instruments. }  If the prime source of water vapor is {\\new a tilted topographic feature on the surface such as a scarp or wall}, and the outburst occurs at sunrise, a slight slowing down of the comet's rotation should become evident in the course of time.  The momentum of evaporating water along the direction of rotation is $\\sim N(mkT)^{1/2}$, where $N$ is the total number of water molecules released over a period of time, $T\\sim 300$ K is the surface temperature at sunrise and $m$ is the molecular weight of water. With a lower limit to the number of water molecules vaporized from the nucleus during the 2005 perihelion passage of  {\\new $N\\sim 10^{35}$ } (Sect.\\ \\ref{qwat-monitoring}), the net angular momentum transferred to the comet nucleus through evaporation at radius $r \\sim 3$ km should have been roughly $N(mkT)^{1/2}r\\sim 3\\times 10^{22}$g cm${^2}$ s$^{-1}$. The angular momentum  of the comet is $\\sim \\pi Mr^2/P$, where $P\\sim 40.83\\pm0.33$ hr is the rotation {\\new period}, $M\\sim 4\\pi \\rho r^3/3$ and $\\rho\\sim 0.6$\\,g\\,cm$^{-3}$ is the density of the nucleus cited by \\cite{AHearn2005}. For comet \\tempelfull, $\\pi Mr^2/P\\sim 2\\times 10^{23}$\\,g\\,cm$^2$ s$^{-1}$. While the fractional mass lost by the comet due to the sublimation of ice is quite small, the velocity of sublimated molecules is several thousand times higher than the rotational surface velocity of the comet. A systematic change in the rotational period of the comet could, therefore, amount to as much as {\\new several} percent during a single perihelion passage.  Whether a slowdown in the comet's rotational period by up to a few percent during the 2005 perihelion passage could be culled out from the data is not clear. \\cite{Belton2005} give the pre-perihelion rotation period of the comet at heliocentric distances greater than 4 AU as $41.85\\pm 0.10$ hours, although it is not clear whether this significantly differs from the 40.83 hours measured at perihelion. If so, the comet's rotation rate appears to have sped up during perihelion passage. Either way, it would be useful to re-measure the post-perihelion rotation period at large heliocentric distances to see whether significant changes in the rotation period of Tempel 1 can be established.        \\ack Support of the SWAS mission is provided by NASA through SWAS contract NAS5-30702. Additional support for the SWAS 9P/Tempel 1 observing campaign was obtained by the Deep Impact team. The help and cooperation of NASA HQ is greatly appreciated.  \\label{lastpage}"
    },
    "0606/gr-qc0606059_arXiv.txt": {
        "abstract": "We present a new method for the calculation of black hole perturbations induced by extended sources in which the solution of the nonlinear hydrodynamics equations is coupled to a perturbative method based on Regge-Wheeler/Zerilli and Bardeen-Press-Teukolsky equations when these are solved in the frequency domain. In contrast to alternative methods in the time domain which may be unstable for rotating black-hole spacetimes, this approach is expected to be stable as long as an accurate evolution of the matter sources is possible. Hence, it could be used under generic conditions and also with sources coming from three-dimensional numerical relativity codes. As an application of this method we compute the gravitational radiation from an oscillating high-density torus orbiting around a Schwarzschild black hole and show that our method is remarkably accurate, capturing both the basic quadrupolar emission of the torus and the excited emission of the black hole. ",
        "introduction": "\\label{intro}  It is well-known that matter falling into a black hole, or orbiting around it, is a source of gravitational radiation.  The detection of these waves, and of the imprint they carry on the excitation of the black hole quasi normal modes (QNM), would provide an invaluable test of General Relativity in the strong-field regime.  This process can be studied using the theory of black hole perturbations, the origin of which dates back to 1957 when Regge and Wheeler~\\cite{RW} showed that the axial perturbations of a Schwarzschild black hole can be studied by solving a radial equation in the frequency domain, the angular part of the perturbation being known in terms of a suitable expansion in tensor spherical harmonics.  Among the many dedicated to this subject, two papers can be considered as milestones in the theory of black hole perturbations. The first one is by F. Zerilli~\\cite{Zer}, who derived the radial equation governing the polar perturbations of a Schwarzschild black hole, thus completing the work of Regge and Wheeler (see~\\cite{nr05} for a recent review). The second one is by S. Teukolsky~\\cite{Teu0}, who achieved the formidable task of reducing the equations for the perturbations of a Kerr black hole to a master radial equation.  Both the Zerilli and the Teukolsky equations were derived in the frequency domain.  The theory of black hole perturbations has been developed largely to study the stability properties of black holes and to determine the gravitational signals emitted by point particles moving near black holes on different orbits (see for instance~\\cite{NOK} and references therein). In addition, it has been used to evaluate back-reaction effects on the orbital motion of a point particle moving around a nonrotating black hole~\\cite{CKP}.  Although it clearly represents only a toy-model, the point-particle approximation has been very useful to study a number of phenomena that occur in the neighborhood of black holes and it has also been applied to study stellar perturbations~\\cite{noi1,noi2}. However, if one wants to investigate realistic astrophysical processes, in which finite-size effects and internal dynamics need to be properly taken into account, the point-particle approximation is certainly too crude. Therefore, several attempts have been made towards a more realistic description. In ref.~\\cite{DS}, for instance, the point particle orbiting a black hole was treated as having a spin angular momentum, while in ref.~\\cite{HSW} the infalling matter was assumed to be a finite size sphere of dust. Furthermore, in ref.~\\cite{SW} a semianalytic method was developed to study the radial infall of a shell of dust; while this latter approach has the advantage of being easily generalizable to sources with arbitrary shapes, it has the serious drawback that it can be applied only to sources in which the matter configurations are freely falling onto the black hole. It is worth stressing that in all cases mentioned above the black hole is assumed to be nonrotating.   Alternative and more general approaches to treat extended sources have recently been developed by Papadopulous and Font~\\cite{PF99} and by Nagar et al.~\\cite{NFZD}, by combining perturbative approaches with fully multidimensional hydrodynamical calculations. In ref.~\\cite{NFZD}, in particular, the polar perturbations of a Schwarzschild black hole induced by an oscillating, thick accretion disk (a torus) have been studied by solving the inhomogenous Zerilli equation in the {\\it time domain}. The source terms were computed self-consistently by using the results of two-dimensional (2D) simulations in which the relativistic hydrodynamical equations were solved for a torus oscillating in the Schwarzschild, unperturbed background. The purpose of that work was to compare the gravitational signal found by solving the Zerilli equation, with the one found using the standard quadrupole formalism applied to the oscillating torus.  In this paper we develop a new approach to study the perturbations of a Schwarzschild black hole excited by extended matter sources, by solving the inhomogeneous Bardeen-Press-Teukolsky (BPT) equation~\\cite{BP,Teu} in the {\\it frequency domain}. Also in this case the source for the perturbative equations is the result of independent 2D hydrodynamical simulations in the time domain. Because of the combined use of both time and frequency domains approaches, we refer to this as to a {\\it hybrid approach}.  There are at least three good reasons why our hybrid approach may be preferrable over a pure time-domain approach as the one used in~\\cite{NFZD}. Firstly, the Regge-Wheeler and the Zerilli equations (hereafter RWZ equations) cannot be generalized to rotating black holes, whereas the BPT equation admits such a generalization, {\\it i.e.,} the Teukolsky equation (We recall that the Teukolsky equation is not separable in the time domain because the eigenvalues of the angular functions are frequency dependent; this is not the case in the frequency domain where the equation is indeed separable.). Secondly, although the BPT equation is intrinsically unstable when integrated in the time domain, its integration does not offer any numerical difficulty when performed in the frequency domain. Thirdly, it is a common experience that by removing the error in the time integration of a system oscillating around its equilibrium state, a frequency-domain approach improves the overall accuracy. The work presented here serves therefore to develop a new procedure for the solution of the BPT equation with extended matter sources, whose generalization will allow to reach a long-standing goal: {\\it i.e.,} the study of the gravitational-wave emission from an extended source perturbing a {\\em rotating} black hole.  While the mathematical apparatus behind our hybrid approach is well-know, our numerical implementation is more innovative and to test it against an astrophysically realistic and computationally complex configuration we have here also considered the gravitational-wave emission from a high-density torus oscillating in its orbital motion around a Schwarzschild black hole. This is the same source considered in ref.~\\cite{NFZD} and thus offers the possibility of comparing for the same test case our hybrid approach with one which is instead in the time-domain only.  To this scope, we recall that the possibility that extremely dense and massive accretion tori form around a black hole has been considered in recent years, initially as a possible explanation of $\\gamma$-ray bursts~\\cite{Woo}.  In these systems, a torus of mass $M_t\\sim 0.1\\!-\\!1M_\\odot$ and mean density as high as $\\sim10^{11}\\!-\\!10^{12}$ g/cm$^3$, forms around a black hole of a few solar masses. The torus can be located within a few tens of km from the horizon and could also be subject to a dynamical instability which would induce its accretion onto the black hole on a dynamical timescale (see~\\cite{font:02,ZRF} and references therein). High-density tori could be generated in many astrophysical scenarios (see~\\cite{Woo,ZRF} for a general overview): in the coalescence of black hole-neutron star or neutron star binaries, as shown by several numerical simulations~\\cite{st06} in the collapse to a black hole of a rapidly rotating supramassive neutron star~\\cite{VS,detal_06}, as a byproduct of collapsars, that is, massive stars in which iron core-collapse does not produce an explosion, but forms a black hole~\\cite{MW}.  As discussed in~\\cite{ZRF}, these systems can be created with event rates comparable to that of core collapse supernovae. In addition, because the tori have an intrinsically large quadrupole moment, even small-amplitude oscillations would produce gravitational signals that may be detectable by ground based interferometers Virgo and LIGO~\\cite{ZRF}. As a result, the calculation of the gravitational waves from these tori does not only represent a useful testbed for our novel approach, but it also offers the opportunity for a more accurate calculation of the emission from a realistic and intense source of gravitational radiation.  The gravitational radiation emitted by a torus oscillating around a black hole was first computed in ref.~\\cite{ZRF} within the Newtonian quadrupole formalism.  However, the quadrupole formalism may not be able to describe a system with such a high energy density, and located so close to a black hole horizon with sufficient accuracy. In particular, the possible excitation of the black hole QNM cannot be revealed by the quadrupole approach.  The same system has been studied in~\\cite{NFZD}, where the perturbed equations have been integrated in the time domain and no excitation of the black hole quasi-normal modes (QNMs) was found. Conversely, as we will later discuss, our hybrid method applied to the same system allows to catch the excitation of the black hole quasi normal modes, though the peak they produce in the energy spectrum is much smaller than that corresponding to the torus oscillation.  The plan of the paper is as follows: in Section~\\ref{approach} we review the main equations which describe the perturbations of a Schwarzschild black hole in the RWZ and in the BPT approaches when these are expressed in the frequency domain. Furthermore, we also recall the expression of the gravitational energy flux computed within the Newtonian quadrupole formalism. In Section~\\ref{hybr} we present the analytical and numerical setup of our approach for the solution of both the perturbative and general relativistic hydrodynamics equations. Finally, in Section~\\ref{results} we discuss the results of the numerical integrations for the representative case of an oscillating torus, while in Section~\\ref{conclusions} we draw our conclusions. A number of Appendices contain details on the form of the source functions employed for the solution of the perturbative equations. ",
        "conclusions": "\\label{conclusions}  We have presented a new {\\it hybrid} approach to study the perturbations of a Schwarzschild black hole excited by extended matter sources in which the solution in the {\\it time-domain} of the relativistic hydrodynamical equations in a multidimensional spacetime is coupled to the solution in the {\\it frequency-domain} of either the Regge-Wheeler/Zerilli equations or of the Bardeen-Press-Teukolsky equations.  We believe that our hybrid approach may be preferrable over a pure time-domain approach based on the RWZ equations for at least three different reasons. Firstly, the Regge-Wheeler and the Zerilli equations cannot be generalized to rotating black holes, whereas the BPT equation admits such a generalization. Secondly, the BPT equation is intrinsically unstable when integrated in the time domain but its solution is regular computed in the frequency domain. Thirdly, for a system undergoing small oscillations around an equilibrium state, a frequency-domain approach is expected to be significantly more accurate than one based in the time-domain, especially at high frequencies, where the black hole contributions are expected to emerge. Within this framework, therefore, the work presented here serves as a first step towards the study of the gravitational-wave emission from an extended source perturbing a rotating black hole.  As a test of its effectiveness and to provide a close comparison with what done so far in perturbative approaches based in the time-domain, {\\it i.e.,} ref.~\\cite{NFZD}, we have considered as extended source to the perturbation equations the oscillations of a high-density torus orbiting around a Schwarzschild black hole. The dynamics of the torus has been followed numerically using a 2D numerical code solving the general relativistic hydrodynamics equations using HRSC schemes in spherical polar coordinates.  Overall, the results of our study show that the hybrid approach (either with the RWZ equations or with the BPT one) works remarkably well, producing an energy spectrum of the emitted gravitational radiation which contains both the basic quadrupolar emission of the torus as well as the excited emission of the black hole.  The latter feature was completely washed out in the complementary time-domain approach of ref.~\\cite{NFZD}.  Being totally generic, the formalism and methodology presented here is could be used under more general conditions and in particular also with sources coming from 3D numerical relativity codes as long as an accurate evolution of the matter sources is possible. Future work in this line of research will therefore comprise both the extension of the present results to 3D codes, as well as the generalization of the perturbative formalism to the Teukolsky equation in order to model the gravitational emission of rotating black holes when perturbed by extended sources."
    },
    "0606/astro-ph0606263.txt": {
        "abstract": "{This paper describes the general characteristics of raw data from fiber-fed spectrographs in general and fiber-fed IFUs in particular. The different steps of the data reduction are presented, and the techniques used to address the unusual characteristics of these data are described in detail.  These techniques have been implemented in a specialized software package, R3D, developed to reduce fiber-based integral field spectroscopy (IFS) data.  The package comprises a set of command-line routines adapted for each of these steps, suitable for creating pipelines. The routines have been tested against simulations, and against real data from various integral field spectrographs (PMAS, PPAK, GMOS, VIMOS and INTEGRAL).  Particular attention is paid to the treatment of cross-talk. } ",
        "introduction": "In fiber-fed spectrographs the slit is not placed at the focal plane of the telescope. Instead, the light from different locations on the sky is conducted into the spectrograph using optical fibers.  At one end, the optical fibers are located in the focal plane of the telescope and at the other they are rearranged along the slit of the spectrograph, in the so-called pseudo-slit. Fiber-fed spectrographs are normally used for multi-object spectroscopy (MOS), where the fiber ends in the focal plane of the telescope can be placed at different positions on the sky, or for Integral Field Spectroscopy (IFS), where the fibers are packed in bundles or coupled with arrays of lenses.  Integral Field Spectroscopy (IFS) is a technique for obtaining multiple spectra simultaneously of a (more or less) contiguous area of the sky. Although fiber-fed spectrographs are widely used, there are other techniques to conduct the light from different positions in the sky to the spectrograph, e.g. lens arrays and image slicers.  These different implementations have produced a set of instruments, which, while sharing the basics of the technique, produce very different representations of the spectra at the detectors. This apparent diversity has led to the development of reduction techniques and/or packages for each individual instrument (e.g., P3d,\\cite{beck01}; \\cite{roth05}; VIPGI, \\cite{vimos05}; GEMINI, \\cite{turn06}). Together with the inherent complexity of this technique, this has reduced the use of IFS for decades to a handful of specialists, each usually working with a specific instrument.   IFS developers became aware of this handicap, and have started to produce standard techniques and tools valid for any integral field unit (IFU). In a recent effort, the Euro3D RTN (\\cite{net02}) has created a standard data format (\\cite{kiss04}), a coding platform (\\cite{peco04}) and a visualization tool (\\cite{sanc04}), useful for any of the existing IFU. All these tools are freely distributed to the community. However most of them are for working with reduced data, while the reduction itself has still not been addressed with a generalistic approach.  The astronomical community is confronted with a jungle of different instrument-specific pipelines, whose procedures and parameters differ greatly, and whose outputs are difficult to compare.  Building on earlier work, we present here a summary of the most important steps in the reduction of data from fiber-fed spectrographs in general, and fiber-fed IFS in particular. All of them have been implemented in {\\tt R3D}, a reduction package able to reduce data from any fiber-fed IFU, whose main characteristics we also present.  The structure of this article is as follows. Section \\ref{sec:code} describes the coding platform used to develop {\\tt R3D}. Section \\ref{sec:raw} summarizes the major characteristics of the raw data from fiber-fed spectrographs. The different steps required for a proper data reduction are explained in Section \\ref{sec:red}. In Section \\ref{sec:com} we compare {\\tt R3D} with other packages in existence.  Section \\ref{sec:con} summarizes the conclusions. The data examples shown throughout this article are from PMAS (\\cite{roth05}) in the PPAK mode (\\cite{kelz06}), using 2$\\times$2 pixel binning, although {\\tt R3D} is not limited to this instrument/setup. For examples of its use with another instruments, please consult S\\'anchez \\& Cardiel (2005). ",
        "conclusions": "\\label{sec:con}  This article has explained the particular characteristics of the raw data and data reduction of fiber-fed spectrographs in general and IFUs in particular. We presented a specialized package for reducing this kind of data, {\\tt R3D}, explaining in detail how the various steps in the data reduction process have been implemented. The effect of cross-talk and its possible solutions were discussed in detail, with simulations that illustrate them. {\\tt R3D} has been tested over several IFUs, including PMAS in its Lensarray and PPAK modes (which differs substantially), GMOS, VIMOS and INTEGRAL. Throughout this article we used PPAK data to illustrate how the data {\\it evolve} after passing through each reduction step. The basic adopted solutions for the reduction steps were compared with those of the IRAF task {\\tt dofibers}, which is widely used to reduce this kind of data.  Integral Field Spectroscopy is evolving fast from a technique for specialists to a common user technique, and generalist tools, like {\\tt R3D}, are needed for the astronomical community.  Thus we distribute the package freely.  It provides a common package for reducing data from any kind of fiber-fed spectrograph in general and IFUs in particular. We plan to extend it to other types of IFUs (such as TIGER-based lensarrays or Image Slicers). Use of {\\tt R3D} will help soften the learning curve for working with IFS data obtained with different IFUs."
    },
    "0606/astro-ph0606228_arXiv.txt": {
        "abstract": "We show that the primordial curvature perturbations may originate not from the quantum fluctuations of inflaton during inflation but from isocurvature perturbations which are amplified during preheating after inflation. We consider a simple preheating model, whose potential is given by $V(\\phi, \\chi)=\\frac{1}{4}\\lambda \\phi^4+\\frac{1}{2}g^2 \\phi^2 \\chi^2$ with $g^2/\\lambda=2$, as a possible realization of generating curvature perturbations during preheating. We make use of the $\\delta N$ formalism which requires only knowledge of the homogeneous background solutions in order to evaluate the evolution of curvature perturbations on super-horizon scales. We solve the background equations numerically and find that the amplitude and the spectral index of curvature perturbations originating from preheating can be tuned to the observed values if the isocurvature perturbations at the end of inflation are not suppressed on super-horizon scales. We also point out that the tensor to scalar ratio in $\\frac{1}{4}\\lambda \\phi^4$ inflation model can be significantly lowered, hence letting the $\\frac{1}{4}\\lambda \\phi^4$ model, which is ruled out by the Wilkinson Microwave Anisotropy Probe (WMAP) data combined with SDSS data, get back into the observationally allowed region. ",
        "introduction": "The primordial curvature perturbations which are coherent over the horizon scales at recombination epoch are seeds for large scale structure. A simple scenario to explain the origin of curvature perturbations is that vacuum fluctuations of inflaton driving inflation convert into classical perturbations at around the time of horizon crossing, which are almost adiabatic, almost scale-invariant and almost Gaussian. These perturbations are constant until the scales of interest re-enter the horizon.  Recently alternative mechanisms were proposed which generate the primordial curvature perturbations after inflation, not from the inflaton fluctuations. In curvaton scenario \\cite{Linde:1996gt, Moroi:2001ct, Lyth:2001nq}, quantum fluctuations of light scalar field during inflation which are not necessarily related to the inflaton, convert into curvature perturbations after inflation ends. In inhomogeneous reheating scenario \\cite{Dvali:2003em, Kofman:2003nx, Matarrese:2003tk}, spatial fluctuations of the decay rate of the inflaton to radiation generate curvature perturbations. Both of these mechanisms are based on the formula \\cite{Lyth} \\begin{equation} {\\dot {\\cal R}_c}=\\frac{H}{\\rho+P}\\delta P_{\\rm nad},  \\label{int1} \\end{equation} which is valid on super-horizon scales. Here ~${\\dot {}}$ ~is a derivative with respect to the cosmological time, $H \\equiv \\dot{a}/a$ represents the Hubble parameter, $\\rho$ is the energy density, $P$ is the pressure, ${\\cal R}_c$ is the curvature perturbation on comoving hypersurface and $\\delta P_{\\rm nad} \\equiv \\delta P-\\frac{{\\dot P}}{ {\\dot \\rho}} \\delta \\rho$ is called the non-adiabatic part of the pressure perturbation or isocurvature perturbation. If $\\delta P_{\\rm nad}$ is non-zero after inflation, which is the case of the curvaton and the inhomogeneous reheating scenarios, ${\\cal R}_{c}$ on super-horizon scales is generated after inflation.  As another candidate for the mechanism of generating curvature perturbations after inflation, preheating model has also been considered~\\cite{Bassett:1999mt, Bassett:1998wg, Zibin:2000uw, Tanaka:2003ck, Kolb:2004jm}. Preheating \\cite{Traschen:1990sw, Kofman:1994rk, Yoshimura:1995gc, Fujisaki:1995dy} is a process in which the coherent oscillation of the inflaton excites the fluctuations of fields coupled to the inflaton by parametric resonance. This process was investigated in detail \\cite{Yoshimura:1995gc, Fujisaki:1995dy, Son:1996uv, Boyanovsky:1996sq, Prokopec:1996rr, Baacke:1996kj, Khlebnikov:1996mc, Khlebnikov:1996wr, Khlebnikov:1996zt, Boyanovsky:1997cr, Kofman:1997yn, Greene:1997fu, Baacke:1997rs, Greene:1997ge, Greene:1998nh, Suyama:2004mz} without metric perturbations, which is valid in dealing processes occurring well within the horizon scale. It was then recognized that particle production in preheating occurs much more efficiently than in the old reheating scenario where inflaton decays perturbatively to other particles.  Naively, fluctuations of such fields coupled to the inflaton seem to be excellent source of isocurvature perturbations. This means there is a possibility of generating curvature perturbations on super-horizon scales during preheating epoch. However, whether preheating can affect the evolution of curvature perturbations on super-horizon scales or not \\cite{Bassett:1999mt, Bassett:1998wg} is not completely understood yet. A necessary condition that curvature perturbations on super-horizon scales are affected by parametric resonance is that isocurvature perturbations are not suppressed on super-horizon scales \\cite{Gordon:2000hv}. Though this condition is not satisfied for large class of two-field preheating models, there remains some models which satisfy this condition. A simple example is a massless self-coupled inflaton $\\phi$ coupled to another scalar field $\\chi$, i.e. with potential \\cite{Bassett:1999cg} \\begin{equation} V(\\phi, \\chi)=\\frac{\\lambda}{4}\\phi^4+\\frac{g^2}{2}\\phi^2 \\chi^2, \\label{int2} \\end{equation} with $g^2/\\lambda = 1\\sim 3$.  For the model given in Eq.~(\\ref{int2}) with $g^2/\\lambda=2$, where the $k=0$ mode has the largest growth rate($k$ is the wavenumber), it was suggested in Ref.~\\cite{Zibin:2000uw} that large-scale curvature perturbations are much amplified and exceed the current observational upper bound by using the mean field approximation.   On the other hand, in Ref.~\\cite{Tanaka:2003ck}, by making use of the $\\delta N$ formalism \\cite{Sasaki:1995aw, Sasaki:1998ug, Wands:2000dp}, it was suggested that the amplification of curvature perturbations during preheating is less efficient than is expected in \\cite{Bassett:1999cg, Zibin:2000uw}.  In the $\\delta N$ formalism, we can evaluate the evolution of curvature perturbations only by the knowledge of the homogeneous background solutions and as was emphasized in \\cite{Tanaka:2003ck}, unphysical acausal energy transfer does not occur in $\\delta N$ formalism while this energy transfer is not inhibited in the mean field approximation used in \\cite{Zibin:2000uw}.  However in Ref.~\\cite{Tanaka:2003ck}, whether the generation of curvature perturbations on super-horizon scales takes place during preheating or not remains an open question because the number of solving background equations during preheating varying the initial value of the isocurvature perturbations at the end of inflation is not enough.  In this paper, we numerically solve background equations for the model given by Eq.~(\\ref{int2}) during preheating and show that scale invariant curvature perturbations with amplitude observed in the temperature anisotropy of Cosmic Microwave Background(CMB) can be generated during preheating. \\footnote{Generating the primordial curvature perturbations in preheating was also discussed in \\cite{Kolb:2004jm}. They considered instant preheating model \\cite{Felder:1998vq} in which case the resulting curvature perturbations are proportional to the initial isocurvature perturbations. As we will show in Sec. IV, the resulting curvature perturbations for the case we consider are very sensitive to the tiny differences of the initial isocurvature perturbations.}  Before we end this section, we define the power spectrum of any field $f({\\vec x})$, which we denote as ${\\cal P}_f$, by \\begin{equation} \\langle f_{\\vec k_1} f_{\\vec k_2} \\rangle =\\frac{2\\pi^2}{k^3_1} {\\cal P}_f (k_1) \\delta ( {\\vec k_1}+{\\vec k_2} ). \\label{int3} \\end{equation} Here $\\langle \\cdots \\rangle$ is the ensemble average and $\\vec{k}$ is the wave number vector.  The paper is organized as follows. In Sec. II we briefly review the $\\delta N$ formalism which we use in a later section in order to evaluate curvature perturbations generated during preheating. In Sec. III we give the basics of preheating and give homogeneous background equations which we numerically solve. In Sec. IV we solve the background equations numerically and discuss how we can obtain the power spectrum of curvature perturbations from the solutions of the background equations. Then we give the resulting power spectrum after preheating. Sec. V is a summary. ",
        "conclusions": "The main purpose of this paper is to examine whether the generation of the curvature perturbations occurs or not during preheating and, if it occurs, whether they could be seeds for large-scale structure in the universe.  To do so, we applied $\\delta N$ formalism to two-field preheating model whose potential is given by $\\frac{\\lambda}{4}\\phi^4+\\frac{g^2}{2}\\phi^2 \\chi^2$. Then we have analyzed the dependence of the $e$-folding number during preheating on the isocurvature perturbations at the end of inflation which corresponds to perturbations of $\\chi$-field. We verified that the $e$-folding number during preheating is very sensitive to the initial value of $\\chi$ due to the chaotic nature of the background trajectory in field space, which has been already observed in \\cite{Tanaka:2003ck}. Though the $e$-folding number seems almost random as a function of $\\chi_{ini}$, we found that the fluctuation of the smoothed $e$-folding number over a large number of initial values of $\\chi$ exceeds the statistical fluctuation. Hence the $e$-folding number is not completely randomized.  This deviation of the smoothed $e$-folding number from the statistical fluctuation provides the generation of curvature perturbations on large scales. We numerically found that the amplitude and the spectral index of the curvature perturbations originating from preheating at current cosmological scales can be fixed to the observed values by tuning the spectral index of $\\chi$-field and the infrared cutoff. An interesting point is that this consequence is derived only from the requirement $g^2/\\lambda=2$. Hence by lowering $\\lambda$ which determines the energy scale of inflation, we can lower the tensor to scalar ratio to the observationally allowed value."
    },
    "0606/astro-ph0606702_arXiv.txt": {
        "abstract": "We present $ugriz$ photometry and optical spectroscopy for 28 DB and DO white dwarfs with temperatures between 28,000K and 45,000K.  About 10 of these are particularly well-observed; the remainder are candidates. These are the hottest DB stars yet found, and they populate the ``DB gap'' between the hotter DO stars and the familiar DB stars cooler than 30,000K. Nevertheless, after carefully matching the survey volumes, we find that the ratio of DA stars to DB/DO stars is a factor of 2.5 larger at 30,000~K than at 20,000~K, suggesting that the ``DB gap'' is indeed deficient and that some kind of atmospheric transformation takes place in roughly 10\\% of DA stars as they cool from 30,000~K to 20,000~K. ",
        "introduction": "\\label{sec:intro}  Existing surveys of white dwarfs have shown a peculiar absence of stars with helium atmospheres with temperatures between roughly 30,000K and 45,000K.  Currently, the hottest published DB temperature is for the star PG~0112+104 near 30,000~K \\citep{lie86}. The most accurate determination is probably $T_{\\rm eff} = 30,783\\pm269$~K, $\\logg = 7.78\\pm0.02$ \\citep{bea95} using fits to He~I lines.  At this temperature and below, only He~I lines appear in the spectrum.  Hydrogen often appears as a trace constituent, and these spectra are classified DBA.  Beauchamp analyzed 51 DB and DBA stars, none hotter than PG~0112+104. \\citet{The91} analyzed the far ultraviolet spectrophotometry of 11 of the hottest known DB stars but generally showed these to have lower $T_{\\rm eff}$ than previously-published values. More recently \\citet{koe01} analyzed 19 more DB stars, but the hottest $T_{\\rm eff}$ determination was 25,400~K.  The He~II 4686\\AA\\ line should start appearing at about 40,000~K (38,000~K for a high signal-to-ratio spectrum), so spectra at this temperature and above would be classified DO.  However, the coolest published temperature for a DO star before any publications from the Sloan Digital Sky Survey (SDSS; York et al. 2000) is the rather crude estimate of $47,500\\pm2,500$~K for PG~1133+489 \\citep{wes85}.  Unfortunately, this star was not included in the more precise analysis of DO stars by \\citet{dre96}.  These authors included DO stars discovered in later surveys such as the Hamburg Schmidt survey \\citep{heb96}, but did not discover any cooler objects, thus supporting the existence of a ``DB gap''.  One or two stars classified DO from the SDSS may be similar in temperature to PG~1133+489 \\citep{Krz04}, but none of the stars discussed in that paper are considerably cooler.  One suggested cause of this ``DB gap'' has been that as the DO stars cool, small residual amounts of hydrogen diffuse to their surfaces, so that 50,000~K helium atmosphere stars become 40,000~K hydrogen DA white dwarfs \\citep{fon87}.  \\citet{fon87} also proposed that the reappearance of helium atmosphere stars below 30,000~K is due to convective mixing of the H-rich atmosphere into the much more massive helium envelope when the outer convection boundary reaches high enough to penetrate the hydrogen.  An unpleasant aspect of this theory is that the hydrogen layer mass evidently has to be quite thin, of the order of $10^{-15}$ of the stellar mass, such as was found for PG~1305--017 \\citep{ber94}. If the white dwarf progenitor's asymptotic giant branch phase is not terminated with a special episode of mass loss, the remaining hydrogen envelope is predicted by models to be of the order $10^{-4}\\msun$.  Asteroseismological studies of several pulsating DAV white dwarfs appear to confirm that these objects have hydrogen layer masses at least within a few orders of magnitude of this high value \\citep{cle95,bra98}.  Since the majority of white dwarfs within the DB temperature range remain DA, these apparently conflicting results may perhaps be consistent if a small fraction of the DAs retain only ultrathin outer hydrogen layers that can be convectively mixed below 30,000~K.  The problem of whether a DB gap exists is complicated by the known existence of several peculiar DAB, DBA, or DAO stars believed to lie in the 30--45,000~K range that (1) show evidence of spectrum variability and/or (2) do not fit simple atmospheric models.  Simple atmospheric models may be either homogeneous (completely mixed) in H and He throughout the atmosphere or stratified with the hydrogen all in a very thin, upper layer.  An example near 45,000~K is PG1210+533, which shows variable H, He I, and He II line strengths, probably modulated with an (unknown) rotational period like the magnetic star Feige 7 \\citep{lie77,ach92}. PG~1210+533 is not known to be magnetic.  An example near 30,000K is GD~323, which shows a peculiar DAB spectrum that has not been fit successfully with either a homogeneous or stratified model atmosphere \\citep{lie84,koe94}.  This star has a variable spectrum \\citep{per05} but is again not known to be magnetic.  In contrast to these are two stars in the middle of the gap which are well fit with homogeneous atmospheres models. HS~0209+0832 is fitted at 36000~K with a 2\\% helium abundance \\citep{Jor93} and PG~1603+432 at 35,000K with 1\\% helium \\citep{Ven04}. The peculiarity of some of these stars might be related to the edges of the DB gap, but in any case their uncertain relationship to the DB sequence has led workers to exclude them, perhaps too quickly, from the statistics of the DB gap. The possible roles of stellar wind mass loss, accretion and convective dredge-up in the spectral evolution of DAB stars is discussed in \\citet{Mac91}.  It must be emphasized, however, that the total number of well-analyzed DB stars prior to the SDSS was rather small, only slightly more than the 70 listed in the first paragraph.  Thus, the validity and significance of the alleged ``DB gap'' has not been clear.  However, SDSS has included many such stars in its spectroscopic program, resulting in a substantial increase in the census of DBs (as well as many more DAs).  Here we report the discovery of 6 DB white dwarfs with high signal-to-noise ratio photometry and spectroscopy that place their temperatures in the previously empty range of 30,000~K to 45,000~K.  We also discuss 4 other well-observed stars at the edge of the range, and 18 less well-observed candidates also likely to be in the DB gap. ",
        "conclusions": ""
    },
    "0606/astro-ph0606158_arXiv.txt": {
        "abstract": "This paper addresses the origin of the far-infrared (FIR) continuum of QSOs, based on the Quasar and ULIRG Evolution Study (QUEST) of nearby QSOs and ULIRGs using observations with the {\\it Spitzer Space Telescope}. For 27 Palomar-Green QSOs at z$\\lesssim$0.3, we derive luminosities of diagnostic lines (\\ne212m , \\nevIR, \\oivIR ) and emission features (\\pahIR\\ emission which is related to star formation), as well as continuum luminosities over a range of mid- to far-infrared wavelengths between 6 and 60$\\mu$m. We detect star-formation related PAH emission in 11/26 QSOs and fine-structure line emission in all of them, often in multiple lines. The detection of PAHs in the average spectrum of sources which lack individual PAH detections provides further evidence for the widespread presence of PAHs in QSOs.  Similar PAH/FIR and [NeII]/FIR ratios are found in QSOs and in starburst-dominated ULIRGs and lower luminosity starbursts.  We conclude that the typical QSO in our sample has at least 30\\% but likely most of the far-infrared luminosity ($\\sim 10^{10\\ldots 12}\\Lsun$) arising from star formation, with a tendency for larger star formation contribution at the largest FIR luminosities.  In the QSO sample, we find correlations between most of the quantities studied including combinations of AGN tracers and starburst tracers. The common scaling of AGN and starburst luminosities (and fluxes) is evidence for a starburst-AGN connection in luminous AGN. Strong correlations of far-infrared continuum and starburst related quantities (PAH, low excitation [NeII]) offer additional support for the starburst origin of far-infrared emission. ",
        "introduction": "The infrared properties of luminous active galactic nuclei (AGN) and ultraluminous infrared galaxies (ULIRGs) hold clues for the understanding of galaxy formation, the star formation history of the universe and the connection between black hole and galaxy formation and evolution. Physical insight into the co-evolution of AGN and star formation can be gained by study of the low redshift members of these populations. The Infrared Astronomical Satellite (\\IRAS) and the Infrared Space Observatory (\\ISO) have allowed the first comprehensive far-infrared (FIR) studies of extragalactic sources and some high quality mid-infrared (MIR) observations on the brighter sources \\citep[see e.g. the review of ][]{genzel00}. Issues discussed in \\IRAS\\ and \\ISO\\ studies include the nature of ultraluminous infrared galaxies and their starburst and AGN energy sources \\citep{genzel98,klaas01}, and the SEDs of QSOs and their evolutionary implications \\citep{haas03}. A  2.5 to 45 $\\mu$m spectral inventory of starburst and AGN prototypes \\citep{sturm00} and detailed mid-infrared spectroscopic investigations of local lower luminosity AGN \\citep[][hereafter S02]{clavel00,sturm02} have addressed AGN unification and the starburst-AGN connection.  The {\\it Spitzer Space Telescope} (hereafter \\spitzer ) has enabled us to build upon previous work with significantly improved sensitivity. This work is part of a series of papers describing the results of the Quasar and Ulirg Evolution STudy (QUEST), focussing on mid-infrared spectroscopy of a sample of 54 QSOs and ULIRGs at redshifts z$\\lesssim$0.3 with the Infrared Spectrograph \\IRS\\ onboard \\spitzer .  The sample was selected with the aim of investigating the possible connections between those two groups of luminous (L$_{Bol}\\gtrsim 10^{12}\\Lsun$) objects in the nearby universe, and is closely connected to optical and near-infrared studies of the morphology and dynamical properties of these two populations \\citep[e.g.][]{dasyra06a,dasyra06b,veilleux06}. The purpose of the present paper is to study starburst signatures in QSO host galaxies and to show that most, and in some cases all, of the far-infrared luminosity in QSOs showing strong PAH features is due to starburst activity. A forthcoming complementary paper addresses the spectral energy distribution (SED) of QSOs including the clear far-infrared starburst contribution to the $\\lambda > 30 \\mu m$ continuum of QSOs.  To investigate the link between AGN activity and star formation and the extent to which they occur simultaneously, it is important to quantify the star formation activity in QSO hosts.  Such measurements are made difficult, however, by the observational problems of detecting star formation tracers in the presence of extremely powerful AGN emission. SED studies based on the \\IRAS\\ and \\ISO\\ space missions have established QSOs as sources of (sometimes) strong far-infrared emission. \\citep[e.g.][]{neugebauer86,haas03}. In addition to a nonthermal continuum that is detectable in the infrared only in flat spectrum radio-loud QSOs \\citep[e.g.][]{haas98}, a strong far-infrared emission component is often observed, at varying levels with respect to the strong AGN mid-infrared continuum. Due to its steep falloff in the submillimeter regime, the origin of this far-infrared emission must be thermal emission of optically thin dust \\citep{chini89,hughes93}. While the warmer T$\\sim$200K dust, which dominates the mid-infrared SEDs of QSOs, is generally accepted to be predominantly AGN heated, there is still considerable dispute about the origin of the cooler T$\\sim$50K emission often dominating the far-infrared. Direct heating by the powerful AGN, but at distances ensuring sufficently low temperatures, is one possibility \\citep[e.g.][]{sanders89,haas03,ho05}. Other models prefer an origin in vigorous star formation in the QSO host \\citep[e.g.][]{roro95}. \\cite{roro95} used radiative transfer modelling to infer an SED of AGN heated dust that, in $\\nu$L$_\\nu$ units, peaks in the mid-IR and decays towards the far-infrared, a feature shared by many other such models. In the QSO SEDs that are often flat over a wide wavelength range including the far-infrared, the far-infrared component is then plausibly ascribed to a component with an SED similar to that of a star-forming galaxy, in accordance with evidence for coexistence of star formation and AGN in spatially resolved lower luminosity AGN. \\Citet{roro95} found a tight correlation of AGN optical emission and mid-infrared continuum and a weaker correlation between optical and far-infrared emission, which is supporting the view that the far-infrared does not result directly from AGN heating but that there is a connection between AGN and starburst luminosities in the QSOs.  Our goals are to quantify star formation in QSO hosts and to estimate its contribution to the the far-infrared emission. In the mid-infrared, the contrast between the emission from possibly dust-obscured star formation and from the central AGN is favourable, and established star formation tracers are available. We use two such tracers:  (1) The mid-infrared broad aromatic `PAH' emission features arise in regions of the interstellar medium of a galaxy where their aromatic carriers are present, and where their transient excitation is made possible by a non-ionizing ($<13.6$eV) soft UV radiation field. This is the case in the photodissociation regions (PDRs) that accompany Galactic star formation regions \\citep[e.g.][]{verstraete96}, as well as in the diffuse interstellar medium where they are excited by the general interstellar UV radiation field that has leaked from its OB star origins to large scales \\citep[e.g.][]{mattila96}.  PAHs have been used as a quantitative tracer of star formation in galaxies \\citep[e.g.][]{genzel98, foerster04, calzetti05}. Metallicity above 0.2 solar is a prerequisite for strong PAH emission \\citep{engelbracht05}, a condition that is probably safely met for local QSO hosts. Destruction of the PAH carriers in energetic environments but survival in starburst PDRs (though not in \\hii\\ regions proper) is key for the use of PAH features as diagnostic. The PAH features are absent from the hard radiation environment of AGNs according to both empirical \\citep[e.g.][]{roche91,lefloch01,siebenmorgen04a} and theoretical \\citep{voit92} studies. The latter suggest that PAH molecules hit by single energetic EUV/X-ray photons can be efficiently destroyed by photo-thermo dissociation or Coulomb explosion. Since AGN are copious emitters of hard photons, PAH molecules near AGN will be destroyed unless shielded by a large obscuring column.  (2) The low excitation fine-structure emission lines like \\ne212m \\ are among the dominant emission lines of HII regions. Observations of starburst galaxies, as well as a combination of evolutionary synthesis and photoionization modelling, show \\ne212m \\ to be stronger than higher excitation mid-infrared lines ([NeV], [OIV], [NeIII]) in typical ionized regions excited by young stellar populations \\citep{thornley00,verma03}. Use of low excitation lines as star formation tracers requires, however, the consideration of possible contributions from the AGN Narrow Line Region (NLR) which can be significant despite the generally higher excitation of such regions \\citep[e.g.][]{spinoglio92,alexander99} compared to starburst \\hii\\ regions.   Section 2 of the paper describes the sample, observations and data reduction used to obtain the line and continuum fluxes in our sources. Emission lines that are relevant to the present study are tabulated for all sources. In Section 3 we discuss the widespread presence of PAH emission and its relation to other components of the QSO spectra. Finally, Section 4 addresses the issue of star formation in host galaxies of QSOs, shows the importance of this process and compares our results with earlier findings. In a forthcoming paper, we discuss the implications of our results for QSO SEDs in general. We adopt $\\Omega_m =0.3$, $\\Omega_\\Lambda =0.7$ and $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "Sensitive \\spitzer\\ mid-infrared spectroscopy reveals the widespread presence of aromatic `PAH' emission features in z$\\lesssim 0.3$ QSOs from the Palomar-Green sample, indicating the presence of powerful  (\\nulnu60 $\\sim 1.7\\times 10^{10}$ to 2.5$\\times 10^{12}\\Lsun$) star formation activity in these systems. Starburst and AGN activity are connected in QSOs up to these high luminosities. By comparing the ratios of \\pahIR, \\ne212m , and far-infrared emission in QSOs with starbursts we conclude that for the average QSO in our sample at least 30\\% and likely most of the QSO far-infrared emission is due to star formation. The data suggest a trend with the star formation contribution being the largest in the most FIR-luminous QSOs."
    },
    "0606/astro-ph0606472_arXiv.txt": {
        "abstract": "{Molecular hydrogen observations towards Herbig-Haro objects provide the possibility of studying physical processes related to star formation.} {Observations towards the luminous IRAS source IRAS11101-5928 and the associated Herbig-Haro objects HH135/HH136 are obtained to understand whether high-mass stars form via the same physical processes as their low-mass counterparts. } {Near-infrared imaging and spectroscopy are used to infer H$_2$ excitation characteristics. A theoretical H$_2$ spectrum is constructed from a thermal ro-vibrational population distribution and compared to the observations. } {The observations reveal the presence of a well-collimated, parsec-sized  H$_2$ outflow with a total H$_2$ luminosity of about $2L_\\odot$. The bulk of the molecular gas is characterized by a ro-vibrational excitation temperature of $2000\\pm200$~K. A small fraction (0.3\\%) of the molecular gas is very hot, with excitation temperatures around 5500~K. The molecular emission is associated with strong [FeII] emission. The H$_2$ and [FeII] emission characteristics indicate the presence of fast, dissociative J-shocks at speeds of $v_\\mathrm{s} \\approx$ 100 km s$^{-1}$. Electron densities of $n_\\mathrm{e}$ = 3500-4000 cm$^{-3}$ are inferred from the [FeII] line ratios.} {The large H$_2$ luminosity combined with the very large source luminosity suggests that the high-mass protostar that powers the HH135/HH136 flow forms via accretion, but with a significantly increased accretion rate compared to that of low-mass protostars. } ",
        "introduction": "Near-infrared studies of molecular hydrogen emission in Herbig-Haro (HH) objects provide a powerful tool to gain insight in the physical processes that occur during the early phases of low-mass star formation. The total H$_2$ luminosities are proportional to the accretion rates in the early phases of the protostellar evolution (Stanke \\cite{stanke}; Froebrich et al. \\cite{froebrich}), and the H$_2$ ro-vibrational population distribution allows us to determine the physical scenarios that are at work in the supersonic jets that form during the accretion phase (McCoey \\cite{mccoey}, and references therein).  A complementary tool to study such regions is available via the near-infrared emission lines of [FeII] and other atomic emission lines, that allow us to measure electron densities in the shocked material (e.g., Nisini et al. \\cite{nisini}).  In a number of HH objects, and most prominently in HH objects that show pronounced [FeII] emission, a temperature stratification is inferred from the H$_2$ observations, where part of the molecular gas reach temperatures above 5000~K (Giannini et al. \\cite{giannini04}). The combined H$_2$ and [FeII] emission has been explained in terms of fast J-type shocks, where the [FeII] emission is produced in dissociative parts of the shocks, and where H$_2$ arises in the slower, non-dissociative regions (e.g., Gredel \\cite{gredel94}). Detailed model calculations of J-type shocks with magnetic precursors confirm such a model, and produce population distributions where the rotational excitation temperatures increase with increasing vibrational state of H$_2$ (e.g., Flower et al. \\cite{flower}). C-type shocks at largely different physical conditions produce similar H$_2$ population distributions, however, and the H$_2$ population distribution alone does not allow us to distinguish between both scenarios (Flower et al. \\cite{flower}). The presence of [FeII] emission is generally interpreted in terms of dissociative J-type shocks, although it is conceivable that C-type shocks produce [FeII] as well (Le Bourlot et al. \\cite{lebourlot}).  Near-infrared studies of outflows from intermediate- and high-mass star forming regions are rare, either because such outflows are not very frequent {\\tt per se} or because high-mass star-forming regions are deeply embedded in general. Some outflows that are observed from high-luminous sources such as Orion ($10^5 L_\\odot$) lack the high degree of collimation that is typical for outflows from low-mass star-forming regions. Other flows that emerge from luminous IRAS sources, such as HH80/81 ($2 10^4 L_\\odot$, Mart\\'i, Rodr\\'iguez, \\& Reipurth \\cite{marti}), IRAS 16547-4247 ($6 10^4 L_\\odot$, Brooks et al. \\cite{brooks03}), or IRAS 18151-1208 ($2 10^4 L_\\odot$, Davis et al. \\cite{davis04}), are highly collimated. It is not clear whether intermediate- and high-mass stars form in a similar way to low-mass stars, but with enhanced accretion rates, or whether different processes, such as the merging of low-mass protostars, are at work. Enlarging the sample of near-infrared observations of regions of intermediate- and high-mass star formation and comparing the general properties of their H$_2$ and [FeII] emission with those of low-mass star-forming regions is therefore desirable.  In the following, a study of the molecular outflow \\object{HH135}/\\object{HH136}, which is powered by the cold and very luminous ($10^4 L_\\odot$) IRAS source \\object{IRAS11101-5928}, is presented.  The pair of Herbig-Haro objects HH135/HH136 was discovered by Ogura \\& Walsh (\\cite{ogura}) in an objective prism survey and is located in the bright rimmed cloud No. 64 of Sugitani \\& Ogura (\\cite{sugitani}) in the Eastern Carina region. The HII region, also known as Gum 36, is believed to be excited by the open cluster Stock~13, for which photometric distances of 2.7 kpc are available (Steppe \\cite{steppe}). The general morphology of HH135/HH136 indicates that the two objects are formed at the opposite flow directions of a bipolar flow, which is driven by IRAS11101-5829.  The velocity field of the HH135/HH136 region is complex, however, and the fact that the main part of the emission from both HH135 and HH136 is blue shifted led Ogura \\& Walsh (\\cite{ogura}) to conclude that HH135 and HH136 form two different, independent flows.  Near-infrared JHK polarimetric observations of the associated reflection nebula carried out by Tamura et al. (\\cite{tamura}) showed that IRAS11101--5829 is the only illuminating source of the nebula that is associated with HH135/HH136. A more recent millimetre study by Ogura et al. (\\cite{ogura98}) presented  a model that explains the observed velocity features and where HH135 and HH136 form part of a single, bipolar flow driven by IRAS11101-5829. This view is supported by the very recent polarimetric observations by Chrysostomou et al.  (\\cite{chrysostomou}), who proposed that a strong helical magnetic field threading though HH135/HH136 maintains the strong collimation of the flow.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=12cm]{4746fig1.ps} \\caption{[SII] image of HH135/HH136 and part of the G~36 region in Carina. Optical knots in HH136 are labeled following the notation by Ogura \\& Walsh (\\cite{ogura}). {The crossed square south of knot J marks the position of the IRAS source IRAS11101-5829. } North is up and east is left. } \\label{img_SII} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=-90,width=12cm]{4746fig2.ps} \\caption{Gunn z image of HH135/HH136. } \\label{img_gunnz} \\end{figure}    The purpose of this paper is to investigate the molecular hydrogen emission of HH135/HH136, and to study how the emission characteristics differ from those of low-mass star-forming regions. The observations are presented in Sect.~\\ref{observations}, which also includes a more detailed discussion of the data reduction method and errors. The imaging results are given in Sect.~\\ref{imaging} and the spectroscopy  in Sect.~\\ref{spectroscopy}. Section~\\ref{sumspectra} contains an analysis of the global properties of the molecular hydrogen emission, and Sect.~\\ref{FeII} summarises the results of an analysis derived from the [FeII] lines, which are observed towards HH135/HH136. The main conclusions of this study are summarized in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} The observations presented above and the main conclusions are summarized as follows:  \\begin{enumerate} \\item The images obtained in the (1,0) S(1) line of molecular hydrogen reveal the presence of a well-collimated molecular outflow that extents over a scale of about 1 pc. \\item {A number of the H$_2$ emission line knots are associated with faint, underlying continuum emission. } \\item The ro-vibrational excitation temperatures of H$_2$ of the various knots in the HH135/HH136 flow are remarkably constant, and are well characterized by a narrow range of $T_\\mathrm{ex} = 2000 \\pm 200$K. \\item The molecular part of the shocked gas contains a small fraction of some 0.3\\% of hot H$_2$ at a ro-vibrational excitation temperature of $T_\\mathrm{ex} = 5500 \\pm 200$K. \\item Very strong emission lines from [FeII] occur towards various knots in HH135/HH136, and emission from HeI, hydrogen recombination lines, [CI], [SII], and [NII] are present as well. The [FeII] line ratios indicate the presence of a fast J-type shock at a speed of $v_s \\approx 100$ km s$^{-1}$. Electron densities are of the order of $n_e = 3500 - 4000$ cm$^{-3}$ and electron temperatures are $T_e \\approx 3000$~K.  A more comprehensive analysis of the atomic emission is hampered by the low spectral resolution and the occurrence of a very large number of line blends. \\item The ro-vibrational population distribution of H$_2$, together with the presence of strong [FeII] emission, which is spatially displaced from the H$_2$ emission, indicate that the emission lines arise form in the cooling regions of fast, dissociative J-type shocks, where the [FeII] emission traces the fast, dissociative parts of the shocks and where the H$_2$ emission emerges from regions of oblique shocks where the shock speeds are lower. The presence of emission from re-forming H$_2$ molecules in the cooling post-shock region is not ruled out. \\item The large H$_2$ luminosity of 2 $L_\\odot$ suggests that the intermediate- to high-mass protostar that powers the HH135/HH136 outflow forms via a significantly increased accretion rate, compared to the accretion rates of low-mass protostars.  \\end{enumerate}"
    },
    "0606/astro-ph0606191_arXiv.txt": {
        "abstract": "{In this paper we consider a large sample of optically selected clusters, in order to elucidate the physical reasons for the existence of X-ray underluminous clusters. For this purpose we analyze the correlations of the X-ray and optical properties of a sample of 137 spectroscopically confirmed Abell clusters in the SDSS database.  We search for the X-ray counterpart of each cluster in the ROSAT All Sky Survey. We find that 40\\% of our clusters have a marginal X-ray detection or remain undetected in X-rays. These clusters appear too X-ray faint on average for their velocity disperiosn determined mass, i.e. they do not follow the scaling relation between X-ray luminosity and virial mass traced by the other clusters. On the other hand, they do follow the general scaling relation between optical luminosity and virial mass. We refer to these clusters as the X-ray-Underluminous Abell clusters (AXU clusters, for short) and designate as 'normal' the X-ray detected Abell systems. We examine the distributions and properties of the galaxy populations of the normal and the AXU clusters, separately. The AXU clusters are characterized by leptokurtic (more centrally concentrated than a Gaussian) velocity distribution of their member galaxies in the outskirts ($1.5 < r/r_{200} \\leq 3.5$), as expected for the systems in accretion. In addition, the AXU clusters have a higher fraction of blue galaxies in the external region and show a marginally significant paucity of galaxies at the center. Our results seem to support the interpretation that the AXU clusters are systems in formation undergoing a phase of mass accretion. Their low X-ray luminosity should be due to the still accreting Intracluster gas or to an ongoing merging process.} \\authorrunning{P. Popesso et al.} ",
        "introduction": "Clusters of galaxies are extremely important astrophysical tools. They are the most massive gravitationally bound systems in the universe. Since they sample the high mass end of the mass function of collapsed systems, they can be used to provide tight constraints on cosmological parameters such as $\\Omega_m$, $\\sigma_8$ and $\\Lambda$ (Eke at al 1996, Donahue \\& Voit 1999). Moreover they are extremely powerful laboratories to study galaxy formation and evolution. To investigate the global properties of the cosmological background it is necessary to construct and study a large sample of clusters (Borgani \\& Guzzo 2001).  Several techniques exist to build cluster samples, each based on different clusters properties. The first attempts at a large, homogeneous survey for galaxy clusters was conducted by Abell (1958) with the visual identification of clusters on the Palomar Observatory Sky Survey (POSS) photographic plates. Similar cataloges were constructed by Zwicky and collaborators (Zwicky et al. 1968). Since then, a large number of optically selected samples have been constructed with automated methods: EDCC (Edimburgh Durham Cluster Catalogue: Lumdsen et al. 1992), APM (Automatic Plate measuring; Dalton et al. 1994), PSCS (Palomar Distant Cluster Survey; Postman et al. 1996), EIS (ESO Imaging Cluster Survey; Olsen et al. 1999), ENACS (ESO Nearby Abell Cluster Survey, Katgert et al. 1996, Mazure et al. 1996), RCS (Red sequence Cluster Survey; Gladders \\& Yee 2000) and the samples derived from the Sloan Digital Sky Survey (Goto at al. 2002; Bahcall et al. 2003). The advantage of using optical data is that in general it is relatively easy to build large optically selected cluster catalogs, which allow one to investigate cluster properties with a statistically solid data-base. On the other hand, the main disadvantage of the optical selection is that the selection procedure can be seriously affected by projection effects. Only a very observationally expensive spectroscopic campaign can confirm the overdensities in 3 dimensions.  In 1978, the launch of the first X-ray imaging telescope, the \\emph{Einstein} observatory, began a new era of cluster discovery, as clusters proved to be luminous ($\\ge 10^{42-45}$ ergs $\\rm{s}^{-1}$), extended ($\\rm{r} \\ga 1$ Mpc) X-ray sources, readily identified in the X-ray sky.  Therefore, X-ray observations of galaxy clusters provided an efficient and physically motivated method of identification of these structures.  The X-ray selection is more robust against contamination along the line-of-sight than traditional optical methods, because the X-ray emission, unlike galaxy overdensities, is proportional to the square of the (gas) density. The ROSAT satellite with its large field of view and better sensitivity, allowed to a leap forward in the X-ray cluster astronomy, producing large samples of both nearby and distant clusters (Castander et al. 1995; Ebeling et al. 1996a, 1996b; Scharf et al. 1997; Ebeling et al. 2000; B\\\"ohringer et al. 2001; Gioia et al. 2001; B\\\"ohringer et al. 2002; Rosati et al. 2002 and references therein). The disadvantage of X-ray cluster surveys is their lower efficiency and higher observational cost as compared to optical surveys.  It is clear that understanding the selection effects and the biases due to the different cluster selection techniques is crucial for interpreting the scientific results obtained from such different cluster samples. Castander et al. (1994) used ROSAT to observe cluster candidates in the redshift range 0.7-0.9 from the 3.5 square degree subsample of Gunn et al.'s (1986) optical cluster catalog and found surprisingly weak X-ray emission. Bower et al. (1994) undertook ROSAT X-ray observations of optically selected clusters from Couch et al.'s (1991) 46 $\\rm{deg}^2$ catalog. Bower et al. (1994) selected a random subset of the full catalogue, in the redshift range 0.15--0.66.  The X-ray luminosity of almost all the selected clusters was found to be surprisingly low, suggesting, on the one hand, substantial evolution of the X-ray luminosity function between redshift $z=0$ and $z \\sim 0.4$, and, on the other hand, overestimated velocity dispersions for the nearby X-ray underluminous clusters, perhaps as a consequence of the contamination by galaxy filaments and of radial infall of field galaxies into the clusters.  Similar results were obtained by Holden et al. (1997).  With the ROSAT Optical X-ray Survey (ROXS), Donahue et al. (2002) concluded that there is little overlap between the samples of X-ray-selected and optically-selected galaxy clusters.  Only $\\sim 20$\\% of the optically selected clusters were found in X-rays, while $\\sim 60$\\% of the X-ray clusters were also identified in the optical sample. Furthermore, not all of the X-ray detected clusters had a prominent red-sequence, something that could introduce a selection bias in those cluster surveys based on colour information (Goto et al. 2002, Gladders \\& Yee 2000). Ledlow at al. (2003) analyzed the X-ray properties of a sample of nearby bright Abell clusters, using the ROSAT All-Sky Survey (RASS). They found an X-ray detection rate of 83\\%. Gilbank at al. (2004) explored the biases due to optical and X-ray cluster selection techniques in the X-ray Dark Cluster Survey (XDCS). They found that a considerable fraction of the optically selected clusters do not have a clear X-ray counterpart, yet spectroscopic follow-up of a subsample of X-ray underluminous systems confirmed their physical reality. Lubin et al. (2004) analyzed the X-ray properties of two optically selected clusters at $z \\ge 0.7$, with XMM-Newton. They found the two clusters are characterized by X-ray luminosities and temperatures that are too small for their measured velocity dispersion. Similar results were obtained in the XMM-2dF Survey of Basilakos et al. (2004). They found many more optical cluster candidates than X-ray ones. Deeper XMM data confirmed that their X-ray undetected cluster candidates have intrinsically very low X-ray luminosities.  In this paper we consider a large sample of optically- and X-ray-selected clusters, in order to elucidate the physical reasons for the existence of underluminous optical/X-ray clusters.  The starting point of this work is the analysis we conducted on a sample of X-ray selected clusters sample (Popesso et 2005a, Paper III of this series).  90\\% of those systems are taken from the REFLEX and NORAS catalogs, which are X-ray flux-limited cluster catalogs entirely built upon the ROSAT-All-Sky Survey (RASS). The remaining 10\\% of that sample are groups or faint clusters with X-ray fluxes below the flux limits of REFLEX and NORAS. In Paper III  we found an optical counterpart for each of the X-ray selected clusters of the RASS. Using Sloan Digital Sky Survey (SDSS, see, e.g., Abazajian et al.  2003) optical data for these clusters, we then studied the scatter of the correlations between several optical and X-ray cluster properties (X-ray and optical luminosities, mass, velocity dispersion and temperature). In this paper we extend our analysis to a sample of {\\em optically} selected clusters.  The paper is organized as follows. In section 2 we describe the data and the sample of optically selected clusters used for the analysis. We also describe how we measure the optical luminosity, the velocity dispersion, the mass and the X-ray luminosity of the clusters. In section 3 we analyze the correlation of both the X-ray and the optical cluster luminosities with their masses. In section 4 we describe the optical properties of the Abell clusters without clear X-ray detection and compare them with those of normal X-ray emitting Abell systems. In section 6 we discuss our results and give our conclusions.  We adopt a Hubble constant $\\rm{H}_0=70 \\; \\rm{h} \\; \\rm{km} \\; \\rm{s}^{-1} \\; \\rm{Mpc}^{-1}$, and a flat geometry of the Universe, with $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout this paper. ",
        "conclusions": "We have studied the X-ray and optical properties of 137 isolated Abell clusters. Each object has a confirmed three-dimensional overdensity of galaxies. We have looked for the X-ray counterpart of each system in the RASS data. Three classes of objects have been identified, where the classification is based on the quality of the X-ray detection. 86 clusters out of the 137 Abell systems have a clear X-ray detection and are considered normal X-ray emitting clusters (the 'normal Abell clusters'). 27 systems have a X-ray detection of low significance (less the $3\\sigma$) and 24 do not have clear X-ray detection (a rough estimate of $L_X$ is provided but with huge statistical errors).  The normal Abell clusters follow the same scaling relations observed in the X-ray selected RASS-SDSS clusters. The $24+27$ Abell clusters with unsecure X-ray detection appear to be outliers in the $L_X-M_{200}$ relation determined for X-ray luminous clusters. Their X-ray luminosity is on average one order of magnitude fainter than would be expected for their mass . A careful analysis of the 3D galaxy overdensity of these systems reveals that the individual galaxy velocity distributions in the virial region are gaussian in 90\\% of the clusters and are not ascribable to the superposition of smaller interacting systems. We conclude that these Abell cluster with unsecure X-ray detection in RASS are not spurious detections in the redshift distribution, but are a distinct class of objects. Due to their location with regard to the RASS-SDSS $M-L_X$ relation we call them 'Abell X-ray underluminous clusters' or AXU clusters for short. Several AXU clusters are confirmed to be very faint X-ray objects in the literature. Their X-ray flux is probably too low to be detected in the RASS survey.  Yet, AXU clusters are not outliers from the $L_{op}-M_{200}$ relation, i.e. they have a normal optical luminosity given their mass. Hence, the distinctive signature of AXU clusters seems to lie in an X-ray luminosity which is unexpectedly low.  We have looked for other properties of AXU clusters that make them different from normal Abell clusters. We have shown that AXU clusters do not have more substructures than normal Abell clusters. The galaxy luminosity functions within the virial region of the two cluster samples are very similar to each other. Rather similar are their galaxy number density profiles, even if the AXU clusters seem to lack galaxies near the core, relative to normal Abell clusters (but the significance of this result is low). The fractions of blue galaxies in the two kinds of clusters are only marginaly different, AXU clusters being characterized by a higher fraction.  The main difference between the two classes of objects lies in the velocity distribution of their member galaxies. The galaxy velocity distribution of the normal Abell clusters is perfectly fitted by a Gaussian both in the inner, virialized region ($\\le 1.5 \\, r_{200}$), and also in the external region ($1.5 \\, r_{200} \\le r \\le 3.5 \\, r_{200}$). The AXU clusters instead have a Gaussian velocity distribution only within the virial region. In the external region, their velocity distribution is significantly more peaked than a Gaussian. The analysis of its shape by comparison with dynamical models available in the literature (van der Marel et al. 2000), suggests a radially anisotropic galaxy orbital distribution. However, the galaxies in this external region need not be in dynamical equilibrium with the cluster potential. As a matter of fact, a leptokurtic shape of the velocity distribution is a typical signature of the external, infall regions of dark matter haloes (Wojtak et al. 2005).  The analysis of the velocity distribution of the AXU clusters in their outer regions hence suggests the presence of an unvirialized component of the galaxy population, still in the process of accretion onto the cluster. This infalling population would be mainly composed of field, hence blue, galaxies, which could then explain the excess of blue galaxies in AXU clusters, relative to normal Abell clusters.  On the other hand, the Gaussian velocity distribution in the inner region suggests that there the galaxy population is dynamically more evolved, and probably virialized.    By a similar analysis on a different sample of X-ray underluminous clusters, Bower et al. (1997) came to propose two different scenarios. AXU clusters could be severely affected by projection effects arising from surrounding large-scale structure filaments elongated along the line-of-sight.  Their velocity dispersion, and hence their virial masses would then be severely overestimated by interlopers in the filaments. In the alternative scenario AXU clusters could be clusters not yet formed, or in the phase of forming, or, at least, caught at a particular stage of their evolution, while they are undergoing a rapid mass growth.   Should the former of the two scenarios apply, we would expect AXU clusters to be X-ray underluminous for their mass, but they could still be optically luminous because we partly see the light of the filament projected onto the cluster. However, contamination by interlopers does affect the optical luminosity estimate, but not so much as the virial mass estimate, and not so much in the $i$ band, where contamination by the field (hence blue) galaxies should be small. Therefore, in such scenario it would be surprising that the clusters obey so well the $L_{op}-M_{200}$ relation, which requires that the effects of the filament on the dynamical mass estimate and the optical light in the aperture both conspire not to produce an offset from the relation. It would also be surprising that the AXU clusters show a galaxy LF perfectely consistent with the steep LF found in galaxy clusters (see Popesso et al. paper II and IV) and not the flat LF observed in the field (Blanton et al. 2005). Instead AXU clusters are not outliers from the $L_{op}-M_{200}$ relation. If anything, AXU clusters are overluminous in the optical for their mass. In fact, the biweight-average (see Beers et al. 1990) $i$-band mass-to-light ratios of normal Abell clusters and AXU clusters are $150 \\pm 10 \\, M_{\\odot}/L_{\\odot}$, and $110 \\pm 10 \\, M_{\\odot}/L_{\\odot}$, respectively.    As a further test, we have re-calculated the virial masses of all clusters by considering only red cluster members belonging to the red sequence in the $u-i$ vs. $i$ color-magnitude diagram. In this way contamination by interlopers is strongly reduced (see, e.g., Biviano et al. 1997).  Masses computed using all cluster members are compared to masses computed using only red-sequence members in Fig. \\ref{lowlx}. The cluster masses do not change significantly when only red-sequence members are used to calculate them, suggesting a low level of contamination by interlopers.  The results of our analyses therefore support Bower et al.'s alternative scenario, namely AXU clusters are systems in the stage of formation and/or of significant mass accretion.  If AXU clusters are still forming, the intra-cluster gas itself may still be infalling or have not yet reached the virial temperature. In addition, for AXU clusters undergoing massive accretion, it is to some degree possible that the continuous collisions of infalling groups is affecting the gas distribution, lowering its central density (such as in the case of the so called 'bullet cluster', see Barrena et al. 2002 and Clowe et al. 2004). In both cases the X-ray luminosity would be substantially lower than predicted for the virial mass of the system, because of its dependence on the square of the gas density. We note however that a virialized cluster undergoing a strong collision with an infalling group would show up as a substructured cluster, yet the AXU clusters do not show an increased level of substructures when compared to normal Abell clusters. In summary, we know that the X-ray emission is very much dominated by the central region whereas the optical properties are more global.  Therefore it could well be that we see a rough relaxation on the large scale (within $1.5r_{200}$) of the galaxy system reflected by a rough Gaussian galaxy velocity distribution, while the central region has not yet settled to reach the high density and temperatures of the luminous X-ray clusters.  In order to explore this further, we need much more detailed information on the distribution of the density and temperature of the intracluster gas in AXU clusters, something that cannot be done with the RASS data, but requires the spatial resolution and sensitivity of XMM-Newton.   Our results give supports to the conclusion of Donahue et al. (2002) concerning the biases inherent in the selection of galaxy clusters in different wavebands. While the optical selection is prone to substantial projection effects, also the X-ray selection is not perfect or not simple to characterize. The existence of X-ray underluminous clusters, even with large masses, makes it difficult to reach the needed completeness in mass for cosmological studies. Moreover, as discussed in Paper III, the relation between the X-ray luminosity and mass is not very tight even for the X-ray bright clusters, and the relation between cluster masses and optical luminosities is as tight or perhaps even tighter. Clearly, a multi-waveband approach is needed for optimizing the completeness and reliability of clusters samples.  On the other hand, it becomes clear that for precision cosmology we also need a more observationally oriented prescription of cluster selection from theory, rather than a mere counting of \"relaxed\" dark matter halos.  Predicted distribution functions closer to the observational parameters like temperature or velocity dispersion distribution functions and their relations to X-ray and optical luminosity are needed.    \\vspace{2cm} We thank the referee F. Castander for the useful comments, which helped in improving the paper. We thank Alain Mazure for useful discussion.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P.  Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S.  Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.  The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0606/astro-ph0606644_arXiv.txt": {
        "abstract": "We present high-resolution, high signal to noise absorption-line observations of CN, \\ion{Ca}{2}, \\ion{Ca}{1}, CH$^+$, and CH along twenty lines of sight toward members of the Pleiades. The acquired data enable the most detailed study to date of the interaction between cluster stars and the surrounding interstellar gas. Total equivalent widths are consistent with previous investigations except where weaker features are detected owing to our greater sensitivity. Mean $b$-values for the molecular species indicate that toward most of the Pleiades CH is associated with the production of CH$^+$ rather than CN. An analysis of radial velocities reveals a kinematic distinction between ionized atomic gas and molecular and neutral gas. Molecular components are found with velocities in the local standard of rest of either $\\sim$ +7 km s$^{-1}$ or $\\sim$ +9.5 km s$^{-1}$, with the higher-velocity components associated with the strongest absorption. Atomic gas traced by \\ion{Ca}{2} shows a strong central component at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +7 km s$^{-1}$ exhibiting velocity gradients indicative of cloud-cluster interactions. Gas density estimates derived from measured CH/CH$^+$ column density ratios show good agreement with those inferred from H$_2$ rotational populations, yielding typical values of $n\\sim50$ cm$^{-3}$. Our models do not include the important time-dependent effects on CH$^+$ formation which may ultimately be needed to extract physical conditions in these clouds. ",
        "introduction": "The interstellar medium (ISM) in the vicinity of the Pleiades is a rich environment for the study of processes that result from the interactions between stellar photons and the gas and dust clouds of interstellar space. The stars of this cluster were not formed out of the surrounding material visible as reflection nebulosity. Rather, the spatial association of the stars and the interstellar gas is the result of a chance encounter between the cluster and one or more approaching clouds (White 2003). Such collisions precipitate numerous radiative processes including the photoionization of atomic and molecular species, the photodissociation of molecules such as CH and H$_2$, and the photoelectric heating of diffuse gas by dust grains stimulated by ultraviolet radiation, all of which may help to explain some peculiarities of the ISM near the Pleiades.  Two such peculiarities are the great strength of CH$^+$ absorption lines observed toward many cluster members (White 1984\\emph{a}) and the large amount of H$_2$ in rotationally excited states (Spitzer, Cochran, \\& Hirshfeld 1974). While these anomalies may be related, a rigorous chemical model of CH$^+$ production in the Pleiades, and elsewhere in the ISM, remains to be developed. CH$^+$ is unable to form at the low temperature of diffuse clouds because the reaction leading to its formation, C$^+$ + H$_2$ $\\to$ CH$^+$ + H, is endothermic with an activation energy of $\\Delta E/k=4640$ K. Elitzur and Watson (1978, 1980) showed that the above reaction can produce sufficient amounts of CH$^+$ if the gas is heated by an interstellar shock. However, subsequent observations failed to detect a corresponding overabundance of OH molecules resulting from a similar endothermic reaction (Federman et al. 1996\\emph{a}). Several groups (Draine \\& Katz 1986; Pineau des For\\^ets et al. 1986) employed magnetohydrodynamic (MHD) shocks to lessen the problem with the overproduction of OH, but even their models yielded OH column densities in excess of observations. The velocity shifts between CH and CH$^+$ absorption lines predicted by these shock models have also not been detected (e.g., Gredel, van Dishoeck, \\& Black 1993). Alternate theories for a non-thermal origin of CH$^+$ have since been proposed. Federman et al. (1996\\emph{b}) considered the motions of C$^+$ ions influenced by the passage of Alfv\\'en waves against a static background of cold neutral gas. Their model predicts CH$^+$, CH, and OH column densities in line with observations, but requires an additional mechanism to account for observed HCO$^+$ abundances. More recently, non-equilibrium chemistry was investigated by Joulain et al. (1998) and Falgarone et al. (2005). These authors ascribe the transient heating of localized regions of the cold diffuse medium to intermittent bursts of turbulent dissipation by either MHD shocks or coherent small-scale vortices. Such models may be appropriate for many Galactic sight lines, yet the observational evidence in the Pleiades favors heating either by H$_2$ dissociation or by dust photoelectron emission, rather than shocks, as the agent responsible for CH$^+$ formation (White 1984\\emph{b}).  \\tabletypesize{\\scriptsize} \\begin{deluxetable*}{lcccccccc} \\tablecolumns{9} \\tablewidth{0.8\\textwidth} \\tablenum{1} \\tablecaption{Stellar and Observational Data} \\tablehead{\\colhead{HD} & \\colhead{Name} & \\colhead{Type} & \\colhead{$E(B-V)$} & \\colhead{$B$} & \\colhead{$\\alpha$ [J2000]} & \\colhead{$\\delta$ [J2000]} & \\colhead{$d$} & \\colhead{$\\tau_{exp}$} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{($^h$, $^m$, $^s$)} & \\colhead{($^{\\circ}$, $^{\\prime}$, $^{\\prime\\prime}$)} & \\colhead{(pc)} & \\colhead{(s)} } \\startdata 23288 & 16 Tau & B7 IV & \\phs 0.10 & 5.41 & 03 44 48.22 & +24 17 22.1 & 103 $\\pm$ 11 & $3 \\times 1200$ \\\\ 23302 & 17 Tau & B6 III & \\phs 0.05 & 3.61 & 03 44 52.54 & +24 06 48.0 & 114 $\\pm$ 12 & $2 \\times \\phn 600$ \\\\ 23324 & 18 Tau & B8 V & \\phs 0.05 & 5.58 & 03 45 09.74 & +24 50 21.3 & 113 $\\pm$ 11 & $2 \\times 1800$ \\\\ 23338 & 19 Tau & B6 IV & \\phs 0.04 & 4.20 & 03 45 12.49 & +24 28 02.2 & 114 $\\pm$ 14 & $3 \\times \\phn 600$ \\\\ 23408 & 20 Tau & B7 III & \\phs 0.07 & 3.81 & 03 45 49.61 & +24 22 03.9 & 110 $\\pm$ 13 & $2 \\times \\phn 600$ \\\\ 23410 & & A0 V & \\phs 0.08 & 6.92 & 03 45 48.82 & +23 08 49.7 & 103 $\\pm$ 11 & $4 \\times 1800$ \\\\ 23432 & 21 Tau & B8 V & \\phs 0.07 & 5.73 & 03 45 54.48 & +24 33 16.2 & 119 $\\pm$ 13 & $2 \\times 1800$ \\\\ 23441 & 22 Tau & B9 V & \\phs 0.06 & 6.42 & 03 46 02.90 & +24 31 40.4 & 109 $\\pm$ 11 & $2 \\times 1800$ \\\\ 23480 & 23 Tau & B6 IV & \\phs 0.10 & 4.11 & 03 46 19.57 & +23 56 54.1 & 110 $\\pm$ 13 & $3 \\times \\phn 600$ \\\\ 23512 & & A1 V & \\phs 0.35 & 8.47 & 03 46 34.20 & +23 37 26.5 & $\\ldots$ & $8 \\times 1800$ \\\\ 23568 & & B9.5 V & \\phs 0.07\\tablenotemark{a} & 6.85 & 03 46 59.40 & +24 31 12.4 & 150 $\\pm$ 23 & $4 \\times 1800$ \\\\ 23629 & 24 Tau & A2 V & $-0.03$ & 6.30 & 03 47 21.04 & +24 06 58.6 & $\\ldots$ & $3 \\times 1800$ \\\\ 23630 & $\\eta$ Tau & B7 III & \\phs 0.04 & 2.81 & 03 47 29.08 & +24 06 18.5 & 113 $\\pm$ 13 & $2 \\times \\phn 300$ \\\\ 23753 & & B8 V & \\phs 0.04 & 5.38 & 03 48 20.82 & +23 25 16.5 & 104 $\\pm$ 10 & $2 \\times 1200$ \\\\ 23850 & 27 Tau & B8 III & \\phs 0.04 & 3.54 & 03 49 09.74 & +24 03 12.3 & 117 $\\pm$ 14 & $2 \\times \\phn 600$ \\\\ 23862 & 28 Tau & B8 Vp & \\phs 0.03 & 4.97 & 03 49 11.22 & +24 08 12.2 & 119 $\\pm$ 12 & $3 \\times \\phn 900$ \\\\ 23873 & & B9.5 V & \\phs 0.02\\tablenotemark{a} & 6.59 & 03 49 21.75 & +24 22 51.4 & 125 $\\pm$ 14 & $3 \\times 1800$ \\\\ 23923 & & B8 V & \\phs 0.06\\tablenotemark{a} & 6.12 & 03 49 43.53 & +23 42 42.7 & 117 $\\pm$ 13 & $2 \\times 1800$ \\\\ 23964 & & B9.5 Vp & \\phs 0.11 & 6.80 & 03 49 58.06 & +23 50 55.3 & 159 $\\pm$ 39 & $4 \\times 1800$ \\\\ 24076 & & A2 V & \\phs 0.03\\tablenotemark{a} & 7.02 & 03 50 52.43 & +23 57 41.3 & 102 $\\pm$ 10 & $4 \\times 1800$ \\\\ \\enddata \\tablenotetext{a}{Reddening values for these stars were updated in White (2003).} \\end{deluxetable*}  \\begin{figure} \\centering \\includegraphics[width=0.4\\textwidth]{f1.eps} \\caption[Distribution of observed Pleiades stars.]{Distribution of observed Pleiades stars in relative equatorial coordinates centered on $\\eta$ Tau [$\\alpha(2000)=03^h47^m29^s.08$, $\\delta(2000)=+24^{\\circ}06^{\\prime}18^{\\prime\\prime}.5$]. The label next to each symbol is either the last three digits of the HD number or the Flamsteed or Bayer designation. Decreasing symbol size denotes three magnitude ranges: $B < 5.0$, $5.0 < B < 6.5$, and $6.5 < B$. This key is used in Figures 6, 7, and 8.} \\end{figure}  Precise measurements of CN, CH$^+$, and CH column densities in the Pleiades interstellar gas can offer new insight into the chemical reaction networks active in these diffuse clouds. At the same time, highly sensitive observations of \\ion{Ca}{2} and, to a lesser extent, \\ion{Ca}{1} can be used to trace the kinematic effects of the interaction occurring between interstellar gas and the stellar radiation field in great detail and on large spatial scales. The earlier investigation by White (1984\\emph{a}), comprising spectra for 15 Pleiades members, analyzed interstellar absorption from the above atomic and molecular species, but the moderate velocity resolutions ($\\sim$ 3$-$8 km s$^{-1}$) allowed detections of only one component per sight line. Later studies capable of detecting many line-of-sight components, such as the ultra-high resolution survey by Crane et al. (1995) of interstellar CH$^+$ and CH and the high-resolution survey by Welty et al. (1996) of \\ion{Ca}{2}~K, included only a few of the brightest Pleiades members in their sample. Most recently, high-resolution observations were made by White et al. (2001) of 36 Pleiades stars, both members and nonmembers, in the \\ion{Na}{1}~D lines and of 12 of these stars in the \\ion{Na}{1} ultraviolet doublet. The full analysis of these data (White 2003) revealed considerable complexity and star-to-star variation in the atomic gas traced by neutral sodium, leading to an extensive spatial schematic of cloud-cluster interactions in which two clouds were presumed to be interacting with the UV radiation field of the Pleiades and also with each other. One of the primary motivations for the present investigation was to obtain equally high quality data for the other important optical tracers of the ISM toward a large number of targets in the Pleiades so that a complete picture of the interaction between interstellar gas clouds and the stars of the cluster may be constructed.  In this paper, high-resolution, high signal to noise observations of CN, \\ion{Ca}{2}, \\ion{Ca}{1}, CH$^+$, and CH toward 20 Pleiades members allow us to study cloud-cluster interactions and the implications for the chemistry of the local ISM with a precision unmatched by previous investigations. We describe our observations and detail the process of data reduction in \\S{} 2. In \\S{} 3, we present the results of Gaussian fitting the observed profiles and compare our measurements to those from the literature. The analysis of our observations appears in \\S{} 4, with particular attention paid to the spatial distribution of atomic and molecular velocity components (\\S{} 4.1). We examine \\ion{Na}{1}/\\ion{Ca}{2} column density ratios in \\S{} 4.2 and derive physical conditions in \\S{} 4.3. In \\S{} 5, we interpret our findings in light of the conceptual schematic of cloud-cluster interactions offered by White (2003) and summarize our principal results in \\S{} 6.  \\begin{figure*} \\centering \\includegraphics[width=0.7\\textwidth, angle=90]{f2.eps} \\caption[Spectra of Ca~{\\scriptsize II}.]{Observed Ca~{\\scriptsize II} spectra arranged to approximate the spatial relationships among sight lines. Spectra toward the left (right) of the figure correspond to sight lines in the eastern (western) part of the cluster. Likewise, spectra toward the top (bottom) of the figure correspond to sight lines in the north (south). See Figure 1 for the actual distribution of sight lines. The vertical tick marks indicate velocity components.} \\end{figure*} ",
        "conclusions": "In \\S{} 4.1, the relatively straightforward distribution of molecular gas in the Pleiades as traced by the molecules of CH$^+$, CH, and CN is contrasted with the rather complicated structures seen in the ionized atomic gas traced by \\ion{Ca}{2}. The largest columns of molecular gas are associated with a velocity component at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +9.5 km s$^{-1}$ concentrated in the western region of the cluster. Clearly, some of this material is related to the molecular cloud seen along the line of sight toward HD 23512, particularly that detected in CH and CN absorption in this direction. Indeed, the velocities of the CH and CN components here are consistent with the cloud velocity determined from molecular line emission by Federman \\& Willson (1984) and, although the density of the cloud derived from our CH and CN column densities ($n > 1600$ cm$^{-3}$) is significantly larger than their predicted value ($\\sim$ 300$-$500 cm$^{-3}$), the inclusion of dark-cloud chemistry should correct this discrepancy. Still, whether or not all molecular components that share the cloud velocity have a common origin remains elusive. The general conclusions of the expansive study by White (2003) favor a three body encounter in which gas associated with the Taurus dust clouds and redshifted by the interaction with the Pleiades accounts for some of the red components and a separate cloud approaching from the west and also interacting with the cluster accounts for others. In support of the latter, the strongest components of CH$^+$ absorption are located northwest of the cluster center, with a position and orientation similar to that of the western cloud described by White (2003). The absence of any corresponding CH components in this region indicates the low density of the material. Gas toward 17 Tau and 23 Tau with a velocity of +10 km s$^{-1}$ likely probes the low-density envelope of the molecular cloud which lies just to the south of these sight lines. The rest of the CH$^+$ components near +9 km s$^{-1}$ may also trace this envelope, but the data do not preclude a connection with redshifted Taurus material.  The other, weaker constituent of molecular gas in the Pleiades has a velocity of $\\sim$ +7 km s$^{-1}$, associating it with the central component of atomic gas and with the Taurus clouds. All sight lines with detectable amounts of CH, except that toward HD 23512, and the majority of CH$^+$ sight lines exhibit this component. It is found in both species along sight lines to three stars in the center of the cluster ($\\eta$ Tau, 27 Tau, and 28 Tau) and these show similar CH/CH$^+$ column density ratios (0.2 for $\\eta$ Tau and 0.1 for 27 Tau and 28 Tau). Toward HD 23512, where both species are detected at $\\sim$ +9.5 km s$^{-1}$, the ratio is much larger (0.9). That the column densities of CH and CH$^+$ are well correlated for the central component strengthens the claim that CH toward most of the Pleiades is CH$^+$-like. The larger ratio toward HD 23512 results from the addition of CN-like CH to the total CH column and reinforces the idea that this sight line effectively probes the molecular cloud.  The patterns exhibited in the velocity components of \\ion{Ca}{2} (Figure 8) offer unique insight into the complex interactions between the ISM and the stars of the Pleiades. For instance, the central component at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +7 km s$^{-1}$ seen toward all of the stars in our survey does not have a uniform velocity across the cluster. The slight blueshifts and redshifts apparent in Figures 8\\emph{d} and 8\\emph{f}, respectively, are a strong indication that a dynamical interaction is occurring in which cloud material passing through the UV radiation field of the cluster is deflected toward and away from our line of sight. That the pattern seems to repeat itself at larger and smaller velocities (Figures 8\\emph{c} and 8\\emph{g}) suggests that the encounter of this central component can account for at least some of the observed red and blue components. Gas detected at velocities higher than $\\sim$ +10 km s$^{-1}$ likely derives from multiple sources. The feature seen in Figure 8\\emph{h} may contain components missing from the gas in Figure 8\\emph{g} which have been further redshifted by cluster interactions, but may also contain independent components from a chance meeting with a second approaching cloud. The latter is the probable source of the extreme-velocity components in Figure 8\\emph{i} due to the proximity of these sight lines to the molecular cloud. The shocked atomic components in Figures 8\\emph{a} and 8\\emph{b} show additional indications of a velocity gradient from the center of the cluster to the northwestern edge, despite the noticeable gap toward 17 Tau, 20 Tau, and 22 Tau. In general, the various components of \\ion{Ca}{2} are more pervasive in number and extent than those in the analysis of \\ion{Na}{1} observations by White (2003), as expected if \\ion{Ca}{2} is the most widely distributed species as in Figure 6 of Pan et al. (2005). Clearly, highly sensitive observations of \\ion{Ca}{2} can trace the intricate cloud-cluster and cloud-cloud interactions of the Pleiades in greater detail and on a larger scale than \\ion{Na}{1} observations of equal quality.  An important result of our analysis of \\ion{Na}{1}/\\ion{Ca}{2} ratios (\\S{} 4.2) was that they predicted the proper relationship between velocity components and the expected physical conditions of the associated clouds, strengthening our confidence in the column densities derived from our observations. \\ion{Na}{1}/\\ion{Ca}{2} ratios of order a few to 10 were measured for most positive-velocity components indicating the diffuse nature of much of the interstellar gas near the cluster where calcium is not efficiently depleted onto the surfaces of dust grains. A ratio of order 100 was observed in the +9.2 km s$^{-1}$ component toward HD 23512 since this gas lies in a denser region near the molecular cloud, an environment which results in higher calcium depletion. Final confirmation comes from the +6.2 km s$^{-1}$ component toward this star whose \\ion{Na}{1}/\\ion{Ca}{2} ratio of $\\sim$ 10 rightly places it among the diffuse components.  The precise density estimates of \\S{} 4.3 can place meaningful constraints on the chemical processes active in the diffuse clouds of the Pleiades. A comparison of the values obtained from CH/CH$^+$ column density ratios with those from rotationally excited H$_2$ column densities offers particular insight. We find generally good agreement in the densities obtained from the two methods by adopting a molecular fraction of $f=0.07$. This value of $f$ yields typical densities of $\\sim50$ cm$^{-3}$, significantly below earlier predictions which favored densities between 100 cm$^{-3}$ (Gordon \\& Arny 1984) and 400 cm$^{-3}$ (White 1984\\emph{a}, \\emph{b}). A possible discrepancy with our more precise estimates comes from the recent study by Zsarg\\'o and Federman (2003) who inferred a density upper limit of $n \\leq 3.5$ cm$^{-3}$ from the column densities of the fine-structure levels of \\ion{C}{1} toward $\\eta$ Tau. This determination is an order of magnitude lower than our value toward $\\eta$ Tau, $n=62$ cm$^{-3}$, obtained from both CH/CH$^+$ ratios and H$_2$ rotational populations. Yet, based on their density estimate, Zsarg\\'o and Federman (2003) predict a CH column density upper limit toward $\\eta$ Tau of $N$(CH) $\\leq 7.3 \\times 10^{11}$ cm$^{-2}$, consistent with our measured value of $4.8 \\times 10^{11}$ cm$^{-2}$. To derive this predicted column density, the authors incorporated non-thermal motions of ions and neutrals due to the propagation of Alfv\\'en waves in a transition zone between cloud and intercloud material (Federman et al. 1996\\emph{a}). Without incorporating turbulence, the predicted column density along this sight line is $N$(CH) $\\leq 1.6 \\times 10^{9}$ cm$^{-2}$, much lower than the measured value. Non-thermal models hold great promise for diffuse-cloud chemistry, particularly in explaining the abundance of the CH$^+$ radical as well as the abundances of molecules like CH which are tied to its formation.  The question remains, then, why the density predicted by Zsarg\\'o and Federman (2003) is so much lower than our value. Two possibilities immediately present themselves. First, observations of H$_2$, CH, and CH$^+$ may sample different cloud depths due to the stratification of molecular species along the line of sight. This scheme suggests that rotationally excited H$_2$ molecules probe denser portions of the cloud closer to the stars, in the main portion of the photodissociation region (PDR), where the UV radiation is strongest and is responsible for populating the higher $J$ levels through UV pumping. CH and CH$^+$, as well as \\ion{C}{1}, would then trace the less dense outer portions of the cloud away from the PDR. A major drawback to this scenario is it requires $f=0.5$ or higher to yield low enough densities from our measured CH/CH$^+$ column density ratios to agree with the \\ion{C}{1} result, which would subsequently increase our H$_2$-derived densities for this diffuse gas to greater than 400 cm$^{-3}$, a value more typical of the Pleiades molecular cloud. A more suitable alternative is that the fine-structure levels of \\ion{C}{1} trace regions of various extent along the line of sight. If the $J=1$ level, for instance, was confined to a region of space 1/10 the size of that occupied by the $J=0$ level, the density upper limit would be more in line with our determination. Finally, the inclusion of time-dependent effects into a chemical model of CH$^+$ formation may still prove to be important since the timescale for H$_2$ dissociation is comparable to the characteristic time in which the cloud is interacting with the cluster. Because the initiating reaction leading to CH production requires H$_2$, a decreasing presence due to photodissociation would demand a higher density to account for the observed columns of CH. Such a requirement could be accommodated by our simple steady-state model by, again, raising slightly the value of $f$(H$_2$).  The most general conclusion of this investigation is that detailed observations of both atomic and molecular species essentially confirm the scenario of cloud-cluster interactions constructed by White (2003). A pervasive foreground gas cloud at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +7 km s$^{-1}$ seen toward nearly every star in both strong atomic and weaker molecular absorption is flowing through the cluster and dynamically interacting with the stellar UV radiation field. The clearly symmetric patterns of both blueshifted and redshifted \\ion{Ca}{2} components extending across the cluster demonstrate the large-scale nature of these interactions. A second cloud at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +10 km s$^{-1}$ is also likely to be interacting with the cluster due to the strong CH$^+$ absorption near this velocity northwest of the cluster center and the existence of a diffuse molecular cloud also at $v_{\\mathrm{\\scriptstyle LSR}}$ $\\sim$ +10 km s$^{-1}$ to the southwest. High velocity \\ion{Ca}{2} components with no direct connection to redshifted foreground gas may also require the existence of a second interacting cloud."
    },
    "0606/astro-ph0606408_arXiv.txt": {
        "abstract": "We explore a model of interacting dark energy where the dark energy density is related by the holographic principle to the Hubble parameter, and the decay of the dark energy into matter occurs at a rate comparable to the current value of the Hubble parameter. We find this gives a good fit to the observational data supporting an accelerating Universe, and the model represents a possible alternative interpretation of the expansion history of the Universe. ",
        "introduction": "Observations of type Ia supernova indicate that the Universe is making a transition from a decelerating phase to an accelerating one. The conventional explanation for this behavior is that a dark energy component is coming to dominate over and the matter-dominated phase is giving way to a phase dominated by a dark energy component. In fact a good fit is obtained from the observational data by assuming the cosmological model ($\\Lambda$CDM) involving a cosmological constant $\\Lambda$ and cold dark matter in about the ratios $\\Omega_{m,0}=0.3$ and $\\Omega_{\\Lambda,0}=0.7$.  In this paper we show that one can describe the cosmological data with a model of dark energy that includes an interaction that effects a transition (decay) of dark energy into matter. The model incorporates a holographic principle\\cite{'tHooft:1993gx,Susskind:1994vu} to determine the dark energy density of the Universe. The principle relates the dark energy scale to the Hubble horizon which has been the subject of speculation for applying holographic ideas to cosmology\\cite{Easther:1999gk,Veneziano:1999ts,Kaloper:1999tt,Brustein:1999ua,Bak:1999hd,Cohen:1998zx,Thomas:2002pq,Bousso:2002ju}. This is a common choice for imposing holography on cosmology, and it has the most natural thermodynamic interpretation. However, this conditon imposed on the dark energy yields an equation of state that is matter-like\\cite{Hsu:2004ri} and therefore inconsistent with observational data. Subsequently efforts have been made to impose a holographic constraint based on some other physical horizon such as the particle horizon\\cite{Fischler:1998st,Bousso:1999xy} or the future event horizon\\cite{Li:2004rb,Huang:2004wt,Zhang:2005yz,Myung:2005sv,Myung:2005pw,Zhang:2005hs,Kim:2005at,Li:2006ci} or some other physical condition\\cite{Horava:2000tb,Danielsson:2004xw,Horvat:2004vn,Guberina:2005mp,Elizalde:2005ju,Guberina:2006fy,Ke:2004nw}.  A suitable evolution of the Universe is obtained when, in addition to the holographic dark energy, an interaction (decay of dark energy to matter) is assumed, and the decay rate should be set roughly equal to the present value of the Hubble parameter for a good fit to the expansion history of the Universe as determined by the supernova and cosmic microwave background (CMB) data.  Interacting dark energy has been studied previously\\cite{Amendola:1999qq,Amendola:2000uh,Zimdahl:2000zm,Balakin:2003tk,Wang:2004cp,Cai:2004dk,Pavon:2005yx,Zimdahl:2005bk,Wang:2005jx,Wang:2005ph,Pavon:2005kr,Berger:2006db,Hu:2006ar} primarily with a goal of explaining the cosmic coincidence problem. A survey of the possible dynamics of dark energy can be found in Ref.~\\cite{Copeland:2006wr}. In the presence of an interaction the dark energy can achieve a stable equilibrium that differs from the usual de Sitter case or the approach to the stable equilibrium can be made more gradual. Such models offer the hope of solving the coincidence problem that exists in the $\\Lambda$CDM model where there is no obvious reason why the transition from matter domination to domination by the dark energy is occuring during the current epoch.  This paper presents a simple model that gives an acceptable expansion history of the Universe in terms of a holographic condition on the dark energy that relates its size to the Hubble scale. This is a common choice for imposing holography on cosmology, and it has the most natural thermodynamic interpretation. The effective equations of state of matter and dark energy coincide and behave like cold dark matter (CDM) at early times. The transition to behavior like a cosmological constant is effected by simply assuming there is a constant decay of dark energy into matter. The coincidence problem appears in the interacting dark energy model as the choice of the scale for the dark energy interaction which must be close to the present value of the Hubble parameter. ",
        "conclusions": "The $\\Lambda$CDM model provides an explanation for all observational data. However there remain a number of important issues it must confront. Given our lack of understanding of the dark energy, one can ask if there are other simple physical properties that dark energy might have that could equally well account for the data.  In this paper we have shown that a model of dark energy with a holographic condition relating the dark energy to the Hubble parameter and a constant interaction of size roughly equal to the Hubble constant $H_0$ can give a good fit to data. This model is characterized by a constant ratio of $\\Omega _m$ to $\\Omega_\\Lambda$. Since the ratio of matter to dark energy is constant due to the holographic condition and the effective equations of state are equal, one can view the evolution of the model as one comprised of one component whose effective equation of state varies between $0$ when the decay of the dark energy into matter is negligible to an asymptotic value of $-1$ when the interaction become important. The transition between one regime and the other occurs when the interaction rate is comparable to the Hubble parameter.  We have shown that the expansion history of the Universe can be reproduced with a model of interacting dark matter. For the specific model presented here the functional dependence $H(z)$ differs slightly between the $\\Lambda$CDM and the model discussed in this paper. Future data from SNAP may be useful in discriminating between them. Finally information that goes beyond the recorded expansion history of the Universe may be useful for ruling out this kind of model."
    },
    "0606/astro-ph0606314_arXiv.txt": {
        "abstract": "% We present an analysis of optical lightcurves of Small Magellanic Cloud (SMC) Be-type stars. Observations show that (1) optical excess flux is correlated with near-IR excess flux indicating a similar mechanism and (2) the lightcurves can trace out ``loops'' in a colour-magnitude diagram. A simple model for the time dependence of bound-free and free-free (bf-ff) emission produced by an outflowing circumstellar disk gives reasonable fits to the observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606064_arXiv.txt": {
        "abstract": "We have undertaken a multiwavelength project to study the relatively unknown properties of groups and poor clusters of galaxies at intermediate redshifts.  In this paper, we describe the \\textit{XMM-Newton} observations of six X-ray selected groups with $0.2<z<0.6$.  The X-ray properties of these systems are generally in good agreement with the properties of low-redshift groups.  They appear to follow the scaling relations between luminosity, temperature, and velocity dispersion defined by low-redshift groups and clusters.  The X-ray emission in four of the six groups is also centered on a dominant early-type galaxy.  The lack of a bright elliptical galaxy at the peak of the group X-ray emission is rare at low-redshifts, and the other two groups may be less dynamically evolved.  We find indications of excess entropy in these systems over self-similar predictions out to large radii.  We also confirm the presence of at least one X-ray luminous AGN associated with a group member galaxy and find several other potential group AGN. ",
        "introduction": "Most galaxies in the universe are members of groups of galaxies, making groups an important environment for the study of galaxy evolution (e.g. Tully 1987).  In addition to being a more common environment for galaxies, different processes are at work in groups versus rich clusters.  With their relatively low velocity dispersions, groups are ideal sites for galaxy-galaxy mergers (Barnes 1985; Merritt 1985).  Many groups are also found to contain diffuse X-ray emission (e.g. Mulchaey et al. 1993; Ponman \\& Bertram 1993; Mulchaey et al. 2003; Osmond \\& Ponman 2004).  This emission is extended on scales of hundreds of kiloparsecs, and group spectra indicate that the emission mechanism is thermal bremsstrahlung and line emission.  The diffuse group medium is therefore analogous to the intracluster medium and indicates a deep potential well in these systems.  However, here again groups may differ importantly from clusters, as non-gravitational heating and cooling may have a larger effect in groups.  Studies of low-redshift groups have found a steepening of the $L_X-T_X$ relationship in the group regime (Helsdon \\& Ponman 2000 a,b; Ponman et al. 1996).  The label ``group'' is used to describe a very diverse population from loose associations of a few galaxies through poor clusters.  X-ray emission, indicating a deep potential well, is found almost exclusively in groups with a significant fraction of early-type galaxies and generally in groups with a central, dominant early-type galaxy (Mulchaey \\& Zabludoff 1998; Mulchaey et al. 2003; Osmond \\& Ponman 2004).  This observation offers both clues to galaxy evolution and a possible connection to rich clusters.  In this paper, we concentrate on this X-ray emitting group population.  It is an open question how groups and the galaxies in them have changed with time.  From an optically-selected group sample based on the CNOC2 survey, Wilman et al. (2005 a,b) find that number of star forming galaxies in groups increases with redshift to $z\\sim0.5$.  On the X-ray side, Jones et al. (2002) studied a few intermediate-redshift groups with ROSAT and did not find evidence for evolution in the group X-ray luminosity function.  Willis et al (2005) find a similar lack of evolution using early data from the XMM Large-Scale Structure Survey.  We have undertaken a program to study in detail a sample of moderate-redshift ($0.2<z<0.6$), X-ray selected groups.  This study includes deeper X-ray data than previous studies of groups at these redshifts as well as HST imaging and ground based spectroscopy.  X-ray selection provides both a method of finding groups at higher redshifts where their sparse galaxy populations are difficult to recognize and a well-defined selection criteria.  Our sample and the initial spectroscopic follow-up are described elsewhere (Mulchaey et al. 2006; hereafter Paper I).  Here we present the results of \\textit{XMM-Newton} observations of six of our groups.  Throughout the paper, we assume a cosmology of $H_0=70h_{70}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Lambda = 0.73$. ",
        "conclusions": ""
    },
    "0606/astro-ph0606587_arXiv.txt": {
        "abstract": "{} {The development of powerful infrared observational technics enables the study of very extincted objects and young embedded star forming regions. This is especially interesting in the context of massive stars which form and spend a non negligible fraction of their life still enshrouded in their parental molecular cloud. Spectrophotometric calibrations are thus necessary to constrain the physical properties of heavily extincted objects.} {Here, we derive UBVJHK magnitudes and bolometric corrections from a grid of atmosphere models for O stars. Bessel passbands are used. Bolometric corrections ($BC$) are derived as a function of \\teff\\ and are subsequently used to derive $BC -$ spectral type ($ST$) and Absolute Magnitudes $- ST$ relations.} {Infrared magnitudes and, for the first time, bolometric corrections are given for the full range of spectral types and luminosity classes. Infrared colors are essentially constant. $(H-K)_{0}$ is 0.05 mag bluer than previously proposed. Optical calibrations are also provided and are similar to previous work, except for $(B-V)_{0}$ which is found to be at minimum -0.28 for standard O stars, slightly larger (0.04 mag) than commonly accepted.} {We present a consistent set of photometric calibrations of optical and infrared magnitudes and bolometric corrections for Galactic O stars as a function of \\teff\\ and spectral type based on non-LTE atmosphere models including winds and line-blanketing.} ",
        "introduction": "\\label{intro}  Massive stars are known to play an important role in various fields of astrophysics, from stellar physics to ISM studies and chemical evolution of galaxies, and to cosmological issues such as the reionisation of the Universe. In particular, the connection between massive stars and star formation is very tight: as a result of their short lifetimes, massive stars are associated with star forming events, and their feedback effects (radiation, winds) have a strong impact on star formation processes. Moreover, their ionising fluxes are responsible for nebular emission lines such as Ly$_{\\alpha}$ or H$_{\\alpha}$, two lines usually used to trace star formation \\citep{kennicutt,russeil}. However, the details of the formation of \\textit{individual} massive stars is still a matter of debate: a standard accretion process faces the problem of the strong radiative pressure generated by the luminosity of young massive proto-stars, so that the mass growth can be stopped at around 10 \\msun. Although progress has been recently made \\citep{ys02,krumholz}, another scenario in which massive stars form through mergers of low mass protostars in dense clusters was proposed by \\citet{bbz98}. This key question of the formation of the most massive stars has triggered a number of observational studies aimed at obtaining constraints on the properties of the youngest objects \\citep[e.g.][]{pauluchii,bik}. Due to the short evolutionary timescale of massive stars, heavily extincted young star forming regions have to be probed, which requires the use of infrared spectrophotometry.  Although in principle using only spectroscopy allows a derivation of spectral types and luminosity classes (LC), photometry can be useful. This is the case when spectra have to be corrected for nebular emission always present in star forming regions, rendering the line strength/shape uncertain. As a result, spectral classification and luminosity classes determinations are difficult. Moreover the luminosity is usually derived from observed magnitudes, extinction and bolometric corrections. Estimates of extinction often rely on intrinsic colors of stars while the knowledge of bolometric corrections requires atmosphere models. Hence, accurate intrinsic photometry is crucial to get access to luminosities. Although such photometry is usually available in the optical range \\citep{kur79,sk82,cgm86}, this is not the case in the infrared where calibrations are incomplete. The widely used intrinsic colors of \\citet{kor83} are only given for O6 to O9.5 dwarfs and the latest O supergiants. \\citet{johnson66} covers the same range of spectral types / luminosity class.   In this context, the recent development of reliable atmosphere models for massive stars is certainly welcome. Indeed, the inclusion of line-blanketing in such models now allows realistic prediction of atmospheric structures and emergent spectra which are used to get quantitative constraints on the properties of massive stars \\citep{paul02,hil03,jc03,martins04,repolust04,martins05}. The grid of models computed by \\citet{msh05} (herafter MSH05) and the associated SEDs are especially interesting since they can be used to compute optical and, most importantly, near infrared photometry for the whole range of O stars. Together with effective temperature scales, calibrations of magnitudes and bolometric corrections as a function of spectral type can thus be produced.  In this paper, we have used the SEDs of MSH05 to calculate UBVJHK photometry. In Sect.\\ \\ref{synth_phot} we present our method and gives the results which are discussed in Sect.\\ \\ref{comp_previous} and summarised in Sect.\\ \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have derived calibrations of UBVJHK photometry of O stars as a function of effective temperature and spectral type using the recent grid of atmosphere models of \\citet{msh05}. UBVJHK photometry was computed as in \\citet{bcp98}, using the system of \\citet{bb88} (near IR) and \\citet{bessel90} (optical).  We provide the first calibrations of near-IR photometry, including bolometric corrections, covering the whole range of spectral types and luminosity classes of O stars. Infrared colors are almost constant. $(H-K)_{0}$ is found to be -0.10, slightly lower (0.05 mag) than the value of \\citet{kor83}.  Optical photometry is consistent with recent studies. One exception is the minimum value of $(B-V)_{0}$ for standard O stars (i.e. with \\logg $\\lesssim$ 4.0) which is found to be -0.28, slightly larger than previously accepted (-0.32). This is important when estimating reddening and distances of OB associations since an error of 0.04 mag in color excess amounts to an error of $\\sim$ 0.1 mag in distance modulus (or $\\sim 0.05$ dex in luminosity).  These calibrations will be useful to study young massive stars embedded in star forming regions and to better understand their formation process."
    },
    "0606/astro-ph0606252_arXiv.txt": {
        "abstract": "{  The method and results of the computation of the model atmospheres and spectral energy distributions of chemically peculiar stars, are discussed. The models are computed with a special consideration of the particular problems encountered when computing model atmospheres for M and C-giants, and of hydrogen deficient stars. We present some computed model atmospheres for Sakurai's object, giants of globular clusters, and C-giants. ",
        "introduction": "The computations of model atmospheres of and theoretical spectra of stars is an essential part of many modern astrophysical projects. An extended grid of model atmospheres and spectra were computed recently using the most complete sets of opacities (see Kurucz (1993, 1999), Hauschildt et al. (1999). In most of the computations the solar abundances (Anders \\& Grevesse 1989) or solar abundances scaled by the metallicity factor [Fe/H] are used. However, the metal abundances in the atmospheres of evolved stars can significantly differ from the solar abundance ratios. The reason for this is because convection and other mixing processes dredge up the products of the nucleosynthesis from the stellar interior.   The atmospheres of the most evolved stars (R CrB, Sakurai's object, etc) present another interesting problem. They are helium and carbon rich, therefore the opacity due to H$^-$ absorption is not as important. The temperature structure of the model atmospheres of these stars is the different from the case of the ``normall''  abundances due to the changes of opacity. ",
        "conclusions": "Due to the high sensitivity of computed model atmospheres and synthetic spectra of evolved stars on the adopted input data, the analysis of their spectra should be done in the framework of the self-consisted approach. Model atmospheres should correspond with the obtained abundances and other physical parameters. Even in the case of the ``normal'' red giants the temperature structure should be computed with the correct abundances. Otherwise, the abundance determination results might be affected by significant errors ($>$ 0.2-0.3 dex, see Pavlenko \\& Yakovina 1994 for more details). Generally speaking, existed grids of the model atmospheres can be used as zero-approach tool to be refined in the process of analysis.   Despite a substantial efforts to develop more sophisticated model atmosphere, a few problems provide the real challenge for the modern theory of stellar astrophysics:  -- atmospheres of the most evolved stars do not exist in hydrostatic equilibrium. In some cases effects of sphericity can be essential.  -- Convection in the extended photospheres can be properly described only in the 3D approach.   -- The further computation of line lists and other physical parameters (dissociation energies, partition functions, etc) of polyatomic molecules is absolutely essential for cool stellar atmosphere modelling to progress. In particular little, no or only poor data exists for species such as C$_3$, CH$_4$ and C$_2$H$_2$. These species provide substantial opacity in the coolest atmospheres. Even the existing line lists of diatomic molecules, such as CN, C$_2$, CH, are not accurate enough to be used for the high resolution analysis of observed spectra.  Naturally, these problems can be solved one at a time. However, it is only possible when paying special attention to theoretical support of current observations."
    },
    "0606/astro-ph0606593_arXiv.txt": {
        "abstract": "The role of asymmetry on the evolution of prebiotic homochirality is investigated in the context of autocatalytic polymerization reaction networks. A model featuring enantiometric cross-inhibition and chiral bias is used to study the diffusion equations controlling the spatiotemporal development of left and right-handed domains. Bounds on the chiral bias are obtained based on present-day constraints on the emergence of life on early Earth. The viability of biasing mechanisms such as weak neutral currents and circularly polarized UV light is discussed. The results can be applied to any hypothetical planetary platform. ",
        "introduction": "The emergence of biomolecular homochirality in prebiotic Earth is a crucial step in the early history of life \\cite{AG93, Bonner}. It is well-known that chiral selectivity plays a key role in the biochemistry of living systems: amino acids in proteins are left-handed while sugars are right-handed. However, laboratory syntheses produce racemic results. This is somewhat surprising, given that statistical fluctuations of reactants will invariably bias one enantiometer over the other \\cite{Blackmond04}:  even though every synthesis is {\\it ab initio} asymmetric \\cite{Dunitz}, the enantiometric excess is nevertheless erased as the reactions unfold. An important exception is the reaction by Soai and  coworkers, where a small initial enantiometric excess is effectively amplified in the autocatalytic alkylation of pyrimidyl  aldehydes with dialkylzincs \\cite{Soai}. As stressed by Blackmond \\cite{Blackmond04}, Soai's reaction succeeds because it features the needed autocatalytic behavior proposed originally by Frank \\cite{Frank53} with enantiometric cross-inhibition catalysed by dimers.  It is unlikely that the specific chemistry of the Soai reaction occurred in early-Earth. However, it displays the relevant signatures of a realistic homochirality-inducing reaction network: autocatalysis, enantiometric cross-inhibition, and enzymatic enhancement performed by dimers or by larger chirally-pure chains.  In the present work, we will investigate the spatiotemporal dynamics of a reaction network recently proposed by Sandars which shares these features  \\cite{Sandars03}. To it, we will add an explicit chiral bias, in order to investigate the efficacy of intrinsic and extrinsic biasing mechanisms proposed in the literature. A first step in this direction can be found in the work by Brandenburg {\\it et al.}, who studied an extension of Sandar's model including bias but no spatial dependence \\cite{BAHN}. By intrinsic we mean either biasing effects related to fundamental physics, such as parity-violating weak neutral currents (WNC) \\cite{WNC0,KN83} or -- given that we know little of early-Earth's prebiotic chemistry and even less of other possible life-bearing planetary platforms  \\cite{Orgel98} --  to some as yet unknown chemical process. By extrinsic we mean possible environmental influences, such as circularly-polarized UV light (CPL) from, for example, active star-formation regions \\cite{Lucas05} or direct seeding of chiral compounds by meteoritic bombardment \\cite{Engel97, Cronin98}. The spatiotemporal dynamics of the reaction network will be shown to be equivalent to a two-phase system undergoing a symmetry-breaking phase transition characterized by the formation of competing domains of opposite chirality. The evolution of the domain network is sensitive to Earth's early environmental history and to the magnitude of the chiral bias.  Using the time-scale associated with the emergence of life on Earth it is possible to obtain a lower bound on the bias.  In particular, it will be shown that the very small bias from WNC is inefficient to generate homochiral conditions. For CPL the situation is less clear due to  uncertainty in the nature and duration of sources, but still highly unlikely. The formalism is set up to be applicable to any planetary platform.  The paper is organized as follows: In section 2 we introduce the biased polymerization model and its pertinent rate equations. In section 3 we describe the reduced ($n=2$) model and how the net chirality can be interpreted as a continuous order parameter satisfying an effective potential. This allows us to introduce explicitly spatial dependence in the study of biased polymerization. In section 4 we describe the dynamics of homochirality using techniques from the theory of phase transitions with mean-field Ginzburg-Landau models in the limit of no bias. In section 5 we generalize our results to include a small bias, describing in detail the wall dynamics in this case and the time-scales associated with the development of homochirality in early-Earth. We also investigate if the onset of prebiotic homochirality on early-Earth could have been the result of a nucleation event.We conclude in section 6 with a summary of our results and an outlook to future work. ",
        "conclusions": "The spatiotemporal evolution of prebiotic homochirality was investigated in the context of an autocatalytic reaction network featuring enantiometric cross-inhibition catalysed by dimers and chiral bias. Domains of opposite chirality, separated by thin interfaces, compete for dominance. It was shown that the dynamics of the domain network is determined by the percolation properties of its initial distribution and subsequently by the two main forces acting on the interfaces, surface tension and chiral bias. Small biases of $g\\leq 10^{-6}$ were shown to be inefficient to drive  reasonably-sized reactor pools toward homochirality within presently-accepted time-scales for the origin of life on Earth.  As a consequence, the present calculations indicate that WNCs cannot explain the observed homochirality of life's  biomolecules. CPL remains a remote possibility, albeit current sources do not look promising either in magnitude and duration. Also, it should be noted that unpolarized UV may destroy any early enantiometric excess.  The results obtained assume that the polymerization dynamics can be captured by truncating the reaction network to $n\\leq 2$ and that the dynamics of dimers is enslaved by that of monomers (the adiabatic approximation). These approximations imply that complete chiral separation can only occur with perfect fidelity $f=1$. Going beyond involves solving the complete network of spatiotemporal rate equations for larger values of $n$ and for varying fidelity, a computer-intensive, but not impossible, task. However,  given that the formation rate of higher $n$ polymers will necessarily be slower, we believe that the results obtained here capture at least qualitatively the essentials of the more general case. It remains a challenging open question whether a simple transformation could be found to reduce the biased higher-$n$ system to an effective field theory as done here for $n=2$.  What other possible sources of bias could have driven Earth's prebiotic chemistry toward homochirality? We cannot rule out the possibility that some unknown chemical bias satisfying the above bound might have been  active. Another, highly unnatractive, possibility is that an unlikely large statistical fluctuation towards one  enantiometer did occur and established the correct initial conditions.  Possible bombardment from meteors contaminated with chiral compounds could also have jump-started the process, although one still needs to explain how the chiral excess formed in the meteors in the first place.  It seems to us that the answer to this enigma will be found in the coupling of the reaction network to the environment. We note again that the results obtained here are within the diffusive, and hence ``gentle,'' evolution towards homochirality. Early-Earth, however, was a dramatic environment. Given the nonlinear properties of the spatiotemporal equations describing the evolution towards homochirality, environmental disturbances, such as meteoritic impacts or volcanic eruptions, must have played a key role in early-Earth's prebiotic chemistry.  These disturbances, if violent enough, would certainly affect the evolution of the chiral domain network and possibly change the bounds obtained in the present work. Gleiser and Thorarinson proposed to model the coupling to an external disturbance stochastically \\cite{GT}. Within their framework, results will depend on how the amplitude of the external ``noise'' compares with the critical value described in section 4.1. Preliminary results indicate that large enough noises (modelling external influences) may redirect the direction of homochirality entirely, erasing any previous evolution toward either handedness. Further work along these lines is underway. In a different approach, Brandenburg and Multam\\\"aki suggested that hydrodynamic turbulence could have sped up the march toward homochirality \\cite{BM}. In either case, it is clear that the evolution toward homochirality, as that of life itself, cannot be separated from Earth's early environmental history.  The author thanks Gustav Arrhenius, Jeffrey Bada, Freeman Dyson, Leslie Orgel, and Joel Thorarinson for stimulating discussions. He also thanks Joel Thorarinson for producing figure 3."
    },
    "0606/astro-ph0606070_arXiv.txt": {
        "abstract": "\\lsi\\ is one of the  few X-ray binaries with Be star companion  from which both radio  and high-energy gamma-ray emission  have been observed.  We present \\xmm\\ and \\intgr\\ observations which reveal variability of the X-ray spectral index of the system. The X-ray spectrum is hard (photon index $\\Gamma\\simeq 1.5$) during the orbital phases of both high and low X-ray flux. However, the spectrum softens at the moment of  transition from high to low X-ray state. The spectrum of the system in the hard X-ray band  does not reveal the presence of a cut-off (or, at least a spectral break)  at 10-60~keV energies,  expected if the compact object is an accreting neutron star.  The observed spectrum and spectral variability can be explained if the compact object in the system is a rotation powered pulsar. ",
        "introduction": "The Be star binary \\lsi\\ is one of the few X-ray binaries (along with PSR B1259-63 and LS 5039) from which radio  and very high-energy gamma-ray emission is observed. Radio emission from the system is highly variable and shows periodicity of $T=26.4960$ days, which can be associated with the binary orbital period \\citep{gregory02}. Radio observations reveal the presence of 100~AU-scale jet in the system which places \\lsi\\ among the Galactic micro-quasars \\citep{massi93,massi04}. The system is also a Galactic \"micro-blazar\"  due to its  association  with 100 MeV gamma-ray source 2CG 135+01 \\citep{tavani98} visible up to TeV energies \\citep{albert06}.  \\citet{mendelson89} found an optical modulation in the V-band with the same periodicity as the radio periodicity. This is confirmed by Paredes et al. (1994); the optical modulation in the V-band is about 0.15 mag.  Optical data allow to constrain the orbital parameters of the system revealing the eccentricity of the orbit, $e\\simeq 0.7$ \\citep{casares05}. However, the measurements are not sufficient to determine the nature of the  compact object (neutron star or black hole), because the inclination of the orbit is poorly constrained.  The X-ray emission from \\lsi\\ is also variable  and has nonthermal spectrum. Since the discovery, the system was twice monitored simultaneously in the X-ray and radio bands over a single orbital cycle \\citep{taylor96,harrison00}. These simultaneous observations  show that X-ray emission peaks almost half an orbit before the radio. In both X-ray and radio bands the orbital phases of maxima of the flux \"drift\" from orbit to  orbit.  Two major types of models of radio-to-X-ray activity of \\lsi\\ were proposed in the literature. Models of the first type, first introduced by \\citet{taylor84} assume that activity of the source is powered by accretion onto the compact object. In the second class of models, first proposed by \\citet{maraschi81}, the activity of the source is explained by interactions of a young rotation powered pulsar with the wind from the companion Be star.  If the system is an accreting neutron star or black hole, one expects to find a cut-off powerlaw spectrum in the hard X-ray band.   The cut-off energy is normally at $10-60$~keV \\citep{white83,filippova05} for neutron stars and at $\\sim 100$~keV for  black holes \\citep{mcclintock03} (the cut-off can move to even higher energies in the \"high/soft\" state see e.g. \\citet{belloni}).   Of course, it is possible that the emission from the \"central engine\" of the  micro-blazar \\lsi\\ is  \"masked\" by the emission from the jet, which can be beamed toward an observer on Earth,  similarly to what is observed in blazars.   In this case the hard X-ray spectrum of the source is a superposition of the powerlaw  jet emission and emission from the accretion disk. If the jet and accretion contributions to the X-ray spectrum are comparable, then emission from the accretion disk should at least produce an observable  spectral feature (e.g. a bump, a break or turnover) in the  10-100~keV energy band.  Below we present a study of \\lsi\\ in the 0.5-100~keV  energy band with \\xmm\\ and \\intgr\\ which shows that the spectrum of the source is well fit by a simple powerlaw, without any signature of high-energy cut-off. Moreover, the X-ray powerlaw matches smoothly to the higher-energy spectrum in the 100~keV -- 10~MeV band found with the {\\it CGRO} instruments OSSE \\citep{tavani96} and COMPTEL \\citep{vandijk96}. The broad band spectrum of the source does not exhibit a cut off up to at least GeV energies. The spectrum of \\lsi\\ in the X-ray -- soft gamma-ray band is thus very different from the typical  spectra of  accreting neutron stars and black holes.   Contrary to accreting compact object models, a  featureless powerlaw keV~--~MeV spectrum is expected in the \"rotation powered pulsar\" scenario. Unfortunately, the population of X-ray binaries with a rotation powered pulsars as compact objects is not so rich as that of accreting pulsars. In fact, the only firmly established example of X-ray binary with a young radio pulsar as a compact object, PSR B1259-63 \\citep{johnston92}, exhibits a featureless powerlaw spectrum in the X-ray band which smoothly continues to the higher energies, exactly as in \\lsi. Moreover, the overall spectral energy distributions and variability properties of the two systems are qualitatively very similar: large radio outbursts which occur once per orbit, similar X-ray spectra and high-energy gamma-ray emission  extending up to the TeV energy band. Recent observations of PSR B1259-63 during its 2004 periastron passage \\citep{chernyakova06} show that the X-ray emission from the system  is most probably produced via   inverse Compton scattering of the UV radiation from the Be star by the pulsar wind electrons, the same mechanism  as proposed by \\citet{maraschi81}  for  \\lsi  The \\xmm\\ and \\intgr\\ observations of \\lsi\\, reported here,  reveal the variability of the X-ray spectrum of the system during the transition from high to low X-ray state. We find that the spectra of the system in high and low X-ray states are very similar, but during the transition to the low flux state the spectrum softens. It turns out that explanation of such behavior of the system is not trivial and, in fact, the shape and variability of the spectrum put tight constraints on the injection mechanism of high-energy electrons responsible for the X-ray emission as well as on the geometry of the X-ray emission region.  This paper is organized as follows: in Section 2 we describe the details of the \\intgr\\ and \\xmmsp  data analysis.  The results (imaging, spectral and timing analysis)  are presented in Section 3. In Section 4 we discuss the physical implications of the results for the synchrotron -- inverse Compton broad band model of the source. This gives us the possibility to constrain physical parameters of the system and develop a physical model of the source which we study in more details in Section 5.  Finally, we summarize our findings in Section 6.  \\section[]{Observations and Data Analysis}  \\subsection{\\intgr\\ observations}  Since the launch of \\intgr\\ \\citep{winkler} on  October 17, 2002, \\lsi\\ was several times in the field of view of the main instruments   during the routine Galactic plane scans and pointed observations  (see Table \\ref{intdata} for details).  Most of the times the distance of the source  from the center of the field of view  was too large to use  the X-ray monitor JEM-X.  Therefore IBIS/ISGRI \\citep{lebrun} is the only instrument we can use in our analysis of this source. In this analysis we have used the version 5.1 of the Offline Science Analysis (OSA) software distributed by the ISDC \\citep{courvoisier03}.  In our analysis we have used all available public data spread over the period from the January 2003 (rev 25)  until March 2005 (rev 288). Overall we have analyzed 600 science windows which resulted in an effective vignetting corrected exposure of 273~ksec.  In order to study the spectral variations of the source we have grouped the data into three parts covering orbital phases 0.4 -- 0.6, 0.6 -- 0.8 and  0.8 -- 0.4 (see Table \\ref{intdata}). The spectral analysis was done with the use of the \\texttt{mosaic\\_spec} tool, which  extracts spectra for a given sky position from mosaic sky images.    \\begin{table} \\caption{Journal of the \\intgr\\  observations of \\lsi. $^*$ \\label{intdata}} \\begin{tabular}{c|c|c|c|c|c} \\hline Data & Orbital phase &  Effective   &20-60 keV Flux&$\\Gamma$ \\\\ Set &               &     Exposure  (ks)  &$10^{-11}$ergs s$^{-1}$cm$^{-2}$\\\\ \\hline I1  & 0.4 -- 0.6  &50& 3.8$\\pm 0.6$ &1.7$\\pm 0.4$ \\\\ I2  & 0.6 -- 0.8  &23& 3.0$\\pm 1.0$&3.6$^{+1.6}_{-1.1}$ \\\\ I3  & 0.8 -- 0.4  & 200&2.4$\\pm0.3$ &1.4$\\pm 0.3$ \\\\ I$_{\\mbox{av}}$& &273&2.5$\\pm 0.3$ &1.6$\\pm 0.2$ \\\\ \\hline \\end{tabular} $^*$ Given errors represent 68\\% confidence interval uncertainties. \\end{table}  \\subsection{\\xmm\\ observations}  \\xmm\\ has observed \\lsi\\  with the EPIC instruments  five times during 2002. Four observations have been done during the same orbital cycle, and the fifth one  has been done seven months later. The log of the \\xmm\\ data analyzed in this paper is presented in Table~\\ref{data}.  \\begin{table} \\caption{Journal of \\xmm\\ observations of \\lsi \\label{data}} \\begin{tabular}{c@{\\,}c@{\\,}ccc@{\\,}c} \\hline Data& Observational& Date &    MJD&  Orbital&Exposure (ks)\\\\ Set&   ID         &      & (days)&    Phase  &MOS1/MOS2/PN     \\\\ \\hline X1& 0112430101& 2002-02-05&52310  & 0.55  &5.8/5.8/5.0  \\\\ X2& 0112430102& 2002-02-10&52315  & 0.76  &5.8/5.8/5.0  \\\\ X3& 0112430103& 2002-02-17&52322  & 0.01  &1.5/1.5/5.0  \\\\ X4& 0112430201& 2002-02-21&52326  & 0.18  &7.0/7.0/5.1  \\\\ X5& 0112430401& 2002-09-16&52533  & 0.97  &6.2/6.2/6.0  \\\\ \\hline \\end{tabular} \\end{table}   The \\xmmsp Observation Data Files (ODFs) were obtained from the on-line Science Archive\\footnote{http://xmm.vilspa.esa.es/external/xmm\\_data\\_acc/xsa/index.shtml}; the data were then processed and the event-lists filtered using {\\sc xmmselect} within the Science Analysis Software ({\\sc sas}) v6.0.1. In all observations the source was observed with the PN detector in the Small Window Mode, and with MOS1 and MOS2 detectors in the Full Frame Mode.  In all observations a medium filter was used.  The event lists for spectral analysis were extracted from a 40$^\\prime$$^\\prime$ radius circle at the source position for  MOS1,  from a 30.0$^\\prime$$^\\prime$ radius circle for MOS2, and for PN  a region of 25$^\\prime$$^\\prime$ around the source position was chosen.  Background photons were collected from a region located in the vicinity of the source with the same size area as the one chosen for the source photons.  For the spectral analysis, periods affected by soft proton flares need to be filtered out. To exclude them we have used script \\texttt{xmmlight\\_clean.csh} \\footnote{http://www.sr.bham.ac.uk/xmm2/xmmlight\\_clean.csh}. The source countrate in 2  -- 10 keV energy range  varied from about 0.03  to 0.06 cts/s for MOS  instruments, and from 0.05 to 0.10 cts/s for PN. With the help of the \\texttt{GRPPHA} FTOOL we have re-binned the spectra, so that  each energy bin has at least 30 counts.  Data from MOS1, MOS2 and PN detectors were combined in the spectral analysis to achieve better statistics. ",
        "conclusions": "In this paper we have presented the \\intgr\\ and \\xmm\\ observations of  \\lsi. We have found that the overall spectrum of the system in 0.5-100~keV band is well fit by a featureless powerlaw with the photon index $\\Gamma_{ph}\\simeq 1.5$ both in the high and low flux states. The 0.5-100~keV powerlaw spectrum matches smoothly the higher-energy spectrum of the source in 100~keV -- GeV band. The powerlaw with the same photon index seems to continue without a cut-off up to 1-10~MeV energies. Non-observation of a cut-off or a break in the spectrum at 10-100~keV energies, typical for the accreting neutron stars and black holes, implies that the accretion in the source proceeds at a very low (if not zero) rate.  This favors the scenario in which the compact object is a rotation-powered pulsar.  We have discovered that the spectrum is hard in both high and low X-ray states, but  softens during the transition from the high flux to the low flux state.  The spectral characteristics of the system favor the model in which electrons are initially  injected at high energies and subsequently cool to lower energies forming the characteristic \"cooling\" spectrum with the spectral index $\\Gamma_e\\simeq 2$. We have proposed a model in which  cold relativistic pulsar wind (electron and/or proton loaded) with the bulk Lorentz factor $\\Gamma\\sim 10^2\\div 10^3$ provides a  source of X-ray emitting electrons. Such relativistic wind penetrating into the Be star disk can also be responsible for the formation of jet-like outflow from the system."
    },
    "0606/astro-ph0606300_arXiv.txt": {
        "abstract": "We present the first 6.7~GHz methanol maser linear polarization map of the extended filamentary maser structure around the compact {\\sc Hii} region W3(OH). The methanol masers show linear polarization up to $\\sim 8$~per~cent and the polarization angles indicate a magnetic field direction along the North-South maser structure. The polarization angles are consistent with those measured for the OH masers, taking into account external Faraday rotation toward W3(OH), and confirm that the OH and methanol masers are found in similar physical conditions. Additionally we discuss the Zeeman splitting of the 6.7~GHz methanol transition and present an upper limit of $\\sim 22$~mG for the magnetic field strength in the maser region.  The upper limit is fully consistent with the field strengths derived from OH maser Zeeman splitting. ",
        "introduction": "Through polarization observations, maser are unique probes of the magnetic fields that are thought to play an important role in star-forming regions (SFRs). Observations of the Zeeman splitting of several different transitions of the OH masers indicate mG level magnetic fields, and the linear polarization of the masers has also been used to determine the structure of the magnetic field \\citep[e.g.][]{Etoka05, Bartkiewicz05}. In addition to the OH masers, the H$_2$O masers at 22~GHz have also been used to determine the magnetic field strength and structure in the more dense areas of SFRs \\citep[e.g.][]{Sarma01, Imai03, Vlemmings06}. These observations have revealed magnetic field strengths of several tens to hundreds of mG.  The 6.7~GHz methanol maser is a strong tracer of high-mass star-formation \\citep{Menten91} and is typically one of the brightest masers in those regions. The masers occur in regions with similar densities and temperatures to the ground state OH masers, where enhanced densities of both molecular species are to be expected as a result of material from grain mantles being evaporated under the influence of weak shocks \\citep{Hartquist95}. Theoretical modelling has shown that both masers can be pumped simultaneously \\citep{Cragg02}. Similar to H$_2$O, methanol is a non-paramagnetic molecule. As a result, the Zeeman splitting is only a small fraction of the typical maser line width. Thus, the determination of the magnetic field strength using 6.7~GHz methanol masers requires sensitive and high spectral resolution observations. As yet, no successful methanol Zeeman experiment has been reported. Only recently have the first measurements of linear polarization of the 6.7~GHz methanol maser been made. Observations by \\citet{Ellingsen02} indicate a fractional linear polarization up to $10$~per~cent, although no polarization maps have been produced until now.  In this letter the polarization properties of the 6.7~GHz methanol masers in the massive star-forming region W3(OH) are presented. The polarization has been determined using MERLIN\\footnote{MERLIN is a national facility operated by the University of Manchester at Jodrell Bank Observatory on behalf of PPARC} observations presented by \\citet[][~hereafter HSC06]{HarveySmith06}, where the focus has been on the astrometry and distribution of the methanol masers.  W3(OH) is one of the most intensively studied star-forming regions and contains one of the most luminous ultracompact {\\sc Hii} (UC{\\sc Hii}) regions. It is located in the Perseus spiral arm of the Galaxy at a distance of $1.95\\pm0.04$~kpc as determined using maser astrometry \\citep{Xu06}. It is the site of several young high- and intermediate-mass stars, the most prominent of which is the UC{\\sc Hii} region itself, which is ionized by an embedded zero-age main-sequence O-star \\citep[e.g][]{Dreher81}. The W3(OH) region displays strong emission from a wide variety of maser species \\citep[e.g][~and references therein]{Wright04a, Etoka05}.  Strong OH and methanol masers are seen projected onto the UC{\\sc Hii} region and strong H$_2$O masers have been found towards the related Turner-Welch (TW) Object \\citep{Turner84}. ",
        "conclusions": "MERLIN has been used to produce the first maps of the polarization of 6.7~GHz methanol masers in the star-forming region W3(OH). Linear polarization of up to $\\sim 8$~per~cent, similar in level to that found from the 12.2~GHz methanol masers at the same velocities, has been found. We determined an upper limit of $22$~mG to the magnetic field strength through the RHC-LHC cross-correlation method. As the 6.7~GHz methanol masers are much less affected by internal and external Faraday rotation than the lower frequency OH masers, they are excellent probes of the magnetic field direction in the maser region. After correction for Faraday rotation we find that the magnetic field is orientated parallel to the extended N-S filamentary structure detected in methanol and OH maser observations."
    },
    "0606/astro-ph0606136_arXiv.txt": {
        "abstract": "% We propose a solution to the problem of astrometric and photometric calibration of coronagraphic images with a simple optical device which, in theory, is easy to use.  Our design uses the Fraunhofer approximation of Fourier optics. Placing a periodic grid of wires (we use a square grid) with known width and spacing in a pupil plane in front of the occulting coronagraphic focal plane mask produces fiducial images of the obscured star at known locations relative to the star. We also derive the intensity of these fiducial images in the coronagraphic image. These calibrator images can be used for precise relative astrometry, to establish companionship of other objects in the field of view through measurement of common proper motion or common parallax, to determine orbits, and to observe disk structure around the star quantitatively. The calibrator spots also have known brightness, selectable by the coronagraph designer, permitting accurate relative photometry in the coronagraphic image. This technique, which enables precision exoplanetary science, is relevant to future coronagraphic instruments, and is particularly useful for ``extreme'' adaptive optics and space-based coronagraphy. ",
        "introduction": "Stellar coronagraphs suppress a star's point spread function (PSF) in order to enable the imaging of faint objects or structure, such as planets or proto-planetary disks, in orbit about the  star \\citep{Lyot39, Sivaramakrishnan01}. Accurate relative astrometry between the star and its apparent companions or extended disk structure is essential to establish physical association through confirmation of common proper motion or common parallax \\citep{Oppenheimer018pcsurvey}, to determine orbital parameters, and to observe disk structure quantitatively.  Relative photometry between the star and these other objects in the field of view is also necessary to conduct studies of disk or companion physics and chemistry. However, a perfect coronagraph will completely remove the core of the star's PSF, meaning that measuring its intensity or location in the scientific data is impossible. \\citet{LS05} show that decentering the on-axis star results in complex PSF morphology --- tilt errors can even introduce artifacts that look like point sources within a few resolution elements of the center of the occulting focal plane mask of the coronagraph, even in simple theoretical simulations.  We have also observed these effects in real data from the optimized, diffraction-limited Lyot Project coronagraph \\citep{Sivaramakrishnan01,  OppenheimerAMOS03, OSM03, OppenheimerSPIE04, Makidon05}, which is deployed at the Air Force AEOS 3.6\\ m telescope \\citep{LRCN02}. This coronagraph is the first to operate in the regime of  ``extreme adaptive optics'' (ExAO), with Strehl ratios in the 70-85\\% range, and it has played a role in developing advanced techniques of exoplanetary science. (The NICMOS coronagraph on the Hubble Space Telescope is not optimized for coronagraphy --- the function of its focal plane occulter is to reduce the total flux on the detector, since there is no coronagraphic suppression of the stellar background away from the PSF core.) It brings early results to a burgeoning field of coronagraphic imaging of nearby stars. In \\citet{Digby06} we discuss some of the problems of coronagraphic astrometry uncovered when analyzing real data from the Lyot Project coronagraph, including  the fact that such data is often far more complex than simple theory might predict \\citep{LS05, SSSLOM05}. Over exposures a few seconds to several minutes long  objects are smeared out on the detector because of slow drifts in the AO system, differential refraction effects and field rotation.  In addition, the dynamic range of the detector prevents simultaneous non-coronagraphic imaging of faint companions with the AO target star.  The image is further complicated by the motion of speckles, which occurs on all timescales present in the obervation. This paper presents a solution to some of the problems discussed in \\citet{Digby06}.   We explain how to create controlled fiducial spots in the coronagraphic image which pinpoint the location of the star, in addition to providing calibration for accurate relative photometry, thus solving the fundamental problem of coronagraphic image calibration. By inserting a reticulate grid of wires into a pupil ahead of the focal plane occulting mask in the coronagraph we create predictable images of the central star at desired locations and chosen brightness in the coronagraphic image plane. We extend the analytical theory of segmented aperture telescope point-spread functions (PSFs) for direct imaging \\citep{Chanan98, SivaramakrishnanGSMT01, Yaitskova02} as well as that of coronagraphic imaging \\citep{SivaramakrishnanGSMT01, SY05} to describe the effects of our proposed reticulate grid of wires in the coronagraphic case.  Our methods, though relevant to any coronagraphic observation, are particularly applicable to space-based high dynamic range coronagraphy dedicated to the extremely challenging task of finding and characterizing Earth-like exoplanets, as our approach does not disturb the extremely smooth, well-corrected wavefront that such searches demand. Our technique does not require exquisite calibration of deformable mirrors, for instance, nor does it require any change in a deformable mirror shape before, during, or after an observation.  We note a simultaneous but independent study by \\citet{Marois06}, where a pupil plane before the focal plane mask is also modified in order to create images suitable for photometric and astrometric purposes.  \\S 2 elucidates the theoretical basis for our calculations of telescope PSFs. \\S 3 describes the perfect coronagraph used to illustrate the utility of this new theoretical construct, as well as the difficulties in proper numerical simulations of the reticulate pupil mask. In \\S 4, we extend the theory to a more practical level, addressing the effects of finite bandwidth in astronomical imaging, imprecise positioning of optics, residual uncorrected wavefront errors, and atmospheric differential refraction. \\S 5 describes the practical application of the theory, including some details on how to design reticulate pupil masks, and how to calibrate and reduce coronagraphic data to derive precision relative astrometry and photometry. ",
        "conclusions": ""
    },
    "0606/astro-ph0606699_arXiv.txt": {
        "abstract": "The clustering properties of dark matter halos are a firm prediction of modern theories of structure formation.  We use two large volume, high-resolution N-body simulations to study how the correlation function of massive dark matter halos depends upon their mass, assembly history, and recent merger activity. We find that halos with the lowest concentrations are presently more clustered than those of higher concentration, the size of the effect increasing with halo mass; this agrees with trends found in studies of lower mass halos. The clustering dependence on other characterizations of the full mass accretion history appears weaker than the effect with concentration. Using the integrated correlation function, marked correlation functions, and a power-law fit to the correlation function, we find evidence that halos which have recently undergone a major merger or a large mass gain have slightly enhanced clustering relative to a randomly chosen population with the same mass distribution. ",
        "introduction": "The observed Universe contains order on all scales we can probe. It is generally believed that the largest structures arose via the amplification of primordial (quantum) fluctuations during a period of accelerated expansion, processed by the subsequent $13\\,$Gyrs of gravitational instability. The pattern of clustering of objects on large scales is a calculable prediction of cosmological models and thus comprises one of the fundamental cosmological statistics.  Within modern theories of structure formation, the clustering of rare, massive dark matter halos is enhanced relative to that of the general mass distribution \\citep{Kai84,Efs88,ColKai89,MoWhi96,SheTor99}, an effect known as bias.  The more massive the halo, the larger the bias. As a result, the mass of halos hosting a given population of objects is sometimes inferred by measuring their degree of clustering -- allowing a statistical route to the notoriously difficult problem of measuring masses of cosmological objects (e.g.~\\citet{CooShe}).  Since halos of a given mass can differ in their formation history and large-scale environment\\footnote{ The large-scale environment of a halo refers to the density, smoothed on some scale larger than the halos, e.g.~$5-10\\,h^{-1}$Mpc.}, a natural question arises: do these details affect halo clustering? In currently viable scenarios for structure formation, objects grow either by accretion of smaller units or by major mergers with comparable-sized objects. The formation history of a halo can thus be characterized by its mass accumulation over time, such as when it reached half of its mass, had a mass jump in a short time, or last underwent a (major) merger.  Theoretically, the simplest descriptions of halo growth and clustering \\citep{bondetal,bower,lc93,lc94,kit-sut,kit-sut2} do not give a dependence upon halo formation history \\citep{Whi93,SheTor04b,FurKam05,Har06}. To reprise these arguments: pick a random point in the universe and imagine filtering the density field around it on a sequence of successively smaller scales.  The enclosed density executes a random walk, which in the usual prescription is taken to be uncorrelated from scale to scale.  The formation of a halo of a given mass corresponds to the path passing a certain critical value of the density, $\\delta_c$, at a given scale. The bias of the halo is the `past' of its random walk and its history the `future' of the walk. All halos of the same mass at that time correspond to random walks crossing the same point, and thus have the same bias. (Note that the derivation, using sharp $k$-space filtering, does not match the way the prescription is usually applied, and this has been suggested by some of the above authors as a way to obtain history dependence.  Introducing an environmental dependence through e.g.~elliptical collapse will also give a history dependence.)  The lack of dependence on halo history in the simplest descriptions does not close the discussion theoretically or otherwise. While these analytic methods work much better than might be expected given their starting assumptions, the Press-Schechter based approaches still suffer many known difficulties (e.g.~\\citet{ShePit97}, \\citet{BenKamHas05}). Other analytical ways of estimating the clustering of mergers have been explored.  For example, \\citet{FurKam05} defined a merger kernel (not calculable from first principles) and assumed that all peaks within a certain volume eventually merged.  Such an ansatz implies that recently merged halos are more clustered for $M>M_*$ and less clustered for $M<M_*$, with some dependence upon predecessor mass ratios and redshifts. (Here $M_*(z)$ is the mass at which $\\sigma(M)$, the variance of the linear power spectrum smoothed on scale $M$, equals the threshold for linear density collapse $\\delta_c(z)$, see e.g.~\\citet{Peacock}.) Using close pairs as a proxy for recently merged halos, they found a similar enhancement of clustering for $M>M_*$ and reduction for $M<M_*$ in several (analytic) clustering models. To foreshadow our results: the signals we see are consistent with this trend.   Simple analytic models cannot be expected to capture all of the complexities of halo formation in hierarchical models, and full numerical simulations are required to validate and calibrate the fits. Fortunately, numerical simulations are now able to produce samples with sufficient statistics to test for the dependence of clustering on formation history. Early work by \\citet{LemKau99} showed that the properties of dark matter halos, in particular formation times, are little affected by their large-scale environment if the entire population of objects is averaged over. They interpreted this as evidence against formation history and environment affecting clustering.  As emphasized by \\citet{SheTor04b}, however, this finding --- plus the well known fact that the typical mass of halos depends on local density --- implies that the clustering of halos of the same mass must also depend on formation time. Using a marked correlation function, \\citet{SheTor04b} found that close pairs tend to have earlier formation times than more distant pairs, work which was extended and confirmed by \\citet{Har06}. \\citet{GaoSprWhi05} found that later forming, low-mass halos are less clustered than typical halos of the same mass at the present; a possible explanation of this result was given by \\citet{WanMoJin06}. \\citet{Wec06} found a similar dependence upon halo formation time, showing that the trend reversed for more massive halos and that the clustering depended on halo concentration.  However, in order to probe to higher masses these authors assumed that the mass dependence was purely a function of the mass in units of the non-linear mass, then used earlier outputs to probe to higher values of this ratio. It should be noted that scaling quantities by $M/M_*$ gives a direct equality only if clustering is self-similar.  Since $P(k)$ is not a power-law and $\\Omega_{\\rm mat}\\ne 1$, a check of this approximation, as is done here, is crucial.   These formation time dependencies are based on (usually smooth) fits to the accretion history of the halo. However, halo assembly histories are often punctuated by large jumps from major mergers that have dramatic effects on the halos. Major mergers can be associated with a wide variety of phenomena, ranging from quasar activity \\citep{KauHae00} and starbursts in galaxies \\citep{MihHer96} to radio halos and relics in galaxy clusters (see e.g.~\\citet{Sar04} for phenomena associated with galaxy cluster mergers).  Major mergers of galaxy clusters are the most energetic events in the universe. It follows that major merger phenomena can either provide signals of interest or can cause noise in selection functions that depend upon a merger-affected observable.  If recently merged halos cluster differently from the general population (merger bias), and this is unaccounted for, conclusions drawn about halos on the basis of their clustering would be suspect.  The question of whether such merger bias exists remains unresolved, as previous work to identify a merger bias through N-body simulations and analytic methods yields mixed results \\citep{Got02,Per03,ScaTha03,FurKam05}.  In this paper we consider the clustering of the most massive dark matter halos, measured from two large volume ($1.1\\,h^{-1}$Gpc)${}^3$ N-body simulations described in \\S\\ref{sec:sims}.  We concentrate on massive halos, as most previous simulations did not have the volume to effectively probe this end of the mass function, and furthermore, for the largest mass halos the correspondence between theory and observation is particularly clean. We first examine the long-term growth history of halos, calculating the ``assembly bias'' as a function of growth history in \\S\\ref{sec:history}, extending previous results mentioned above to higher masses. We then look to short-term history effects (i.e. events), measuring the ``merger bias'' as a function of recent major merger activity or large mass gain in \\S\\ref{sec:merger}, where we find a weak, but statistically significant, signal for both cases. We conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The large-scale structure of the Universe is built upon a skeleton of clustered dark matter halos.  For the past two decades we have known that rarer, more massive dark matter halos cluster more strongly than their lower mass counterparts. Halos of a fixed mass, however, can differ in their formation history and large-scale environment, and recent work on halos smaller than galaxy clusters has shown that this can lead to further changes in their clustering.  In this paper we have used two large-volume, high resolution N-body simulations to study the clustering of massive halos as a function of formation history. We confirmed earlier results that the lower concentration massive halos are more clustered than the population as a whole; extending these results to higher masses (and thus lower redshifts) than had been probed previously \\citep{Wec06}. (Previous work had looked at similar regimes of $M/M_*$ but for smaller $M$ and thus higher redshift; note again that exact scaling with $M/M_*$ is not expected for non power law $P(k)$ and $\\Omega_m \\ne 1$.) Similarly, we confirmed the enhanced clustering of halos with later formation times, though the signal was not as strong as for concentration. The signal for bias based on a halo reaching half of its mass is weaker than that seen in \\citet{GaoSprWhi05} (again for higher $z$), and not statistically significant in our case.  We also investigated whether recent merger activity affected the clustering of massive halos --- a topic with a muddied history in the literature. While we found statistically significant ($>2\\sigma$) merger effects on clustering in many cases we considered, both for recent major mergers and large mass gain, in most cases this signal was weak: a 5--10\\% increase in bias. Our strongest signal came from using a likelihood fit of the correlation function to a power law, particularly for major mergers within $0.4\\,h^{-1}$Gyr of the present, where we saw a typical merger bias of up to $20\\%$. This bias signal is not necessarily at odds with the lack of signal in previous work, which looked for larger bias than that seen on average here.  Even with a $(1.1\\,h^{-1}{\\rm Gpc})^3$ volume, massive halos remain very rare objects and small changes in their correlations are difficult to detect. We were plagued by the competing effects that increasing the severity of the merger (and hence underlying signal) decreases the number of pairs, worsening the statistics. General trends remain elusive, since changing various criteria (e.g.~merger definition, minimum mass, time step) generally changed the number of halos involved, thus changing the errors. However, we did find that the strength of the merger bias typically increased with increasing merger ratio, i.e.~more major mergers are more strongly biased. Finally, we note that the correlations found between the last large (20\\%) mass gain and the different definitions of formation redshifts provide a connection between the assembly bias studied in \\S\\ref{sec:history} and the merger bias in \\S\\ref{sec:merger}. This bias is not expected from direct application of extended Press-Schechter theory, and it provides a phenomenon that a more precise analytic model of mergers should reproduce.  We thank J. Bullock, R. Croft, G. Jungman, P. Norberg, G. Rockefeller, R. Sheth, E. Scannapieco, R. Wechsler and A. Zabludoff for enlightening conversations and especially R. Sheth and R. Wechsler, who also provided useful comments on the draft. J.D.C., D.E.H. and M.W. thank the staff of the Aspen Center for Physics for their hospitality while this work was being completed. The simulations and analysis in this paper were carried out on supercomputers at Los Alamos National Laboratory and NERSC. J.D.C. was suported in part by the NSF. M.W. was supported in part by NASA. D.E.H. gratefully acknowledges a Feynman Fellowship from LANL."
    },
    "0606/astro-ph0606650_arXiv.txt": {
        "abstract": "{} { We have carried out a survey of 60 Herbig Ae/Be stars in the 3 micron wavelength region in search for the rare spectral features at 3.43 and 3.53~micron. These features have been attributed to the presence of large, hot, hydrogen-terminated nanodiamonds. Only two Herbig Ae/Be stars, \\object{HD~97048} and \\object{Elias~3$-$1} are known to display both these features.} { We have obtained medium-resolution spectra ($R \\sim 2500$) with the ESO near-IR instrument ISAAC in the 3.15$-$3.65~micron range. } { In our sample, no new examples of sources with prominent nanodiamond features in their 3~micron spectra were discovered. Less than 4\\% of the Herbig targets show the prominent emission features at 3.43 and/or 3.53~$\\mu$m. Both features are detected in our spectrum of HD~97048. We confirm the detection of the 3.53~$\\mu$m feature and the non-detection of the 3.43~$\\mu$m feature in MWC~297. Furthermore, we report tentative 3.53~$\\mu$m detections in V921~Sco, HD 163296 and T~CrA. The sources which display the nanodiamond features are not exceptional in the group of Herbig stars with respect to disk properties, stellar characteristics, or disk and stellar activity. Moreover, the nanodiamond sources are very different from each other in terms of these parameters. We do not find evidence for a recent supernova in the vicinity of any of the nanodiamond sources. \\\\ We have analyzed the PAH~3.3~$\\mu$m feature and the \\Pfd hydrogen emission line, two other spectral features which occur in the 3~micron wavelength range. We reinforce the conclusion of previous authors that flared-disk systems display significantly more PAH emission than self-shadowed-disk sources. The \\pfd line detection rate is higher in self-shadowed-disk sources than in the flared-disk systems.} { We discuss the possible origin and paucity of the (nano)diamond features in Herbig stars. Different creation mechanisms have been proposed in the literature, amongst others in-situ and supernova-induced formation. Our data set is inconclusive in proving or disproving either formation mechanism.} ",
        "introduction": "Herbig Ae/Be (HAEBE) stars are pre-main-sequence objects of intermediate (1.5--8 M$_\\odot$) mass. These sources exhibit large infrared (IR) excesses, for the greater part due to thermal emission of circumstellar dust. The geometry of this dust surrounding the late-B, A and F-type members of this class is disklike \\citep[e.g.][]{mannings97, testi, pietu, fuente, natta04}. Also for early-B-type stars, evidence for disks has been found \\citep[e.g.][]{vink02}. However, there is growing evidence that the structure of the circumstellar matter in early-type B stars is fundamentally different from Herbig Ae and T~Tauri stars \\citep[e.g.][]{fuente,bik04}.  The chemistry and mineralogy of the dust component of circumstellar disks around HAEBE stars has been studied with unprecedented precision based on near-to-mid-IR spectra (2--200~$\\mu$m) provided by the \\textit{Infrared Space Observatory} \\citep[ISO,][]{kessler}. The launch of this satellite in 1995 was the start of a brand-new era for the research of protoplanetary disks. The spectra revealed a large variety in dust properties and species, from small carbonaceous molecules like polycyclic aromatic hydrocarbon molecules (PAHs) to silicate dust. Moreover, some sources were shown to contain crystalline grains, similar to those found in comets in our own solar system \\citep{waelkens96,malfait, malfait99,vandenancker00a, vandenancker00b, meeus01}. All available ISO 2--15~$\\mu$m spectra of HAEBE stars have been studied as a whole by \\citet[][hereafter AV04]{ackeiso}.  A handful of HAEBE sources \\citep[HD~97048 and Elias~3$-$1,][and references therein]{vankerckhoven} and the post-AGB binary \\object{HR~4049} \\citep{geballe89} display peculiar spectral features at 3.43 and 3.53~$\\mu$m. Comparison with laboratory spectra has convincingly shown that the carriers of these features are hydrogen-terminated nanometer-sized diamonds \\citep[hereafter nanodiamonds,][]{guillois}. Near-IR observations of HD~97048 by \\citet{habart04a} have spatially resolved the emission region of the strong 3.43 and 3.53~$\\mu$m features on a scale of $0.18^{\\prime\\prime} \\times 0.18^{\\prime\\prime}$. The observations prove beyong doubt that the emission region is confined to the inner 15~AU of the circumstellar {\\em disk}. More tentative detections of the features were reported in the literature \\citep[e.g. in HD~142527 and HD~100546,][respectively]{waelkens96,malfait98b}, but only one additional 3.53~$\\mu$m feature was discovered persuasively \\citep[in MWC~297,][]{terada}. Summing up, only two HAEBE stars display both the 3.43 and the 3.53~$\\mu$m feature, and one HAEBE object shows the 3.53~$\\mu$m feature. Considering that the ISO sample contains 45 targets, only a minor fraction of the HAEBE stars appear to have such a spectacular 3 micron spectrum (AV04).  The currently operating successor to the ISO satellite, Spitzer, is not equipped with a spectrograph which can observe in the 3~micron region of the spectrum. Ground-based observations are therefore the only alternative. With the present paper, we intend to enlarge the sample of HAEBE stars that are observed in the 3 micron region and possibly identify new targets which display nanodiamond emission. Therefore, we have employed the near-IR ISAAC instrument (ESO$-$VLT) and focused on the wavelength region around 3.3 (polycyclic aromatic hydrocarbons, hereafter PAHs), 3.43 and 3.53~$\\mu$m (nanodiamonds). In Sect.~\\ref{dataset}, we present the sample and the data reduction procedure. The detected features are discussed in Sect.~\\ref{observations}. ",
        "conclusions": "}  We summarize the two main conclusions of the present paper.  \\begin{itemize} \\item Our survey of 3~micron spectra of HAEBE stars has not revealed a new spectacular nanodiamond source like HD~97048 or Elias~3$-$1. This negative result increases the rarity of these near-IR bands in the group of Herbig stars. Within the current sample, only 4\\% of the targets are surrounded by detected nanodiamonds. We find a large range in nanodiamond luminosity, which spans 1$-$2 orders of magnitude. No clear correlation is found between the presence of the nanodiamond features and the disk and stellar parameters, although a weak link between nanodiamond luminosity and stellar luminosity may be present. Furthermore, there does not appear to be a specific disk/star property which the nanodiamond sources have in common. We have shown that our findings are inconclusive with respect to the influence of the surroundings of the targets. We do not find any evidence for a correlation of nanodiamond emitters with the locations of known recent supernova remnants, nor do we find evidence for an enrichment in heavy elements of the photosphere of the system with the most prominent nano-diamond features, HD~97048. \\item Most to all Herbig stars with a flared disk are PAH emitters. This conclusion once again confirms the suggestion of \\citet{meeus01} that PAH molecules can only be excited in systems where they can capture direct stellar UV photons. In the case of self-shadowed disks, the puffed-up inner rim casts a shadow over the outer disk, preventing the PAH molecules from being heated by the central star. Our findings reinforce the results of AV04 and \\citet{habart}. The latter studies were based on ISO data. Due to the large aperture sizes of the ISO instruments, it is difficult to differentiate between PAH emission emanating from the disk surface, and contamination from the wider surroundings of the target. The present ISAAC data set, which clearly confirms the correlation, is considerably more reliable in this sense: both the narrow slit and the chop-nod technique, which allows for an efficient subtraction of constant background emission, reduce the possibility of confusion with extended PAH emission drastically. \\end{itemize}  Laboratory experiments have been instrumental in the identification of the carriers of the 3~micron features \\citep[e.g.][]{guillois,sheu02,jones04,mutschke04}. Currently, the astronomical spectra can be reproduced from the absorption spectra of films of hydrogen-terminated diamonds \\citep{sheu02}. The spectral features in the 3~micron range are due to CH-bond stretches (as is the PAH~3.3~$\\mu$m feature). The prominent peaks around 3.43 and 3.53~$\\mu$m observed in HAEBE spectra can only be reproduced when the diamonds are sufficiently large. Typical sizes for the diamonds in HAEBE stars ($\\sim$50-100~nm) are 2 orders of magnitude larger than the nanodiamonds found in our solar system \\citep[e.g.][]{jones04}. The spectral features hence indicate the presence of large, hot hydrogen-terminated diamonds.  The question remains whether there are indeed no diamonds in the majority of disks. It may be possible that nanodiamonds of small sizes are present, but remain undetected, because their spectral signature is less pronounced \\citep[e.g.][]{sheu02}. Also diamonds which are not hydrogen-terminated are not detectable in the 3~micron region. Diamonds that are stripped from hydrogen atoms are hence overlooked by our survey. A third possibility why (nano)diamonds can remain undetected is their exact location |and hence temperature| in the disk. Cool hydrogen-terminated diamonds will not radiate significantly in the near-IR. The absence of the 3~micron features in the spectrum of a HAEBE system hence does not necessarily mean that no (nano)diamonds are present.  \\subsection*{In-situ formation or extraneous origin?}  Diamond was the first presolar mineral to be identified in meteorites \\citep{lewis87} and it has the highest relative abundance among carbonaceous grains. The diamond has been classified as presolar based on the presence of a significant overabundance of very heavy (Xe-H) {\\em and} very light (Xe-L) Xe isotopes (together Xe-HL). The overabundance of this noble-gas carrier is expected to be produced by $r$- and $p$-process nucleosynthesis in supernovae. The link between the presence of nanodiamonds and the xenon overabundance led to the hypothesis that the diamonds are formed in, and injected into the system by a supernova outflow \\citep{clayton89}. Isotope anomalies on other heavy elements such as Te and Pd also point to supernova nucleosynthesis \\citep{richter98, maas01}.  Recent analysis of $^{60}$Fe isotopes in chondrites has clearly demonstrated that a nearby type~II supernova has contributed material to the natal cloud from which our own solar system formed \\citep{mostefaoui05,tachibana06}. This process may have been important for planet formation, as the the radioactive decay of $^{60}$Fe into $^{60}$Ni was an important heat source for the early planetary melting and differentiation and keeping asteroids thermally active for much longer than would be possible from the decay of $^{26}$Al alone.  Infrared spectroscopy has shown that solar-system diamonds contain N and O, probably in chemical functional groups on their surfaces \\citep{lewis89,mutschke95,andersen98,jones04}. However, the similarity of C- and N-isotopes of these diamonds and the solar system as a whole supports the idea that not all solar-system diamonds originate from supernovae and that the supernova contribution to the diamonds in our own solar system is probably not very large. It thus remains unclear how the majority of the solar-system diamonds has formed.  Although advocated by several recent studies \\citep{kouchi05,binette05}, the presence of nanodiamonds in the interstellar medium remains controversial. The absence of observable 3.43 and 3.53~$\\mu$m features in the ISM suggests that hydrogenated hydrocarbons cannot be more abundant than $\\approx$~0.1 ppm in the interstellar medium \\citep{tielens87}. \\citet{vankerckhoven} showed that in HD~97048 and Elias~3$-$1, 1 nanodiamond part per billion relative to hydrogen is required to reproduce the 3~micron spectra observed by ISO, making an ISM origin of nanodiamonds in principle a viable option. This hypothesis has problems in explaining the paucity of detections of the 3.43 and 3.53~$\\mu$m feature in Herbig stars within our sample, however.  We found no evidence for the presence of a supernova remnant (SNR) in the proximity of any of our sample sources. We made use of the catalog of galactic SNRs published by \\citet{green04}. Only R~Mon and V590~Mon are potentially located close to a SNR; the Monoceros Nebula (G205.5$+$0.5). These sources display no detected nanodiamond features, however. We note that the absence of a SNR in the vicinity of our sample stars in the above-mentioned catalog does not necessarily imply that no supernova has actually occured. Most of the sample targets are located in massive-star-forming regions. It may very well be that the most massive members in such a region have recently gone off as a supernova. In a relatively crowded star forming region, the ejecta of such an event likely affect a number of surrounding young stellar objects or molecular clouds. In this scenario, the presence of diamonds in the circumstellar disk is expected to be common to all disk systems in the vicinity of the supernova. Unfortunately, we have no spectra of disk sources close to Elias~3$-$1, HD~97048 or MWC~297 to check this hypothesis.  The parent star of a disk which contains supernova-injected diamonds must be be polluted by supernova material as well, under the supernova hypothesis. Enrichment of the stellar photosphere by a supernova outflow is expected to alter the photospheric abundance pattern of the star significantly. An abundance analysis of the strongest diamond source HD~97048 has produced upper limits on the abundances of a few heavy elements. The Sr abundance is at least one order of magnitude less than the solar abundance of this element. No evidence supporting the supernova hypothesis is found.  Except for HD~97048 and Elias~3$-$1, both nanodiamond features have also been observed in the post-AGB binary HR~4049 \\citep{geballe89}. The post-AGB phase is very short ($\\sim$10,000~yr) compared to the lifetime of an intermediate-mass pre-main-sequence star. Due to the short timescales, it is unlikely that a nearby supernova is the cause of the presence of nanodiamonds in this system. However, the oxygen isotopes in the circumbinary disk of HR~4049 display peculiar behaviour: the relative $^{17}$O and $^{18}$O abundances are two orders of magnitude larger than the surface abundances in other evolved stars \\citep[$^{16}$O/$^{17}$O~=~8.3$\\pm$2.3 and $^{16}$O/$^{18}$O~=~6.9$\\pm$0.9,][]{cami01}. These exceptional values are most likely related to the binary nature of HR~4049, and possibly to a nova outburst. The presence of nanodiamonds in the circumstellar environment may be related to such an event as well.  The extraneous origin of the circumstellar diamonds is disputed, however. \\citet{dai02nature} have stated that, if the supernova hypothesis is true, nanodiamonds should be abundant in solar-system comets as well. These objects are believed to have formed further out in the early solar system and are likely more pristine than meteorites. The authors have investigated fragile, carbon-rich IPD particles which enter the Earth's atmosphere with speeds in the range of cometary bodies, which suggests that the grains originate from these objects. Nanodiamonds are absent or strongly depleted in such IPD grains, which indicates that diamonds are not uniformly abundant in the solar system. This may support the hypothesis that the detected meteoritic nanodiamonds have formed {\\em in situ}, and are not of presolar origin.  \\citet{goto00} has suggested that the X-ray hardness of the radiation field of the central star may play a decisive role in the formation of diamond. However, our analysis has shown no obvious link between the source's X-ray luminosity and presence/absence of nanodiamonds. Furthermore, the strength of the few detected diamond features appears to be independent from the X-ray strength of the central star as well.  The extraneous origin of nanodiamonds is an attractive hypothesis, because it naturally explains the paucity of nanodiamond sources and the seemingly very different properties of these targets. No evidence was found, however, to substantiate the claim of an external triggering source in the vicinity of the targets. In-situ formation of nanodiamonds is a viable alternative. However, it remains unclear why only very few disks can/could provide the specific conditions needed to produce diamond. The disk and stellar parameters of the diamond sources are not uncommon, nor do the targets form a consistent group within the Herbig stars. Selection effects in terms of the size, temperature or hydrogenation of the diamonds may nonetheless be present. Such effects could prevent the detection of features in the 3~micron region in spectra of the majority of circumstellar disks. A larger sample of systems displaying diamond emission is needed to distinguish between in-situ formation and an extraneous origin of the (nano)diamonds in Herbig Ae/Be stars and in our own solar system."
    },
    "0606/astro-ph0606185_arXiv.txt": {
        "abstract": "{ We used our database of VLT-UVES quasar spectra to build up a sample of 70 Damped Lyman-$\\alpha$ (DLA) or strong sub-DLA systems with total neutral hydrogen column densities of $\\log N($H\\,{\\sc i}$)\\ga 20$ and redshifts in the range $1.7<z_{\\rm abs}<4.3$. For each of the systems, we measured in an homogeneous manner the metallicities relative to Solar, [X/H] (with ${\\rm X}={\\rm Zn}$, or S or Si), and the velocity widths of low-ionization line profiles, $\\Delta V$. We provide for the first time evidence for a correlation between DLA metallicity and line profile velocity width, which is detected at the $6.1\\sigma$ significance level. This confirms the trend previously observed in a much smaller sample by \\citet{1998ApJ...494L..15W}. The best-fit linear relation is $[{\\rm X}/{\\rm H}]=1.55(\\pm 0.12)\\log\\Delta V -4.33(\\pm 0.23)$ with $\\Delta V$ expressed in km s$^{-1}$. The slope of the DLA velocity-metallicity relation is the same within uncertainties between the higher ($z_{\\rm abs}>2.43$) and the lower ($z_{\\rm abs}\\le 2.43$) redshift halves of our sample. However, the two populations of systems are statistically different. There is a strong redshift evolution in the sense that the mean metallicity {\\sl and} mean velocity width increase with decreasing redshift. We argue that the existence of a DLA velocity-metallicity correlation, over more than a factor of 100 spread in metallicity, is probably the consequence of an underlying mass-metallicity relation for the galaxies responsible for DLA absorption lines. Assuming a simple linear scaling of the galaxy luminosity with the mass of the dark-matter halo, we find that the slope of the DLA velocity-metallicity relation is consistent with that of the luminosity-metallicity relation derived for local galaxies. If the galaxy dynamical mass is indeed the dominant factor setting up the observed DLA velocity-metallicity correlation, then the DLA systems exhibiting the lowest metallicities among the DLA population should, on average, be associated with galaxies of lower masses (e.g., gas-rich dwarf galaxies). In turn, these galaxies should have the lowest luminosities among the DLA galaxy population. This could explain the difficulties of detecting high-redshift DLA galaxies in emission. ",
        "introduction": "Over the past decade, significant progress in our understanding of early galaxy evolution has been made with large samples of high-redshift galaxies drawn from deep multi-band imaging \\citep[][and references therein]{2003ApJ...592..728S}. However, even before the first surveys for Lyman-Break Galaxies (LBGs) had begun, samples of DLA absorbers observed on the lines-of-sight to distant quasars had been constructed \\citep{1986ApJS...61..249W,1995ApJ...454..698W}. These absorbers were thought at the time to be the best carriers of information on the population of high-redshift galaxies, but, despite many attempts to identify the galaxies responsible for DLA absorption lines [hereafter called DLA galaxies], only very few could be detected in emission \\citep[see, e.g.,][]{1993A&A...270...43M,1996Natur.382..234D,1999MNRAS.303..711L,1999MNRAS.305..849F,2002ApJ...574...51M,2004A&A...422L..33M,2004A&A...417..487C,2005MNRAS.358..985W}. However, there is little doubt that DLA systems arise from the densest regions of the Universe and are closely associated with galaxies. It is therefore crucial to establish the connection between the absorption-selected DLA systems and emission-selected galaxies. In addition, the detailed information that becomes available only through the combination of morphology, colour and luminosity, with QSO absorption-line spectroscopy, makes these galaxy/absorber associations unique laboratories to study the physical processes at work during galaxy formation \\citep[see][]{1999ApJ...522..604P}.  Progress in this field has been slow. Firstly, a huge amount of work is needed to derive important parameters in DLA systems such as gas kinematics, metallicity, or dust and molecular fractions \\citep[see, e.g.,][]{1997ApJ...486..665P,1997ApJ...487...73P,1998ApJ...507..113P,1999ApJS..121..369P,2001ApJ...552...99P,1998A&A...337...51L,2003MNRAS.346..209L}. Secondly, as mentioned above, the known high-redshift DLA systems have proved to be very difficult to detect in emission. This has caused some confusion and for a while suggestions were put forward that DLA absorbers may not be related to high-redshift galaxies at all. \\citet{1998MNRAS.295..319M} and \\citet{1998ApJ...495..647H,2000ApJ...534..594H} resolved this issue showing that the difficulty of detecting high-redshift DLA systems in emission is an unavoidable consequence of the absorption cross-section selection which tends to reveal faint galaxies because they have an integrated cross-section larger than that of bright galaxies \\citep[see also][]{1999MNRAS.305..849F}.  Recently, \\citet{2004A&A...422L..33M} tentatively suggested that, if a galaxy luminosity-metallicity relation similar to that observed at $0\\la z\\la 1$ \\citep[e.g.,][]{2002ApJ...581.1019G,2003ApJ...599.1006K,2004MNRAS.350..396L,2004ApJ...613..898T} was already in place at high redshifts, then it would be possible to significantly increase the DLA galaxy detection probability by carefully selecting DLA systems with the highest metallicities. In fact, the few DLA galaxies that have to date been identified in emission do give support to the conjecture that a luminosity-metallicity relation was already in place at $z\\approx 2-3$, although the result is only marginally statistically significant \\citep[][see also Christensen et al., in prep.]{2004A&A...422L..33M}. This is in line with the near-Solar or even super-Solar metallicities derived for bright Lyman-break or bright K-band selected galaxies at similar redshifts \\citep{2004ApJ...612..108S,2004ApJ...608L..29D}. A mass-metallicity relation has recently been put into evidence for UV-selected star-forming galaxies at $z\\sim 2.3$ by \\citet{2006astroph0602473E}.  In this paper, we provide for the first time evidence for the existence of a velocity-metallicity correlation for high-redshift DLA galaxies that could be the consequence of an underlying mass-metallicity relation for the galaxies responsible for DLA absorption lines. From the observation of a sample of 17 DLA systems at $z_{\\rm abs}<3$, \\citet{1998ApJ...494L..15W} previously showed that the DLA systems exhibiting the largest line profile velocity widths span a narrow range of high metallicities. However, these authors also suggested that systems with small velocity widths span a wide range of metallicities. Recently, \\citet{2003MNRAS.345..480P} found a hint of an increase of the mean DLA metallicity with increasing velocity width, but the statistical significance of their result is low. In this paper, we use our database of VLT-UVES quasar spectra to build up a sample of 70 DLA or strong sub-DLA systems with total neutral hydrogen column densities of $\\log N($H\\,{\\sc i}$)\\ga 20$ and redshifts in the range $1.7<z_{\\rm abs}<4.3$. We present new, homogeneous measurements of DLA metallicities and line profile velocity widths in Sect.~\\ref{sample} and the observed velocity-metallicity relation in Sect.~\\ref{relation}. We discuss in Sect.~\\ref{discussion} the use of the DLA gas kinematics as a proxy for the mass of DLA galaxies and the possibility of the existence of a mass-metallicity relation for high-$z$ DLA galaxies. We conclude in Sect.~\\ref{conclusions}. ",
        "conclusions": "Using a sample of 70 DLA or strong sub-DLA systems with total neutral hydrogen column densities of $\\log N($H\\,{\\sc i}$)\\ga 20$ and redshifts in the range $1.7<z_{\\rm abs}<4.3$, we have shown that there is a correlation between metallicity ([X/H]) and line profile velocity width ($\\Delta V$) at the $6.1\\sigma$ significance level. The best-fit linear relation is $[{\\rm X}/{\\rm H}]=1.55(\\pm 0.12)\\log\\Delta V -4.33(\\pm 0.23)$ with $\\Delta V$ expressed in km s$^{-1}$. We argued that the existence of a DLA velocity-metallicity correlation, over more than a factor of 100 spread in metallicity, is probably the consequence of an underlying mass-metallicity relation for the galaxies responsible for DLA absorption lines. Assuming a simple linear scaling of the galaxy luminosity with the mass of the dark-matter halo, we found that the slope of the DLA velocity-metallicity relation is consistent with that of the luminosity-metallicity relation derived for local galaxies. If the galaxy dynamical mass is indeed the dominant factor setting up the observed DLA velocity-metallicity correlation, then the DLA systems exhibiting the lowest metallicities among the DLA population should, on average, be associated with galaxies of lower masses.  Eq.~\\ref{eq5} implies that the more than two orders of magnitude spread in DLA metallicity could reflect a more than ten magnitudes spread in DLA galaxy luminosity. Even though low-mass galaxies, i.e., gas-rich dwarf galaxies, can undergo periods of intense star formation activity and, in this case, have high luminosities in the UV, it is a fact that, on average, they show lower star-formation rates than more massive galaxies (\\citealt{2004MNRAS.351.1151B}; see also \\citealt{2004ApJ...603...12O}). This may well explain the difficulty of detecting high-redshift DLA galaxies in emission \\citep[e.g.,][]{2000ApJ...536...36K,2001ApJ...551...37K}. Furthermore, the non-detection of Ly$\\alpha$ emission from the galaxies responsible for low-metallicity DLA systems, down to Ly$\\alpha$ fluxes fainter than most of the Ly$\\alpha$ emitters from the deep survey of \\citet{2003A&A...407..147F}, could be a consequence of their low masses and their correspondingly, on average, low star-formation activity.  A significant fraction of DLA galaxies could be actively forming stars for some period of time, but due to their small masses the bursts of star-formation would not be powerful enough and/or would have too short life-times. Ly$\\alpha$ emission would simply be too faint or would be strong for a too short spell of time to be detected by the current generation of 8-10\\,m class telescopes. Conversely, the existence of a DLA mass-metallicity relation can explain the recent, tentative result by \\citet{2004A&A...422L..33M} that the few DLA systems with detected Ly$\\alpha$ emission have higher than average metallicities (see also Christensen et al., in prep.). This should be confirmed by additional deep imaging of the fields of QSOs with carefully selected DLA absorbers.  The DLA velocity-metallicity correlation relation studied in this paper also needs to be investigated in the context of new high-resolution smoothed-particle hydrodynamics simulations including the effects of feedback in a self-consistent manner \\citep[see, e.g.,][]{2004MNRAS.348..435N}."
    },
    "0606/astro-ph0606466_arXiv.txt": {
        "abstract": "We have obtained EUV spectra between 90 and 255 \\AA\\/ of the comets C/2002 T7 (LINEAR), C/2001 Q4 (NEAT), and C/2004 Q2 (Machholz) near their perihelion passages in 2004 with the Cosmic Hot Interstellar Plasma Spectrometer (CHIPS). We obtained contemporaneous data on Comet NEAT Q4 with the $Chandra$ X-ray Observatory ACIS instrument, marking the first simultaneous EUV and X-ray spectral observations of a comet. The total CHIPS/EUV observing times were 337 ks for Q4, 234 ks for T7, and 483 ks for Machholz and for both CHIPS and $Chandra$ we calculate we have captured all the comet flux in the instrument field of view. We set upper limits on solar wind charge exchange emission lines of O, C, N, Ne and Fe occurring in the spectral bandpass of CHIPS.  The spectrum of Q4 obtained with $Chandra$ can be reproduced by modeling emission lines of C, N O, Mg, Fe, Si, S, and Ne solar wind ions. The measured X-ray emission line intensities are consistent with our predictions from a solar wind charge exchange model. The model predictions for the EUV emission line intensities are determined from the intensity ratios of the cascading X-ray and EUV photons arising in the charge exchange processes.  They are compatible with the measured limits on the intensities of the EUV lines. For comet Q4, we measured a total X-ray flux of 3.7$\\times 10^{-12}$ ergs cm$^{-2}$ s$^{-1}$, and derive from model predictions a total EUV flux of 1.5$\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$. The CHIPS observations occurred predominantly while the satellite was on the dayside of Earth. For much of the observing time, CHIPS performed observations at smaller solar angles than it was designed for and EUV emission from the Sun scattered into the instrument limited the sensitivity of the EUV measurements. ",
        "introduction": "Soft X-ray emission has been detected from at least 20 comets since 1996.  After the consideration and rejection of several possible emission mechanisms, a consensus has been reached that the primary mechanism is charge exchange collisions between highly charged solar wind minor ions and neutral atoms and molecules of the cometary atmospheres, as suggested by \\citet{Cravens97}. The observations have been reviewed by \\citet{Cravens02}, \\citet{Kras04a} and \\citet{Lisse04}. The solar wind charge exchange (SWCX) process also occurs in the heliosphere where it contributes to the soft X-ray background \\citep{Cox88, Cravens00, Rob03, Pepino04, Lalle04} and may be responsible for some fraction of the short and long term enhancements seen at 250 eV \\citep{Sno95, Sno04}. These recent, interesting observations indicate it is important to understand both the microphysics of the charge-exchange process and the phenomenology of where, when and how this process occurs within the solar system.  Theoretical models of the X-ray and EUV emission lines produced in comets by charge exchange with the solar wind ions have been constructed by \\citet{Weg98}, \\citet{schwad00}, \\citet{khar00}, \\citet{khar01}, \\citet{khar03}, \\citet{Beiers03}, and \\citet{Lisse04}.  Observations of diffuse SWCX emission other than from comets have been carried out using $Chandra$ \\citep{War04} and XMM-Newton \\citep{Sno04}. \\citet{War04} observe the optically dark side of the moon and attribute the X-rays they detect to charge exchange collisions between solar wind ions and neutral hydrogen in the Earth's geocorona. \\citet{Sno04} ascribe the emission they see as coming from the Earth's exosphere or the heliosphere. In both cases, the emission is characterized by strong variability and correlation with the solar wind particle flux. X-rays generated by SWCX have been detected from the atmospheres of Mars and Venus \\citep{mur96, Den02, holm04, fok04, Gun04}.  In summary, the highly-ionized, heavy ions that stream out in the solar wind are a potential source of soft X-ray and EUV emission from within the solar system.  Charge exchange emission may occur involving neutral atoms and molecules from comets entering the inner solar system, planetary exospheres, and neutral atoms in the inter-planetary medium. A highly-ionized solar wind ion may give rise to several photons as the ion captures electrons sequentially in its flight from the Sun.  A detailed ionization and density map of the heavy solar wind ions as a function of time and position is required for a complete description of the X-ray and EUV emissions in the solar system. Measurements of SWCX emission from comets (and other measurements) can provide a useful probe of the density and ion composition of the solar wind.  Many SWCX lines are predicted to appear at EUV energies below the low-energy limit of $Chandra$ and XMM-Newton (Kharchenko \\& Dalgarno 2000, 2001, Kharchenko et al. 2003). Some of these EUV lines arise from ions not conclusively identified from emission at higher energies. These lines may be detected with CHIPS, a spectrometer sensitive in the EUV with resolution of several hundred. The lines appear to be present in the ROSAT PSPC spectrum of C/1990 K1 (Levy) (\\citep{Den97}) which rises steeply at energies below 200 eV. \\citet{Kras01} observed comet C/1996 B2 (Hyakutake) at low resolution at EUV wavelengths and claimed detection of lines from O, C, He and Ne. At lower energies, \\citet{khar01} predict that emission lines from the \\ion{O}{6} doublet at 1032 and 1038 A should be the brightest FUV lines SWCX lines. A sensitive search for EUV and FUV  emission lines cometary SWCX was conducted with the Far-Ultraviolet Spectroscopic Explorer (FUSE) observatory \\citep{Weaver02, Feldman05}. \\citet{Weaver02} report a a tentative detection of 1032 \\AA\\/ emission line from comet C/2000 WM1 (LINEAR).   An alternative interpretation of this feature has been suggested by \\citet{Feldman05} from in FUSE observations of the brighter comet C/2001 Q4 (NEAT); they attribute this feature to H$_2$ fluorescent emission line pumped by radiation from \\ion{N}{3} and \\ion{O}{6} ions in the solar corona. The non-detection of the brightest FUV \\ion{O}{6} lines from SWCX may be because the relatively small field of view of FUSE allowed observation of only about 6000-7000 km around the cometary nucleus. For the Q4 FUSE observations, the peak brightness of the charge-exchange EUV and X-ray emissions was shifted by more than this amount from the comet center, hence the FUSE aperture may have missed it \\citep{Feldman05}.  The favorable perigee passages of the bright Comets C/2001 Q4 NEAT and C/2002 T7 LINEAR in May 2004 and Machholz in January 2005 provided an excellent opportunity to use the unique low energy spectroscopic capabilities of CHIPS to study SWCX in comets at EUV wavelengths.  Due to a combination of high solar activity and their close approaches, Comets C/2001 Q4 NEAT and C/2002 T7 LINEAR are the brightest X-ray targets available since Comet C/Hyakutake in 1996 and Comet 153P/2002 C1 (Ikeya-Zhang) in 2002.  In this paper, we discuss the observational data from CHIPS and $Chandra$, describe the procedures by which we measure the line strengths and discuss the implications for the cometary charge exchange process. ",
        "conclusions": "We have measured EUV emission from three comets and provided the first-ever simultaneous EUV and X-ray measurement and modeling of cometary SWCX emission.  The X-ray data for comet Q4 show clear emission lines that are satisfactorily interpreted as solar wind charge exchange lines from ions expected in the solar wind. We have used the same model to identify and predict the strengths of EUV lines.  The measured EUV line strengths are consistent with the theoretical model that reproduces the X-ray line intensities. The CHIPS data and model results show that the EUV lines from all three comets are quite faint.  Nonetheless, the high spectral resolution measurements of the EUV lines have now set useful limits on many emission lines in this band that may serve as a guide for future EUV observations of comets."
    },
    "0606/astro-ph0606716_arXiv.txt": {
        "abstract": "We present results of simulations performed with the Geant4 software code of the effects of Galactic Cosmic Ray impacts on the photoconductor arrays of the PACS instrument. This instrument is part of the ESA-Herschel payload, which will be launched in late 2007 and will operate at the Lagrangian L2 point of the Sun-Earth system. Both the Satellite plus the cryostat (the shield) and the detector act as source of secondary events, affecting the detector performance. Secondary event rates originated within the detector and from the shield are of comparable intensity. The impacts deposit energy on each photoconductor pixel but do not affect the behaviour of nearby pixels. These latter are hit with a probability always lower than 7\\%.\\\\ The energy deposited produces a spike which can be hundreds times larger than the noise. We then compare our simulations with proton irradiation tests carried out for one of the detector modules and follow the detector behaviour under 'real' conditions.\\\\ ",
        "introduction": "The European Space Agency Herschel satellite, scheduled for launch on August 2007 with an Ariane-5 rocket, will operate at the Lagrangian L2 point of the Sun-Earth system (see ESA web page: {\\it www.rssd.esa.int/SA-general/Projects/Herschel}). Herschel is the fourth cornerstone mission of ESA Horizon 2000 program and will perform imaging photometry and spectroscopy in the far infrared and submillimetre part of the spectrum. The Herschel payload consists of two cameras/medium resolution spectrometers (PACS and SPIRE) and a very  high resolution heterodyne spectrometer (HIFI).  The PACS (Photo-conductor Array Camera and Spectrometer) instrument will perform efficient imaging and photometry in three wavelength bands within 60-300$\\mu$m and spectroscopy and spectroscopic mapping with spectral resolution between 1000 and 2000 over the entire wavelength range (60-210$\\mu$m) or short segments.\\\\ PACS is made of four sets of detectors: two Ge:Ga photoconductor arrays for the spectrometer part and two Si-bolometer arrays for the photometer part. On each instrument side each detector covers roughly half of the PACS bandwidth (see Poglitsch et al., 2004). More about Herschel - PACS can be found in the PACS Web pages: {\\it pacs.mpe-garching.mpg.de} and {\\it pacs.ster.kuleuven.ac.be}.  It is well known that cosmic rays may influence strongly the detector behaviour in space. The performances of detectors such as ISOCAM and ISOPHOT on board of the Infrared Space Observatory (ISO) were largely affected for time periods long enough to corrupt a large amount of data. Our goal, in the present paper, is to exploit our present knowledge on detectors and cosmic environment to understand how the detector performances change from the expectations.  In this paper we focus our attention on the GeGa photoconductor arrays (hereafter PhC) of the PACS instrument, following its response under simulated cosmic rays irradiation. The detectors are similar to those on board of ISO and we expect therefore that they will be strongly affected by cosmic ray hitting.  In addition we compare proton irradiation tests, performed on a single module of the PhC, with Geant4 simulations and draw conclusions about their performances.  The paper is organized as follows: in Sections \\ref{simulations} through \\ref{PL}, the used simulation toolkit, the input detector design, the Galactic Cosmic Ray spectra and the Physics List are outlined. In Section \\ref{results} we report the results. Section \\ref{IPN} deals with the irradiation tests which were performed at UCL-CRC on one detector module and the comparison with our simulations. In the final Section \\ref{discussion} we summarize the results reported in this draft. ",
        "conclusions": "\\label{discussion} \\subsection{The photoconductors in the space environment} Our simulations of the effects of Galactic Cosmic Ray impacts on the photoconductor arrays show that both the satellite plus the cryostat (the shield) and the detector act as source of secondary events, affecting the detector performance. Secondary event rates originated within the detector and from the shield are of comparable intensity. The impacts deposit energy on each photoconductor pixel but do not affect the behaviour of nearby pixels. These latter are hit with a probability always lower than 7\\%.\\\\ The energy deposited produces a spike which can be hundreds times larger than the noise. The present simulations are not able to follow the temporal behaviour of each pixel and cannot be used to determine the shape of the output signal as a function of time.  \\subsection{The photoconductors in the lab test} We have simulated the experiment carried out at UCL-CRC and compared the simulation outputs to the measurements reported in Groenewegen \\& Royer (2005). We find similar rates and events (see Tables \\ref{MGvsG4}) on each pixel. The simulated energy distribution differs substantially from that of the measurements.  We tried to ascribe these differences to the following causes: \\begin{enumerate} \\item As far as the event number and rate is concerned, it may be possible that the chosen cuts are rather large. But due to the complexity of the experimental hall it is hard to tune them optimally. We would need information on the intermediate passage of the beam through matter. A simple correction could be a slight increase of the Ge cuts. A clear benchmark would be the comparison between the energy distributions. \\item The presence of a primary and a secondary peak seems to be the only feature common between experimental and simulated data. What it is not clear is why, due to the crossing through matter, the Gaussian beam is ({\\it should be}) transformed as in Figure \\ref{estop05}, that is first a tail, then a peak, whereas experimental data favour the tail next to the secondary peak. Either Ge has some physical properties which are not included in the G4 Physics List, or the deglitching algorithm used is not correct. \\end{enumerate}  \\begin{table} \\begin{tabular}{rrr} \\hline & \\cite{MGb}    & Geant4    \\\\ Files       & Events    & Events\\\\ \\hline L26-L29     &   755 &  853  \\\\ H3-H6       & 29008 &29430  \\\\ \\hline \\end{tabular} \\caption[]{Measured versus predicted values.}\\label{MGvsG4} \\end{table}"
    },
    "0606/astro-ph0606520_arXiv.txt": {
        "abstract": "The coincidence problem of late cosmic acceleration is a serious riddle in connection with our understanding of the evolution of the Universe. In this paper we show that an interaction between the dark energy component (either phantom or quintessence) and dark matter can alleviate it. In this scenario the baryon component is independently conserved. This generalizes a previous study [S. del Campo, R. Herrera, and D. Pav\\'{o}n, Phys. Rev. D \\textbf{71}, 123529 (2005)] in which neither baryons nor phantom energy were considered. ",
        "introduction": "According to the conventional picture, the present accelerated expansion of the Universe is driven by the negative pressure of an unknown and unclustered component (dubbed ``dark energy\") that currently contributes about $70$\\% of the total density. The remaining $30$ per cent is shared between cold dark matter ($\\rho_{dm} \\sim 25$\\%) and cold baryons ($\\rho_{b} \\sim 5$\\%) \\cite{accel,wmap,hzt,snls,wmap3y}. The two latter, being pressureless, redshift with expansion faster than the dark energy component. Then, the cosmic coincidence problem arises, ``Why are the densities of matter and dark energy  of precisely the same order today?\" \\cite{coincidence}. Clearly, this is one outstanding riddle in our understanding of the Universe. To solve it one is forced to  adopt an evolving dark energy field (either quintessence or phantom energy) or accept an incredibly  tiny cosmological constant and admit that the ``coincidence\" is just a coincidence that hopefully might be somewhat alleviated with the help of the anthropic idea \\cite{reviews}. Here we take the view that before resorting to the second option we should further explore the first one. Yet, an evolving dark energy cannot solve the problem either unless a suitable interaction (coupling) with matter is allowed \\cite{interaction,recent}. Note that the coupling alters the rate at which  both matter and dark energy redshift with expansion; this is why it can potentially alleviate the aforesaid problem.  Interestingly, rather than in connection with the coincidence problem, which was not even formulated at the moment, this interaction was first proposed as a mechanism to reduce the value of the cosmological constant \\cite{wetterich}. Notice that in the absence of underlying symmetry that would suppress the coupling matter--dark energy there is no a priori reason to dismiss it. In the last years, various proposals at the fundamental level for the coupling leading to a constant ratio matter/dark energy at late times were advanced (see e.g., \\cite{federico} and references therein), and specific phenomenological models have been built and contrasted with observational data (high redshift supernovae and CMB anisotropies) and seen to pass the tests \\cite{tests}. Further, the Akaike \\cite{akaike} and Bayesian informative criteria \\cite{bic} when applied to high redshift supernovae data suggest  a transfer of energy from the phantom component to the matter component; yet, the conventional $\\Lambda$CDM model is still preferred \\cite{szydlowsky}.  While a constant ratio matter/dark--energy at late times (including the present one) would clearly alleviate the coincidence problem it should be noted that a much less strong condition would serve. It would suffice that nowadays the aforesaid ratio varies only slowly (i.e., less faster than the scale factor) and be of order unity (``soft coincidence\"). Recently, it was found within this approach that the quintessence scenario was favored over the tachyon scenario for the latter would imply an excessive amount of pressureless matter today \\cite{recent}. However, in order to circumvent the tight constraints on long--range forces \\cite{tight} the baryon component was left aside altogether. This may be justified because the analysis was restricted to times about the present one which, as said above, is characterized, among other things, by a low value to the baryon density. In this paper we generalize the analysis by including baryons in the energy density budget as an independently conserved component and extend the study to earlier times -though not to the radiation dominated era. As dark energy components quintessence and phantom are separately considered.  The outline of the paper is as follows. In section II we present the model. In section III we constrain it with recent high redshift supernovae data. Finally, in section IV we discuss and summarize our findings. As usual, a zero subscript or superscript attached to any quantity means that it should be evaluated at the present epoch. ",
        "conclusions": "We studied a model of late cosmic acceleration by assuming that the dark matter and dark energy components are coupled to each other so that there is a transfer of energy from the latter to the former, while the baryon component is conserved. By suitably choosing the interaction term, $Q$, the ratio between both dark sources of gravity is seen to evolve, at present time, less faster than the scale factor. This considerably alleviates the coincidence problem albeit it does not solve it in full. Clearly, to achieve the latter one should derive the present value of the aforesaid ratio, or at least show that it has to be of order unity. For the time being, $r_{0}$ ought to be understood as an input parameter. This also holds for a handful of key observational quantities such as the present value  of the cosmic background radiation temperature, the cosmological constant $H_{0}$, or the ratio between the number of baryons and photons.  The transition deceleration--acceleration occurs recently, at redshifs lower than unity (see Fig. \\ref{fig:deceleration}). This contrasts with other interacting models in which the transition is predicted to take place much earlier, at redshifhts as hig as $10$ \\cite{transition}.  It is noteworthy that the expression for the potential, Eq. (\\ref{Vphia}), is identical irrespective of whether the dark energy component is a phantom or a quintessence field.  Both when the dark energy is a quintessence or a phantom field the model fits rather well the HZT data set, and in the latter case clearly better than the concordance model $\\Lambda$CDM does ($\\chi^{2}_{phantom} = 173$, $\\chi^{2}_{quintessence} = 177$, $\\chi^{2}_{\\Lambda} = 178$). Yet, the Bayesian information criterion (BIC) \\cite{bic}, given by the formula BIC$= \\chi^{2} +% p\\, \\ln N$, where $p$ is the number of free parameters of the model ($2$ for the $\\Lambda$CDM model, $4$ for dark energy models) and $N$ the number of data points, distinctly favors the $\\Lambda$CDM model for it yields a lower figure (BIC$_{quintessence} \\simeq 197$, BIC$_{phantom} \\simeq 193$, BIC$_{\\Lambda} \\simeq 188$).  The fits to the SNLS data are rather similar ($\\chi^{2}_{phantom} = 111.02$, $\\chi^{2}_{quintessence} = 111.09$, $\\chi^{2}_{\\Lambda} = 111.03$). However, the Bayesian information criterion again sides with the $\\Lambda$CDM model (BIC$_{quintessence} \\simeq $ BIC$_{phantom} \\simeq 128$, BIC$_{\\Lambda} \\simeq 119$).  We have not used the  Akaike criterion because it favors models with larger number of parameters \\cite{favors}.  Thus, from the point of view of the Bayesian information criterion the $\\Lambda$CDM model is  preferred, and  the soft coincidence model for quintessence seems clearly disfavored if not discarded. As for the phantom model, the situation appears somewhat undecided. While, regarding the BIC, it lags behind the concordance model it shows a better fit to the HZT sample of supernovae and alleviates the coincidence problem. We may summarize by saying that the SNIa data are not yet conclusive thereby a richer supernovae statistics, especially at redshifts larger than unity, is needed to come to a verdict.  Unfortunately, the supernovae constraints on the $c^{2}$ parameter is rather poor. This is coherent with the fact that the interaction dark energy--dark matter affects the luminosity distance only at third order in the redshift \\cite{grg1-2}. Nevertheless, we hope to be able to break the degeneracy by submitting the model to further tests such as the CMB temperature anisotropies and the distribution of matter at cosmological scales. This will be the subject of a future work."
    },
    "0606/astro-ph0606430_arXiv.txt": {
        "abstract": "Several of the current and next-generation cosmic microwave background (CMB) experiments have polarimetric capability, promising to add to the finesse of precision cosmology.  One of the contaminating Galactic foregrounds is thermal emission by dust.  Since optical interstellar polarization is commonly seen, from differential extinction by aligned aspherical dust particles, it is expected that thermal emission from these grains will be polarized.  Indeed, in the Galactic plane and in dark (molecular) clouds, dust emission in the infrared and submillimetre has been measured to be polarized.  It seems likely that the faint diffuse cirrus emission, of more relevance to CMB experiments, will be polarized too.  We discuss how well the amount of polarization of this component can be predicted, making use of what is known about optical (and infrared and ultraviolet) interstellar polarization and extinction.  Some constraints on the alignment of the carrier of the dust-correlated anomalous microwave emission can be made as well. ",
        "introduction": "\\label{int}  Several of the current and next-generation cosmic microwave background (CMB) experiments have the capability of measuring linear polarization, a complementary probe for precision cosmology (e.g., Laureijs \\cite{lau06}; Polenta \\etl\\ \\cite{pol05}; Beno\\^it \\etl\\ \\cite{ben04}; Prunet \\cite{pru06}). Among these experiments are Archeops (Beno\\^it \\etl\\ \\cite{ben04}), BOOMERanG (B2K in 2003; Polenta \\etl\\ \\cite{pol05}; Netterfiled \\cite{net06}), and the Planck Surveyor (Laureijs \\cite{lau06}; Prunet \\cite{pru06}; Lammare \\etl\\ \\cite{lam03}). All-sky surveys by IRAS from 12 to 100~\\micron\\ revealed faint diffuse emission (``cirrus'') everywhere, even at high Galactic latitudes (Low \\etl\\ \\cite{low84}). One of the diffuse foregrounds contaminating the CMB signal is thermal emission by this diffuse interstellar dust (e.g., de Oliveira-Costa \\etl\\ \\cite{deo04}; Tucci \\etl\\ \\cite{tuc05}), expected to correspond most closely spatially to the 100~\\micron\\ emission (\\S~\\ref{comp}). This becomes the dominant foreground above about 100 GHz ($\\lambda < 3$~mm $=3000$~\\micron), and as the CMB fades toward higher frequencies, the rising dust spectrum becomes the dominant signal (Fig.~\\ref{contam}). Where polarization is concerned, some further caution is needed in defining where the dust dominates.  The frequency above which the polarized dust intensity exceeds the polarized synchrotron background is probably somewhat higher because the synchrotron component is likely to have a higher degree of polarization.  Of course, it is the spatial fluctuations, which will vary from component to component, that are in the end important to CMB contamination. Multifrequency observations through the transitional range are therefore a prerequisite for separating these signals (component separation, not a specific topic of this paper; see e.g., Jaffe \\etl\\ \\cite{jaf04}).  \\begin{figure} \\includegraphics[width=5cm]{giardlagachefig2_3000.eps} \\qquad \\includegraphics[width=5cm]{olivfig1_3000.eps} \\caption{Left: Foreground components contaminating the CMB at two representative latitudes. From www.planck.fr/heading136.html Giard \\& Lagache.  Right: Alternative presentation from de Oliveira-Costa \\etl\\ (\\cite{deo04}) showing the anomalous emission as well. } \\label{contam} \\end{figure}  Some CMB experiments like B2K minimize cirrus contamination by concentrating measurements in regions of low foreground column density (Netterfield \\cite{net06}). One strategy for cirrus mitigation could therefore be to mask out regions of bright cirrus.  But only 20\\% of the sky has an HI column density below $10^{20}$~cm${-2}$ (Miville-Desch\\^enes, private communication) and even that produces a non-negligible foreground ($\\sim 1$~MJy~sr${-1}$ at 100~\\micron).  The precision sought by instruments like the Planck Surveyor depends on broad sky coverage (Prunet \\cite{pru06}), and of course Planck will in any case produce data for the entire sky. Thus, as an alternative strategy, one needs to measure the properties of cirrus at higher frequencies where the CMB is not important, and then extrapolate to lower frequencies where one does have to address component separation.  To inform/confirm the extrapolation, the cirrus properties at the lower frequencies can be assessed empirically in regions of brighter cirrus.  The spectral properties of cirrus are of course an important diagnostic of the interstellar dust.  Since optical polarization is commonly seen, from differential extinction by aligned aspherical dust particles, it is expected that thermal emission from these grains will be polarized (Stein \\cite{ste67}). Indeed, in the Galactic plane and in dark (molecular) clouds, dust emission in the infrared and submillimetre has been measured to be polarized (Vaillancourt \\cite{vai06}; Hildebrand \\etl\\ \\cite{hil00}; Vaillancourt \\etl\\ \\cite{val03}).  In these dense regions, the interpretation of the polarization is complicated by many issues: beam dilution, distortions in the magnetic field topology, changes in the degree of alignment within clouds, grain evolution, and a range of grain temperatures and optical depth which affect which grains dominate the emission and polarization in various parts of the infrared to submillimetre spectrum.  These are not considered here.  It seems likely that the faint diffuse cirrus emission will be polarized too, and this will not be affected by the above-mentioned complications.  Diffuse dust polarization is of more relevance to CMB experiments.  Polarimetric observations in the transitional high frequency range (see Table~\\ref{angular}) were made by Archeops at 353 GHz (Beno\\^it \\etl\\ \\cite{ben04}) and B2K and MAXIPOL at lower frequencies (Polenta \\etl\\ \\cite{pol05}; Johnson \\etl\\ \\cite{jon03}), precursors to Planck HFI. Archeops has reported detecting the dust polarization (\\S~ref{sobs}). From the point of view of best examining the dust polarization, Planck has no polarimetric capability at 545 and 857~GHz where the dust polarized intensity would be the strongest.  However, balloon-borne experiments like PILOT (Bernard \\cite{ber05}) and BLAST-pol will map the diffuse polarization with great sensitivity at higher frequencies.  This paper attempts to quantify the degree to which the diffuse dust emission is polarized, a contribution toward both component separation and dust physics.  \\begin{table}[!b] \\caption{Angular Resolution$^1$ of High-Frequency CMB Polarization Experiments} \\begin{center} \\begin{tabular}{lccccccccc} \\hline \\\\ $\\nu$ (GHz) &      100  &  143 & 150  & 217  & 240  & 345 & 353  & 545 & 1250 \\\\ $\\lambda$ (\\micron) & 3300 & 2100 & 2000 & 1380 & 1250 & 870 & 850  & 550 & 240 \\\\ \\hline \\\\ Archeops &&&&&&& 13 \\\\ % BOOMERanG 2K &&& 10 && 7 & 7 \\\\  % MAXIPOL & & 10 && & 10 \\\\ PILOT & &&&&&&& 2.3 & 1 \\\\ Planck HFI & 9.2 &7.1 && 5 &&& 5\\\\ \\hline \\\\ \\end{tabular} \\end{center} {$^1$ FWHM in arc minutes} \\\\ \\label{angular} \\end{table}  WMAP (Kogut \\etl\\ \\cite{kog03}) and Planck have polarimetric capability at lower frequencies, down to 23 GHz and 30 GHz respectively, which can be used to examine the polarization of two other CMB foregrounds, synchrotron emission and the dust-correlated anomalous microwave (hereafter simply ``anomalous'') emission (Draine \\& Lazarian \\cite{dra98}; Finkbeiner \\etl\\ \\cite{fin99}; Lazarian \\& Finkbeiner \\cite{laz03}; Finkbeiner \\cite{fin04}; Finkbeiner \\etl\\ \\cite{finl04}; de Oliveira-Costa \\etl\\ \\cite{deo04}; Davies \\cite{davies06}; Davis \\cite{davis06}). Some constraints on the alignment of a dust carrier of the latter emission, which would be diagnostic of the emission mechanism, are made here as well (\\S~\\ref{anomalous}).  \\subsection{A Legacy from Optical Polarization} \\label{legacy}  Both optical interstellar polarization and polarization of thermal emission depend on aspherical (flattened or elongated) grains that are aligned. The orientation of the electric vector of the emitted radiation is along the long axis of the mean grain profile projected on the plane of the sky, while that of the optical polarization, being of transmitted light suffering differential extinction, is orthogonal to this. Particular ingredients in alignment theories are that grains are rapidly spinning about their axis of largest moment of inertia and that this is on average aligned with the Galactic magnetic field, so that consistent with observations the electric vector of optical polarization is along the direction of the projected magnetic field, while it is perpendicular for emission (see Fig.~\\ref{dustrad}). There are several potentially viable alignment mechanisms as discussed elsewhere (Lazarian \\& Finkbeiner \\cite{laz03}; Lazarian \\& Yan \\cite{laz04}; Roberge \\cite{rob04}; Vaillancourt \\cite{vai06}) and in the Appendix. Nevertheless, it is still difficult to predict the degree of alignment with certainty {\\it ab initio}, or as will become clear, gauge other important factors, and so a focus of this paper is to assess how well the degree of polarization of the diffuse cirrus component can be predicted semi-empirically.  \\begin{figure}[!b] \\includegraphics[width=8cm]{dust_radiation_3000.eps} \\caption{The relationship of the polarization of emitted and transmitted radiation to the mean grain profile, aligned with respect to the Galactic magnetic field.  In the far-infrared and submillimetre, there is more radiation emitted with polarization E-vector along the long axis.  This is also the axis along which there is more extinction, and so the E-vector of the net transmitted radiation is orthogonal to this axis, and to the emitted polarization and the magnetic field.  Graphic from Ponthieu \\& Lagache at www.planck.fr/article263.html. } \\label{dustrad} \\end{figure}  To predict the degree of polarization of the dust-related CMB foreground(s), we draw on what can be deduced about alignment from optical (and infrared and ultraviolet) interstellar polarization and extinction, concentrating on diffuse dust associated with atomic gas. Still, it should be cautioned that the optical measurements are by necessity largely for column densities with $E_{B-V} > 0.1$ and so $N_H > 6 \\times 10^{20}$~cm$^{-2}$. This corresponds to lines of sight with 100-\\micron\\ cirrus brightness $I_{100} > 6$~MJy~sr$^{-1}$.  While there will be many such directions in which to examine the polarization of the low-frequency foreground cirrus with forthcoming experiments, for CMB applications one needs to work at even lower column densities too. Another issue for any intercomparison of optical polarization and the polarization of low-frequency dust emission is that the Galaxy is optically thin for the latter.  Therefore, we see the whole Galaxy, or right out of it, unlike probes with stars which rely on differential extinction along that path.  But at high latitude, most relevant to the CMB, the effective paths might not be too dissimilar if the stars used were sufficiently distant.  A lot of attention goes into the degree of polarization $p = P/I$, the ratio of the polarized intensity $P$ to the total emission $I$. Non-polarizing grains (or other emission components entirely) can contribute to $I$, potentially confusing the interpretation of any spectral dependence of $p$.  We therefore also emphasize the importance of examining $P$ and its spectral dependence directly, as a diagnostic of the aligned grains. $P$ is of course what is relevant to the calculation of the power spectrum of the $E$ and $B$ modes of the polarized intensity. We do not attempt to predict these (Prunet \\etl\\ \\cite{pru98}; Tucci \\etl\\ \\cite{tuc05}) since they depend in addition on the spatial variation of the degree and orientation of the alignment, whose statistical properties are not known (i.e., the statistics of $P$ are different than the statistics of $p_{em}I$).  However, the dust contribution to the power spectrum should scale with the same frequency dependence as $P^2$ (from dust), unless hoped-for ``spatial -- electromagnetic frequency decoupling'' breaks down, i.e., there were different populations of grains with different alignment statistics whose relative contributions to $P$ (and $I$) changed with frequency. A $TE$ measurement of the dust emission at 353~GHz by Archeops (Ponthieu \\etl\\ \\cite{pon05}) provides a basis from which the potential contamination at lower frequencies can be estimated more directly. ",
        "conclusions": "\\label{results}  The results from the different steps of this recipe are collected in Table~\\ref{predictt}.  There $pa_{em}$ is evaluated at 350~GHz, but this choice is inconsequential (Fig.~\\ref{figpem}).  \\begin{table}[!b] \\caption{Prediction of the Net Submillimetre Polarization} \\begin{center} \\begin{tabular}{lccccccc} \\hline \\\\ &\\multispan{4}\\, Oblate\\hfil &\\multispan{3}  Prolate\\hfil\\\\ Axial ratio      & $\\sqrt{2}$ & 2    & 4    & 6    & 2$^1$& 2$^2$& 4$^2$\\\\ \\hline \\\\ $(p/\\tau)_{ex}$   & 0.04 & 0.08 & 0.12 & 0.12 & 0.05 & 0.09 & 0.14 \\\\ $R$               & 0.67 & 0.36 & 0.23 & 0.23 & 0.53 & 0.30 & 0.19 \\\\ $pa_{em}$$^3$     & 0.31 & 0.56 & 0.83 & 0.90 & 0.39 & 0.56 & 0.91 \\\\ $R pa_{em}$       & 0.21 & 0.20 & 0.19 & 0.21 & 0.21 & 0.17 & 0.17 \\\\ $d R pa_{em}$ (\\%)& 9.3  & 8.9  & 8.7  & 9.4  & 9.3  & 7.5  & 7.8  \\\\ \\hline \\\\ \\end{tabular} \\end{center} {$^1$ perfect spinning alignment} \\\\ {$^2$ picket fence alignment} \\\\ {$^3$ evaluated at 350~GHz} \\\\ \\label{predictt} \\end{table}  For perfectly aligned oblate silicate particles ($R \\sim 1$ in this case), the axial ratio needs to be no higher than 1.4 to produce the maximum $(p/\\tau)_{ex}$ observed.  For larger axial ratios, grains must be somewhat disaligned by a quantifiable amount ($R < 1$) to produce the same $(p/\\tau)_{ex}$. The value of $pa_{em}$ is large, and as expected higher for the more extreme axial ratios (Fig.~\\ref{figpem}).  The application of the recipe to the models for axial ratios closer to unity is probably more reliable, and these ratios are possibly the most realistic as well. But it turns out that the product $R pa_{em}$ is fairly robust, depending little on the shape and axial ratio.  \\subsection{Dilution}\\label{dilute}  Non-aligned grains contribute to the thermal emission and dilute the maximum polarization expected.  The Kim \\& Martin models used involve silicates and (large) graphite particles (Kim \\etl\\ \\cite{kim94}), following Mathis \\etl\\ (\\cite{mat77}) and Draine \\& Lee (\\cite{dra84}).  The relative contributions to the submillimetre emission can be judged from calculations by Draine \\& Anderson (\\cite{dra85}), which do produce about the right amount of far-infrared emission per unit column density.  This suggests quite a reduction, by $d \\sim 1/3$.  A more recent variant of this model, with slightly different size distributions and apportionment of material (Li and Draine \\cite{lid01}), gives $d \\sim 0.58$ at 353~GHz (their Fig.~9).  We adopt the mean, $d \\sim 0.45$, and note another major contribution to the systematic uncertainty.  Note that the dilution of $(p/\\tau)_{ex}$ in the optical by extinction by unaligned graphite (carbonaceous) particles has resulted in a larger $R_{ex}$, and so this low $d$ can be considered as payback, compensating in the product $dR$.  Likewise, if there were optical polarization by large carbonaceous particles, there would be an accompanying decrease in $R$ but an increase in $d$, and so this change too would not produce a greatly different net $p_{em}$ in the submillimetre.  \\subsection{Predicted Net Polarization $p_{em}$}\\label{predict}  This self-consistent model then predicts a maximum net polarization $d R pa_{em}$ as listed in Table~\\ref{predictt} and summarized as the median $p_{em} (\\%) = 8.9 \\pm 0.7 \\pm 3.5 $. The result is not very dependent on the shape or axial ratio of the grains, as reflected in the modest rms 0.7\\% among the different models. The systematic uncertainty 3.5\\% % is much more substantial, with the rough estimate based on actual application of the recipe to the models: calculating $\\tau_{ex}$ appropriately (step (ii); \\S~\\ref{disor}), the (frequency dependent) dilution $d$ near 353~GHz (step (vi); \\S~\\ref{dilute}), and encompassing the assumption $R_{em} = R_{ex}$ (step (v); \\S~\\ref{disor}).  Thus the maximum polarization to be expected is quite large, even without the possible boost accompanying spectral flattening of the amorphous silicate emissivity through 350~GHz.  But when averaged over large regions with (i) non-uniform alignment (beam dilution) or (ii) less than the optimal alignment, including the effects changes in the direction of the magnetic field, or unfavorable field orientation (something that changes systematically on a large scale in the Galaxy), or (iii) alignment which changes along the line of sight, then typically half of this might be expected, judging from Fig.~\\ref{smf} which shows the depolarization from (ii) and (iii). This would still leave $p_{em} \\sim 5$\\%.  \\subsection{Observations}\\label{sobs}  Beno\\^it \\etl\\ (\\cite{ben04}) present Archeops measurements of polarization in the Galactic plane at 353~GHz.  They find $p_{em} \\sim 4 - 5 \\%$ for the diffuse emission averaged over several square degrees and even higher values in some large clouds. The orientation of the $E$ vector roughly perpendicular to the plane is reassuringly as expected. However, no clear correlation with optical interstellar polarization (for which the path lengths to the background stars would be shorter than for the emission) has been established and the data are not of high enough signal to noise to follow this at high resolution. Because of the long lines of sight and high column densities, the Galactic plane, even its diffuse emission, is probably not the best place to evaluate the polarization of the higher latitude cirrus foreground of the CMB. The amount of polarization observed is actually close to what is predicted by the analysis above.  Although such estimates were actually made in 2003, before the Archeops results, this can hardly be taken as a triumph, given the large systematic uncertainties and the extra uncertainties of modeling the emission in the Galactic plane.  Nevertheless it does support the view that the polarization of the cirrus at higher latitude will be significant.  Although the variations in alignment along and across the line of sight at high latitude are not known, it seems possible (because of the shorter paths and thus simpler geometry being integrated) that these variations will be less than in the Galactic plane, resulting in less depolarization.  \\subsection{Spectral Behaviour of $p_{em}$ and $P$} \\label{spectp}  As noted above, the spectral index of the amorphous silicate emissivity might be less steep than the $\\beta \\sim 2$ inherent in our calculations.  A modest change was introduced by Li \\& Draine (\\cite{lid01}, their eq.~1 and Fig.~9) to fit the frequency dependence of the diffuse emission.  Boudet \\etl\\ (\\cite{bou05}) suggest even stronger variations with frequency over the range 150 -- 3000~GHz. In this case, the factor $d$ would be tend to increase with decreasing frequency as the relative importance of large silicates to the submillimetre emission rose.  In models where it is the silicates that produce the polarization (subscribed to here), this frequency dependence of $d$ would result in an increase in the net $p_{em}$ at the lower frequencies and our estimate at 353~GHz could be low by a factor $\\sim 1.6$ (not included in the systematic errors below). It will be interesting to learn the results from B2K (and Planck) which measure within this interesting range of frequencies.   Because of different weighting of different components, the spectral dependence of the polarized intensity $P$ can be different than that of total emission, $I$ (\\S~\\ref{emission}). Through dilution, non-aligned grains lower the net polarization. If the frequency dependence of emission for the diluting component is different, then this introduces a frequency dependence into $p_{em} = P/I$.  We have just mentioned one example.  To characterize the intrinsic emission properties of the carriers of the polarization more directly, this confusion from dilution can be avoided in principle by examining $P$, which would isolate the spectral dependence of the polarizing emitters.  For frequency ranges over which $pa_{em}$ for this component could be assumed to be constant, this would directly give $\\beta$, and its variations, for the spectral emissivity ($I$) of this aligned-grain component.  $P$ is also the quantity on which to base statistical evaluation of the spatial variations.  \\begin{figure}[!b] \\includegraphics[width=7cm, angle=90] {toymodel.ps} \\caption{Frequency dependence of observables $p_{em}$, $I$, and $P$ (solid curves, from top) in an illustrative model in which the value of $\\beta$ for the amorphous silicate component varies with frequency (see text).  The silicate and graphite constributions to $I$ are shown by the dashed and dot-dash curves, respectively (for normalization, see text). } \\label{toymodel} \\end{figure}  A toy model for illustration is presented in Figure~\\ref{toymodel}. Here it has been assumed that $\\beta$ for the aligned silicate component changes significantly with frequency, in the spirit of what has been inferred from astronomical observations and measurments in the laboratory (Boudet \\etl\\ \\cite{bou05}; J.-P.\\ Bernard, private communication), although to be sure this is not (yet) well constrained. Specifically the model uses $\\beta = 1.7$ for $\\lambda < 1$~mm ($\\nu > 300$~GHz), flattening (perhaps too abruptly in this illustrative model) to 0.5 through 2~mm (150~GHz), and then steepening to 1.5 (again for visual emphasis) beyond 5~mm ($\\nu < 60$~GHz).  For the silicate grains we take $T = 17$~K and $pa_{em}$ = 0.2 (0.089/0.45 as in \\S~\\ref{predict}.  In a more realistic model there would be a slight increase in $pa_{em}$ with decreasing frequency as a resultof impliedchanges in the complex refractive index of the silicate.  There is also an unpolarized graphite component with $T = 20.5$~K with relative emission such that $d=0.45$ at 353~GHz (consistent with Table~\\ref{predictt}).  For the graphite component, $\\beta = 2$.  It should be kept in mind that the ``graphite'' might alternatively be amorphous carbon, in which some frequency dependent changes in $\\beta$ might also be expected.  The ``observables'', $I$, $P$ and $p_{em}$, are shown by the solid curves.  in Figure~\\ref{toymodel}.  The model illustrates several important features. The spectral index measured from $I$ is a function of frequency and would not be representative of either grain material. On the other hand, the spectral index measured from $P$ would track that of the aligned amorphous silicates, even though this component is diluted in $I$. Toward lower frequencies silicates provide a more dominant contribution to the total emission ($d$ is larger) and so even though $pa_{em}$ does not rise significantly for this component, the net $p_{em}$ does increase to lower frequencies.  It appears that the frequency dependence could be quite complicated, such that it would not be correct to extrapolate at constant $\\beta$ and/or $p_{em}$ from meaurements at a single high frequency toward the CMB frequncies.  Rather, to disentangle this most convincingly, multi-frequency measurements of both $p_{em}$ and $P$ (and $I$) will need to be examined in regions of relatively bright foregrounds to produce a self-consistent model. Such an assessment of the dust polarization on the brighter cirrus, uncontaminated by the CMB at lower frequencies, would provide a basis for understanding the fainter cirrus contaminating the CMB. While this is not possible with Archeops, B2K and Planck have polarization capability at three frequencies and PILOT will extend this to higher frequencies (Table~\\ref{angular}). It would be advantageous to make use of polarization measurements below 100~GHz too (e.g., WMAP and Planck), but at some point synchrotron emission, which probably has a higher intrinsic polarization, would no longer be negligible (Fig.~\\ref{contam}), and there is the anomalous emission to considere too..  \\subsection{Alignment of Small Grains and Polarization of Anomalous Emission}\\label{anomalous}  As discussed in \\S~\\ref{polarization}, $(p/\\tau)_{ex}$ is very low in the ultraviolet, where the extinction comes from small grains (VSGs, PAHs). What polarization there is (Fig.~\\ref{pandtau}) is consistent with a fading contribution coming from big grains (Kim \\& Martin \\cite{kim95}).  However, it is possible that there is a very low level of residual alignment of small grains, such as from the Davis-Greenstein process, could be present (Kim \\& Martin \\cite{kimm95}; Wolff \\etl\\ \\cite{wol97}; Lazarian \\& Finkbeiner \\cite{laz03}).  Tiny grains, including if not exclusively PAHs, are the ones that could spin most rapidly to produce low frequency ($\\sim 20$~GHz) emission (Draine \\& Lazarian \\cite{dra98}), possibly accounting for the dust-correlated anomalous microwave emission (Finkbeiner \\cite{fin04}; Finkbeiner \\etl\\ \\cite{finl04}; de Oliveria-Costa \\etl\\ \\cite{deo04}; Davies \\cite{davies06}; Davis \\cite{davis06}). Since these small particles should not be well aligned, the degree of polarization of that component would not be high (tiny compared to the degree of polarization of the synchrotron component), a diagnostic feature that might be used to advantage by WMAP and Planck LFI. Still, the spectre of a few percent polarization for spinning dust could be a serious contamination for precise CMB polarization measurments (Lazarian \\& Draine \\cite{laz00}; Lazarian \\& Finkbeiner \\cite{laz03}).  We can be somewhat more quantitative in a semi-empirical if model-dependent way, starting from the Li \\& Draine (\\cite{lid01}) model, in which PAHs are responsible for most if of the strong 2175~\\AA\\ feature in the extinction curve (as reasonable and economical assumption).  We then use the tight limits placed on the polarization of the 2175~\\AA\\ feature (Martin \\etl\\ \\cite{mar99}; \\S~\\ref{polarization}).  The relative change in polarization, compared to extinction, at the bump gives for this component $p/\\tau_{ex} < 0.002$.  The same quantity for perfectly aligned PAHs could be calculated using small-particle formulae as in equations \\ref{qspheroid} and \\ref{pem} (the extinction is pure absorption), given a shape and refractive indices.  However, there are no refractive indices since only absorptivity has been modeled (Li \\& Draine \\cite{lid01}), and for these large molecules it might not even be particularly appropriate.  Nevertheless, since the PAHs are probably highly flattened, this quantity would be near unity, just as it is for tiny graphite spheroids of even modest asphericity (Martin \\etl\\ \\cite{mar95}; Wolff \\etl\\ \\cite{wol97}).  Therefore, $R_{ex} < 0.002$.  Similarly, for the rotational dipole emission of well aligned tiny grains, $pa_{em}$ would also be unity.  It would be reasonable to assume similar effects from disorientation: $R_{em} \\sim R_{ex}$. Thus for the dust-related anomalous emission, we would predict a degree of polarization less than 0.2\\%, probably very much less given that polarization of the 2175~\\AA\\ feature is unusual.  This is more than an order of magnitude, probably nearer two, less than for the degree of polarization of the thermal dust emission, and also very much less than expected for synchrotron emission.  The real situation is more subtle, however, because polarization of the rotational emission depends on the alignment of the angular momentum vector with respect to the magnetic field rather than the body axis with respect to the field, needed for the ultraviolet polarization (Lazarian \\& Draine \\cite{laz00}). If there is only partial alignment of the angular momentum vector and the body axis, then the low frequency polarization would be higher and with detection of polarization of the anomalous emission one might learn something more about the alignment.  To complicate things further, there is an alternative model for the anomalous microwave emission, magnetic-dipole emission from magnetic grains (Draine \\& Lazarian \\cite{dra99}; Lazarian \\& Finkbeiner \\cite{laz03}), which would have a very distinctive frequency-dependent polarization signature.  Not enough is known about the properties of these grains or their alignment to make firm predictions of the degree of polarization.  However, since magnetic inclusions might be a factor in grain alignment, this emission might be more closely related to ``normal-sized'' grains which are better aligned, and therefore might be more highly polarized than emission from spinning dust.  Following similar bootstrapping arguments such as above, the order of magnitude might be like that given by $R pa_{em}$ in Table~\\ref{predictt}, thus $\\sim 20$\\%, reduced by other diluting contributions to the microwave emission.  Again, synchrotron polarization would be a consideration."
    },
    "0606/astro-ph0606740_arXiv.txt": {
        "abstract": "Spectra have been obtained with the low-resolution modules of the Infrared Spectrograph (IRS) on the Spitzer Space Telescope ($Spitzer$) for 58 sources having f$_{\\nu}$(24\\,\\um) $>$ 0.75\\,mJy.  Sources were chosen from a survey of 8.2 deg$^{2}$ within the NOAO Deep Wide-Field Survey region in Bo\\\"{o}tes (NDWFS) using the Multiband Imaging Photometer (MIPS) on $Spitzer$.  Most sources are optically very faint ($I$ $>$ 24\\,mag).  Redshifts have previously been determined for 34 sources, based primarily on the presence of a deep 9.7\\,\\um silicate absorption feature, with a median z of 2.2.  Spectra are presented for the remaining 24 sources for which we were previously unable to determine a confident redshift because the IRS spectra show no strong features. Optical photometry from the NDWFS and infrared photometry with MIPS and the Infrared Array Camera on $Spitzer$ (IRAC) are given, with $K$ photometry from the Keck I telescope for some objects.  The sources without strong spectral features have overall spectral energy distributions (SEDs) and distributions among optical and infrared fluxes which are similar to those for the sources with strong absorption features.  Nine of the 24 sources are found to have feasible redshift determinations based on fits of a weak silicate absorption feature. Results confirm that the \"1 mJy\" population of 24\\,\\um $Spitzer$ sources which are optically faint is dominated by dusty sources with spectroscopic indicators of an obscured AGN rather than a starburst.  There remain 14 of the 58 sources observed in Bo\\\"{o}tes for which no redshift could be estimated, and 5 of these sources are invisible at all optical wavelengths. ",
        "introduction": "Imaging surveys at infrared wavelengths with the Spitzer Space Telescope ($Spitzer$) have the potential to reveal populations of sources that were previously unknown.  In particular, surveys with the Multiband Imaging Photometer for $Spitzer$ (MIPS) \\citep{rie04} at 24\\,\\um should reveal sources which are luminous because of emission from dust, and which may contain sufficient dust to make them very faint or invisible optically.  The Infrared Spectrograph on $Spitzer$ (IRS)\\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.} \\citep{hou04} is sufficiently sensitive to obtain low-resolution spectra of these MIPS sources at flux density levels of f$_{\\nu}$(24\\,\\um) $<$ 1 mJy.  As part of initial efforts to characterise this \"1 mJy\" population, we surveyed 8.2 deg$^{2}$ within the Bo\\\"{o}tes field of the NOAO Deep Wide-Field Survey (NDWFS) \\citep{jan99} with MIPS to produce a catalog of mid-infrared sources.  The MIPS data were obtained with an effective integration time at 24\\,\\um of $\\sim$90\\,s per sky pixel, reaching a 5 $\\sigma$ detection limit of $\\sim$ 0.3\\,mJy for unresolved sources. This field was chosen because the deep and well calibrated optical imagery in $B_W$, $R$, and $I$ bands makes possible the identification of infrared sources with very faint optical counterparts and allows confident selection of infrared sources lacking optical counterparts to very deep optical limits.  We then selected for subsequent spectroscopic observations with $Spitzer$  those sources which are the faintest optically (typically $I$ $>$ 24) while also bright enough ($>$ 0.75\\,mJy at 24\\,\\um) for spectroscopy with the IRS within integration times of $\\sim$ 1 h. Other than the mid-infrared flux limit, the only selection criterion was optical faintness, because we were primarily interested in understanding sources which would not have been identified in previous optical studies.  Our first results reported the discovery of a significant population of optically obscured, high redshift sources (\\citet{hou05}; hereinafter H05).  To date, we have observed 58 Bo\\\"{o}tes sources with the IRS.  Continuum was detected in all 58 objects observed, confirming that all mid-infrared sources are real even when they have no optical counterparts and also demonstrating that the MIPS-derived positional uncertainties are less than $\\sim$ 0.5'' rms. Redshifts for 17 sources determined by fitting templates of known local objects were reported in H05; an additional 17 sources with redshifts include 14 in Higdon et al. (in preparation), 2 in \\citet{des06}, and one in \\citet{dey05}.  The most important result is that these sources are generally at high redshift. Of the 34 sources with redshifts, the median redshift (z) is 2.2.  The spectra of these sources with redshifts are dominated by strong silicate absorption centered at rest frame 9.7 $\\mu$m. Only two sources are best fit by strong PAH emission features not requiring silicate absorption.  There remain 24 spectra which we have observed with the IRS but for which redshifts have not yet been reported because there are no strong spectral features.  We discuss these sources in the present paper.  All Bo\\\"{o}tes sources with 1.9 $<$ z $<$ 2.7 from IRS redshifts have measurable redshifts because of the deep silicate absorption feature centered at rest frame 9.7\\,\\um.  Because of the accessible wavelength range of the IRS, redshifts beyond about 2.8 could not be measured because this absorption feature would be longward of the IRS wavelength limit. This allows two possible interpretations of the 24 sources for which we have not been able to derive redshifts: either they are a category of sources at redshifts z $\\la$ 2.8 having weak spectral features, or they are at z $\\ga$ 2.8.  It is crucial to know which conclusion is more representative of the \"no-z\" sample.   If features are weak, that is important for interpreting the nature of the sources, and for understanding why they are optically faint.  If the \"no-z\" sources are at higher redshifts than for the measured sources, this would be evidence for an obscured population at even higher redshifts and luminosities than the sources already identified.  Because this population of optically-faint infrared sources would not have been selected using criteria available before $Spitzer$, it is necessary to consider whether all of the targets do indeed represent a distant, extragalactic population, as we have assumed.  Our initial selection ruled out Solar System objects by verifying that the infrared source showed no proper motion between the two epochs of the MIPS observations; the selection at 24\\,\\um insures that any Solar System objects would be in the main asteroid belt or closer and hence would have measurable proper motion between $Spitzer$ observations.  All sources were also detected at the same positions with the $Spitzer$ Infrared Array Camera (IRAC) \\citep{faz04} with observations  obtained at a different epoch.  Known stellar populations are ruled out because the extreme infrared to optical flux ratios imply very cool source temperatures for a black body.  In addition, we report new $K$ band observations with high spatial resolution which resolve at least 4 of the 24 sources. All indications, therefore, are that these sources are an extragalactic population, but locating them in the universe requires the determination of redshifts.  In this paper, we give all available data for the 24 sources with \"featureless\" spectra and compare the overall spectral energy distributions for this sample of sources to those sources with previously determined redshifts.  We illustrate all of the spectra in order to discuss which ones might have weak features and suggest new redshifts for 9 sources based on possible but weak spectral absorption features. ",
        "conclusions": "We present IRS spectra and optical and infrared photometric data for 24 optically faint sources with f$_{\\nu}$ (24\\,\\um) $>$ 0.8\\,mJy selected from within the 8.2 deg$^{2}$ $Spitzer$ Bo\\\"{o}tes survey within the NOAO Deep Wide-Field Survey. The Bo\\\"{o}tes sample with IRS spectra comprises a set of 58 sources with high infrared to optical ratios ($\\nu$f$_{\\nu}$(24 \\ums)$/ \\nu$f$_{\\nu}$($I$) = IR/opt $>$ 50) and typical f$_{\\nu}$(24\\,\\um) = 1 mJy.  Redshifts have been previously determined for 34 of the 58 sources, with a median z = 2.2.  The 24 sources discussed in the present paper are the remaining sources which do not show sufficiently strong spectral features in the IRS low-resolution spectra for confident redshift determination by fitting with templates containing either strong silicate absorption or strong PAH emission.  There are no significant systematic differences in overall SEDs or fluxes for the sample with redshifts (usually having strong silicate absorption at rest frame 9.7\\,\\um) and the sample without redshifts (having weak or no absorption or emission features).  16 of the 24 sources show power-law SEDs determined from photometry through infrared and optical wavelengths.  Sources could be at z $\\la$ 2.8 with weak spectral features, or could show no features in the spectrum because z $\\ga$ 2.8 and features are redshifted out of the IRS spectral range. 10 of the 24 sources have log [f$_{\\nu}$(24 \\ums)/f$_{\\nu}$(8 \\ums)] $>$ 1.0, a value expected for starbursts, but none of the sources show PAH emission features in the infrared spectra normally associated with a luminous starburst.  Possible redshifts are suggested for 9 of the 24 sources based on fitting a profile of weak silicate absorption.  Photometric redshift estimates are given for 2 sources whose SEDs show evidence of a stellar component exceeding the dust continuum, although neither of these sources shows PAH emission, and the photometric redshift does not agree with the weak-silicate redshift for the one source with both measured. These suggested new redshifts provide evidence in favor of the explanation for the majority of the \"no-redshift\" sources that they are similar in nature and luminosity to the more heavily absorbed sources, but with weaker absorption features, and would bring the total number of redshifts determined for faint Bo\\\"{o}tes sources to 44 of the 58 sources observed, counting one photometric redshift.  Even if the suggested new redshifts are correct, there remain 14 sources from the Bo\\\"{o}tes sample of 58 with no redshift estimates.  With our present data, it is not possible to reach any conclusions regarding the nature of these remaining sources.  Five of these 14 sources without redshifts are not detected at any optical wavelength in the NDWFS.  Whether these represent a dusty population at redshifts higher than previously measured, or whether they represent a population of optically-obscured sources at z $<$ 2.8 with weak spectral features remains ambiguous. Because such sources are a significant fraction of the optically faint sources in the Spitzer \"1 mJy\" population of 24\\,\\um sources, continued efforts to determine their nature are important."
    },
    "0606/astro-ph0606576_arXiv.txt": {
        "abstract": "We have observed the bipolar post-AGB candidate OH 231.8+4.2, using the mid-infrared interferometer MIDI and the infrared camera with the adaptive optics system NACO on the Very Large Telescope. An unresolved core ($<$200~mas in FWHM) is found at the center of the OH\\,231.8+4.2 in the 3.8~$\\mu$m image.  This compact source is resolved with the interferometer. We used two 8-meter telescopes with four different baselines, which cover projected baseline lengths from 62 to 47 meters, and projected position angles from 112 to 131 degrees that are almost perpendicular to the bipolar outflow.  Fringes from 8 to 9~$\\mu$m and from 12 to 13.5~$\\mu$m were clearly detected, whilst the strong silicate self-absorption allows only marginal detection of visibilities between 9 and 12~$\\mu$m.  The fringes from the four baselines consistently show the presence of a compact circumstellar object with an inner radius of 30--40~mas, which is equivalent to 40--50 AU at 1.3~kpc.  This clearly shows that the mid-infrared compact source is not the central star (3~AU) but circumstellar material.  The measured size of the circumstellar material is consistent with the size of such disks calculated by hydrodynamic models, implying the circumstellar material may have a disk configuration. ",
        "introduction": "Low and medium mass stars ($\\sim$1--8 $M_{\\sun}$ on the main sequence) experience an intensive mass-loss phase during the Asymptotic Giant Branch (AGB) phase.  Typically, the AGB wind is spherically symmetric. However, during the next evolutionary stages, i.e., the post-AGB phase and the planetary nebula (PN) phase, a high fraction of stars show asymmetric shapes in their circumstellar envelopes, such as elliptical and bipolar.  One of the hypotheses about the formation of the bipolar shape invokes a binary disk scenario \\citep{Balick02, VanWinckel03}.  Part of the material lost during the intensive AGB mass-loss wind is trapped in the binary system, and a circumbinary disk is formed in the plane of the binary orbit. The disk restricts the direction of the low density but high velocity post-AGB and PN wind in the equatorial plane, and focuses the wind towards two poles.  The size of the binary disk will be small \\citep[less than 100 AU (80~mas for our target);][]{Mastro99}, and so requires interferometric observations to be resolved.  \\object{OH 231.8+4.2} (IRAS~07399$-$1435; RA 07h42m16.83s Dec $-$14d42m52.1s; hereafter OH~231) is one of the well studied post-AGB candidates.  TiO bands are detected from the central region, suggesting the central star exhibits a M9 spectral type  \\citep{Cohen81}. \\citet{SanchezContreras04} claimed the presence of a spectroscopic binary from optical spectra, because in addition to TiO and VO bands from the M-type star, Balmer lines and continuum excess are detected.  OH~231 is probably located in the open cluster, M~46 \\citep{Jura85}, thus the distance is relatively well determined (1.3 kpc).  The outflow is strongly bipolar, and bubbles and shocked regions are found in the outflow \\citep[e.g.][]{Bujarrabal02}. L-, N- and Q-band seeing-limited images show an unresolved core at the center of this object \\citep{Kastner92, Jura02}.  The infrared color of this compact source is extremely red, and it is believed to be a dusty disk \\citep{Jura02}.  The velocity structure of SiO masers also suggest the presence of a rotating disk around this compact source \\citep{SanchezContreras02}. The OH masers appear to be associated with an expanding torus \\citep{Zijlstra01}.  In this paper we present both high resolution infrared (IR) images and and mid-IR interferometeric visibilities of the central compact source, so as to resolve the compact source at the center and so as  to determine if this source is a disk. ",
        "conclusions": "Our L'-band and N-band images using the AO system find an unresolved compact object in the center with  dimension less than 200~mas.  This object has a red color and consists of circumstellar material.  This circumstellar material is responsible for the fringes detected by mid-infrared interferometer.  There are other emission mechanisms which can cause correlated flux in the mid-infrared, such as (1) the central star and (2) the central star + binary companion(s), but we can exclude these possibilities. First, the correlated flux of a single star with a radius 3~AU \\citep{SanchezContreras02} at 1.3~kpc (2.3~mas) is almost 100~\\% of the total flux (i.e. unresolved) with current baseline lengths. The observed fringes indicate a much larger source size.  Using the parameters of the central star from \\citet{Jura02} ($T_{\\rm eff}=2500$~K, radius of $4.6\\times10^{13}$~cm), the flux of the central star is $\\sim$6~Jy at 10~$\\mu$m, thus the correlated flux should also be $\\sim$6~Jy on any baseline.  The measured values are inconsistent with the expected correlated flux; the flux drops to 0.1--0.2~Jy for the longest baseline. Thus, the measured visibilities are not from the central star. Second, the silicate absorption is detected in correlated flux, which shows the circumstellar origin of the mid-infrared compact source.  Although the presence of a binary companion has been suggested by \\citet{SanchezContreras04}, it is unlikely that we measured the orbit of the binary companion around the central star, nor the radius of the binary companion.  A sharp angular dependence of visibilities would be expected in the orbit case, which is not detected (Fig.~\\ref{Fig-gauss}).  The possibility of the binary companion radius is ruled out because the flux would be 1~mJy level at 1.3~kpc if the companion is an A0V star \\citep{SanchezContreras04}.  Therefore, the MIDI visibilities are due to dusty circumstellar material. The absence of stellar emission in the correlated flux suggests this material is still optically thick at 10~$\\mu$m.  \\citet{Jura02} analyzed the spectral energy distribution of this object and argued that the mid-infrared unresolved source is a disk. \\citet{SanchezContreras02} measured the velocity distribution of SiO maser lines. The SiO masers show the presence of a rotating disk, within 4~mas from the star. \\citet{Zijlstra01} find that OH velocity structure could be viewed as an expanding torus, with an additional component from a bipolar, ballistic outflow.  Our NACO images also suggest the presence of a disk or torus.  The central region in the L'-band shows a trapezium shape: this may be interpreted as a flared outflow from a torus or a disk.  The trapezium shape is also seen in the NB2.12 image.  The small blobs seen in the trapezium cloud of NB2.12 image, is probably due to non-uniform extinction within the flared disk, or an illusion caused by scattered light.  The north part of the `trapezium region' is brighter than other regions in NB2.12, possibly because the disk is slightly inclined with the northern part nearer to the earth.   The obtained size of the dense circumstellar material is at least 40--50 AU.  The optical depth from our preliminary analysis using {\\it DUSTY} \\citep{Ivezic97} suggests $\\tau_{\\rm{8\\mu m}}$=1.6 and $\\tau_{\\rm{13~\\mu m}}$=2.7, showing the actual inner radius is smaller than the values which we measured.  Nevertheless, if the density distribution is $\\rho=\\rho_0 (r/r_0)^{-\\alpha}$ where $\\alpha=1$--2, the optical depth increases dramatically for closer inner radii, and the measured size need not be that different from the actual inner radius of the dusty shell or torus.  The hydrodynamic model of \\citet{Mastro99} shows that an accretion disk with a radius of 40--50 AU can be formed in a binary disk.  The co-incidence of the radii from the theory and the observations implies that the circumstellar material is shaped as a disk or torus in this object.   On the other hand, \\citet{Jura02} assumed a binary companion with an orbit of 3--~5~AU for OH 231, and the inner radius of the disk might be $\\sim 1.7$ larger than the companion orbit, which should be 5--9~AU. Our measured inner radius is much larger than \\citet{Jura02}'s assumption.  This may be because the companion is actually further out than expected because the central star is larger than expected by \\citet{Jura02}, or because our mid-infrared measurements observe the radius at which disk has cooled down enough to allow dust to condense out of the gas phase. An alternative solution could be that the disk is gradually expanding and losing momentum, and at the time of the formation, the disk could be much smaller than the current size.  In conclusion, we measured the angular size of the circumstellar material in OH~231to be 40--50~AU, and this material is probably in a disk or toroidal configuration.  Future MIDI observations with different baseline angles are required to confirm the presence of a disk like structure and to measure its inclination angle.   The OH masers are located to the south and east, within the darker lane seen in the L'-band image, and trace the waist of the trapezium, allowing for orientation effects.  Only foreground (blue-shifted) OH is seen: the red emission appears to be more extended and is not picked up by the MERLIN interferometer, while the blue emission consist of more compact components. The velocity field shows a clear gradient in the southern part of the maser spots, with $v_{\\rm LRS}$=20--30\\,km\\,s$^{-1}$ on the east and west edge and the bluest component ($v_{\\rm LRS}$=10\\,km\\,s$^{-1}$) in the middle. It is inconsistent with an expanding torus, but is consistent if we observe the blue-shifted rim of a biconical outflow tilted towards us. Thus, the dynamics of the source change from a rotating disk in the very centre, within a few stellar radii, to a conical outflow at $\\sim$1000~AU. The stellar position estimated by the the SiO maser tracing an equatorial rotating torus \\citep{SanchezContreras02} coincides with the central position of the blue-shifted ($\\sim$0--5\\,km\\,s$^{-1}$) biconical distribution seen in OH.   All the components of the binary disk hypothesis \\citep{Balick02, VanWinckel03} may therefore be present in OH~231.8+4.2: a rotating SiO maser disk very close to the central star, a compact circumstellar material at $\\sim$40--50~AU which may have a disk-like distribution, and a bipolar outflow."
    },
    "0606/astro-ph0606210_arXiv.txt": {
        "abstract": "We report the discovery of a new Einstein cross at redshift $z_S = 2.701$ based on \\lya\\, emission in a cruciform configuration around an SDSS luminous red galaxy ($z_L = 0.331$).  The system was targeted as a possible lens based on an anomalous emission line in the SDSS spectrum.  Imaging and spectroscopy from the W.M.~Keck Observatory confirm the lensing nature of this system.  This is one of the widest-separation galaxy-scale lenses known, with an Einstein radius $\\theta_{\\rm E}\\simeq 1.84$\\,arcsec.  We present simple gravitational lens models for the system and compute the intrinsic properties of the lensed galaxy.  The total mass of the lensing galaxy within the $8.8 \\pm 0.1$ kpc enclosed by the lensed images is $(5.2 \\pm 0.1) \\times 10^{11} M_{\\sun}$.  The lensed galaxy is a low mass galaxy (0.2$L_*$) with a high equivalent-width Ly$\\alpha$ line ($EW_{{\\rm Ly}\\alpha}^{\\rm rest} = 46 \\pm 5$ \\AA).  Follow-up studies of this lens system can probe the mass structure of the lensing galaxy, and can provide a unique view of an intrinsically faint, high-redshift, star-forming galaxy at high signal-to-noise ratio. ",
        "introduction": "Strong gravitational lensing is a powerful tool for the measurement of lensing galaxy masses and for the detailed study of magnified high-redshift sources.  Multiple lensed images can directly constrain models for the distribution of mass in the lens on the scale of the Einstein radius $\\theta_{\\rm E}$.  In well-studied cases, these observations directly test theories for the central dark-matter profile in both galaxies \\citep[e.g.,][]{kt03, dye_warren_05} and galaxy clusters (\\eg \\citealt{kneib_2218_95, kneib_2218_96}; \\citealt*{abdelsalam_98a, abdelsalam_98b}).  Simultaneously, lensed background galaxies can be magnified by factors of up to several tens, providing data of a quality which would not otherwise be possible.  Such data have been used to constrain the faint end of the high-redshift galaxy luminosity function \\citep[e.g.,][]{santos_2004}, study distant galaxies at wavelengths which would be impractical without magnification \\citep*[e.g.,][]{chary_2005}, and probe detailed properties of high-redshift galaxies \\citep{pettini_cb58}.  Systematic searches for new galaxy-scale strong gravitational lenses are traditionally based on imaging detections \\citep[e.g.,][]{maoz_snapshot, gregg_2000, wisotzki_2002, inada_0924, morgan_ctq327, richards_snap, pindor_2004}, although a handful of lenses have been identified in spectroscopic observations \\citep{huch85, warren_0047_96, john03}. The Sloan Lens ACS (SLACS) survey (\\citealt{slacs1, slacs2, slacs3}) efficiently identifies new gravitational lenses by searching the spectroscopic database of the Sloan Digital Sky Survey (SDSS; \\citealt{york_sdss}) for systems consisting of low-redshift ($z \\sim 0.1$--$0.4$) luminous elliptical galaxies superposed with moderate-redshift ($z \\sim 0.3$--$1$) star-forming galaxies (see also \\citealt{bolton_speclens}, \\citealt*{whw_2005}, and \\citealt{whwdm_2006}). This technique relies upon the detection of {\\em multiple} anomalous nebular emission lines in the SDSS fiber spectra of lensing early-type galaxies.  In principle, the anomalous emission line technique can also identify lenses with higher redshift, \\lya-emitting source galaxies \\citep{warren_0047_96, h00, wil00}.  In practice this approach has been less productive because strong, high-redshift \\lya\\, emitters are less numerous on the sky than moderate-redshift galaxies showing oxygen and Balmer emission lines to the SDSS line-flux limits, and the increase in lensing cross section with source redshift does not overcome this effect.  This {\\em Letter} reports the discovery of a new spectroscopically-selected strong gravitational lens with a \\lya-emitting galaxy as its source, SDSS~J101129.49$+$014323.3 (hereafter \\einstein).  The lensed galaxy forms a highly symmetric Einstein cross which, by virtue of its large physical scale, provides leverage for determining the dark-matter halo mass of the lensing elliptical. The system also provides a highly magnified view of an intrinsically faint, compact, and high-redshift star-forming galaxy.  For all calculations, we assume a universe with $(\\Omega_{\\mathrm{M}},\\Omega_{\\Lambda},h) = (0.3,0.7,0.7)$. ",
        "conclusions": "We present the discovery of a new high-redshift Einstein-cross gravitational lens, \\einstein.  The lens galaxy is a bright elliptical at $z_{\\rm lens} = 0.331$, while the lensed source is a $\\sim 0.2 L_*$ \\lya -bright, star-forming galaxy at redshift $z_{\\rm source} = 2.701$. Though such high-$z$ \\lya\\ lenses appear to be much less numerous in the SDSS than lenses with lower redshift ($z \\la 1$) oxygen- and Balmer-line emitting sources \\citep{bolton_speclens, whw_2005, slacs1, whwdm_2006}, the discovery of \\einstein\\ demonstrates their presence.  Depending on the luminosity function slope of the \\lya -emitting source population, deeper spectroscopic surveys could yield an appreciable sample of such lens systems.  As with all gravitational lenses, this system offers a powerful tool for measuring the mass in the lensing galaxy.  \\einstein\\ is of particular interest for its relatively wide ($\\sim 4\\arcsec$) image separation, probing the lens galaxy at a radius within which the contributions of luminous and dark matter are expected to be comparable.  As with the sample of lenses discovered by the SLACS survey \\citep{slacs1}, the image of the lens galaxy is not overwhelmed by lensed-quasar images and is thus accessible to accurate photometric and dynamical observations. The high degree of symmetry of the image configuration suggests that a largely model-independent test of the relative degree of flattening between the mass and light distributions in the early-type lens galaxy will be enabled by high-resolution imaging. A modest Cycle~15 {\\it HST}/ACS program to image this system has been awarded two orbits (P.I. Moustakas).  \\einstein\\ also offers a highly magnified view of a sub-$L_*$ starforming galaxy at high redshift.  This discovery is thus complementary to the more luminous Lyman-break galaxy MS1512$-$cB15 \\citep{yee_cb58}, which is also strongly magnified by gravitational lensing \\citep{seitz_cb58}.  The magnification of \\einstein\\ suggests this system as a target for deep spectroscopic studies of the high-redshift IGM that would otherwise be infeasible due to the intrinsic source faintness."
    },
    "0606/astro-ph0606026_arXiv.txt": {
        "abstract": "{We  revisit the evolutionary scenario for Hot Flasher low-mass structures, where mass loss delays the He flash till the initial phases of their White Dwarf cooling sequence.} {Our aim has been to test the theoretical results vis-a-vis different assumptions about the efficiency of mass loss.} {To this purpose, we present evolutionary models covering a fine grid of masses, as obtained assuming a single episode of mass loss in a Red Giant model of 0.86 $M_{\\odot}$ with Z=0.0015.} {We find a reasonable agreement with previous evolutionary investigations, showing that for the given metallicity late Hot Flashers are predicted to cover the mass range M=0.4975 to M= 0.4845 ($\\pm$0.0005) $M_{\\odot}$, all models igniting the He-flash with a mass of the H-rich envelope as given by $M_e$=0.00050 $\\pm$0.00002 $M_{\\odot}$. The ignition mechanism is discussed in some details, showing the occurrence of a bifurcation in the evolutionary history of stellar structures at the lower mass limit for He ignition. Below such a critical mass, the structures miss the He ignition, cooling down as a Hot Flasher-Manqu\\'e He White Dwarf.  We predict that these structures will cool down, reaching the luminosity $logL/L_{\\odot}$=-1 in a time at the least five times longer than the corresponding cooling time of a normal CO White Dwarf.} {On very general grounds, one expects that old stellar clusters with a sizeable population of Hot Flasher should likely produce at least a similar amount of slow-cooling  He White Dwarfs. According to this result, in a cluster where 20\\% of Red Giants escape the He burning phase, one expects roughly twice as White Dwarfs than in a normal cluster where all Red Giants undergo their He flash}  \\keywords {Stars: evolution, Stars: White Dwarfs, Stars: mass loss} ",
        "introduction": "Over the last decades the evolution of low-mass stellar structures has been the subject of a large amount  of investigations, aimed at constraining the evolutionary status of stars in  old stellar systems, such as Galactic Globular Clusters. Since long time we know that present  Globular Cluster stars are expected to leave their Main Sequence to climb along the Red Giant Branch (RGB) till the onset of the He-flash. After the phase of central (Horizontal Branch) and shell (Asymptotic Giant Branch) He burning phases, they will eventually cool down under the form of Carbon-Oxygen  (CO) White Dwarfs (WDs).  In this context,  the occurrence of extended \"Blue Tails\" in the Horizontal Branches (HB) of several Galactic Globulars has already been understood in terms of RGB structures which have lost the large majority of their H-rich  envelope before igniting He to become HB stars. Castellani \\& Castellani (1993; but see also Castellani, Degl'Innocenti \\& Pulone 1995) found that, for extreme mass loss, there are stellar models which fail to ignite He at the tip of the RGB, but undergo a late He-flash during the contraction toward their He-WD structure or in the early stages of the WD cooling sequence. Similar structures are now known in the literature as \"Hot He-Flashers\" (HFs). Even larger mass loss will prevent the He ignition, definitively producing He White Dwarfs.  Hot Flashers have been extensively investigated by several authors. D'Cruz et al. (1996) made use of Reimers (1975, 1977) formula for mass loss, taking  the efficiency parameter $\\eta_R$ as a free parameter to explore the range of mass-loss producing HFs for selected assumptions about the star metallicity. Sweigart (1997) discovered that when the He-flash occurs along the WD cooling sequence (\"late\" HFs), then convection  is expected to reach the H-rich envelope, enhancing  He and Carbon abundances in the stellar atmosphere and driving strong H-flashes.  Brown et al. (2001) adopted again the Reimers formalism to explore the occurrence of late HF for the metal abundance Z=0.0015, in connection with observational evidence for extremely hot HB stars in the Galactic globular NGC2808. Quantitative estimates of the mixing driven by the He-flash have been finally presented by Cassisi et al. (2003), who were able, for the first time, to follow in detail the growth of such an instability in late HFs.  In this paper we revisit the  HF theoretical scenario but adopting different assumptions concerning the mass-loss mechanism. On this basis we will present and discuss  new evolutionary results, focusing the attention on the stellar structures marking the transition between late HF and bona fide He WD. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[width=8cm]{Age_LL.eps} \\caption {The cooling ages versus the WD luminosity for our most massive He WD model ($0.4840 M_{\\odot}$: heavy line) compared with the less massive model $0.4600 M_{\\odot}$, the  $0.449 M_{\\odot}$ with element diffusion (Serenelli et al. 2002) and a $0.5 M_{\\odot}$ CO WD (Prada Moroni \\& Straniero 2002). {\\it L} is in fraction of $L_{\\odot}$ and {\\it AGE} is in years.} \\label{f:fig4} \\end{figure}   In this paper we have addressed the problem of Hot Flasher, investigating in details  the predicted evolutionary behavior of low mass stars  with Z=0.0015 after an episode of mass loss during their RG evolution. We found that stellar masses in the range $0.485 \\le M \\le 0.497 M_{\\odot}$ experience the He flash (and the explosive H-reignition) during their WD cooling phase, when the residual H shell burning has reduced the H-rich stellar envelope down to $M_e \\sim 0.0005 M_{\\odot}$, independently of the mass of the model. Such a result appears in reasonable agreement with theoretical predictions given by Brown et al (2001) for the same  metallicity, the small difference ($M \\sim 0.0005$ against $0.0006 M_{\\odot}$) being likely the effect of small differences in the  adopted input physics. Supporting, in turn, the evidence that the mass of the H-rich envelope plays a critical role in the onset of the  delayed He flashes.  According to this result, we are also predicting the structural parameters needed to evaluate the  evolutionary times of structures below the lower mass limit for   He ignition, which will definitely cool down as He WD. We find that these He WD will reach the luminosity $logL/L_{\\odot}$=-1 in a time about 5 times longer than normal Carbon-Oxigen WDs do, giving a detectable contribution to the abundance of WD above such a luminosity if and when a not marginal fraction of RG stars escape the He ignition.  For the sake of completeness, one has finally to advise that the current scenario slightly depends on the luminosity of the RG models undergoing the mass loss episode, since peeling off RG structures either before (as in the previously  reported computations) or after the first dredge up gives HF progenitors with different He abundances in the stellar envelopes. Numerical experiments for our $0.86 M_{\\odot}$ models  have shown, e.g., that when the models is stripped after the dredge up ($logL/L_{\\odot}= 1.48, M_c=0.24M_{\\odot}$) the lower mass limit for LHF moves from 0.485 to 0.491 $M_{\\odot}$, with  LHF models characterized by slightly smaller H-rich envelopes, as given by $M_e^f=0.00048 M_{\\odot}$. As a whole, these appear as marginal differences, not affecting the theoretical scenario we are dealing with."
    },
    "0606/astro-ph0606160_arXiv.txt": {
        "abstract": "% Emission line spectra from H~{\\sc ii} regions are often used to study properties of the gas in star-forming regions, as well as temperatures and luminosities of the ionising sources. Empirical diagnostics for the interpretation of observational data must often be calibrated with the aid of photoionisation models. Most studies so far have been carried out by assuming spherical or plane-parallel geometries, with major limitations on allowed gas and dust density distributions and with the spatial distribution of multiple, non-centrally-located ionising sources not being accounted for. We present the first results of our theoretical study of geometric effects, via the construction of a number of 3D photoionisation models using the {\\sc mocassin} code for a variety of spatial configurations and ionisation sources. We compare integrated emission line spectra from such configurations and show evidence of systematic errors caused by the simplifying assumption of a single, central location for all ionising sources. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606483_arXiv.txt": {
        "abstract": "We derive the effects of dark matter (DM) decays and annihilations on structure formation. We consider moderately massive DM particles (sterile neutrinos and light DM), as they are expected to give the maximum contribution to heating and reionization.  The energy injection from DM decays and annihilations produces both an enhancement in the abundance of coolants (H$_2$ and HD) and an increase of gas temperature. We find that for all the considered DM models the critical halo mass for collapse, $m_{crit}$, is generally higher than in the unperturbed case. However, the variation of $m_{crit}$ is small. In the most extreme cases, i.e. considering light DM annihilations (decays) and halos virializing at redshift $z_{vir}>30$ ($z_{vir}\\sim{}10$), $m_{crit}$ increases by a factor $\\sim{}$4 ($\\sim{}$2). In the case of annihilations the variations of $m_{crit}$ are also sensitive to the assumed profile of the DM halo.  Furthermore, we note that the fraction of gas which is retained inside the halo can be substantially reduced (to $\\approx 40$ per cent of the cosmic value), especially in the smallest halos, as a consequence of the energy injection by DM decays and annihilations. ",
        "introduction": "One of the fundamental questions concerning the formation of first structures is the minimum halo mass (critical mass, $m_{crit}$) for collapse at a given redshift (Silk 1977; Rees \\& Ostriker 1977; White \\& Rees 1978; Couchman 1985; Couchman \\& Rees 1986; de Araujo \\& Opher 1988, 1991; Haiman, Thoul \\& Loeb 1996).  Tegmark et al. (1997; T97) thoroughly addressed such question, pointing out how $m_{crit}$ crucially depends on the abundance of H$_2$, the main coolant present in the metal free Universe. Subsequent studies (Abel et al. 1998; Fuller \\& Couchman 2000; Galli \\&{} Palla 1998, 2002; Ripamonti 2006) refined the model of T97, accounting also for minor effects, such as the cooling induced by HD molecules.  The production of molecules and $m_{crit}$ are sensitive to any physical process which can release energy in the intergalactic medium (IGM). In fact, the injection of energy in the IGM can  either delay the collapse of first halos (because of the increased gas temperature, or of photodissociation of molecules) or favour structure formation (because of the enhancement in the abundance of free electrons, which act as catalysts for the formation of molecules).  For this reason, it is crucial to understand the influence of reionization sources on structure formation. Many studies have shown that massive metal free stars are efficient in dissociating H$_2$ molecules, quenching star formation in the first halos (Haiman, Rees \\& Loeb 1997; Ciardi, Ferrara \\& Abel 2000; Ciardi et al. 2000; Haiman, Abel \\& Rees 2000; Ricotti, Gnedin \\& Shull 2002; Yoshida et al. 2003). Intermediate mass black holes, produced by the collapse of first stars, are thought to efficiently re-heat the IGM, increasing $m_{crit}$ and reducing star formation in the smaller mass halos (Ricotti \\& Ostriker 2004; Ricotti, Ostriker \\& Gnedin 2005; Zaroubi et al. 2006).  Also particle decays and annihilations can be sources of partial reionization and heating (see Mapelli, Ferrara \\& Pierpaoli 2006 and references therein), and could influence structure formation. For example, Shchekinov \\& Vasiliev (2004) investigated the possible effect on $m_{crit}$ due to ultra-high energy cosmic rays (UHECRs) emitted by particles decaying in the early Universe. Biermann \\& Kusenko (2006) considered the impact on structure formation due to sterile neutrino decays. Both these studies found a substantial enhancement on the abundance of molecular coolants (H$_2$ and/or HD). However, they neglected the possible increase of gas temperature due to UHECRs or decays, respectively.  More recently, Stasielak, Biermann \\& Kusenko (2006) evaluated the effect of sterile neutrino decays accounting also for the heating of the gas. However, their single-zone model  is likely to oversimplify the crucial behaviour of gas density during the halo collapse.  In this paper, we consider the influence of dark matter (DM) decays and annihilations on structure formation, taking into account variations induced both in the chemical and in the thermal evolution of the IGM and of the gas inside halos. Furthermore, we substitute the single-zone models, which are commonly adopted in previous papers (Haiman et al. 1996 is an important exception), with more sophisticated 1-D simulations. We focus on relatively low mass DM particles, such as sterile neutrinos and light DM (LDM), as their effect on the IGM is expected to be much more important than that of heavier ($\\gtrsim{}100$ MeV) DM particles (Mapelli et al. 2006).  Sterile neutrinos are expected to decay into active neutrinos and keV-photons (Dolgov 2002), while LDM can either decay or annihilate producing electron-positron pairs (Boehm et al. 2004; Hooper \\& Wang 2004; Picciotto \\& Pospelov 2005; Ascasibar et al. 2006). keV-photons interact with the IGM both via Compton scattering and photo-ionization; instead, the electron-positron pairs undergo inverse Compton scattering, collisional ionizations, and positron annihilations (Zdziarski \\& Svensson 1989; Chen \\& Kamionkowski 2004; Ripamonti, Mapelli \\& Ferrara 2006, hereafter RMF06). RMF06 derived the fraction $f_{abs}(z)$ of energy emitted by sterile neutrino decays and LDM decays or annihilations which is effectively absorbed by the IGM through these processes. In this paper, we adopt the fits of $f_{abs}(z)$ given by RMF06.   In Section 2 we describe the hydro-dynamical code used to derive $m_{crit}$ and the DM models which we adopt. In Section 3 we discuss the effect of DM decays and annihilations on the chemical and thermal evolution of the IGM, giving an estimate of the Jeans mass. In Section 4 we describe the chemical and thermal evolution of the gas inside the halos, deriving $m_{crit}$. In the discussion (Section 5) we address various points, such as the variations in the baryonic mass fraction inside the halos induced by DM decays/annihilations and the influence of the concentration of the DM profile.   We adopt the best-fit cosmological parameters after the 3-yr WMAP results (Spergel et al. 2006), i.e. $\\Omega{}_{\\rm b}=0.042$, $\\Omega{}_{\\rm M}=0.24$, $\\Omega{}_{\\rm DM}\\equiv{}\\Omega{}_{\\rm M}-\\Omega{}_{\\rm b}=0.198$, $\\Omega{}_\\Lambda{}=0.76$, $h=0.73$, $H_0=100\\,{}h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper we derived the effects of DM decays and annihilations on structure formation. We considered only moderately massive DM particles (sterile neutrinos and LDM), as they are expected to give the maximum contribution to heating and reionization (Mapelli et al. 2006). To describe the interaction between the IGM and the decay/annihilation products we followed the recipes recently derived by RMF06.  We accounted not only for the diffuse cosmological contribution to heating and ionization, but also for the local contribution due to DM decays and annihilations occurring in the halo itself. The local contribution results to be dominant in the case of DM annihilations especially for cuspy DM profiles.   The energy injection from DM decays/annihilations produces both an enhancement in the abundance of coolants (H$_2$ and HD) and an increase of gas temperature.  We found that for all the considered DM models (sterile neutrino decays, LDM decays and annihilations) the critical halo mass for collapse, $m_{crit}$ is often higher than in the unperturbed case. This means that DM decays and annihilations tend to delay the formation of structures. However, the variation of $m_{crit}$ is minimal. In the most extreme cases, i.e. considering LDM annihilations (decays) and halos virializing at redshift $z_{vir}>30$ ($z_{vir}\\sim{}10$), $m_{crit}$ increases of a factor $\\sim{}4$ ($\\sim{}$2). In the case of decays, the variations of $m_{crit}$ are almost independent from the assumed concentration of the DM halo, although higher concentrations (corresponding to smaller values of $m_{crit}$) seem to be associated with slightly stronger effects of the DM energy injection. The dependence on concentration is more evident in the case of annihilating particles, where higher concentrations lead to substantially larger effects. This happens because the ``local'' contribution is important.   In summary, the effects of DM decays and/or annihilations on structure formation are quite small, except in some extreme cases (e.g. very high concentration for annihilations). However, the energy injection from DM decays/annihilations has important consequences on the fraction of gas which is retained inside the halo. This fraction can be substantially reduced, especially in the smallest halos ($\\lesssim{}10^6\\,{}M_\\odot{}$).  Finally, we point out that our results are quite different from the conclusions of Biermann \\& Kusenko (2006) {and Stasielak et al. (2006)\\footnote{Biermann \\& Kusenko (2006) and Stasielak et al. (2006) do not calculate the critical mass $m_{crit}$. So, it is quite difficult to make a quantitative comparison between their results and ours.}, who suggest that sterile neutrino decays can favour the formation of first objects.  The discrepancy is likely due to our more complete treatment which includes the hydrodynamics of the collapsing structures. In fact, our hydro-dynamical treatment allows to describe the detailed gas density evolution during the collapse, resulting in markedly different temperature and chemical properties with respect to those found by a simple one-zone model."
    },
    "0606/astro-ph0606103.txt": {
        "abstract": "The unified model of Seyfert galaxies suggests that there are hidden broad-line regions (HBLRs) in Seyfert 2 galaxies (S2s). However, there is increasing evidence for the appearance of a subclass of S2s lacking of HBLR (non-HBLR S2s). An interesting issue arises as to relations of non-HBLR S2s with other types of Seyfert galaxies and whether or not they can be included in the unified model. We assemble two sub-samples consisting of 42 non-HBLR S2s and 44 narrow-line Seyfert 1s (NLS1s) with redshift $z\\le 0.05$ from published literatures to explore this issue. We compare black hole masses in the galactic centers, accretion rates, infrared color ratio ($f_{60 \\mu \\rm m}/f_{25 \\mu \\rm m}$) as a potential indicator of the dusty torus orientation, \\oiii $\\lambda 5007$, radio and far infrared luminosities. We find that non-HBLR S2s and NLS1s have: 1) similar distributions of the black hole masses ($10^6-3\\times 10^7\\sunm$) and the Eddington ratios ($L_{\\rm Bol}/L_{\\rm Edd}\\sim 1$); 2) significantly different distributions of $f_{60 \\mu \\rm m}/f_{25 \\mu \\rm m}$ ratios; 3) similar distributions of bulge magnitudes and luminosities of [O {\\sc iii}], radio, far infrared emission. The similarities and differences can be understood naturally  if they are intrinsically same but non-HBLR S2s are viewed at larger angles of observer's sight than NLS1s. We thus suggest that non-HBLR S2s only have \"narrower\" broad line regions and they are the counterparts of NLS1s viewed at high inclination angles. The absence of the polarized emission line in non-HBLR S2s is caused by the less massive black holes and high accretion rate similar to NLS1s. The implications of the unification scheme of non-HBLR S2s and NLS1s are discussed. % ",
        "introduction": "Seyfert galaxies were traditionally divided into two classes according to the presence or absence of the broad permitted optical lines. With the discovery of the polarized broad lines in NGC 1068, a prototypical Seyfert 2 galaxy, Antonucci \\& Miller (1985) suggested that there should be a geometrically and optically thick \"torus\" surrounding a Seyfert 1 nucleus with broad lines (BLS1s). If the torus is face on, we can \"see\" the broad line region directly and the galaxies appear as BLS1s. Otherwise it appears as a Seyfert 2. This is the basic idea of the unified model (UM) of AGNs (Antonucci 1993). Many authors have reported pieces of evidence for the orientation-based UM (Miller \\& Goodrich 1990; Tran et al. 1992; Mulchaey et al. 1994, Tran 1995; Yong et al. 1996; Heisler, Lumsden \\& Baily 1997, hereafter HLB97; Moran et al. 2000; Lumsden et al. 2001). It has been found that at least $35\\%$ of Seyfert 2 galaxies have broad emission lines in polarized lights (Tran 2001; Moran et al. 2000) and $\\sim 96\\%$ of the objects have column densities ranging from $10^{22}$ to $10^{24}$cm$^{-2}$ (Risaliti et al. 1999; Bassani et al. 1999). However, spectropolarimetric surveys of complete samples of Seyfert 2 galaxies suggest that hidden Seyfert 1 nuclei have not been detected in $>50\\%$ of these objects from the CfA and 12 $\\mu$m samples of Seyfert 2 galaxies (Tran 2001, 2003); and $10-30\\%$ S2 are found unabsorbed in X-rays (Panessa \\& Bassini 2002), even $50\\%$ among {\\em ROSAT}-selected Seyferts (Gallo et al. 2006). Non-HBLR S2s are shown to be systematically weaker than their HBLR counterparts, and can not be explained by different orientations (Tran 2001; 2003, Lumsden \\& Alexander 2001), challenging the unification scheme. Tran (2003) suggested that there are \"true\" Seyfert 2 galaxies. The paradigm of the unification scheme for {\\em all} Seyfert galaxies remains a matter as debate among the literatures (e.g. Miller \\& Goodrich 1990; Kay 1994; Tran 2001, 2003).  We still poorly understand why we can not detect polarized broad lines (PBLs) in some S2s. The absence of PBLs could be  attributed to edge-on line of sight and hidden of electron scattering region (Miller \\& Goodrich 1990; HLB97; Taniguchi \\& Anabuki 1999). Nicastro (2000) related the absence of PBLs to higher Eddington ratios. He suggested that the width of broad emission lines is Keplerian velocity of an accretion disk at a critical distance from the central black hole, which is the transition radius between radiation  and gas pressure-dominated region (see also Nicastro et al. 2003). It has also been suggested that some non-HBLR S2s are intrinsically weak and lack of broad line region (Tran 2001, 2003; Gu \\& Huang 2002; Laor 2003). As argued by Tran (2003), \"it appears that much of the difference between S1s and S2s can be explained solely by orientation, it would be difficult for the same model to apply among the HBLR and non-HBLR S2s without invoking intrinsic physical differences\", what are the physical meanings for the absence of polarized broad lines in Seyfert 2 galaxies? Do they really not have \"broad\" line region?  It has been suggested that some of Seyfert 2 galaxies without PBLs as a new subclass are probably  lack of broad line region (Tran 2001, 2003). Tran (2003) confirmed that polarized hidden broad-line region S2s share many similar large-scale characteristics with BLS1s, while non-HBLR S2s do not. Deluit (2004) analyzed {\\em Beppo}-SAX data of Seyfert 2s and found that non-HBLR S2s are different in hard X-rays (15$-$136keV) from those with hidden BLR Seyfert 2. There is growing evidence for that not {\\em all} Seyfert 2 galaxies might be intrinsically similar in nature. As we discuss in detail in \\S4 (see Table 1. for possible types of Seyfert 2 galaxies), some non-HBLR S2s may result from fuel-depleting Seyfert 1 and 2 galaxies if the dusty torus is supplying matter onto the black holes (Krolik \\& Begelman 1988). These objects could be characterized by low or absent absorption in X-ray band, they thus might be the progenitors of the optically-selected unabsorbed Seyfert 2 galaxies defined by Panessa \\& Bassini (2002)\\footnote{The roles of the gas-dust ratio has been discussed by Maiolino et al. (2001) and Gallo et al. (2006)}. This definitely makes it more complicate to study the physics of non-HBLR S2s.  However, this paper focuses on the absorbed non-HBLR S2s. If they were powered by less massive black holes and obscured by torus at larger viewing angles (Tran 2003), what are their counterparts at low orientation? This motivates us to explore the relation between non-HBLR S2s and narrow line Seyfert 1 galaxies.   {\\footnotesize \\begin{center}{\\sc Table 1 The possible types of Seyfert 2 galaxies} \\vglue 0.3cm \\begin{tabular}{lcc}\\hline\\hline &  absorbed                           & unabsorbed \\\\ &  ($N_{\\rm H}\\ge 10^{22}$cm$^{-2}$)  & ($N_{\\rm H}< 10^{22}$cm$^{-2}$)\\\\ \\hline non-HBLR    &  $\\surd$                                    & $\\surd$ (unabsorbed Seyfert 2)   \\\\ HBLR        &  $\\surd$(classical Seyfert 2)               & $\\surd$ (Gallo et al. 2006) \\\\ \\hline  \\end{tabular} \\parbox{3.1in} {\\baselineskip 9pt \\noindent {\\sc Note:} The symbol $\\surd$ indicates that this type is generally observed. Polarized spectroscopic measurements are available only for four unabsorbed Seyfert 2 galaxies, NGC 2992 (Rix et al. 1990), NGC 5995 (Lumsden \\& Alexander 2001, Tran 2001), NGC 7590 (Heisler et al. 1997) and NGC 4501 (Tran 2003, Cappi et al. 2006). A polarized broad H$\\alpha$ line has been found in the first two objects, but not in NGC 4501 and NGC 7590. Gallo et al. (2006) find $\\sim 50\\%$ of {\\em ROSAT}-selected Seyfert galaxies are low absorption Seyfert 2s.} \\end{center} }   As a distinct subclass, NLS1s have very narrow Balmer lines [FWHM (H$\\beta$) $\\le 2000~\\rm km~ s^{-1}$], strong Fe {\\sc ii} lines (Osterbrock \\& Pogge 1985), and violent variability in soft X-ray band (Boller et al. 1996). They likely contain less massive black holes at the Eddington limit rates (Boller et al. 1996; Laor et al. 1997) and can be explained by slim disk (Wang et al. 1999; Wang \\& Zhou 1999, Mineshige et al. 2000; Wang \\& Netzer 2003; Wang 2003; Ohsuga et al. 2003, Chen \\& Wang 2004; Collin et al. 2002; Collin \\& Kawaguchi 2004, Kawaguchi et al. 2004). The brighter soft X-ray fluxes favor the pole-on orientation hypothesis since an edge-on thick disk is dimmer than the lower inclination (Madau 1988; Boller et al. 1996; Leighly 1999a, b; see more detail calculations of Watarai et al. 2005). It is thus expected that the soft X-ray selected NLS1s tend to have a pole-on orientation to observers (Boller et al. 1996). This is further supported by the polarization observations showing that most of the soft-X-ray-selected {\\em ROSAT} AGNs have polarization lower than $\\le 1\\%$ and no clear optical reddening (Grupe et al. 1998a). Optically-selected NLS1s from the Sloan Digital Sky Survey (SDSS)  are weak in soft X-ray bands and hence have lower Eddington ratios (William et al. 2004). If the dusty tori generally exist in Seyfert galaxies and their orientations are random,  what are the counterparts of NLS1s viewed at larger angles? The presence of the non-HBLR S2s and NLS1s as new members of Seyfert galaxies have strong impact on the classical unified model. We suggest that non-HBLR S2s are the counterpart of NLS1 viewed at larger angles. If so, the black hole masses, accretion rates (also Eddington ratios) as fundamentally intrinsic parameters should have same distributions.  In this paper, we compared large-scale properties of non-HBLR S2s and NLS1s to quantitatively test the above issues. We find non-HBLR S2s share the potential isotropic characters with NLS1s while they are different greatly in the potential anisotropic properties. This  may suggest that they are basically the same objects but viewed from different angles, adding new ingredient to the classical unification scheme. ",
        "conclusions": "We assemble two sub-samples of consisting of 44 NLS1s and 42 non-HBLR S2s to examine whether they have similar or same central engines. Most of NLS1s in the present sample have a steep soft X-ray spectrum ($\\Gamma_{\\rm SX}>2$) whereas the non-HBLR S2s have strong absorption in X-ray band. We estimate black hole masses and accretion rates of these samples. We find that: 1) the two kinds of Seyfert galaxies have same distributions of black hole masses from FWHM([O {\\sc iii}]) at $p_{\\rm null}=44.8\\%$ in ASURV test; 2) the \\oiii luminosities as an accretion rate indicator are similar and the black holes have accretion rates close to or above the Eddington limit; 3) the $f_{60\\mu}/f_{25\\mu}$ as orientations of the dusty torus are very different. We suggest that non-HBLR S2s are counterpart of NLS1s at edge-on orientation.  This hypothesis is further supported by the comparison with other indicators. NLS1s and non-HBLR S2s can be unified based on orientation, but they have less massive black holes with higher accretion rates different from the broad line Seyfert 1 galaxies and HBLR S2s. This hypothesis sets up a scenario of a population of less massive black holes at {\\em all} orientations, which have higher accretion rates and are evolving to broad line Seyfert galaxies. Future work on unification scheme with the growth of the black holes and electron scattering screen are need to be done in future."
    },
    "0606/astro-ph0606497_arXiv.txt": {
        "abstract": "% The planets capture model for the eruption of V838 Mon is discussed. We used three methods to estimate the location where the planets were consumed. There is a nice consistency for the results of the three different methods, and we find that the typical stopping / slowing radius for the planets is about 1R$_\\odot$. The three peaks in the optical light curve of V838 Mon are either explained by the swallowing of three planets at different radii or by three steps in the slowing down process of a single planet. We discuss the other models offered for the outburst of V838 Mon, and conclude that the binary merger model and the planet/s scenario seem to be the most promising. These two models have several similarities, and the main differences are the stellar evolutionary stage, and the mass of the accreted material. We show that the energy emitted in the V838 Mon event is consistent with the planets scenario. We suggest a few explanations for the trigger for the outburst and for the double structure of the optical peaks in the light curve of V838 Mon. ",
        "introduction": "V838 Mon had an extraordinary multi-stage outburst during the beginning of 2002. Imaging revealed the presence of a spectacular light echo around this object (Bond et al. 2003). The amplitude of the outburst in the optical band was about 9.5 mag. The post-outburst spectroscopic observations of V838 Mon showed that it was very red throughout the eruption and long after it ended (Munari et al. 2002; Banerjee \\& Ashok 2002; Kimeswenger et al. 2002; Evans et al. 2003; Kaminsky \\& Pavlenko 2005; Tylenda 2005). This is inconsistent with an exposed hot white dwarf in novae.  Evans et al. (2003) and Retter \\& Marom (2003) concluded that the progenitor star of V838~Mon probably had a radius of $\\sim 8 R_\\odot$, a temperature of $\\sim 7,300$ K and a luminosity of $\\sim 100-160 L_\\odot$. Tylenda, Soker \\& Szczerba (2005b) presented a detailed analysis of the progenitor. They argued that V838~Mon is likely a young binary system that consists of two $5-10 M_\\odot$ B stars and that the erupting component is a main-sequence or pre-main sequence star. They also estimated for the progenitor a temperature of $\\sim 4,700-30,000$ K and a luminosity of $\\sim 550-5,000 L_\\odot$. Tylenda (2005) adopted a mass of $\\sim 8 M_\\odot$ and a radius of $\\sim 5 R_\\odot$ for the progenitor of V838~Mon. There is additional supporting evidence that the erupting star belongs to a binary system with a hot B secondary star (Munari \\& Desidera 2002; Wagner \\& Starrfield 2002; Munari et al. 2005).  Spectral fitting suggested that V838~Mon had a significant expansion from a few hundreds to several thousands stellar radii in a couple of months during the outburst (Soker \\& Tylenda 2003; Retter \\& Marom 2003; Tylenda 2005; Rushton et al. 2005). Interferometric observations at the end of 2004 with the Palomar Testbed Interferometer confirmed the huge radius of the post-outburst star with an estimate of $1,570 \\pm 400 R_\\odot$ and suggested some asymmetric structure (Lane et al. 2005). There are only very rough estimates on the mass of the ejecta (Rushton et al. 2003; Lynch et al. 2004; Tylenda 2005).  \\subsection{Models for the outburst}  Soon after its outburst, V838~Mon was recognized as the prototype of a new class of stars (Munari et al. 2002; Bond et al. 2003), which currently consists of three objects: M31RV (Red Variable in M31 in 1988; Rich et al. 1989; Mould et al. 1990; Bryan \\& Royer 1992), V4332~Sgr (Luminous Variable in Sgr, 1994; Martini et al. 1999), and V838~Mon (Peculiar Red Variable in 2002), plus three candidates -- CK~Vul, which was identified with an object that had a nova-like event in the year 1670 (Shara \\& Moffat 1982; Shara, Moffat \\& Webbink 1985; Kato 2003; Retter \\& Marom 2003), V1148 Sgr, which had a nova outburst in 1943 and was reported to have a late type spectrum (Mayall 1949; Bond \\& Siegel 2006), and the peculiar variable in Crux that erupted in 2003 (Della Valle et al. 2003).  So far, seven explanations for the eruption of these objects have been supplied. The first invokes a nova outburst from a compact object, which is embedded inside a common red giant envelope (Mould et al. 1990). In the second model, an atypical nova explosion on the surface of a cold white dwarf was suggested (Iben \\& Tutukov 1992; Boschi \\& Munari 2004). Soker \\& Tylenda (2003) proposed a scenario in which a main sequence star merged with a low-mass star. This model was lately revised by Tylenda \\& Soker (2006), and summarized by Soker \\& Tylenda (2006). Van-Loon et al. (2004) argued that the eruption was a thermal pulse of an AGB star. Munari et al. (2005) explained the outburst of V838~Mon by a shell thermonuclear event in the outer envelope of an extremely massive (M $\\sim 65 M_\\odot$) B star. Lawlor (2005, 2006) proposed another mechanism for the eruption of V838~Mon. He invoked the born-again phenomenon to explain the first peak in the light curve and altered the model by adding accretion from a secondary main-sequence star in close orbit to explain the second peak in the optical light curve of V838~Mon.  A promising model for the peculiar eruption of V838~Mon was suggested by Retter \\& Marom (2003) and was further developed by Retter et al. (2006). This paper summarizes this model for V838 Mon and similar objects.  The peculiar and enigmatic outburst of V838 Mon led to a specific meeting dedicated to this phenomenon that was held in La Palma, Spain on 2006 May, in which the first author of this paper presented the planets-swallowing model of V838 Mon. The most important result that was presented in the conference is probably two new very reliable distance estimates that are consistent with a distance of about 6 kpc to V838 Mon (Sparks 2006; Afsar \\& Bond 2006). This is somewhat smaller than what was previously believed (e.g., Bond et al. 2003), and thus it has some impact on the energy emitted in the outburst. ",
        "conclusions": ""
    },
    "0606/astro-ph0606174_arXiv.txt": {
        "abstract": "We suggest that the mass lost during the evolution of very massive stars may be dominated by optically thick, continuum-driven outbursts or explosions, instead of by steady line-driven winds.  In order for a massive star to become a Wolf-Rayet star, it must shed its hydrogen envelope, but new estimates of the effects of clumping in winds from O-type stars indicate that line driving is vastly insufficient.  We discuss massive stars above roughly 40--50 M$_{\\odot}$, which do not become red supergiants, and for which the best alternative is mass loss during brief eruptions of luminous blue variables (LBVs).  Our clearest example of this phenomenon is the 19th century outburst of $\\eta$ Carinae, when the star shed 12--20 M$_{\\odot}$ or more in less than a decade.  Other examples are circumstellar nebulae of LBVs and LBV candidates, extragalactic $\\eta$ Car analogs (the so-called ``supernova impostors''), and massive shells around supernovae and gamma-ray bursters.  We do not yet fully understand what triggers LBV outbursts or what supplies their energy, but they occur nonetheless, and present a fundamental mystery in stellar astrophysics.  Since line opacity from metals becomes too saturated, the extreme mass loss probably arises from a continuum-driven wind or a hydrodynamic explosion, both of which are insensitive to metallicity.  As such, eruptive mass loss could have played a pivotal role in the evolution and ultimate fate of massive metal-poor stars in the early universe. If they occur in these Population III stars, such eruptions would also profoundly affect the chemical yield and types of remnants from early supernovae and hypernovae thought to be the origin of long gamma ray bursts. ",
        "introduction": "Mass loss is a critical factor in the evolution of a massive star.  In addition to the direct reduction of a star's mass, it profoundly affects the size of its convective core, its core temperature, its angular momentum evolution, its luminosity as a function of time, and hence its evolutionary track on the HR diagram and its main-sequence (MS) lifetime (e.g., Chiosi \\& Maeder 1986).  Wolf-Rayet (WR) stars are the descendants of massive stars as a consequence of mass loss in the preceding H-burning phases, during which the star sheds its H envelope (Abbott \\& Conti 1987; Crowther 2006).  While the maximum initial mass of stars is thought to be $\\sim$150 M$_{\\odot}$ (Figer 2005; Kroupa 2005), WR stars do not have masses much in excess of 20 M$_{\\odot}$ (Crowther 2006).\\footnote{By ``WR stars'' we mean H-deficient WR stars (core-He burning phases or later), and not the luminous H-rich WNL stars (Crowther et al.\\ 1995), which are probably still core-H burning.}  Thus, very massive stars have the immense burden of removing 30--130 M$_{\\odot}$ during their lifetime before the WR phase, unless they explode first.  Stellar evolution calculations prescribe $\\dot{M}$($t$) based on semiempirical values, so it is important to know when most of this mass loss occurs.  In this letter we address the question of whether this mass loss occurs primarily via steady stellar winds, or instead through violent, short-duration eruptions or explosions.  Recent studies of hot star winds indicate that mass-loss rates on the MS are much lower than previously thought.  These mass-loss rate reductions are significant enough to affect MS evolution, but they also raise an important question: {\\it If mass loss via stellar winds is insufficient to strip off a star's H envelope and form a WR star, then how and when does it occur}?  Simultaneously, observations of nebulae around luminous blue variables (LBVs) and LBV candidates have revealed very high ejecta masses -- of order 10 M$_{\\odot}$.  In $\\eta$ Car we know that the mass was ejected in a single outburst and is not swept-up ambient material.  Together, these facts suggest that short-duration outbursts like the 19th century eruption of $\\eta$ Car could dominate mass lost during the lives of the most massive stars, and would be critical to form WR stars.  As detailed below, the extreme mass-loss rates of these bursts imply that line opacity is too saturated to drive them, so they must instead be either continuum-driven super-Eddington winds or outright hydrodynamic explosions.  Unlike steady winds driven by lines, the driving in these eruptions may be largely independent of metallicity, and might play a role in the mass loss of massive metal-poor stars (Population III stars). ",
        "conclusions": ""
    },
    "0606/astro-ph0606032_arXiv.txt": {
        "abstract": "We discuss new functionality of the spectral simulation code CLOUDY which allows the user to calculate grids with one or more initial parameters varied and formats the predicted spectra in the standard FITS format.  These files can then be imported into the x-ray spectral analysis software XSPEC and used as theoretical models for observations.  We present and verify a test case.  Finally, we consider a few observations and discuss our results. ",
        "introduction": "X-ray spectrometers typically record photon counts per energy bin.  The photon count for each bin is equal to the integral of the incident spectrum times an instrument response function, which is a function of both photon energy and detector bin.  In general, this integral cannot be reliably inverted to recover the incident spectrum, in part because inversion techniques tend to be unstable to small changes in the photon count.  The solution then is to compare the photon counts recorded by the spectrometer with best-fit theoretical photon counts, calculated by integrating the product of the known instrument response function with theoretical spectra.  XSPEC (last described by Arnaud 1996) is an open-source, X-ray spectral-fitting program, first developed in 1983 and still in development today.  The current \\textit{de facto} standard for X-ray spectral analyis, XSPEC calculates theoretical photon counts from theoretical spectra to find the best fit between the observed and theoretical photon counts.  In order to do this, XSPEC must have not only data files containing the observed photon counts, background spectrum, and instrumental response (all of which can be readily obtained), but also theoretical models of any and all spectra that the user believes will accurately represent the actual source spectrum (as a whole or in part).  XSPEC comes bundled with many such theoretical models, but can also import external models via two methods: an external subroutine or a table file.  The table file option can be used if calculating an individual model is too CPU-intensive. The commands which control the importing of table models are as follows: \\texttt{atable} for additive tabular models (intended to represent sources of emission), \\texttt{mtable} for multiplicative tabular models (for absorption, filtering, and extinction), and \\texttt{etable} for exponential tabular models (for any exponential effects).  The purpose of this work is to greatly expand the pool of theoretical models by creating functionality allowing the existing spectral simulation code CLOUDY to produce predicted spectra in a format which the user can import into XSPEC. CLOUDY predictions have been used in XSPEC before (recent examples include Kraemer et al. 2005 and Turner et al. 2005). However, the new functionality discussed here greatly simplifies the process.  We consider only additive and multiplicative models. ",
        "conclusions": ""
    },
    "0606/hep-ph0606261_arXiv.txt": {
        "abstract": "The neutrino chirality-flip process under the conditions of the supernova core is investigated in detail with the plasma polarization effects in the photon propagator taken into account. It is shown that the contribution of the proton fraction of plasma is essential. New upper bounds on the neutrino magnetic moment are obtained: $\\mu_\\nu < (0.5 - 1.1) \\, \\times 10^{-12} \\, \\mu_{\\rm B}$ from the limit on the supernova core luminosity for $\\nu_R$ emission, and $\\mu_\\nu < (0.4 - 0.6) \\, \\times 10^{-12} \\, \\mu_{\\rm B}$ from the limit on the averaged time of the neutrino spin-flip. The best astrophysical upper bound on the neutrino magnetic moment is improved by the factor of 3 to 7. ",
        "introduction": "Nonvanishing neutrino magnetic moment leads to various chirality-flipping processes when the left-handed neutrinos produced inside the supernova core during the collapse could change their chirality becoming sterile with respect to the weak interaction. These sterile neutrinos would escape from the core leaving no energy to explain the observed luminosity of the supernova.  This process was investigated by several authors. R. Barbieri and R.~N. Mohapatra~\\cite{Barbieri:1988} considered the neutrino spin-flip via both $\\nu_L e^- \\to \\nu_R e^-$ and $\\nu_L p \\to \\nu_R p$ scattering processes in the inner core of a supernova immediately after the collapse. However, they did not consider the essential plasma polarization effects in the photon propagator, and the photon dispersion was taken in a phenomenolical way, by inserting an {\\it ad hoc} thermal mass into the vacuum photon propagator.  Imposing for the $\\nu_R$ luminosity $Q_{\\nu_R}$ the upper limit of $10^{53}$ ergs/s, the authors~\\cite{Barbieri:1988} obtained the upper bound on the neutrino magnetic moment: \\begin{eqnarray} \\mu_\\nu < (0.2-0.8) \\times 10^{-11} \\, \\mu_{\\rm B} \\,. \\label{eq:lim_Barbi} \\end{eqnarray}  Later on, A. Ayala, J.~C. D'Olivo and M. Torres~\\cite{Ayala:1999,Ayala:2000} used the formalism of the thermal field theory to take into account the influence of hot dense astrophysical plasma on the photon propagator. The upper bound on the neutrino magnetic moment was improved by them in the factor of 2: \\begin{eqnarray} \\mu_\\nu < (0.1-0.4) \\times 10^{-11} \\, \\mu_{\\rm B} \\,. \\label{eq:lim_Ayala} \\end{eqnarray} However, those authors considered only the contribution of plasma electrons, and neglected the proton fraction. Thus, the reason exists to reconsider the neutrino spin-flip processes in the supernova core more attentively.  We will show in part, that the proton contribution into the photon propagator is essential, as well as the scattering on plasma protons. ",
        "conclusions": "\\begin{itemize} \\item We have investigated in detail the neutrino chirality-flip process under the conditions of the supernova core. The plasma polarization effects caused both by electrons and protons were taken into account in the photon propagator. The rate $\\Gamma (E)$ of creation of the right-handed neutrino with the fixed energy $E$, the energy spectrum, and the luminosity have been calculated. \\item From the limit on the supernova core luminosity for $\\nu_R$ emission, we have obtained the upper bound on the neutrino magnetic moment $ \\mu_\\nu < (0.5 - 1.1) \\, \\times 10^{-12} \\, \\mu_{\\rm B}\\,. $ \\item From the limit on the averaged time of the neutrino spin-flip, we have obtained the upper bound $ \\mu_\\nu < (0.4 - 0.6) \\, \\times 10^{-12} \\, \\mu_{\\rm B}\\,. $ \\item We have improved the best astrophysical upper bound on the neutrino magnetic moment by the factor of 3 to 7. \\end{itemize}  \\newpage"
    },
    "0606/astro-ph0606048_arXiv.txt": {
        "abstract": "{The most recently celebrated cosmological implications of the cosmic microwave background studies with WMAP (2006), though fascinating by themselves, do, however, create some extremely hard conceptual challenges for the presentday cosmology. These recent extremely refined WMAP observations seem to reflect a universe which was extremely homogeneous at the recombination age and thus is obviously causally closed at the time of the cosmic recombination era. From the very tiny fluctuations apparent at this early epoch the presently observable nonlinear cosmic density structures can, however, only have grown up, if in addition to a mysteriously high percentage of dark matter an even higher percentage of dark energy is admitted as drivers of the cosmic evolution. The required dark energy density, on the other hand, is nevertheless 120 orders of magnitude smaller then the theoretically calculated value. These are outstanding problems of present day cosmology onto which we are looking here under new auspices. We shall investigate in the following, up to what degree a universe simply abolishes all these outstanding problems in case it reveals itself as universe of a constant total energy. As we shall show basic questions like: How could the gigantic mass of the universe of about $10^{80}$ proton masses at all become created? - Why is the presently recognized and obviously indispensable cosmic vacuum energy density so terribly much smaller than is expected from quantumtheoretical considerations, but nevertheless terribly important for the cosmic evolution? -Why is the universe within its world horizon a causally closed system? - , can perhaps simply be answered, when the assumption is made, that the universe has a constant total energy with the consequence that the total mass density of the universe (matter and vacuum) scales with $R_\\mathrm{u}^\\mathrm{-2}$. Such a scaling of matter and vacuum energy abolishes the horizon problem, and the cosmic vacuum energy density can easily be reconciled with its theoretical expectation values. In this model the mass of the universe increases linearily with the world extension $R_\\mathrm{u}$ and can grow up from a Planck mass as a vacuum fluctuation.} ",
        "introduction": "This paper aims at the investigation of possible cosmological consequences of the hypothesis that the total mass density - or equivalently the energy density, of the universe (baryonic and dark matter, vacuum energy, photons) behaves according to a $R_\\mathrm{u}^{-2}$-law where $R_\\mathrm{u}$ denotes the cosmic scale or, if extended to the light horizon, the extension of the universe. In this respect $R_\\mathrm{u}$ can be called the distance between two arbitrary space points comoving with the Robert- son-Walker-like homologous cosmological expansion of the universe, validating the relation $\\dot{R_\\mathrm{u}}/R_\\mathrm{u}=\\dot{S}/S$, if $S$ for instance is the distance between any two freely comoving galaxies (dots on top of quantities mean derivatives with respect to cosmic time. This relation was already suggested by considerations carried out by Kolb (1989), Overduin \\& Fahr (2001, 2003) or Fahr \\& Heyl (2006). A detailed justification for such a scaling behaviour will be provided in section 9 at the end of this paper where we will show that such a density scaling indeed might help to describe a reasonably realistic scenario which correctly delivers the basic parameters of our expanding universe.  In the following we shall work in the frame of General Relativity and shall base ourselves on the Friedmann equations. The only special prerequisites for our investigation are the assumption of a total mass density scaling with $R_\\mathrm{u}^\\mathrm{-2}$ and a topologically flat universe with a curvature parameter $k=0$, which is also strongly suggested nowadays by the recent WMAP data (WMAP 2006).  A short suggestive substantiation for our assumption of a $R_\\mathrm{u}^\\mathrm{-2}$-scaling can be derived from the 1. Friedmann equation  \\setlength{\\mathindent}{0pt} \\begin{equation} H^\\mathrm{2}  = \\left( {\\frac{{\\dot R_\\mathrm{u}}}{R_\\mathrm{u}}} \\right)^\\mathrm{2}  = \\frac{{8\\pi G\\rho _\\mathrm{tot}}}{3}  - \\frac{kc^\\mathrm{2}}{{R_\\mathrm{u}^\\mathrm{2}}}, \\label{1}% \\end{equation} with $H$ the Hubble parameter, $R_\\mathrm{u}$ and $\\dot{R_\\mathrm{u}}$ the extension of the universe and its time derivative, respectively, $G$ the constant of gravitation, $c$ the velocity of light, and $\\rho_\\mathrm{tot}$ the total mass density of all gravitating constituents of the universe (e.g. baryonic and dark matter, equivalent mass of the vacuum energy). For a flat universe, $k=0$, the above equation can be written as  \\setlength{\\mathindent}{0pt} \\begin{equation} \\dot {R_\\mathrm{u}}^\\mathrm{2}  = \\frac{{8\\pi G\\rho _\\mathrm{tot}}}{3} R_\\mathrm{u}^\\mathrm{2} = \\Phi_\\mathrm{eff}, \\label{2}% \\end{equation} where the expression $8\\pi G\\rho _\\mathrm{tot}R_\\mathrm{u}^\\mathrm{2}/3$ can be interpreted as an effective cosmic action potential $\\Phi_\\mathrm{eff}$ of gravitation which is the driving quantity for the dynamics of the universe. Respecting the cosmological principle means, that the potential $\\Phi_\\mathrm{eff}$ is present at each space point of the expanding universe. We now take as our basic assumption that the gravitational energy $\\mathrm{d}m\\Phi_\\mathrm{eff}$, related to each mass element $\\mathrm{d}m$ of our homogeneously matter-filled universe, just equals its rest mass energy, i.e.  \\setlength{\\mathindent}{0pt} \\begin{equation} \\Phi_\\mathrm{eff} \\mathrm{d}m = \\frac{{8\\pi G\\rho _\\mathrm{tot}}}{3} R_\\mathrm{u}^\\mathrm{2} \\mathrm{d}m = c^\\mathrm{2} \\mathrm{d}m. \\label{3}% \\end{equation} Then, this evidently results in a total density $\\rho_\\mathrm{tot}$ given by  \\setlength{\\mathindent}{0pt} \\begin{equation} \\rho_\\mathrm{tot}={\\frac{{3c^\\mathrm{2}}}{{8\\pi GR_\\mathrm{u}^\\mathrm{2}}}} \\propto {\\frac {1}{R_\\mathrm{u}^\\mathrm{2}}}, \\label{4}% \\end{equation} which expresses the assumed $R_\\mathrm{u}^\\mathrm{-2}$-scaling as the basis of our further investigations. A deeper view to this relation is offered in section 9 of this paper. ",
        "conclusions": "The assumed $R_\\mathrm{u}^\\mathrm{-2}$ scaling of matter and vacuum energy densities combined with the assumption of a flat universe with curvature $k=0$ leads to a universe which does not face a horizon problem any longer and thus does not require a cosmic inflation at the beginning. Furthermore, the theoretically calculated and unexplainable high amount of vacuum energy - with a value about $10^\\mathrm{120}$ higher than observed - one perfectly fits into the idea of an universe with a $R_\\mathrm{u}^\\mathrm{-2}$ scaling of the matter and vacuum energy density. In addition, it has been shown, that the present universe might have its origin in a quantum mechanical fluctuation that took place in the Planck era and that the vacuum energy is nothing else but the scaling rest energy associated with half a Planck mass within a Planck volume $V_\\mathrm{Pl}$. Finally, the whole mass of the universe can be explained by the accumulation of half Planck masses up to the present time which are generated as virtual quantum mass releases per Planck time and permitted to become real in the expanding universe. These results have been obtained on the basis of a so-called \"economical\" universe where the rest energy of the released Planck masses are compensated by the negative gravitational binding energy and which expresses one of the most fundamental laws of physics - the conservation of energy."
    },
    "0606/astro-ph0606081_arXiv.txt": {
        "abstract": "It is generally thought that conservative mass transfer in Algol binaries causes their orbits to be wider, in which the less massive star overflows its Roche-lobe. The observed decrease in the orbital periods of some Algol binaries suggests orbital angular momentum loss during the binary evolution, and the magnetic braking mechanism is often invoked to explain the observed orbital shrinkage. Here we suggest an alternative explanation, assuming that a small fraction of the transferred mass forms a circumbinary disk, which extracts orbital angular momentum from the binary through tidal torques. We also perform numerical calculations of the evolution of Algol binaries with typical initial masses and orbital periods. The results indicate that, for reasonable input parameters, the circumbinary disk can significantly influence the orbital evolution, and cause the orbit to shrink on a sufficiently long timescale. Rapid mass transfer in Algol binaries with low mass ratios can also be accounted for in this scenario. ",
        "introduction": "An Algol binary is a semidetached binary system consisting of (1) an early type, main sequence primary component which does not fill its Roche lobe, and (2) a lobe-filling, less massive star that is substantially above the main sequence. The less massive star is cooler, fainter and larger \\citep{Giur81,pete01}. It is believed that star (2) is initially more massive, and evolves first to overflow its Roche lobe to transfer mass to star (1). After the rapid mass exchange, the lobe-filling star (2) becomes less massive in the Algol binary stage.  The evolution of Algol binaries has been investigated extensively (e.g., Plavec 1968; Paczynski 1971; Refsdal \\& Weigert 1969). Although conservative mass transfer seems to roughly reproduce the observed characteristics of a considerable fraction of Algol binaries (e.g. Nelson \\& Eggleton 2001; De Loore \\& van Rensbergen 2004), nonconservative evolution with mass and orbital angular momentum loss taken into account is required for better comparison between theories and observations (Thomas 1977; Refsdal et al. 1974; Sarna 1993). With regard to the orbital evolution, conservative mass transfer in Algol binaries leads to increase of the orbital periods because mass is transferred from the less massive star to the more massive component (Huang 1963). However, there is strong evidence showing that the orbital periods of some Algol binaries are actually decreasing \\citep{qian00a,qian00b,qian01a,qian01b,qian01c,qian02,lloy02}. Although it is uncertain whether the measured period changes are long-term (secular) changes or transient fluctuations, they clearly demonstrate that these systems must undergo mass and angular momentum loss during the evolution. Moreover, \\citet{qian02} find that the decreasing rates $\\dot{P}_{\\rm orb}$ scale with the orbital periods $P_{\\rm orb}$ roughly as $-\\dot{P}_{\\rm orb}\\propto P_{\\rm orb}$.  This orbital decay may be explained by angular-momentum loss via magnetic stellar winds \\citep{verb81}. Sarna, Muslimov, \\& Yerli (1997) and Sarna, Yerli, \\& Muslimov (1998) proposed that dynamo action can occur in mass-losing stars in Algol-type binaries, and produce large-scale magnetic fields, which lead to magnetic braking of the stars. However, there could exist some potential difficulties for the magnetic braking mechanism (see below). Here we suggest an alternative interpretation assuming that a small fraction of the transferred mass would form a circumbinary (CB) disk surrounding the binary system \\citep[see also][]{van94}. Previous works \\citep{taam01,chen06} have already indicated that the CB disk can efficiently remove the orbital angular momentum and accelerate the binary evolution in accreting white dwarf or neutron star binaries. We describe the input physics that is necessary for the evolution model in section 2. Numerically calculated results for the evolutionary sequences of Algol binaries in two cases are presented in section 3. We make brief discussion and conclude in section 4. ",
        "conclusions": "It is an open question why some of the Algol systems show orbital shrinkage while others not. Although it is generally believed that there must be orbital angular momentum loss during the evolution of these binaries, the appropriate mechanism(s) for angular momentum loss has not been verified. Based on theoretical model of \\citet{spru01} and \\citet{taam01}, we suggest that a CB disk may play an important role in determining the orbital change in Algol systems. Note that our CB disk hypothesis does not exclude the effect of (magnetized) mass loss. On the contrary, the formation of the CB disk may be closely related to the mass loss processes in the binary evolution. The existence of the CB disks may be revealed by their infrared radiation, as in GG Tau \\citep{rodd96}. The Monte-Carlo simulation by \\citet{wood99} has reproduced the morphology of the near IR radiation which is in accord with the observations by \\citet{rodd96}.  In this paper we have calculated the evolutionary sequence for Algol binaries with typical masses and orbital periods, taking into account both magnetic braking and the CB disk. The detailed calculations indicate that the orbital evolution of Algol binaries can be significantly affected by the CB disk, especially for Case A mass transfer. With adequate values of $\\delta$, it is possible to account for the decrease of the orbital periods in some Algol binaries. In addition, the existence of CB disk can accelerate the evolution and enhance the mass transfer rates, as already noticed in the evolutionary study of CVs \\citep{taam01}. This suggests a plausible explanation for the rapid mass transfer observed in a few Algol binaries with very low mass ratios \\citep{qian02a}.  Finally, we note that the magnitude of the mass feeding parameter $\\delta$ is highly uncertain, let alone its variation with the mass transfer rates. These uncertainties make it difficult to present direct comparison between observations and theoretical predications for Algol binaries. However, our analysis provides a reasonable mechanism for angular momentum loss in semi-detached binaries including Algol binaries. More detailed multi-wavelength observations can provide stringent tests for the CB disk model, and the theories of the evolution of Algol binaries."
    },
    "0606/astro-ph0606562_arXiv.txt": {
        "abstract": "{We present corrections for the change in the apparent scalelengths, central surface brightnesses and axis ratios due to the presence of dust in pure disk galaxies, as a function of inclination, central face-on opacity in the B-band (${\\tau}^{f}_{\\rm B}$) and wavelength. The correction factors were derived from simulated images of disk galaxies created using geometries for stars and dust which can reproduce the entire spectral energy distribution from the ultraviolet (UV) to the Far-infrared (FIR)/submillimeter (submm) and can also account for the observed surface-brightness distributions in both the optical/Near-infrared and FIR/submm. We found that dust can significantly affect both the scalelength and central surface brightness, inducing variations in the apparent to intrinsic quantities of up to 50$\\%$ in scalelength and up to 1.5 magnitudes in central surface brightness. We also identified some astrophysical effects for which, although the absolute effect of dust is non-negligible, the predicted variation over a likely range in opacity is relatively small, such that an exact knowledge of opacity is not needed. Thus, for a galaxy at a typical inclination of $37^{\\circ}$ and having any ${\\tau}^{f}_{\\rm B}>2$, the effect of dust is to increase the scalelength in B relative to that in I by a factor of $1.12 \\pm 0.02$  and to change the B-I central colour by $0.36\\pm 0.05$ magnitudes. Finally we use the model to analyse the observed scalelength ratios between B and I for a sample of disk-dominated spiral galaxies, finding that the tendency for apparent scalelength to increase with decreasing wavelength is primarily due to the effects of dust. ",
        "introduction": "\\label{intro}  A primary goal of modern studies of star-forming galaxies is to understand how these systems were assembled over cosmic time. If the disks of spiral galaxies grow from the inside out, as predicted by semi-analytical hierarchical models for galaxy formation (e.g. Mo, Mao \\& White 1998), one would predict the stellar populations to be younger and have lower metallicity in the outer disk than in the inner disk, such that local universe galaxies should be intrinsically larger at the shorter wavelengths where light from the young stellar population is more prominent. For the same reason one would expect the intrinsic sizes of spiral disks to be larger at the current epoch than at higher redshift. Observationally, such predictions can be tested in two ways. One way is to compare the spatial distribution of the constituent stellar populations at different wavelengths, for local universe galaxies. Another way is to look for structural differences in galaxies observed at different cosmological epochs, at the same rest frame wavelength. Both methods require an analysis of the surface-brightness distribution of spiral galaxies in the optical and near-infrared (NIR) to quantify the distribution of starlight, for example by deriving the scalelength of the disk. This is done by fitting observed images with models for the surface-brightness distribution of stellar light, whereby the disk component is usually specified by an exponential distribution. The derived exponential scalelengths can either be intercompared between different wavelengths for local universe galaxies (Elmegreen \\& Elmegreen 1984, Peletier et al. 1994, Evans 1994, De Jong 1996a,b, de Grijs 1998, Cunow 1998, 2001, 2004, MacArthur et al. 2003, M\\\"ollenhoff 2004), or between galaxies at different redshifts at a given wavelength (Lilly et al. 1998, Simard et al. 1999, Ravindranath et al. 2004, Trujillo \\& Aguerri 2004, Trujillo et al. 2005, Barden et al. 2005). However, the appearance of disk galaxies is strongly affected by dust and this effect is different at different wavelengths and for different opacities. This has consequences not only for the derivation of the variation of intrinsic scalelength with wavelength, but also for the variation of intrinsic scalelength with cosmological epoch, since the opacity of disk galaxies is expected to have been systematically higher in the past ( e.g. Dwek 1998, Pei et al. 1999).  The effect of dust on the observed scalelengths and central surface brightnesses of disk galaxies has been previously modelled by Byun et al. (1994), Evans (1994) and Cunow (2001). By means of radiative transfer calculations these works investigated the dependence of this effect on star/dust geometry, opacity, inclination and wavelength. Recently, a better knowledge of the star/dust geometry has been obtained through a joint consideration of the direct starlight, emitted in the ultraviolet (UV)/optical/NIR, and of the starlight which is re-radiated in the Far-infrared (FIR)/submillimeter (submm). In a series of papers devoted to modelling the spectral energy distributions (SEDs) we derived geometries of the distribution of stellar light and dust that are successful in reproducing not only the observed integrated SEDs, but also the observed radial profiles both in the optical/NIR and FIR/submm (Popescu et al. 2000; hereafter Paper~I, Misiriotis et al. 2001; hereafter Paper~II, Popescu et al. 2004, see also Popescu \\& Tuffs 2005 \\footnote{A simplified version of this geometrical prescription has been applied by Misiriotis et al. (2004) to fit the FIR SEDs of bright IRAS galaxies.}).  In this paper, the fourth in this series, we use these derived distributions of stars and dust to obtain a new quantitative measure of the effect of dust on the observed photometric parameters in the optical wavebands. Furthermore, because it is no longer necessary to explore a wide range of star/dust geometries, we are also able to systematically explore the full parameter space in opacity, inclination and wavelength, and tabulate the results in a form convenient for the use of the community. We give quantitative measures of the change in the apparent scalelength, central surface brightnesses and inclination of disk galaxies due to the presence of dust.\\footnote{ These corrections are only valid for normal disk galaxies and are not applicable to systems with different star/dust geometries such as starburst or dwarf galaxies.} All these changes are expressed as the ratio of the apparent quantity (i.e. that obtained by fitting images of dusty disks with pure exponential disks) to the intrinsic quantity (i.e. that which would be obtained in the absence of dust). These corrections have been derived from a subset of the simulated images (those in the optical bands) presented in Tuffs et al. (2004; hereafter Paper~III).  In Sect.~2 we give a brief description of the distributions of stars and dust used in the simulated images. In Sect.~3 we specify the fitting procedure used to extract apparent scalelengths, central surface brightnesses and axis ratios from the simulated images. These quantities are tabulated in Sect. 4, where we also describe and explain their dependence on opacity, inclination and wavelength due to the effect of dust. In Sect.~5 we examine the impact of these new results on our ability to derive quantities of astrophysical interest from optical observations and give in Sect.~6 a specific example of the determination of the variation of intrinsic scalelength with wavelength for local universe galaxies. A summary of the paper is given in Sect.~7. ",
        "conclusions": "\\label{conclusion} We have fitted the simulated images of dusty disk galaxies presented in Paper~III with a template brightness distribution corresponding to an inclined infinitely thin dustless exponential disk to obtain apparent scalelengths, central surface-brightness distributions and axis ratios. These are the apparent photometric quantities that an observer would extract from observed images of galaxies. Using the prior knowledge of the corresponding intrinsic quantities input in the simulations, we were able to derive the correction factors (listed in Tables~1--5 of this paper) for the conversion of the apparent to intrinsic quantities due to dust, as a function of inclination $i$, central face-on optical depth in B band  ${\\tau}^{f}_{\\rm B}$, and wavelength.  The apparent to intrinsic scalelength ratio is always greater than unity and can vary up to 50$\\%$. The apparent to intrinsic central surface-brightness ratio expressed in magnitudes can be either positive or negative, depending on whether the dimming due to dust is more or less important than the brightening due to the increase in the column density of stars induced by inclining the disk. This ratio can change up to 1.5 magnitudes due to the effect of dust. The ratio of the apparent to intrinsic axis ratio is not strongly affected by dust, but has an opposite dependence on inclination according to whether the lines of sight through the disk are predominantly optically thin or optically thick. In the former case the vertical distribution of stars makes the galaxy appear progressively rounder with increasing inclination than an infinitely thin disk, whereas in the latter case the opposite trend is seen, because an increasing proportion of the observed light originates from a thin layer of stars above the dust.  Assuming that the basic geometry of dust and stars in local universe spiral galaxies also applies for higher redshift spiral galaxies, our tabulated corrections can be used to correct for the increase in the apparent disk scalelength with increasing redshift due to the expected increase in opacity.  This will allow the intrinsic evolution of disk sizes with cosmological epoch to be investigated. As an example, we show that for a possible variation in opacity between 2 at $z=0$ to 8 or more at $z=1$, the apparent scalelength in the B band  of a galaxy seen at $i=37^{\\circ}$ would increase by a factor of at least 1.24 due to dust.  We used our model to analyse the distribution of observed scalelength ratios between B and I for a sample of disk-dominated spiral galaxies. We found that the predominance of galaxies with larger apparent scalelength in B than in I is primarily due to the effects of dust."
    },
    "0606/astro-ph0606424_arXiv.txt": {
        "abstract": "I review photo-polarimetric and spectropolarimetric observations of V838 Mon, which revealed that it had an asymmetrical inner circumstellar envelope following its 2nd photometric outburst. Electron scattering, modified by pre- or post-scattering H absorption, is the polarizing mechanism in V838 Mon's envelope.  The simplest geometry implied by these observations is that of a spheroidal shell, flattened by at least 10\\% and having a projected position angle on the sky of $\\sim$37$^{\\circ}$.  Analysis of V838 Mon's polarized flux reveals that this electron scattering shell lies interior to the envelope region in which H$\\alpha$ and Ca II triplet emission originates.  To date, none of the theoretical models proposed for V838 Mon have demonstrated that they can reproduce the evolution of V838 Mon's inner circumstellar environment, as probed by spectropolarimetry. ",
        "introduction": "Linear polarimetry can provide powerful diagnostic information regarding the geometry of unresolved astrophysical environments.  Numerous literature resources \\citep{nor92, bjo00} eloquently discuss these diagnostic capabilities; for the purpose of this review I will simply summarize several fundamental principles.  The observed intrinsic polarization of unresolved sources is simply the net integrated polarization of the system.  The density, geometrical distribution, and scattering properties of scatterers in a system are several factors which will influence the strength of the observed intrinsic polarization; systems which either lack an extended envelope of material or are characterized by a symmetrical envelope will exhibit zero net intrinsic linear polarization.  Non-uniform illumination of an extended envelope, by sources such as star-spots and/or binary companions, may also produce a net intrinsic polarization.  Several factors may influence the wavelength dependence of observed intrinsic linear polarization, as discussed by \\citet{nor92} and \\citet{bjo00}.  These factors include: a) the scattering process (i.e. Thompson versus Mie scattering); b) the nature of the illuminating source; c) the dilution of polarized light by the presence of additional unpolarized (i.e. direct) light; and d) the preferential absorption of more scattered light than direct (unpolarized) light.  One is typically is unable to directly measure the intrinsic polarization of astrophysical sources, as the actual observed polarization is comprised of interstellar (time independent) and intrinsic (possibly time dependent) components. Identifying and removing interstellar polarization (ISP) from data is a critical, non-trivial exercise; however, several techniques have proven to be successful in this regard. The field star technique \\citep{mcl79} is one method; successful implementation requires one to identify a suitable number of field stars which are a) intrinsically unpolarized; b) located at a similar distance as the target of interest; and c) located a small angular distance from the target of interest.  If one assumes emission lines in the target of interest are intrinsically unpolarized \\citep{har68}, measuring the polarization in these lines can yield estimates of the ISP, although \\citet{qui97} have shown that this assumption is not always valid.  Finally, the wavelength dependence of ISP is known to follow the empirical Serkowski law \\citep{ser75}; hence, measuring the wavelength dependence of the total observed polarization will, in certain cases, allow one to accurately parameterize the ISP. ",
        "conclusions": "V838 Mon exhibited clear evidence of an intrinsic polarization component beginning at least on 4 February 2002; evidence of this intrinsic component disappeared by 13 February 2002.  The wavelength dependence of this intrinsic polarization suggests that electron scattering, modified by pre- or post-scattering absorption by hydrogen, was the polarizing mechanism.  The presence of an intrinsic component implies that V838 Mon's circumstellar envelope was asymmetrical; one possibly geometry of this envelope is a spheroidal shell flattened by at least 10\\%.  From V838 Mon's polarized flux, we know that this shell is located interior to the region of V838 Mon's envelope responsible for producing emission features at H$\\alpha$ and the Ca II triplet.  Current theoretical efforts to identify the mechanism responsible for V838 Mon's outburst have primarily focussed on explaining its photometric evolution.  To date, none of these models have demonstrated that they can reproduce the evolution of V838 Mon's inner circumstellar environment, as probed by spectropolarimetry."
    },
    "0606/astro-ph0606649_arXiv.txt": {
        "abstract": "We present a survey of the extinction properties of ten lensing galaxies, in the redshift range $z=0.04$--$1.01$, using multiply lensed quasars imaged with the ESO VLT in the optical and near infrared.  The multiple images act as `standard light sources' shining through different parts of the lensing galaxy, allowing for extinction studies by comparison of pairs of images.   We explore the effects of systematics in the extinction curve analysis, including extinction along both lines of sight and microlensing, using theoretical analysis and simulations.  In the sample, we see variation in both the amount and type of extinction.  Of the ten systems, seven are consistent with extinction along at least one line of sight.  The mean differential extinction for the most extinguished image pair for each lens is $\\bar{A}(V) = 0.56\\pm0.04$, using \\cardelli~parametrization.  The corresponding mean $\\bar{R}_V=2.8\\pm0.4$ is consistent with that of the Milky Way at $R_V=3.1$, where $R_V=A(V)/E(B-V)$.  We do not see any strong evidence for evolution of extinction properties with redshift.  Of the ten systems, B1152+199 shows the strongest extinction signal of $A(V)=2.43\\pm0.09$ and is consistent with a \\cardelli~with $R_V=2.1\\pm0.1$.  Given the similar redshift distribution of SN Ia hosts and lensing galaxies, a large space based study of multiply imaged quasars would be a useful complement to future dark energy SN Ia surveys, providing independent constraints on the statistical extinction properties of galaxies up to $z\\sim1$. ",
        "introduction": "The study of extinction curves of galaxies at high redshift has generated a lot of interest in recent years \\citep[see e.g.,][]{riess,  falco1999,goudfrooij,murphy,kann,  goicoechea, york}.  Light reaching us from distant sources is extinguished by dust along its path making it important to correct measurements for the amount and properties of the extinction.  Extragalactic dust extinction can for example affect measurements of Type Ia supernovae (SNe Ia) used to determine various cosmological parameters \\citep[e.g.,][]{riess1998, perlmutter} and the star-formation rates for high redshift starburst galaxies which are used as probes of galaxy evolution \\citep[see e.g.,][]{madau}.  Yet, even if dust properties and thus extinction may vary with redshift and environment,  an average \\cardelli~is often applied when calibrating extragalactic data due to the lack of knowledge of the extinction properties of higher redshift galaxies.  Traditionally, extinction curves are measured by comparing the spectra of two stars of the same spectral type, which have been reddened by different amounts \\citep[see e.g.,][]{massa}.  As it becomes significantly harder to measure spectra of individual stars with distance this method is limited in its application to the Milky Way and the nearest galaxies.  The extinction curves of the Milky Way along different lines of sight have been mapped extensively using this method and have been shown to follow an empirical parametric function which depends only on one parameter, $R_V=A(V)/E(B-V)$, where $A(\\lambda)$ is the total extinction at wavelength $\\lambda$ and $E(B-V)=A(B)-A(V)$ \\citep{cardelli}.  The mean value of $R_V$ in the Milky Way is $3.1$ \\citep{cardelli} but for different lines of sight the value ranges from as low as $R_V\\approx1.8$ toward the Galactic bulge \\citep{udalski} and as high as $R_V\\approx5.6$--$5.8$ \\citep{cardelli,fitzpatrick}.  A lower $R_V$ corresponds to a steeper rise of the extinction curve into the UV, whereas it has little effect on the extinction in the infrared.  Extinction curves have also been obtained for the Small and Large Magellanic Clouds (hereafter, SMC and LMC, respectively) and M31 using this method.  The mean extinction curve of the LMC differs from the \\cardelli~in that the bump at $2175$~$\\AA$ is smaller by a factor of two (as measured by the residual depth of the bump when the continuum has been extracted) and the curve has a steeper rise into the UV for wavelengths shorter than $2200$~$\\AA$ \\citep{nandy81}.  The extinction curve of the SMC is well fitted by  an $A(\\lambda)\\propto\\lambda^{-1}$ curve which deviates significantly from the \\cardelli~and the LMC extinction for $\\lambda^{-1}\\ge4$~$\\mu m^{-1}$ and in particular shows no bump at $2175$ $\\AA$  and a steeper rise into the UV \\citep{prevot84}.  \\citet{bianchi} found that the extinction of M31 follows that of the average \\cardelli.  The various extinction properties shown by these galaxies, especially in the UV and shorter wavelengths, further strengthens the need to find a method to study the extinction curves of more distant galaxies.  A few methods have been proposed for measuring extinction curves for more distant galaxies.  One method is basically an extension of the traditional method of comparing stars of the same spectral type to comparing the SNe Ia \\citep{riess,perlmutter97}.  The extinction is estimated from comparison with unreddened, photometrically similar SNe Ia.  A subset of SNe Ia, with accurately determined extinctions and relative distances, is then used to further determine the relationship between light and color curve shape and luminosity in the full sample.  SN Ia extinction studies usually give lower $R_V$ values than the mean Galactic value of $R_V=3.1$ \\citep{riess1996b, krisciunas, wang2006}.  Quasars with damped Ly$\\alpha$ systems (DLAs) in the foreground have also been studied by \\citet{pei} and were found to be on average redder than those without.   By comparing the optical depths derived from the spectral indices and the ones derived from excess extinction at the location of the \\cardelli~bump they found that their sample of five quasars with DLAs is not consistent with the \\cardelli, marginally compatible with the LMC extinction and fully compatible with SMC extinction. \\citet{murphy} studied a larger sample of the Sloan Digital Sky Survey quasars with damped Ly$\\alpha$ systems in the foreground and found no sign of extinction.  They suggest that the difference between their results and that of \\citet{pei} may be due to the small number statistics in the study by \\citet{pei}.  \\citet{ellison} also found that intervening galaxies cause a minimal reddening of background quasars in agreement with the results of \\citet{murphy} while \\citet{york} found $E(B-V)$ of up to $0.085$ for quasars from the Sloan Digital Sky Survey with Mg II absorption.  In their study \\citet{york} found no evidence of the $2175$~$\\AA$ bump (at variance with \\citet{malhotra}) and found that the extinction curves are similar to SMC extinction.  \\citet{ostman} studied the feasibility of measuring extinction curves by using quasars shining through galaxies.  For the two such systems which survived their cuts, they argued that the extinction curves in the foreground spiral galaxies were consistent with Galactic extinction.  They further suggested a possible evolution in the dust properties with redshift, with higher $z$ giving lower $R_V$ by studying values obtained from the literature in addition to their own.  Extinction curves of high redshift galaxies have also been studied by looking at the spectral energy distribution of gamma-ray bursts (GRBs).  For example, \\citet{palli} fitted a \\cardelli, an SMC and an LMC extinction law to the afterglow of  GRB 030429. The afterglow, at $z=2.66$, was best fit by an SMC like extinction curve with $A(V)=0.34\\pm0.04$. \\citet{kann} studied the extinction of a sample of 19 GRB afterglows and fitted them to various dust extinction models.  They found that the SMC extinction law was preferred by a great majority of their Golden Sample (seven out of eight) while one afterglow was best fit by a \\cardelli~(the other eleven were equally well fit by SMC, LMC and Galactic extinction).  The mean extinction in the $V$-band was $A(V)=0.21\\pm0.04$.  \\citet{goudfrooij} reviewed the dust properties of giant elliptical galaxies and found that they are typically characterized by small $R_V$ if they are in the field or in loose groups, but that if they are in dense groups or clusters their $R_V$ values are close to the mean Galactic $R_V=3.1$.  Early type elliptical galaxies typically have low $A(V)$ \\citep[see e.g.,][who found $A(V)\\lesssim0.35$ for dust lanes in ellipticals]{goudfrooij1994}.  \\citet{nadeau} pointed out that gravitationally lensed quasars could be used to measure the extinction curves of higher redshift galaxies.   \\citet{falco1999}  explored a large sample of 23 lensing galaxies using this method and found that only seven were consistent with no extinction.  This method has also been applied to single systems by e.g. \\citet{jaunsen,toft,motta,wucknitz,munoz,wisotzki,goicoechea} and shows varying extinction properties between different lensing systems.  Here we present a systematic study of the extinction curves of gravitational lenses based on a survey of $10$ lens systems.  We have made a dedicated effort to minimize the number of unknowns and effects that can mimic extinction.  We have broad wavelength coverage in nine different optical and NIR broad bands. An effort was made to minimize the time between the observations for each system in the different bands to minimize the effect of intrinsic quasar variability and microlensing.  All our systems have spectroscopically determined redshifts for both the quasar and the lensing galaxy.  Finally, our systems span the range of $z=0.04$--$1.01$ giving us the possibility to explore possible evolution with redshift.   The outline of this paper is as follows:  In \\S~\\ref{sec:method} we describe the details of the employed method and discuss different sources of systematics and of random errors which may affect our results.  We also present the results of simulations which explore the effects of these errors on data sets similar to those we obtain in the survey.  In \\S~\\ref{sec:data} we present the data and the data reduction of the ESO VLT survey for the 10 lensing systems.  We present the results of our analysis of each individual system in \\S~\\ref{sec:res_ind} and the analysis of the full sample in \\S~\\ref{sec:res_full}.  Finally we summarize our results in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:results}  We open this chapter by presenting in \\S~\\ref{sec:res_ind} the results of our extinction curve analysis for each of the 10 lensing systems. We present the detailed analysis of systems where at least one image pair has a two sigma detection of extinction for one of the three applied extinction laws.  The systems are presented in order of increasing redshift.  We then move on to discussing in \\S~\\ref{sec:res_full} the overall results of our analysis, as well as statistical properties of the full sample.  \\subsection{The individual systems} \\label{sec:res_ind}  \\subsubsection{Q2237+030} Q2237+030 was discovered by \\citet{huchra} and consists of a quasar at redshift $z=1.70$ and a spiral lensing galaxy at $z=0.04$ making it the nearest known lensing galaxy to date.  The system was later resolved into four images forming an Einstein cross with the lensing galaxy in the middle  \\citep{yee88, schneider88}.  \\citet{schneider88} modeled the system in detail and found a predicted time delay of order of one day between the images and amplifications of $4.6,4.5,3.8$ and $3.6$ for images A, B, C and D, respectively.  \\citet{falco96} studied the system with the VLA at radio wavelengths and obtained flux density ratios of $1.00,1.08,0.55$ and $0.77$ for images A, B, C, D, respectively.  Q2237+030 has previously been noted in the literature as having high variability which is uncorrelated between the four images \\citep[see e.g.,][]{irwin, corrigan}.  This system is difficult to interpret as it shows a lot of scatter in the points which cannot be explained by extinction alone nor microlensing (see Figures~\\ref{fig:q2237c} and \\ref{fig:q2237d}).  None of the extinction laws we apply give a good fit to the data for any pair of images (as the redshift of this system is very low, $z=0.04$, we do not expect to see much difference between the quality of the fits for the different extinction laws, see \\S~\\ref{sec:pureext}).  All the image pairs, except C$-$A and D$-$A, yield $A(V)$ consistent with zero.  The data points and fits for C$-$A and D$-$A can be seen in Figures~\\ref{fig:q2237c} and \\ref{fig:q2237d} and the parameters of the fits in Tables~\\ref{tab:q2237ac} and \\ref{tab:q2237ad}.  \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for Q2237+030: C$-$A} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$} & \\colhead{$s$} & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $ 0.53\\pm0.03$ & $>7$  &\\nodata  & \\textit{0.65}& \\nodata & $3.6\\pm0.4$\\\\ \\cardellitab & $0.29\\pm0.06$ & $2.9\\pm1.4$ & \\nodata & $0.85\\pm0.04$ & \\nodata & $3.4\\pm0.4$\\\\ \\cardellitab & $0.27\\pm0.07$ & $2.6\\pm 1.5$ & \\nodata & \\textit{0.65} & $-0.23\\pm0.05$ & $3.4\\pm0.4$\\\\ Power law  & $0.49\\pm0.02 $ & \\nodata & $0.6\\pm0.1$ & \\textit{0.65} & \\nodata & $3.2\\pm0.4$\\\\ Power law  & $0.3\\pm0.5$   & \\nodata & $1\\pm4$ & $0.8\\pm0.5$ & \\nodata & $3.4\\pm0.4$\\\\ Power law  & $0.3\\pm0.2$ & \\nodata & $1\\pm3$ & \\textit{0.65} & $-0.2\\pm0.1$ & $3.4\\pm0.4$\\\\ Linear law& $0.49\\pm0.02$ & \\nodata & \\textit{1.0} & \\textit{0.65} & \\nodata & $3.2\\pm0.4$\\\\ Linear law& $0.35\\pm0.05$ & \\nodata & \\textit{1.0} & $0.78\\pm0.04$ &\\nodata & $3.1\\pm0.4$\\\\ Linear law& $0.34\\pm0.05$ & \\nodata & \\textit{1.0} & \\textit{0.65} & $-0.15\\pm0.05$ & $3.1\\pm0.4$\\\\ \\enddata \\tablecomments{ The extinction of image C compared to image A.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:q2237ac} \\end{deluxetable*}  \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for Q2237+030:  D$-$A} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$} & \\colhead{$s$} & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $ 1.23\\pm0.03$ & $>7$  &\\nodata  & \\textit{0.28}& \\nodata & $7.1\\pm0.4$\\\\ \\cardellitab & $0.35\\pm0.06$ & $3.1\\pm1.5$ & \\nodata & $1.05\\pm0.05$ & \\nodata & $1.9\\pm0.4$\\\\ \\cardellitab & $0.28\\pm0.07$ & $2.1\\pm 1.2$ & \\nodata & \\textit{0.28} & $-0.85\\pm0.05$ & $1.9\\pm0.4$\\\\ Power law  & $1.11\\pm0.03 $ & \\nodata & $0.27\\pm0.04$ & \\textit{0.28} & \\nodata & $1.9\\pm0.4$\\\\ Power law  & $0.4\\pm1.1$   & \\nodata & $0.9\\pm0.4$ & $0.9\\pm1.1$ & \\nodata & $2.0\\pm0.4$\\\\ Power law  & $0.3\\pm0.4$ & \\nodata & $1.3\\pm0.6$ & \\textit{0.28} & $-0.8\\pm0.4$ & $2.0\\pm0.4$\\\\ Linear law& $1.18\\pm0.03$ & \\nodata & \\textit{1.0} & \\textit{0.28} & \\nodata & $5.9\\pm0.4$\\\\ Linear law& $0.41\\pm0.06$ & \\nodata & \\textit{1.0} & $0.97\\pm0.05$ &\\nodata & $1.9\\pm0.4$\\\\ Linear law& $0.39\\pm0.06$ & \\nodata & \\textit{1.0} & \\textit{0.28} & $-0.75\\pm0.05$ & $1.9\\pm0.4$\\\\ \\enddata \\tablecomments{  The extinction of image D compared to image A.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:q2237ad} \\end{deluxetable*}  \\begin{figure} \\epsscale{1.0} \\plotone{f11.eps} \\caption{Q2237+030, C$-$A:  The upper panel shows the data points and the best fit extinction curves.  The lower panel shows the original data points and their shift due to a microlensing signal, along with the best fit through the shifted points.  The parameters of the fits are given in Table~\\ref{tab:q2237ac}.  \\textit{Annotation:}  Filled circles are the original data points with error bars.  The curves correspond to the \\cardelli~(eq. \\ref{eq:car}) with $\\Delta\\hat{m}$ free (dash-dot) and fixed (solid), the power law (eq. \\ref{eq:alpha}) with  $\\Delta\\hat{m}$ free (dash-dot-dot-dot) and fixed (dotted) and the linear law (eq. \\ref{eq:lambda}) with $\\Delta\\hat{m}$ free  (long dash) and fixed (short dash).  Shifted data points due to a microlensing signal are plotted in open boxes (\\cardelli), triangles (power law) and diamonds (linear law).\\label{fig:q2237c}} \\end{figure}  \\begin{figure} \\epsscale{1.0} \\plotone{f12.eps} \\caption{Q2237+030, D$-$A:  The upper panel shows the data points and the best fit extinction curves.  The lower panel shows the original data points and their shift due to a microlensing signal, along with the best fit through the shifted points.  The parameters of the fits are given in Table~\\ref{tab:q2237ad}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview. \\label{fig:q2237d}} \\end{figure}   We use the radio measurements of \\citet{falco96} to fix $\\Delta\\hat{m}$ and in particular for the D$-$A image pair this changes the results significantly.  As the flux density ratios in the radio agree with the model predictions for the D$-$A image pair it is interesting to note that none of the extinction laws we apply give good fits to the data unless we allow for corrections due to achromatic microlensing (see Figures~\\ref{fig:q2237c} and \\ref{fig:q2237d} and Tables~\\ref{tab:q2237ac} and \\ref{tab:q2237ad}).  This might suggest that either D is demagnified by a strong microlensing signal, or that image A is magnified (or both) and the residual `extinction curve' which we are fitting may be effects of chromatic microlensing.  The same effect, but not as strong, is seen in the C$-$A image, again suggesting a slight demagnification of C, a magnification of A or a combination of the two.  In previous microlensing studies, component D has not been seen to have as strong a microlensing signal as A and C have \\citep{irwin, alcalde, gil-merino} so therefore it is perhaps more likely that we are seeing the magnification of image A.  Another explanation for the shift in $\\Delta\\hat{m}$ might be intrinsic variations of the quasar components as discussed in \\S \\ref{sec:lensing}.  Such variability could introduce an overall shift of the data, resulting in inaccurate estimates of $\\Delta\\hat{m}$.  By inspecting the $V$-band lightcurves from \\citet{kochanek2004} we see that component A is indeed in a bright phase at the time of our observations (Julian date of around 2451670).  Component D is fairly stable but component C is getting dimmer climbing down from a peak in its brightness and is still fairly bright, perhaps making the C$-$A shift in $\\Delta\\hat{m}$ less prominent than the D$-$A shift.   \\subsubsection{PG1115+080} PG1115+080 is a  multiply imaged system discovered by \\citet{weynman} as a triply imaged system with the quasar at redshift $z=1.72$.  The A component was later resolved into two separate images, A1 and A2, by \\citet{hege} making the system a quad.  The lensing galaxy was located by \\citet{christian} and is an early type galaxy \\citep{rusin}.  Its redshift and that of three neighboring galaxies were determined to be at $z=0.31$ by \\citet{kundica}.  The time delays between the components were determined by \\citet{schechter}.  \\begin{deluxetable*}{lrrrrr} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for PG1115+080} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $0.12\\pm0.06$ & $3.3\\pm1.9$ & \\nodata & $0.29\\pm0.05$  & $1.2\\pm0.3$\\\\ Power law  & $0.3\\pm2.0$   & \\nodata & $0.5\\pm1.5$ & $0.1\\pm2.0$ & $1.2\\pm0.3$\\\\ Linear law& $0.13\\pm0.03$ & \\nodata & \\textit{1.0} & $0.28\\pm0.04$  & $1.2\\pm0.3$\\\\ \\enddata \\tablecomments{  The extinction of image A2 compared to reference image A1.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:pg1115ab} \\end{deluxetable*}   The system shows very weak differential extinction for all the images (all but A2$-$A1 have $A(V)$ equal to $0$ within two sigma, see Table~\\ref{tab:extinction}).  The data points and fits for A2$-$A1 are shown on Figure  \\ref{fig:pg1115} and the parameters of the fits in Table  \\ref{tab:pg1115ab}.  The low extinction signal is in agreement with the results of \\citet{falco1999}. \\begin{figure} \\epsscale{1.0} \\plotone{f13.eps} \\caption{PG1115+080, A2$-$A1:  Best fit extinction curves for A2$-$A1.  The parameters of the fits can be seen in Table~\\ref{tab:pg1115ab}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.\\label{fig:pg1115}} \\end{figure}  \\subsubsection{B1422+231} B1422+231 is a quadruply imaged system first discovered by \\citet{patnaik} in the JVAS survey and confirmed to be a lensing system by \\citet{lawrence92}.  The lensing system consists of an early type main galaxy \\citep{yee} at $z=0.34$ \\citep{kundic, tonry98} and five nearby galaxies \\citep{remy, bechtold}.  The quasar is at a redshift of $z=3.62$ and the maximum image separation is  $1.3''$  \\citep{patnaik}.  The images show intrinsic variability which has been used to determine the time delay by studying radio light curves \\citep{patnar}.  We only use three of the four images in our analysis as D was too faint to give usable results (all the visible bands gave zero detection).  As for the other components they show very weak differential extinction and give very weak constraints on the differential extinction curves.  All the fits have $A(V)$ consistent with zero (see Table~\\ref{tab:extinction}).  We also get $A(V)$ consistent with zero when we fix $\\Delta\\hat{m}$ (where the values for the $\\Delta\\hat{m}$ are taken to be the average between those deduced by \\citet{patnaik} in the $5~GHz$ and $8~GHz$ bands).  The low differential extinction between the images is in agreement with the results of \\citet{falco1999}.  \\subsubsection{B1152+199} \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for B1152+199} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$} & \\colhead{$s$} & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $ 2.03\\pm0.03$ & $1.61\\pm0.05$ &\\nodata  & \\textit{1.18}& \\nodata & $2.7\\pm0.4$\\\\ \\cardellitab & $2.43\\pm0.09$ & $2.1\\pm0.1$ & \\nodata & $0.85\\pm0.07$ & \\nodata & $2.0\\pm0.4$\\\\ \\cardellitab & $2.41\\pm0.09$ & $2.0\\pm 0.1$ & \\nodata & \\textit{1.18} & $0.32\\pm0.07$ & $2.1\\pm0.4$\\\\ Power law  & $2.01\\pm0.03 $ & \\nodata & $1.98\\pm0.04$ & \\textit{1.18} & \\nodata & $4.0\\pm0.4$\\\\ Power law  & $2.7\\pm0.1$   & \\nodata & $1.45\\pm0.08$ & $0.6\\pm0.1$ & \\nodata & $2.8\\pm0.4$\\\\ Power law  & $2.6\\pm0.1$ & \\nodata & $1.52\\pm0.07$ & \\textit{1.18} & $0.6\\pm0.1$ & $3.2\\pm0.4$\\\\ Linear law& $1.94\\pm0.03$ & \\nodata & \\textit{1.0} & \\textit{1.18} & \\nodata & $10.0\\pm0.4$\\\\ Linear law& $3.57\\pm0.07$ & \\nodata & \\textit{1.0} & $-0.23\\pm0.06$ &\\nodata & $3.4\\pm0.4$\\\\ \\enddata \\tablecomments{  The extinction of image B compared to reference image A.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:b1152} \\end{deluxetable*}  B1152+199 is a doubly imaged system first discovered by \\citet{myers} in the CLASS survey with a background quasar at $z=1.02$, a lensing galaxy at $z=0.44$ and image separation of $1\\farcs56$.  It was observed in radio wavelengths (at frequencies  $1.4,5,8.4$ and $15$~$GHz$) by \\citet{rusin}.  The extinction curve has previously been studied and fitted by a \\cardelli~with $1.3\\leq R_V\\leq2.0$   and $E(B-V)\\sim1$ \\citep{toft} suggesting that it is a heavily extinguished system.  B1152+199 shows a very strong extinction signal as can be seen in Figure~\\ref{fig:b1152} and Table~\\ref{tab:b1152}.  It has the strongest extinction signal of all ten systems with $A(V)=2.43\\pm0.09,2.7\\pm0.1,3.57\\pm0.07$ at $\\chi^2_\\nu=2.0,2.8,3.4$ for the \\cardelli, power law and linear law respectively.  Using the radio measurements of \\citet{rusin} to fix $\\Delta\\hat{m}$ we similarly get $A(V)=2.03\\pm0.03,2.01\\pm0.03,1.94\\pm0.03$ at $\\chi^2_\\nu=2.7,4.0,10.0$.  We also analyze the data with respect to a possible achromatic microlensing signal, keeping $\\Delta\\hat{m}$ fixed.  This yields a non-zero microlensing correction for the \\cardelli~and power law of $s=0.32\\pm0.07,0.6\\pm0.1$ and $A(V)=2.41\\pm0.09,2.6\\pm0.1$ at $\\chi^2_\\nu=2.1,3.2$ respectively.  The best fit for the linear law lies outside the validity of the method with $s>1$ (see \\S~\\ref{sec:micro}) which would correspond to a microlensing signal of $>1$ mag.    \\begin{figure} \\epsscale{1.0} \\plotone{f14.eps} \\caption{B1152+199:  The upper panel shows the data points and the best fit extinction curves to them.  The lower panel shows the original data points and their shift due to a microlensing signal.  The parameters of the fits are given in Table~\\ref{tab:b1152}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview. \\label{fig:b1152}} \\end{figure}  It is clear that in all cases the \\cardelli~provides the best fit to the data suggesting Galactic type dust although the best fit $R_V$ values are lower than those commonly seen in the Milky Way.  It is possible that the measured $R_V$ value is being lowered by a non-zero extinction in the A image provided it has a higher value of $R_V$ (see discussion in \\S~\\ref{sec:extboth}).  However, given the very strong extinction signal this would require very strong extinction along both lines of sight in addition to a strong differential signal.  This is unlikely given the fact that component A is at more than twice the distance from the center of the lensing galaxy than component B with A at $1\\farcs14$ and B at $0\\farcs47$ from the center \\citep{rusin}. Measurements in the X-ray further suggest that the A component is non-extinguished (K. Pedersen et al., 2006, in preparation).  \\subsubsection{Q0142$-$100} \\label{sec:q0142} Q0142$-$100 is a doubly imaged system first discovered by \\citet{surdej} and also known in the literature as UM 673.  The quasar is at a redshift of $z=2.72$ \\citep{macalpine} and the lensing galaxy, which is of early type \\citep{rusin}, is at a redshift of $z=0.49$ \\citep{surdej}.   \\citet{wisotzki} studied this system using spectrophotometric observations and found signs of differential extinction but no microlensing.  The data points and our fits for Q0142$-$100 can be seen in Figure~\\ref{fig:q0142} and the parameters of the fits in Table~\\ref{tab:q0142}.  All the fits give similar $\\chi^2_\\nu$ but the parameters, in particular for the power law, are poorly constrained due to the lack of data points (we did not get any measurements in the infrared for this system).  The extinction is high for an early type galaxy and is not consistent with that found by \\citet{falco1999} who found negligible extinction.  We suspect that the data are being contaminated by the lens galaxy as the B component is located near the galaxy center (at $0\\farcs38$) and the seeing was not optimal for this system (the mean seeing was $0\\farcs87$ compared to $0\\farcs57$ for the full data set). As there are no published radio measurements available we do not have constraints on $\\Delta\\hat{m}$ to analyze the system with respect to a possible microlensing signal.  \\begin{deluxetable*}{lrrrrr} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for Q0142$-$100} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$} &\\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $0.40\\pm0.03$ & $4.8\\pm0.7$ & \\nodata & $1.64\\pm0.04$ & $1.2\\pm0.5$\\\\ Power law  & $4.1\\pm3.8$   & \\nodata & $0.1\\pm0.8$ & $-2.1\\pm3.8$ & $1.1\\pm0.5$\\\\ Linear law& $0.27\\pm0.08$ & \\nodata & \\textit{1.0} & $1.8\\pm0.1$  & $1.1\\pm0.4$\\\\ \\enddata \\tablecomments{  The extinction of image B compared to reference image A.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:q0142} \\end{deluxetable*}  \\begin{figure} \\epsscale{1.0} \\plotone{f15.eps} \\caption{Q0142$-$100:  The plot shows the data points and the best fit to them.  The parameters of the fits can be seen in Table~\\ref{tab:q0142}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview. \\label{fig:q0142}} \\end{figure}  \\subsubsection{B1030+071} B1030+071 is a doubly imaged system first discovered by \\citet{xan1} in the JVAS survey.  They monitored the system in radio wavelengths, finding that the flux density ratios between image A and B range from 12.0 to 18.8 and seem to vary with both time and frequency.  The redshift of the background source was determined to be at $z=1.54$ and the redshift of the lensing object to be at $z=0.60$ \\citep{fassnacht}.  \\citet{falco1999} determined a differential extinction of $E(B-V)=0.02\\pm0.04$ assuming a fixed $R_V=3.1$ \\cardelli.   We were unable to perform an extinction analysis on this system as the deconvolution did not succeed in separating the B component from the main lens galaxy (separated by $0\\farcs11\\pm0\\farcs01$ \\citep{xan1}) making the photometric values unreliable.  For a further study of the extinction of this system higher resolution images would be required.   \\subsubsection{RXJ0911+0551}  \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for RXJ0911+0551: B} \\tablehead{ \\colhead{Image pair} & \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  & \\colhead{$\\chi^2_{\\nu}$}} \\startdata B$-$A & \\cardellitab & $ 0.33\\pm0.06$ & $4.9\\pm0.6$ &\\nodata  & $-0.10\\pm0.05$ &  $1.3\\pm0.4$\\\\ B$-$D  & \\cardellitab & $ 0.20\\pm0.08$ & $4.1\\pm1.1$ &\\nodata  & $-1.00\\pm0.06$ & $1.6\\pm0.5$\\\\ B$-$A & Power law & $0.9\\pm3.7$ & \\nodata & $0.3\\pm0.3$ & $-0.7\\pm3.7$ & $1.3\\pm0.4$ \\\\ B$-$D & Power law & $0.2\\pm1.4$ & \\nodata & $0.8\\pm0.6$ & $-1.1\\pm1.4$ & $1.6\\pm0.5 $\\\\ B$-$A & Linear law & $0.23\\pm0.03$ & \\nodata & \\textit{1.0} & $-0.04\\pm0.04$ & $1.5\\pm0.4$\\\\ B$-$D & Linear law & $0.17\\pm0.03$ & \\nodata & \\textit{1.0} & $-1.00\\pm0.04$ & $1.5\\pm0.4$ \\enddata \\tablecomments{  The extinction of image B compared to reference images A and D.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:rxj0911b} \\end{deluxetable*}   \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for RXJ0911+0551: C} \\tablehead{ \\colhead{Image pair} & \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  & \\colhead{$\\chi^2_{\\nu}$}} \\startdata C$-$A & \\cardellitab & $ 0.20\\pm0.10$ & $3.3\\pm1.5$ &\\nodata  & $0.66\\pm0.07$ &  $1.2\\pm0.4$\\\\ C$-$D  & \\cardellitab & $ 0.09\\pm0.08$ & $2.2\\pm1.5$ &\\nodata  & $0.26\\pm0.06$ & $1.3\\pm0.4$\\\\ C$-$A & Power law & $0.3\\pm2.1$ & \\nodata & $0.9\\pm0.5$ & $0.6\\pm2.1$ & $1.2\\pm0.4$ \\\\ C$-$D & Power law & $0.1\\pm1.8$ & \\nodata & $1.3\\pm1.5$ & $03\\pm1.8$ & $1.2\\pm0.4$\\\\ C$-$A & Linear law & $0.23\\pm0.04$ & \\nodata & \\textit{1.0} & $0.62\\pm0.04$ & $1.1\\pm0.4$\\\\ C$-$D & Linear law & $0.17\\pm0.04$ & \\nodata & \\textit{1.0} & $0.31\\pm0.05$ & $1.2\\pm0.4$ \\enddata \\tablecomments{  The extinction of image C compared to reference images A and D.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:rxj0911c} \\end{deluxetable*}   RXJ0911+0551 is a  multiply imaged system first discovered by \\citet{bade} in the ROSAT All-Sky Survey with the quasar at $z=2.80$.  It was later studied by \\citet{burud} who resolved the system into four images and found that large external shear, possibly due to a cluster, was required to explain the image configuration.  \\citet{kneib} confirmed that the lensing galaxy belongs to a cluster at $z=0.769$.  Observed reddening in at least two (images B and C) of the four images suggest differential extinction by the early type lensing galaxy \\citep{burud}.   \\citet{hjorth} measured the time delay of the system between images A,B,C on the one hand  and D on the other and found the time delay to be $146\\pm8$~days ($2\\sigma$).  We find relatively strong extinction in images B and C compared to images A and D.  Image D also shows some extinction when compared to A but the effect is consistent with zero within two sigmas.  We analyze the extinction curves of B and C compared to A and D.  The data points  and the fits can be seen in Figures~\\ref{fig:rxj0911b} and \\ref{fig:rxj0911c} and the parameters of the fits can be seen in Tables  \\ref{tab:rxj0911b} and  \\ref{tab:rxj0911c}. \\begin{figure} \\epsscale{1.0} \\plotone{f16.eps} \\caption{RXJ0911+0551, B:  The upper panel shows the data points and the fits for the extinction of image B compared to image A.  The lower panel shows the corresponding plot for image B compared to image D.  The parameters of the fits can be seen in Table~\\ref{tab:rxj0911b}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.    \\label{fig:rxj0911b}} \\end{figure}  \\begin{figure} \\epsscale{1.0} \\plotone{f17.eps} \\caption{RXJ0911+0551, C:  The upper panel shows the data points and the fits for the extinction of image C compared to image A.  The lower panel shows the corresponding plot for image C compared to image D.  The parameters of the fits can be seen in Table~\\ref{tab:rxj0911c}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.   \\label{fig:rxj0911c}} \\end{figure}   Assuming image A is completely unextinguished we can estimate the lower limit of the relative extinction of D compared to B and C, $E_D(B-V)/E_{B,C}(B-V)$.  For both B and C, we find that this ratio is around $1/3$ so we expect the extinction curve properties to be affected by both lines of sight (see discussion in \\S~\\ref{sec:extboth}).  That is, we do not expect the extinction curve we get from comparing B and C to D to represent the extinction curve along either line of sight unless their extinction properties are identical.  We see that the $R_V$ value of both B and C are lower when compared to image D than those we get from comparing them to image A suggesting that the extinction properties are indeed different (with image D having a higher $R_V$ value).  We note however that the values of $R_V$ do agree within one sigma for both the differential extinction curves for both B and C.   \\subsubsection{HE0512$-$3329} HE0512$-$3329 is a doubly imaged system first discovered by \\citet{gregg} with an image separation of $0\\farcs644$ and quasar redshift of $z=1.565$.  They estimated a redshift of $z=0.9319$ for the lensing object and found that the lens is most likely a spiral galaxy.  In addition, they estimate the differential reddening assuming negligible microlensing and a standard \\cardelli~with $R_V=3.1$.  This yields $A(V)=0.34$ with the A image being redder than the B image.  \\citet{wucknitz} worked further on disentangling microlensing and differential extinction, and estimated $A(V)= 0.07$ with A being the extinguished image.  This  fit results in an effective $R_V=-2.0$ which can be achieved if the two lines of sight have different $R_V$.  \\begin{deluxetable*}{lrrrrr} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for HE0512$-$3329} \\tablehead{  \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $ 0.14\\pm0.04$ & $1.7\\pm0.4$ &\\nodata  & $-0.67\\pm0.03$ &  $2.1\\pm0.4$\\\\ Power law & $0.23\\pm0.09$ & \\nodata & $1.3\\pm0.3$ & $-0.76\\pm0.09$ & $1.4\\pm0.4$ \\\\ Linear law & $0.35\\pm0.02$ & \\nodata & \\textit{1.0} & $-0.86\\pm0.04$ & $1.4\\pm0.4$ \\enddata \\tablecomments{  The extinction of image A compared to image B.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:he0512} \\end{deluxetable*}   In the case of HE0512$-$3329, it is the brighter image, A, which shows extinction with respect to the B image.  The system is interesting as one of the redshifted data points falls in the range where the $2175$ $\\AA$ bump in the \\cardelli~should lie (see Figure~\\ref{fig:he0512}).  There is however no sign of a bump at $\\lambda=2175$ $\\AA$ and both the power law and the linear extinction law give a much better fit (see Table~\\ref{tab:he0512} for the parameters of the fits).  We redo the fits with no constraints on the $R_V$ values to see if our data could be fit by a negative $R_V$ value but this does not change the result of $R_V=1.7\\pm0.4$.  As there is no radio data available we do not constrain the intrinsic ratio in the fits and we can not constrain the microlensing signal.     \\begin{figure} \\epsscale{1.0} \\plotone{f18.eps} \\caption{HE0512$-$3329:   The plot shows the data points and the best fit to them.  The parameters of the fits can be seen in Table~\\ref{tab:he0512}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.  \\label{fig:he0512}} \\end{figure}   Our results are not in agreement with those of \\citet{wucknitz} who found that fits with the $2175$ $\\AA$ bump better reproduced their data than those without, although the result was not highly significant.  In addition, they found that it is crucial to take microlensing into account when analyzing the extinction curve, which might explain the discrepancy.  However, the detected microlensing signal is only important at wavelengths lower than those we probe, with a small possible effect in the $B$- and $V$-bands.  Therefore, a microlensing signal consistent with the results of \\citet{wucknitz} should not affect our results significantly.  In addition, we note that for their best fitting $R_V$ their fit curves downwards for $\\lambda^{-1}<1$~$\\mu m^{-1}$ which is not consistent with our measurement in the $K$-band (see Figure~\\ref{fig:he0512}).    \\subsubsection{MG0414+0534} MG0414+0534 is a quadruply imaged system first discovered by \\citet{hewitt} with image separation of up to $2^{\\prime\\prime}$.  The quasar, at redshift of $z=2.64$, shows evidence of being heavily reddened by dust in the lensing galaxy \\citep{lawrence95}.  The lens, which has early type spectrum, is at redshift $z=0.9584$ \\citep{tonry} and was modeled by \\citet{falco1997} who found the brightness profile to be well represented by a de Vaucouleurs profile which is characteristic of an elliptical galaxy. \\citet{falco1999} studied the extinction curve of this system and fitted it to a \\cardelli~giving a best fit of $R_V=1.5$, assuming that all lines of sight have the same $R_V$.  \\citet{angonin} studied the origin of the extinction and found, that while the differential extinction is likely due to the lensing galaxy, then there is also evidence for significant reddening which is intrinsic to the source.  \\citet{katz} did an extensive radio survey of the system and found that there was no sign of variability in the radio flux ratios between their measurements and those of \\citet{hewitt} except for the C/B image ratio.  For MG0414+0534, components A1 and A2 show extinction when compared to images B or C with A2 being the more strongly extinguished image (see Table~\\ref{tab:extinction}).  For the A1 image we find different effective extinction laws depending on whether we compare with image B or C (see Table~\\ref{tab:mg0414a1} and Figure~\\ref{fig:mg0414a1}).  In all cases the power law gives the best fit and the linear law the worst (we use the radio measurements of \\citet{katz} to fix $\\Delta\\hat{m}$).  For the \\cardelli~the $R_V$ values do not agree suggesting that perhaps the extinction of images B and C is affecting the differential extinction curve.  We also note though, that we would expect $A(V)$ for A1$-$B to be around 0.3, to be consistent with the other values in Table~\\ref{tab:extinction}, but the best fitting values give a lower value.  We therefore perform another fit where we fix $A(V)=0.3$ in the fits for the A1$-$B pair and this gives $R_V=2.1\\pm0.2;1.8\\pm0.1$ for $\\Delta\\hat{m}$ fixed and free respectively, which are marginally consistent with the results compared to the C image.  However, the $\\chi^2_\\nu=2.6;2.0$ of these fits are significantly worse than those of the original fits.  We do not see any evidence for microlensing except in the case of SMC-like linear extinction which still results in a worse fit than the other two extinction laws.  \\begin{figure} \\epsscale{1.0} \\plotone{f19.eps} \\caption{MG0414+0534, A1:  The upper panel shows the data points and the fits for the extinction of image A1 compared to image B.  The lower panel shows the corresponding plot for image A1 compared to image C.  The parameters of the fits can be seen in Table~\\ref{tab:mg0414a1}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.  \\label{fig:mg0414a1}} \\end{figure}   For image A2 the \\cardelli~gives the best fit when $\\Delta\\hat{m}$ is kept fixed but otherwise the different extinction laws give similar results (see Table~\\ref{tab:mg0414a2} and Figure~\\ref{fig:mg0414a2}).  The parameters of the \\cardelli~are consistent when compared with images B and C suggesting that either A2 dominates the extinction signal or that B and C have similar extinction properties.  There is no evidence for microlensing except in the case of the linear extinction law.  We note that the absolute extinction of image A2, which must be greater or equal to the differential extinction in Table~\\ref{tab:mg0414a2}, is very high given that the lens is an early type galaxy (\\citet{goudfrooij1994} find $A(V)\\lesssim0.35$ for their sample of early type galaxies).  \\begin{figure} \\epsscale{1.0} \\plotone{f20.eps} \\caption{MG0414+0534, A2:  The upper panel shows the data points and the fits for the extinction of image A2 compared to image B.  The lower panel shows the corresponding plot for image A2 compared to image C.  The parameters of the fits can be seen in Table~\\ref{tab:mg0414a2}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview. \\label{fig:mg0414a2}} \\end{figure}  \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{8} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for MG0414+0534: A1} \\tablehead{ \\colhead{Images} & \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  &\\colhead{$s$}& \\colhead{$\\chi^2_{\\nu}$}} \\startdata A1$-$B & \\cardellitab & $ 0.07\\pm0.02$ & $0.4\\pm0.1$ &\\nodata  & $-0.99\\pm0.03$ & \\nodata& $1.5\\pm0.5$\\\\ A1$-$B  & \\cardellitab & $ 0.15\\pm0.03$ & $0.9\\pm0.2$ &\\nodata  & \\textit{-1.07} & \\nodata&$1.6\\pm0.5$\\\\   A1$-$C & \\cardellitab & $ 0.29\\pm0.04$ & $1.5\\pm0.2$ &\\nodata  & $-2.00\\pm0.04$ &  \\nodata&$1.7\\pm0.5$\\\\ A1$-$C  & \\cardellitab & $ 0.27\\pm0.04$ & $1.4\\pm0.2$ &\\nodata  & \\textit{-2.0} & \\nodata&$1.5\\pm0.4$\\\\   A1$-$B & Power law & $0.10\\pm0.05$ & \\nodata & $3.1\\pm0.6$ & $-1.04\\pm0.05$ & \\nodata&$0.9\\pm0.4$ \\\\ A1$-$B & Power law & $0.13\\pm0.03$ & \\nodata & $2.8\\pm0.4$ & \\textit{-1.07} & \\nodata&$1.0\\pm0.4$\\\\  A1$-$C& Power law & $0.15\\pm0.07$ & \\nodata & $2.7\\pm0.6$ & $-1.93\\pm0.05$ & \\nodata&$1.0\\pm0.4$ \\\\ A1$-$C& Power law & $0.24\\pm0.04$ & \\nodata & $2.1\\pm0.3$ & \\textit{-2.0} & \\nodata&$1.1\\pm0.4$\\\\   A1$-$B & Linear law & $0.53\\pm0.04$ & \\nodata & \\textit{1.0} & $-1.37\\pm0.05$ & \\nodata&$2.2\\pm0.5$\\\\ A1$-$B & Linear law & $0.31\\pm0.02$ & \\nodata & \\textit{1.0} & \\textit{-1.07} & \\nodata&$3.4\\pm0.4$\\\\ A1$-$B & Linear law & $0.49\\pm0.03$ & \\nodata & \\textit{1.0} & \\textit{-1.07} & $0.28\\pm0.05$&$2.3\\pm0.4$\\\\  A1$-$C & Linear law & $0.61\\pm0.04$ & \\nodata & \\textit{1.0} & $-2.27\\pm0.05$ & \\nodata&$1.9\\pm0.5$\\\\ A1$-$C & Linear law & $0.42\\pm0.02$ & \\nodata & \\textit{1.0} & \\textit{-2.0} & \\nodata&$2.9\\pm0.4$\\\\ A1$-$C & Linear law & $0.58\\pm0.03$ & \\nodata & \\textit{1.0} & \\textit{-2.0} & $0.26\\pm0.05$&$1.9\\pm0.5$  \\enddata \\tablecomments{  The extinction of image A1 compared to reference images B and C.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:mg0414a1} \\end{deluxetable*}   \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{8} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for MG0414+0534: A2 } \\tablehead{ \\colhead{Images} & \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$}  & \\colhead{$s$}& \\colhead{$\\chi^2_{\\nu}$}} \\startdata  A2$-$B  & \\cardellitab & $ 0.87\\pm0.05$ & $2.7\\pm0.2$ &\\nodata  & $-1.08\\pm0.04$ & \\nodata & $1.8\\pm0.5$\\\\ A2$-$B  & \\cardellitab & $ 0.69\\pm0.03$ & $2.2\\pm0.2$ &\\nodata  & \\textit{-0.93} & \\nodata & $1.8\\pm0.5$\\\\  A2$-$C & \\cardellitab & $ 0.91\\pm0.04$ & $2.6\\pm0.1$ &\\nodata  & $-1.96\\pm0.04$ &  \\nodata & $1.8\\pm0.5$\\\\ A2$-$C  & \\cardellitab & $ 0.81\\pm0.04$ & $2.3\\pm0.2$ &\\nodata  & \\textit{-1.89} & \\nodata & $1.6\\pm0.5$\\\\  A2$-$B & Power law & $1.6\\pm1.2$ & \\nodata & $0.7\\pm0.2$ & $-1.8\\pm1.2$ & \\nodata &$1.7\\pm0.5$ \\\\ A2$-$B & Power law & $0.65\\pm0.03$ & \\nodata & $1.49\\pm0.07$ & \\textit{-0.93} &\\nodata & $2.3\\pm0.5$\\\\  A2$-$C& Power law & $1.3\\pm0.5$ & \\nodata & $0.9\\pm0.2$ & $-2.4\\pm0.5$ & \\nodata & $1.8\\pm0.5$ \\\\ A2$-$C& Power law & $0.75\\pm0.03$ & \\nodata & $1.42\\pm0.07$ & \\textit{-1.89} & \\nodata & $1.9\\pm0.5$\\\\  A2$-$B & Linear law & $1.11\\pm0.04$ & \\nodata & \\textit{1.0} & $-1.34\\pm0.05$ & \\nodata & $1.7\\pm0.5$\\\\ A2$-$B & Linear law & $0.81\\pm0.02$ & \\nodata & \\textit{1.0} & \\textit{-0.93} & \\nodata & $4.0\\pm0.4$\\\\ A2$-$B & Linear law & $1.07\\pm0.04$ & \\nodata & \\textit{1.0} & \\textit{-0.93} & $0.39\\pm0.05$ & $1.6\\pm0.5$\\\\  A2$-$C & Linear law & $1.16\\pm0.04$ & \\nodata & \\textit{1.0} & $-2.25\\pm0.05$ & \\nodata & $1.6\\pm0.4$\\\\ A2$-$C & Linear law & $0.91\\pm0.02$ & \\nodata & \\textit{1.0} & \\textit{-1.89} & \\nodata & $3.5\\pm0.4$\\\\ A2$-$C & Linear law & $1.13\\pm0.04$ & \\nodata & \\textit{1.0} & \\textit{-1.89} & $0.35\\pm0.05$ & $1.5\\pm0.4$ \\enddata \\tablecomments{ The extinction of image A2 compared to reference images B and C.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:mg0414a2} \\end{deluxetable*}   \\begin{deluxetable}{llll} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{MG0414+0534: The extinction properties of A2$-$A1} \\tablehead{ \\colhead{} & \\colhead{$\\Delta\\hat{m}$}  &\\colhead{$A(V)$} & \\colhead{$R_V$}} \\startdata Fit & Free  & $0.61\\pm0.11$ & $3.8\\pm0.7$\\\\ Fit & Fixed & $0.53\\pm0.03$ & $3.5\\pm0.4$\\\\ C   & Free  & $0.62\\pm0.06$ & $4.0\\pm0.9$\\\\ C   & Fixed & $0.54\\pm0.06$ & $3.4\\pm1.0$\\\\ B   & Free  & $0.80\\pm0.05$ & $5.4\\pm2.5$\\\\ B   & Free  & $0.54\\pm0.04$ & $3.7\\pm1.3$ \\enddata \\tablecomments{  The first two lines give the results from the \\cardelli~fit to the data.  The last four lines give the extinction properties calculated from eq. (\\ref{eq:rdiff}) using the properties of A2 and A1 compared to images B and C from Tables~\\ref{tab:mg0414a1} and \\ref{tab:mg0414a2}.} \\label{tab:mg0414_rdiff} \\end{deluxetable}    As the extinction of A1 is significant compared to A2 we expect the extinction properties of both lines of sight to affect the A2$-$A1 extinction curve (see \\S~\\ref{sec:extboth}).  The fit of A2$-$A1 for the Galactic extinction curve gives us $R_V=3.5\\pm0.4,3.8\\pm0.7$ at $A(V)=0.53\\pm0.03,0.61\\pm0.11$ when $\\Delta\\hat{m}$ is kept fixed or free respectively.  If we assume that images B and C have zero extinction we can calculate the effective $R_V$ we expect to get from eq. (\\ref{eq:rdiff}).  The results can be seen in Table~\\ref{tab:mg0414_rdiff} and are in good agreement with the results of the fits.  The extinction of MG0414+0534 is high for an early type galaxy.  We can not exclude the possibility that the extinction may be due to an unknown foreground object and not the lensing galaxy itself.  Finally we note that our estimates of the differential extinction agree with those of \\citet{falco1999} which were obtained by assuming standard Galactic extinction with $R_V=3.1$.    \\subsubsection{MG2016+112} \\begin{deluxetable*}{lrrrrrrr} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Extinction curve fit results for MG2016+112} \\tablehead{ \\colhead{Extinction} & \\colhead{$A(V)$}   & \\colhead{$R_V$}    & \\colhead{$\\alpha$} & \\colhead{$\\Delta\\hat{m}$} & \\colhead{$s$} & \\colhead{$\\chi^2_{\\nu}$}} \\startdata \\cardellitab & $ 0.20\\pm0.03$ & $3.0\\pm0.5$ &\\nodata  & \\textit{-0.092}& \\nodata & $1.7\\pm0.4$\\\\ \\cardellitab & $0.1\\pm0.1$ & $1.8\\pm1.0$ & \\nodata & $-0.01\\pm0.10$ & \\nodata & $1.9\\pm0.5$\\\\ \\cardellitab & $0.11\\pm0.09$ & $1.8\\pm 1.3$ & \\nodata & \\textit{-0.092} & $-0.1\\pm0.1$ & $1.8\\pm0.5$\\\\ Power law  & $0.18\\pm0.03 $ & \\nodata & $1.4\\pm0.2$ & \\textit{-0.092} & \\nodata & $1.4\\pm0.4$\\\\ Power law  & $0.11\\pm0.09$   & \\nodata & $1.8\\pm0.6$ & $-0.02\\pm0.10$ & \\nodata & $1.5\\pm0.5$\\\\ Power law  & $0.12\\pm0.09$ & \\nodata & $1.7\\pm0.5$ & \\textit{-0.092} & $-0.06\\pm0.11$ & $1.5\\pm0.5$\\\\ Linear law& $0.23\\pm0.01$ & \\nodata & \\textit{1.0} & \\textit{-0.092} & \\nodata & $1.6\\pm0.4$\\\\ Linear law& $0.28\\pm0.03$ & \\nodata & \\textit{1.0} & $-0.18\\pm0.06$ &\\nodata & $1.6\\pm0.5$\\\\ Linear law& $0.27\\pm0.03$ & \\nodata & \\textit{1.0} & \\textit{-0.092} & $0.08\\pm0.06$ & $1.6\\pm0.5$\\\\ \\enddata \\tablecomments{  The extinction of image B compared to reference image A.  Numbers quoted in italics were fixed in the fitting procedure.} \\label{tab:mg2016} \\end{deluxetable*}  MG2016+112 was discovered by \\citet{lawrence} and has a giant elliptical lensing galaxy at redshift $z=1.01$ \\citep{schneider85, schneider86}.  The system consists of two images, A and B, of the quasar at redshift $z=3.273$ and an additional image C which may be a third image of the quasar with an additional signal from another galaxy and has been challenging to model \\citep{lawrence, lawrence93,nair}.  The flux of images A and B in the radio at 5 GHz was determined by \\citet{garrett} to be $15.8$~mJy and $17.2$~mJy respectively.  This is the highest redshift system in our sample, and is also interesting since one of the data points lands in the range where the $2175$ $\\AA$ bump in the \\cardelli~should be (see Figure~\\ref{fig:mg2016}).  However, the extinction signal is very weak with $A(V)=0.1\\pm0.1,0.11\\pm0.09,0.28\\pm0.03$ at $\\chi^2_\\nu=1.9,1.5,1.6$ for the \\cardelli, power law and linear law respectively (see Table~\\ref{tab:mg2016} for the parameters of the fits).  When we fix $\\Delta\\hat{m}$ we find somewhat higher extinction of $A(V)=0.20\\pm0.03,0.18\\pm0.03,0.23\\pm0.01$ at $\\chi^2_\\nu=1.7,1.4,1.6$.  In both cases a power law or a linear law is marginally preferred to a \\cardelli.  We also analyze the data with respect to a possible microlensing signal but only find a weak microlensing signal (see Table~\\ref{tab:mg2016} and Figure~\\ref{fig:mg2016}).  Finally we note that our results for the \\cardelli~are consistent with the results of \\citet{falco1999}.    \\begin{figure} \\epsscale{1.0} \\plotone{f21.eps} \\caption{MG2016+112:   The upper panel shows the data points and the best fit extinction curves as given in Table~\\ref{tab:mg2016}.  The lower panel shows the original data points and their shift due to a microlensing signal. The parameters of the fits can be seen in Table~\\ref{tab:mg2016}.  See the caption of Figure \\ref{fig:q2237c} for annotation overview.\\label{fig:mg2016}} \\end{figure}  \\subsection{The full sample} \\label{sec:res_full}  In this section we study the properties of the sample as a whole.  We look for correlations between various parameters and, in particular, search for any dependence on the redshift or the morphology of the galaxies.  Furthermore, we discuss the low $R_V$ values found in SN Ia studies and the possible complementarity of lensing extinction curve studies.  We study, on the one hand, a `golden sample' and, on the other hand, we analyze the full sample.  The `golden sample' is defined to include the image pair with the strongest differential extinction for each lens.  In addition Q0142$-$100 is excluded from the `golden sample' (see \\S \\ref{sec:q0142}).  The `golden sample' therefore consists of eight pairs of images, of which seven have strong enough extinction to analyse the extinction curve.  If not otherwise stated, the results apply to the full sample.  \\subsubsection{$A(V)$ as a function of distance from center of the lensing galaxy} To study the distribution of $A(V)$ as a function of distance from the lens galaxy, we analyze the sample using two methods.  First we plot, in Figure~\\ref{fig:A_dist}, the differential $A(V)$ of the image pairs  (from Table~\\ref{tab:extinction}) as a function of the ratio of the distances from the center of the galaxy (from Table~\\ref{tab:lenses}).  We assign a negative value to the $A(V)$, in those cases where the more distant image is the more strongly extinguished one.  One can see, that when the ratio is small, the image which is nearer the center of the galaxy is the more extinguished one.  However, when the ratio approaches one, the $A(V)$ becomes more evenly scattered around zero.  \\begin{figure} \\epsscale{1.0} \\plotone{f22.eps} \\caption{The differential $A(V)$ for a pair of images~vs.~the ratio of the distances from the center of the lensing galaxy.  The differential $A(V)$ is defined as negative if the image closer to the galaxy is less extinguished.  The figure shows that images closer to the galaxy tend to be the more extinguished but that when the ratio approaches $1$ the scatter increases. \\label{fig:A_dist}} \\end{figure}  For the second method, we assume that the image with the weakest extinction signal is indeed non-extinguished.  We define an absolute $A(V)$ for the other images, by taking the differential extinction compared to this reference image, which we plot as a function of distance from the center of the galaxy, scaled by the lens galaxy scale length\\footnote{The scale length is taken to be the effective radius of a de Vaucouleurs profile fit from \\citet{rusin2003}.} (see  Figure~\\ref{fig:A_scale} and Table~\\ref{tab:lenses}).  From the plot we can see that most $A(V)$ values lie in the range of $0$--$0.5$ for distances smaller than around four scale lengths, but drop for more distant images.  \\begin{figure} \\epsscale{1.0} \\plotone{f23.eps} \\caption{$A(V)$ as a function of the distance of the image relative to the scale radius of the lensing galaxy.  We assign an absolute $A(V)$ to the images by assuming that the least extinguished image for each system is non-extinguished.  The quads are symbolized as circles and the four doubles are symbolized as triangles (up and down facing), a diamond and a box.  The non-extinguished reference images are marked on the plot by open symbols. The error bars for $d_g/d_s$ are smaller than the plotted symbols.  We see that the $A(V)$ values mostly lie in the range of $0-0.5$ for $d_g/d_s \\lesssim 4$ and drop for higher values.\\label{fig:A_scale}} \\end{figure}  Both of these results are consistent with the expectation that the more distant image is on average more likely to pass outside the galaxy and thus not be affected by extinction.  When the distances become similar, secondary effects due to the non-symmetric shape of the lens start becoming important, creating a scatter in the  $A(V)$~vs.~distance plots.  This is in particular the case for the quads where the distances tend to be similar.   \\subsubsection{The different extinction laws} We investigate whether our sample shows a preference for one type of extinction law to another and whether the type of extinction depends on the galaxy type.  We also study the correlation between the parameters of the different fits.  We find that when $\\Delta\\hat{m}$ is allowed to vary, our sample does not show a preference for one extinction law over the other (the mean of the $\\chi^2_\\nu$ is $\\bar{\\chi}^2_\\nu=1.7,1.6,1.8$ for the \\cardelli, power law and linear law) although individual systems can show a strong preference.  If we alternatively look at the fits where $\\Delta\\hat{m}$ was fixed we see that the power law and \\cardelli~are preferred over the linear law in the sample as a whole (with  $\\bar{\\chi}^2_\\nu=2.7,2.1,4.3$) but again individual systems can show different behaviors.  There are three late type galaxies in our sample.  One (HE0512$-$3329) shows a clear preference for an SMC linear law extinction, one (B1152+199) shows a preference for a \\cardelli~and the third (Q2237+030) gives equally good fits to all the extinction laws  (which is expected due to its low redshift, see \\S~\\ref{sec:pureext}).  For the early type galaxies there is also no clear preference for one type of extinction law.  Three systems (PG1115+080, Q0142$-$100, RXJ0911+0551) show no preference for one extinction law over the other, one (MG0414+0534) favors a power law with power index $\\alpha=2-3$ for one of the images (which may be affected by extinction along both lines of sight) and one (MG0216+112) shows a weak preference for a power or a linear law over the \\cardelli.  We therefore conclude that there is no evidence for a correlation between galaxy type and type of extinction in our sample.  When confining the analysis to a Galactic extinction, we find that the mean $R_V$ (for the `golden sample') of the late type galaxies ($\\bar{R}_V^{late}=2.3\\pm0.5$) is marginally lower than that of the early type galaxies ($\\bar{R}_V^{early}=3.2\\pm0.6$), however they are consistent within the error bars and the difference may be due to low number statistics.  This is further discussed in \\S \\ref{sec:lowRv}.  \\begin{figure} \\epsscale{1.0} \\plotone{f24.eps} \\caption{The power index $\\alpha$~vs.~$R_V$.  The data points consist of all image pairs where the extinction curves were analyzed (from fits with $\\Delta\\hat{m}$ free). The plot shows a clear correlation between $\\alpha$ and $R_V$, with lower $\\alpha$ giving higher $R_V$, which is consistent with lower $R_V$ giving a steeper rise into the UV.  The point corresponding to $\\alpha=1.0$ (SMC type extinction) and $R_V=2.9$ \\citep[the mean $R_V$ for the SMC, as determined by][]{pei2}, is marked by a diamond.\\label{fig:alpha_R}} \\end{figure}  We also study the correlation of the parameters $R_V$ and $\\alpha$ for each system to demonstrate the consistency of the two approaches.  As expected, we find that there is a strong correlation with larger $R_V$ giving smaller $\\alpha$, as seen in Figure~\\ref{fig:alpha_R}, consistent with smaller $R_V$ giving steeper rise into the UV.  The exact relationship between $R_V$ and $\\alpha$ can be derived by solving equations (\\ref{eq:car}) and (\\ref{eq:alpha}) and is wavelength dependent.  A first order linear fit to our data gives $\\alpha=(2.5\\pm0.2)-(0.51\\pm0.09)R_V$ which is an applicable approximation within the wavelength range of our data.  In addition we check whether the strength of the extinction is correlated to $R_V$, but we find no evidence for such a correlation (see Figure~\\ref{fig:R_A}).   \\begin{figure} \\epsscale{1.0} \\plotone{f25.eps} \\caption{The figure shows $R_V$~vs.~$A(V)$.  The data points consist of all image pairs where the extinction curves were analyzed (filled circles from fits with $\\Delta\\hat{m}$ free and filled boxes from fits with $\\Delta\\hat{m}$ fixed). The figure shows, as expected, no correlation between $R_V$ and the amount of extinction.\\label{fig:R_A}} \\end{figure}  Finally we study the correlation between the values of $A(V)$ which were found corresponding to the different extinction laws.  The results can be seen in Figure~\\ref{fig:A_A}.  We see that when $\\Delta\\hat{m}$ is free, the power law favors higher $A(V)$ than the other two extinction laws (see left bottom and top panels in Fig.~\\ref{fig:A_A}).  When $\\Delta\\hat{m}$ is fixed, the correspondence becomes much better with the power law giving marginally lower values.  The agreement between the $A(V)$ values derived for the \\cardelli~and the linear law are in general good, regardless of whether we keep $\\Delta\\hat{m}$ fixed or free (with the exception that for $A(V)\\gtrsim1$ the linear law gives higher results when $\\Delta\\hat{m}$ is free).  \\begin{figure} \\epsscale{1.0} \\plotone{f26.eps} \\caption{The graphs show the correlation between the values of $A(V)$ derived for the different fits.  The left column shows the distribution for the fits where $\\Delta\\hat{m}$ is a fitted parameter and the right column gives the corresponding distribution when $\\Delta\\hat{m}$ is fixed.  The dashed line corresponds to $x=y$, which would correspond to perfect agreement in $A(V)$ between the fits, and is plotted for reference. \\label{fig:A_A}} \\end{figure}  Figure~\\ref{fig:A_A} also clearly demonstrates that the $A(V)$ values become much better constrained when $\\Delta\\hat{m}$ is fixed in the fitting.  This suggests that it would be valuable for a future extinction survey, to do a simultaneous radio survey for the systems, in order to constrain the intrinsic magnitude ratio of the images.   \\subsubsection{Evolution with redshift} \\label{sec:redshift} Next, we investigate the behavior of our sample as a function of redshift.  The plots of $R_V$ and $\\alpha$ as a function of redshift can be seen in the upper two panels of Figure~\\ref{fig:all_z}.  We do not see any strong correlation, in either the full nor the `golden sample', although the lower values of $R_V$ and the higher values of $\\alpha$ seem to appear at higher $z$.  Our results do not confirm evolution of $R_V$ with redshift, with lower $R_V$ at higher $z$, as suggested by \\citet{ostman}, but do not exclude such evolution either.  We stress that a larger sample would be needed to make any conclusive claims.  \\begin{figure} \\epsscale{0.9} \\plotone{f27.eps} \\caption{ The top panel shows $R_V$ as a function of $z$.  The middle panels shows $\\alpha$~vs.~$z$.  The bottom panels shows $A(V)$ (as given by a \\cardelli~fit)~vs.~$z$.  On all panels, triangles denote late type galaxies and boxes denote early type galaxies where the values are taken from the fits with $\\Delta\\hat{m}$ kept free.  Circles denote the corresponding fits where  $\\Delta\\hat{m}$ was fixed.  Filled symbols correspond to a `golden' sample as defined in \\S \\ref{sec:res_full}.  We see no strong evolution with $z$ but lower values of $R_V$ seem to appear at higher $z$.  \\label{fig:all_z}} \\end{figure}  We also investigate whether there might be an evolution in the amount of dust extinction with redshift.  Again, we use two samples, the `golden sample' as defined above, and a sample consisting of the highest differential extinction deduced for each image.  The resulting plot can be seen in the bottom panel of Figure~\\ref{fig:all_z} and does not show any correlation between $A(V)$ and $z$.  \\subsubsection{Low $R_V$ values and Type Ia SNe} \\label{sec:lowRv} Recent studies of SNe Ia have suggested that $R_V$  values for SN hosts \\citep[which are mostly late type, see][]{sullivan} could be lower than those for the Milky Way \\citep[see e.g.,][]{riess1996b, krisciunas, wang2006} suggesting that SN hosts are systematically different from the Milky Way.  \\citet{wang} however suggest that the reason for the low values of $R_V$ might be due to circumstellar dust around the SNe themselves.  The presence of such dust would cause inaccurate estimates of the dust extinction of the host galaxies of the SNe Ia.  Lensing studies have also seen more extreme $R_V$ values than those in the Milky Way \\citep[see e.g.,][]{motta, wucknitz} but this has been criticized as possibly being due to extinction along both lines of sight.  However, by choosing systems where the extinction of the measured image dominates the extinction of the reference image, this effect can be avoided (see discussion in \\S~\\ref{sec:extboth}).  In our sample the mean $R_V$ value is $\\bar{R}_V=2.8\\pm0.3$ (with RMS scatter of $1.2$) for the full sample, and $\\bar{R}_V=2.8\\pm0.3$ (with RMS scatter of $1.1$) for the `golden sample'.  These values are marginally lower than, but consistent with, the Milky Way mean value of $R_V=3.1$.  If we look at $R_V$ for the late and early type galaxies separately, we find $\\bar{R}_V^{early}=3.2\\pm0.6$ and $\\bar{R}_V^{late}=2.3\\pm0.5$ for the `golden' sample.  These values are consistent with each other within the quoted error bars, but it is interesting that the late type galaxies have a lower mean $R_V$, in agreement with SN Ia studies.  A larger sample would be needed to determine whether this is a real trend, or due to low number statistics.  The mean extinction in our  `golden sample' is $\\bar{A}(V)=0.56\\pm0.04$ (with RMS scatter of $0.80$).  This gives a lower limit to the mean absolute extinction, as the mean $A(V)$ value is lowered if the reference image is also extinguished.  If we remove the highly extinguished system B1152+199 from our sample, the mean of the `golden sample' reduces to $\\bar{A}(V) = 0.29\\pm0.05$ (with RMS scatter of $0.29$).  If we instead take the highest differential extinction for each image for the full sample into account we get $\\bar{A}(V) = 0.33\\pm0.03$ (with RMS scatter of $0.56$ (or $\\bar{A}(V) = 0.21\\pm0.03$ without B1152+199).  All these values are high enough to cause systematic effects in the calibration of SNe Ia.  A lower mean $R_V$ from a lensing study would strengthen the results of low $R_V$ values from SN studies applying to the interstellar medium.  A lower real $R_V$ value than the assumed one would lead to an overestimation of $A(V)$ given a measurement of $E(B-V)$ (as $A(V)=R_V E(B-V)$).  It is therefore important that the extinction properties of higher redshift galaxies, and SN hosts in particular, be further investigated as assuming a mean Galactic extinction with $R_V=3.1$ in the analysis of SNe Ia could affect the cosmological results.  Lensing galaxies and SN Ia hosts are distributed over a similar redshift range (from $z=0$ to $z\\sim1$) and consist of both early type and late type galaxies.  The majority of lensing galaxies are massive early type galaxies \\citep{kochanek} whereas the SN Ia hosts are mostly late type \\citep{sullivan}.  It would however be possible to select sub-samples of either group which would have the same morphology distribution as the other.  Therefore, studies of the extinction properties of lensing galaxies can complement future dark energy SN Ia surveys, providing an independent measurement of the extinction properties of the SN Ia type hosts.   We have presented an imaging survey of the extinction properties of 10 lensing galaxies using multiply imaged quasars observed with the ESO VLT in the optical and the NIR.  We have made a dedicated effort to reduce the number of unknowns and effects which can mimic extinction.  We have explored, analytically and in simulations, the effects of extinction along both sight lines.  We find that it is not crucial for the reference image to have zero extinction, as long as its extinction is small compared to the other image.  We also study the effects of achromatic microlensing and find that to account for such an effect in photometric data, it is crucial to have constraints on the intrinsic magnitude difference of the images.  We were able to study the extinction of 9 out of 10 of the systems in the survey, the last one had to be discarded due to contamination by the lensing galaxy.  Out of the 9 systems, 8 have a two sigma extinction signal for at least one image pair, which was our limit for doing further extinction curve analysis. However, we suspect that one of those is also contaminated by the lensing galaxy (Q0142$-$100) and exclude it from our  `golden sample'.  The mean extinction for the `golden sample' is $\\bar{A}(V)=0.56\\pm0.09$, using \\cardelli~parametrization, and the mean $R_V$ is $\\bar{R}_V=2.8\\pm0.4$ (compared to $\\bar{R}_V=2.8\\pm0.3$ for the full sample), which is consistent with the mean $R_V=3.1$ found for the Galaxy.  The systems show various extinction properties.  There is no strong evidence for a correlation between morphology and extinction properties.  As our sample covers a broad range in redshifts ($z=0.04$--$1.01$) we have also looked for evolution with redshift.  However, our results neither confirm nor refute evolution of extinction parameters with redshift and we stress that a larger sample would be needed to make any conclusive claims.  Finally we wish to point out that large studies of gravitationally lensed quasars are ideal to study the possible evolution of extinction properties as they are spread over a redshift range from $z=0$ to $z\\approx1$.  Furthermore, the quasars do not affect the environment of the galaxy we wish to study as is the case in SN Ia studies.  For further improvements, however, higher resolution and deeper imaging for a larger sample would be required making a dedicated study with space based telescopes of considerable interest.  A simultaneous radio survey, in order to constrain the intrinsic ratio of the images, would also further improve the results.  Such a study could complement future dark energy SN Ia surveys, providing an independent measurement for the extinction properties of SN Ia type host galaxies."
    },
    "0606/astro-ph0606680_arXiv.txt": {
        "abstract": "In the present paper we develop an algorithm allowing to calculate line-of-sight velocity dispersions in an axisymmetric galaxy outside of the galactic plane. When constructing a self-consistent model, we take into account the galactic surface brightness distribution, stellar rotation curve and velocity dispersions. We assume that the velocity dispersion ellipsoid is triaxial and lies under a certain angle with respect to the galactic plane. This algorithm is applied to a Sa galaxy NGC~4594~= M~104, for which there exist velocity dispersion measurements outside of the galactic major axis. The mass distribution model is constructed in two stages. In the first stage we construct a luminosity distribution model, where only galactic surface brightness distribution is taken into account. Here we assume the galaxy to consist of the nucleus, the bulge, the disc and the stellar metal-poor halo and determine structure parameters of these components. Thereafter, in the second stage we develop on the basis of the Jeans equations a detailed mass distribution model and calculate line-of-sight velocity dispersions and the stellar rotation curve. Here a dark matter halo is added to visible components. Calculated dispersions are compared with observations along different slit positions perpendicular and parallel to the galactic major axis. In the best-fitting model velocity dispersion ellipsoids are radially elongated with $\\sigma_{\\theta}/\\sigma_R \\simeq 0.9-0.4$, $\\sigma_z/\\sigma_R \\simeq 0.7-0.4$, and lie under the angles $\\le 30\\degr$ with respect to the galactic equatorial plane. Outside the galactic plane velocity dispersion behaviour is more sensitive to the dark matter density distribution and allows to estimate dark halo parameters. For visible matter the total $M/L_B = 4.5 \\pm 1.2$, $M/L_R = 3.1 \\pm 0.7$. The central density of the dark matter halo is $\\rho_{\\rm DM}(0) = 0.033~\\rmn{M_{\\sun}pc^{-3}}$. ",
        "introduction": "The study of the dark matter (DM) halo density distribution allows us to constrain possible galaxy formation models and large scale structure formation scenarios \\citep*{b66,b51,b41}. For this kind of analysis, it is necessary to know both the distribution of visible and dark matter. Without additional assumptions rotation curve data alone are not sufficient to discriminate between these two kinds of matter \\citep{b29}. It does not suffice either to use additionally velocity dispersions along the major axis.  Realistic mass and light distribution models must be consistent, i.e. the same model must describe the luminosity distribution and kinematics. Three main classes of self-consistent mass distribution models can be discriminated: the Jeans equations based models, the specific phase space density distribution models and the Schwarzschild orbit superposition based models.  Mass distribution models based on solving the Jeans equations have an advantage that the equations contain explicitly observed functions -- velocity dispersions. On the other hand, there are three equations, but at least five unknown functions (three dispersion components, centroid velocity and the velocity dispersion ellipsoid orientation parameter) and thus the system of equations is not closed. In addition, the use of the Jeans equations neglects possible deviations of velocity distributions from Gaussians and does not garantee that the derived dynamical model has non-negative phase density distribution everywhere. However, within certain approximations the Jeans equations are widely used for the construction of mass distribution models. In the case of spherical systems with biaxial velocity dispersion ellipsoids, such models have been constructed, for example, by \\citet{b6}, \\citet{b64}, \\citet{b42}, \\citet{b82}. In the case of flattened systems with biaxial velocity dispersion ellipsoids, a general algorithm for the solution of the Jeans equations was developed by \\citet*{b8}, \\citet{b19}. Another algorithm in the context of the multi-Gaussian expansion (MGE) formalism was developed by \\citet{b34}. An approximation for cool stellar discs (random motions are small when compared with rotation) has been developed by \\citet{b1}.  Dynamical models with a specific phase density distribution have the advantage that velocity dispersion anisotropy can be calculated directly. On the other hand, due to rather complicated analytical calculations, only rather limited classes of distribution functions can be studied. Spherical models of this kind have been constructed by \\citet*{b17}, \\citet{b5}. In the case of an axisymmetric density distribution, velocity dispersion profiles have been calculated for certain specific mass and phase density distribution forms by \\citet*{b84}, \\citet{b37}, \\citet{b25}, \\citet*{b24}, \\citet*{b28}, \\citet{b65}, \\citet{b2} and others. A special case is an analytical solution with three integrals of motion for some specific potentials: an axisymmetric model with a potential in the St\\\"ackel form \\citep{b27}, isochrone potential \\citep{b26}.  Probably the most complete class of dynamical models has been developed on the basis of the Schwarzschild linear programming method \\citep{b72}. Thus, it is not surprising that just for this method most significant developments occured in last decade. \\citet{b68} and \\citet{b22} have developed this method in order to calculate line-of-sight velocity profiles. Thereafter, \\citet{b16} and \\citet{b88} generalized it for an arbitrary density distribution linking it with MGE method. A modification of the least-square algorithm was done by \\citet{b55}. Interesting comparisons of the results of the Schwarzschild method with phase density calculations within a two-integral approximation have been made by \\citet{b86} and \\citet{b55}.  At present, nearly all dynamical models have been applied for one-component systems. However, the structure of real galaxies is rather complicated -- galaxies consist of several stellar populations with different density distribution and different ellipticities. In addition, in different components velocity dispersions or rotation may dominate.  In our earlier multi-componental models \\citep*[see][]{b79,b80,b32} we approximated flat components with pure rotation models and spheroidal components with dispersion dominating kinematics. For spheroidal components mean velocity dispersions were calculated only on the basis of virial theorem for multi-componental systems. These models fit central velocity dispersions, gas rotation velocities and light distribution with self-consistent models.  In the present paper, we construct a more sophisticated self-consistent mass and light distribution model. We decided to base it on the Jeans equations. For all visible components, both rotation and velocity dispersions are taken into account. The velocity dispersion ellipsoid is assumed to be triaxial and line-of-sight velocity dispersions are calculated. Mass distribution of a galaxy is axisymmetric and inclination of the galactic plane with respect to the plane of the sky is arbitrary.  In order to discriminate between DM and visible matter, it is most complicated to determine the contribution of the stellar disc to the galactic mass distribution. Quite often the maximum disc approximation is used. In the present paper, we attempt to decrease degeneracy, comparing calculated models with the observed stellar rotation curve, velocity dispersions along the major axis and in addition, along several cuts parallel to the major and minor axis. In the case of two-integral models for edge-on galaxies this allowed to constrain possible dynamical models \\citep{b63}.  First measurements of velocity dispersions along several slit positions were made by \\citet{b53} and \\citet{b46}. Later, similar measurements were performed by \\citet{b8}, \\citet*{b38}, \\citet{b77}, \\citet{b16}. In recent years, with the help of integral field spectroscopy, complete 3D velocity and dispersion fields have been measured already for several tens of galaxies.  We apply our model for the nearby spiral Sa galaxy M~104 (NGC~4594, the Sombrero galaxy). This galaxy is suitable for model testing, being a disc galaxy with a significant spheroidal component. The galaxy has a detailed surface brightness distribution and a well-determined stellar rotation curve. M~104 has a significant globular cluster (GC) subsystem. And as most important in our case, the line-of-sight velocity dispersion has been measured along the slit at different positions parallel and perpendicular to the projected major axis.  We construct the model in two stages. First, a surface brightness distribution model is calculated. Here we distinguish stellar populations and calculate their structure parameters with the exception of masses. In the second stage, we calculate line-of-sight velocity dispersions and the stellar rotational curve and derive a mass distribution model.  Sections~2 and 3 describe the observational data used in the modelling process and construction of the preliminary model. In Section~4 we present the line-of-sight dispersion modelling process. Section~5 is devoted to the final M~104 modelling process. In Section~6 the discussion of the model is presented.  Throughout this paper all luminosities and colour indices have been corrected for absorption in our Galaxy according to \\citet*{b71}. The distance to M~104 has been taken 9.1~Mpc, corresponding to the scale 1~arcsec~= 0.044~kpc \\citep*{b39,b61,b81}. The angle of inclination has been taken 84\\degr. ",
        "conclusions": "In the present paper we developed an algorithm, allowing to construct a self-consistent mass and light distribution model and to calculate projected line-of-sight velocity dispersions outside galactic plane. We assume velocity dispersion ellipsoids to be triaxial and thus the phase density is a function of three integrals of motion. The galactic plane may have an arbitrary angle with respect to the plane of the sky. The developed algorithm is applied to construct a mass and light distribution model of the Sa galaxy M~104. In the first stage a luminosity distribution model was constructed on the basis of the surface brightness distribution. The inclination angle of the galaxy is known and the spatial luminosity distribution can be calculated directly with deprojection. Using the surface brightness distribution in $BVRI$ colours and along the major and minor axis, we assume that our components represent real stellar populations and determine their main structure parameters. In the second stage, the Jeans equations are solved and the line-of-sight velocity dispersions and the stellar rotation curve are calculated. Observations of velocity dispersions outside the apparent galactic major axis allow to determine the velocity ellipsoid orientation, anisotropy and to constrain DM halo parameters.  The total luminosity of the galaxy M~104 resulting from the best-fitting model is $L_B = (5.1\\pm 0.6) \\cdot 10^{10} \\rmn{L_{\\sun}}$, $L_R = (7.4\\pm 0.7) \\cdot 10^{10} \\rmn{L_{\\sun}}$. The total mass of the visible matter is $M_{\\rmn{vis}} = (22.9\\pm 3.2) \\cdot 10^{10} \\rmn{M_{\\sun}}$, giving the mean mass-to-light ratio of the visible matter $M/L_B = 4.5 \\pm 1.2~ \\rmn{M_{\\sun}L_{\\sun}^{-1}}$, $M/L_R = 3.1 \\pm 0.7~ \\rmn{M_{\\sun}L_{\\sun}^{-1}}$. The surface brightness distributions in $V$ and $I$ have not sufficient extent to determine the luminosities of the stellar halo and we do not either give galactic total luminosities in these colours and corresponding $M/L$ ratios. Calculated from the model, the $L_B$ coincides well with the total absolute magnitude $M_B = -21.3$ ($=5.2 \\cdot 10^{10}~\\rmn{L_\\odot}$) obtained by \\citet{b39}.  In our model, the mass of the disc is $M_{\\rm disc} = 12\\cdot10^{10}\\rmn{M_{\\sun}}$. This coincides rather well with the disc mass $11.4\\cdot 10^{10} \\rmn{M_{\\sun}}$  calculated with the help of Toomre's stability criterion by \\citet{b83} and with the mass $9.6\\cdot 10^{10} \\rmn{M_{\\sun}}$ derived by \\citet{b34}. On the other hand, \\citet{b34} derived for the bulge mass $> 5\\cdot 10^{11} \\rmn{M_{\\sun}}$, giving $M_\\rmn{disc}/M_\\rmn{bulge} =0.2$. This is similar to the value 0.25 derived by \\citet{b47}, but this is much less than $M_\\rmn{disc}/M_\\rmn{spher} = 1.1$ resulting from our model. An explanation may be that in the models by \\citet{b34} and \\citet{b47} no DM halo was included and hence the extended bulge mass is higher.  In our model, the disc is rather thick ($q= 0.25$). However, disc thickness can be easily reduced to $q=0.15-0.2$ when taking the galactic inclination angle instead of $\\delta = 84\\degr$ to be 83--82$\\degr$. Other parameters remain nearly unchanged. At present this was not done.  Derived in the present model bulge parameters can be used to compare them with the results of chemical evolution models. Our model gives $M/L_V = 7.1 \\pm 1.4~\\rmn{M_{\\sun}L_{\\sun}^{-1}}$ and $(B-V)=1.06$ for the bulge. Comparing spectral line intensities with chemical evolution models, \\citet{b87} obtained for the bulge region the metallicity $Z=0.03$ and the age 11~Gyrs. According to \\citet{b13}, these parameters give $M/L_V = 7-8~\\rmn{M_{\\sun}L_{\\sun}^{-1}}$ and $(B-V) = 1.06-1.08$ for simple stellar population (SSP) models. Bulge parameters from our dynamical model agree well with these values and suggest that our model is realistic.  In our calculations, we corrected luminosities from the absorption in the Milky Way only and did not take into account the inner absorption in M~104. According to \\citet{b35}, absorption in the centre may be at least $A_V\\sim 0.13$~mag and thus $M/L_V = 6.3$ for the bulge. This is slightly too small when compared with the \\citet{b13} SSP models. However, decreasing the bulge age to 10.5~Gyrs allows to fit the results.  Rather sophisticated models of M~104 have been constructed by \\citet{b34,b35} and \\citet{b33}. Due to our different approaches, it is difficult to compare our components and their parameters with those of \\citet{b34,b35}. On the basis of the data used by us, we had no reason to add an additional inner disc or a bar to the bulge region. However, we did not analyse $I$ and $H$ colours and ionized gas kinematics in inner regions as it was done by \\citet{b33}. Modelling of gas kinematics in central regions is beyond the scope of the present paper as gas is not collision-free.  On the basis of velocity dispersion observations only along the major axis it is difficult to decide about the presence of the DM even when dispersions extend up to 2--3 $\\rmn{R_e}$ \\citep{b70}. In the case of M~104, additional dispersion measurements can be used. Velocity dispersions in the case of the slit positioned parallel and perpendicular to the galactic major axis, have been measured by \\citet{b53}. The calculated mass distribution model describes rather well the observed stellar rotation curve and line-of-sight velocity dispersions. Only the two last measured points at a cut 50~arcsec perpendicular to the major axis deviate rather significantly when compared to the model. On the other hand, in addition to stellar velocity dispersion measurements, the mean line-of-sight velocity dispersion of the GC subsystem $\\sigma = 255~\\rmn{km~ s^{-1}}$ was measured by \\citet{b12}. This corresponds to GCs at average distances 5--10 kpc from the galactic center and is in rather good agreement with the dispersions calculated from the model.  In the best-fitting model the DM halo harmonic mean radius $a_0= 40$~kpc and $M=1.8\\cdot 10^{12} {\\rm M_{\\sun}}$ giving slightly falling rotation curve in outer parts of the galaxy (Fig.~\\ref{gasrot}). The central density of the DM halo in our model is $\\rho (0) = 0.033~\\rmn{M_{\\sun} pc^{-3}}$, being also slightly less than it was derived for distant ($z\\sim ~0.9$) galaxies \\citep[$\\rho (0) = 0.012-0.028~\\rmn{M_{\\sun} pc^{-3}}$,][]{b78}. On the other hand, the result fits with the limits derived by \\citet{b9} for local galaxies $\\rho (0) = 0.015-0.050~\\rmn{M_{\\sun} pc^{-3}}$.  An essential parameter in mass distribution determination is the inclination of the velocity dispersion ellipsoid with respect to the galactic plane \\citep[see e.g.][]{b56,b63}. Velocity dispersion ellipsoid inclinations calculated in the present paper are moderate, being $\\le 30\\degr$. In a sense, our approach to the third integral of stellar motion is similar to that by \\citet{b50} -- the local St\\\"ackel fit. In their modelling of the local Milky Way structure, they derived that at $0<z<600$ pc and $6.8<R<8.8$ kpc, the inclination of the velocity dispersion ellipsoid is less than $z/R$, and they studied the corresponding correction in detail. In our model, in the same distance regions (although the Milky Way and M~104 are not very similar objects), inclination correction values are slightly smaller. Variations of the corrections with $R$ and $z$ are qualitatively similar. The largest difference is the variation of the correction value with $z$, for which \\citet{b50} obtained an increase by 0.1, when moving from $z=0$ to $z= 0.6$~kpc, but in our model, corresponding increase was only by 0.01. In our model, a significant increase of the ellipsoid inclination angle begins at larger $z$, which can be explained by higher thickness of the disc component of M~104 ($q=0.25$). In a following paper we intend to construct similar models for other galaxies with velocity dispersion measurements outside the galactic major axis. This may lead to more firm conclusions about the inclination of velocity dispersion ellipsoids outside the galactic plane.  Ratios of the line-of-sight velocity dispersions are given in Fig.~\\ref{aniso}. It is seen that velocity dispersion ellipsoids are quite elongated -- anisotropies in the symmetry plane at outer parts of the galaxy are $<0.5$. \\begin{figure} \\includegraphics[width=84mm]{figure13.eps} \\caption{Dispersion ratios in galactic meridional plane. Coordinate $z$ is in kpc.}\\label{aniso} \\end{figure}  Modelling the disc Sb galaxy NGC~288 within a constant velocity ellipsoid inclination approximation, \\citet*{b43} estimated that the dispersion ratio $\\sigma_z/\\sigma_R = 0.70$. Within epicycle approximation \\citet{b89} derived for Sb galaxy NGC~3982 the dispersion anisotropy $\\sigma_z/ \\sigma_R = 0.73$. These two galaxies are morphologically close to the Sa galaxy modelled in the present paper, and it is seen that dispersion ratios are more anisotropic in our case. In general, a radially elongated dispersion ellipsoid is rather common \\citep*{b75}. On the other hand, there are exceptions -- galaxy NGC~3949 (\\citet{b89} has the dispersion ratio $\\sigma_z/\\sigma_R = 1.18$).  By using the quadratic programming method \\citep{b23}, the distribution function within the three-integral approximation has been numerically calculated for the S0 galaxy NGC~3115 by \\citet*{b36}. Assuming some similarity between S0 and Sa galaxies, it is interesting to compare the derived velocity dispersion behaviour outside the galactic plane. Although detailed comparison is difficult a similar structure of isocurves is seen. But the results disagree in values of $\\sigma_z/ \\sigma_R$. Our ratios are radially elongated, the ratios by \\citet{b36} are vertically elongated. Qualitatively this is in agreement with the general trend that galaxies of earlier morphological type have larger $\\sigma_z/ \\sigma_R$ ratio \\citep{b75}. At larger $R$, the dispersion ratio $\\sigma_z/ \\sigma_R$ decreases. This is in agreement with the decrease of dispersion ratios due to the decrease of the role of interactions with molecular clouds at greater galactocentric distances \\citep[see][]{b48}. The use of this explanation in the case of gas-poor S0 galaxies is not clear. At greater distances $z$ both our and \\citet{b36} dispersion ellipsoids become more spherical.  By using the Schwarzschild method, dispersion ratios for E5--6 galaxy NGC~3377 have been calculated by \\citet*{b20}. Taking into account relation between spherical and cylindrical coordinates, the behaviour of the dispersion ratios as a function of $R$ and $z$ near the galactic plane is in approximate accordance (dispersion ratios by \\citet{b20} are slightly more spherical). Rotational properties of elliptical and Sa galaxies are too different to compare the ratios $\\sigma_{\\theta}/ \\sigma_R$. Unfortunately, it is not possible to compare also orientations of velocity dispersion ellipsoids."
    },
    "0606/astro-ph0606363_arXiv.txt": {
        "abstract": "{\\it XMM-Newton} and {\\it Chandra} have ushered in a new era for the study of dwarf galaxies in the Local Group.  We provide an overview of the opportunities, challenges, and some early results. The large number of background sources relative to galaxy sources is a major theme. Despite this challenge, the identification of counterparts has been possible, providing hints that the same mechanisms producing X-ray sources in larger galaxies are active in dwarf galaxies. A supersoft X-ray source within $2''$ of the supermassive black hole in M32 may be a remnant of the tidal disruption of a giant, although other explanations cannot be ruled out. ",
        "introduction": "Dwarf galaxies constitute the largest number of galaxies residing in groups and clusters. Dwarf galaxies also inhabit the field. They are the building blocks of galaxy formation. Therefore, in the young universe and throughout vast volumes of the present-day universe, dwarf galaxies provide the typical environments for the formation and evolution of X-ray sources (XRSs).  A vast array of multiwavelength observations of Local Group dwarf galaxies allow the detailed star formation histories of individual galaxies to be derived. This is important, because theoretical considerations and observations of larger, more distant galaxies suggest that there are correlations between star formation history and the formation of XRSs (see, e.g., Gilfanov 2004, Gilfanov, Grimm, \\& Sunyaev 2004, and references therein). There is, e.g., a strong correlation between star formation rates and the formation of high-mass X-ray binaries (HMXBs) and supernova remnants (SNRs). There is also a correlation between low-mass X-ray binaries (LMXBs) and the total galaxy mass; the LMXB population in a large galaxy provides a measure of the average star formation history. Detailed information about individual dwarf galaxies can therefore help to both test and refine theory. The Local Group dwarf galaxies in effect form a set of individual laboratories in which the foundations of X-ray astronomy can be studied.   We can best study X-rays from those dwarf galaxies in the Local Group, where our observations can discover sources of relatively low luminosity, and where we can determine source positions well enough to identify counterparts at other wavelengths. {\\it Chandra} and {\\it XMM-Newton} provide new tools. Each can resolve individual systems in the dwarf galaxies of the Local Group, where the X-ray populations are generally not spatially concentrated.  The large effective area of {\\it XMM-Newton} collects enough photons to study many spectral and timing characteristics without prohibitively long exposure times. The superb angular resolution of {\\it Chandra} allows more reliable detection of counterparts, especially in combination with high resolution optical or radio observations; in some cases, structures in extended sources, such as SNRs, may be resolved.    \\vspace{-.3 true in} ",
        "conclusions": ""
    },
    "0606/astro-ph0606155_arXiv.txt": {
        "abstract": "{The M16 nebula is a relatively nearby \\hii ~region, powered by O stars from the open cluster NGC~6611, which borders to a Giant Molecular Cloud. Radiation from these hot stars has sculpted columns of dense obscuring material on a few arcmin scales. The interface between these pillars and the hot ionised medium provides a textbook example of a Photodissociation Region (PDR).} {To constrain the physical conditions of the atomic and molecular material with submillimeter spectroscopic observations.} {We used the APEX submillimeter telescope to map a $\\sim 3'\\times3'$ region in the CO $J=$ 3--2, 4--3 and 7--6 rotational lines, and a subregion in atomic carbon lines. We also observed C$^{18}$O(3--2) and CO(7--6) with longer integrations on five peaks found in the CO(3--2) map. The large scale structure of the pillars is derived from the molecular lines' emission distribution. We estimate the magnitude of the velocity gradient at the tips of the pillars and use LVG modelling to constrain their densities and temperatures. Excitation temperatures and carbon column densities are derived from the atomic carbon lines.} {The atomic carbon lines are optically thin and excitation temperatures are of order 60~K to 100~K, well consistent with observations of other \\hii ~region-molecular cloud interfaces. We derive somewhat lower temperatures from the CO line ratios, of order 40~K. The {\\sc Ci}/CO ratio is around 0.1 at the fingers tips.}{} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606225_arXiv.txt": {
        "abstract": "{} {Based on our well-determined sample of open clusters in the all-sky catalogue \\ascc we derive new linear sizes of some 600 clusters, and investigate the effect of mass segregation of stars in open clusters.} {Using statistical methods, we study the distribution of linear sizes as a function of spatial position and cluster age. We also examine statistically the distribution of stars of different masses within clusters as a function of the cluster age.} {No significant dependence of the cluster size on location in the Galaxy is detected for younger clusters ($<$ 200~Myr), whereas older clusters inside the solar orbit turned out to be, on average, smaller than outside. Also, small old clusters are preferentially found close to the Galactic plane, whereas larger ones more frequently live farther away from the plane and at larger Galactocentric distances. For clusters with $(V - M_V) < 10.5$, a clear dependence of the apparent radius on age has been detected: the cluster radii decrease by a factor of about 2 from an age of 10 Myr to an age of 1 Gyr. A detailed analysis shows that this observed effect can be explained by mass segregation and does not necessarily reflect a real decrease of cluster radii. We found evidence for the latter for the majority of clusters older than 30 Myr. Among the youngest clusters (between 5 and 30 Myr), there are some clusters with a significant grade of mass segregation, whereas some others show no segregation at all. At a cluster age between 50 and 100 Myrs, the distribution of stars of different masses becomes more regular over cluster area. In older clusters the evolution of the massive stars is the most prominent effect we observe.} {} ",
        "introduction": "Open star clusters are gravitationally bound systems of, typically, several hundreds of stars formed together. Primordial conditions during the cluster formation and the location of the parental molecular cloud in the Galaxy play an important role in the fate of a cluster. On the other hand, the stellar content of a cluster evolves with time, and internal and external interactions affect the properties of individual cluster members as well as of the whole cluster as a system. Therefore, the spatial structure and mass distribution that we observe today in a given cluster is the result of the original brand marks and the ongoing evolution.  Numerical simulations of the dynamical evolution predict a mass segregation in open clusters i.e., a different concentration of cluster members with different masses with respect to the cluster centre. This process occurs on approximately the relaxation time-scale and is independent of most of the possible initial conditions (Bonnell \\& Davies~\\cite{boda}). During the dynamical evolution of a cluster, more massive members sink to the centre, whereas less massive stars tend to show a diffuse distribution (Portegies Zwart \\& McMillan~\\cite{pzm}). A relaxed cluster can be thought of as a set of nested spherical clouds of stars of different mass (e.g., see Adams et al.~\\cite{adea01} for illustration). With increasing mass of the stars, the radial density profile becomes steeper and narrower. Due to tidal interactions with the Galactic gravitational field, the cluster can lose stars once they overflow its tidal radius. Due to mass segregation a cluster loses preferentially low mass stars from the cluster corona, which evaporate into the field, up to an entire dissolution of the cluster in the Galaxy (Andersen \\& Nordstr\\\"{o}m~\\cite{andea00}). A sudden mass loss or close passage of giant molecular clouds can considerably disturb a relaxation process and reduce the life-time of open clusters (Kroupa et al.~\\cite{krea01}, Bergond et al.~\\cite{bea01}).  Predicted by simulations (Spitzer \\& Shull~\\cite{spish}), mass segregation was already found in many open clusters. The most reliable results on mass segregation can be expected for nearby clusters like the Pleiades and Praesepe, where cluster members can be observed over a wide range of magnitudes and masses. Compared to the distribution of more massive stars, indications for a flatter density profile of cluster members with $m<1\\, m_{\\odot}$ were obtained in these clusters by Jones \\& Stauffer~(\\cite{jost91}), and more recently by Adams et al.~(\\cite{adea01}, \\cite{adea02}) who used data from 2MASS and USNO-A. Raboud \\& Mermilliod~(\\cite{ramer98a}, \\cite{ramer98b}) found evidence for a continuous flattening of density profiles with decreasing mass of cluster members in the Pleiades and Praesepe, and in a much more distant open cluster, NGC 6231, too. Similar effects were detected by Sagar et al.~(\\cite{sagar88}), who considered 11 distant clusters in the Galactic disk and by de Grijs et al.~(\\cite{deg02a}, \\cite{deg02b}, \\cite{deg02c}), who studied mass segregation in open clusters in the Large Magellanic Cloud.  Additionally to the mass segregation in older clusters due to dynamical evolution, a higher concentration of massive stars to the centre was also found in some very young clusters (e.g., in the Orion Nebula Cluster, Hillenbrand~\\cite{hi97}, Hillenbrand \\& Hartmann~\\cite{hiha}). Due to the youth of these clusters, the central location of massive stars can not be explained by dynamical evolution only. Additional arguments from star formation scenarios and early cluster dynamics have to be considered (Bonnell \\& Davies~\\cite{boda}, Kroupa et al.~\\cite{krea01}).  Since mass segregation has a direct impact on the spatial structure of clusters, the effect should be seen in the {\\it apparent} cluster sizes. In this context, the observed cluster size is an important parameter related to the dynamical state both of the cluster and of the Galactic disk. As open clusters are found over a broad span of ages, a study of global trends including cluster size, if they exist, should be possible from a representative sample of clusters with homogeneous data on the main cluster parameters.  Practically all famous collections (Trumpler, Collinder etc.) of cluster data include estimations of angular sizes of open clusters, but the first systematic determination of apparent diameters was made by Lyng\\aa{}~(\\cite{lyn87}) for about 1000 open clusters from visual inspection of the POSS prints. These estimates are included in the Lund catalogue, 5th edition, together with about other 150 clusters, where an estimate by G. Lyng\\aa{}~ himself was not available, hence taken from the references quoted in the catalogue. For reasons of homogeneity only the diameter of the cluster nucleus (core) was included in the catalogue, although already then it was known that some clusters showed coronae. These data were used by Lyng\\aa{}~(\\cite{lyn82}) and Janes et al.~(\\cite{janea88}) for their studies of properties of the open cluster system. Since that time ages and distances were available for a small fraction of known clusters, their sample included about 400 clusters. No estimations were made on how well this sample represents the Galactic cluster population.  Although structural parameters have been derived for many individual clusters during the last decade, there are only a few studies dealing with a systematic determination of cluster sizes based on objective and uniform approaches for larger cluster samples. Danilov \\& Seleznev~(\\cite{danil}) derived structural parameters for 103 compact distant ($>$ 1~kpc) clusters from star counts down to $B\\approx16$ from homogeneous wide-field observations with a 50-cm Schmidt camera of the Ural university. Based on \\emph{UBV}-CCD observations compiled from literature, Tadross et al.~(\\cite{tad02}) redetermined ages and distances for 160 open clusters. The cluster sizes were estimated visually, from POSS prints, and they are practically identical to the diameters estimated by Lyng\\aa{}. Kharchenko et al.~(\\cite{khea03}) determined radii of about 400 clusters from star counts in \\ascc and USNO-A2.0 catalogues. Nilakshi et al.~(\\cite{nilak}) derived structural parameters of 38 open clusters selected from the Lyng\\aa's~(\\cite{lyn87}) catalogue from star counts in the USNO-A2.0 catalogue. Recently, Bonatto \\& Bica~(\\cite{bonb}) published structural and dynamical parameters of 11 open clusters obtained from star counts and photometric membership based on the 2MASS survey.  Correlations of cluster size with age and Galactic location were found by some of the authors above, though the results are rather controversial (see \\S~\\ref{sec:location} and \\S~\\ref{sec:age} for more details). There are at least two major aspects which must be taken into account in the interpretation of the results. At first, how well does a given sample represent the local population of open clusters in the Galaxy, or which biases can arrise from the incompleteness of the data and influence the results. Second, how homogeneous are data on individual clusters, on their size, age, distance, provided that they are based on observations with different telescopes equipped with different detectors, or if different methods were used for the determination of cluster parameters. The answer is not trivial considering the large set of data compiled from literature, especially.  Using the Catalogues of Open Cluster Data (COCD% \\footnote{\\texttt{ftp://cdsarc.u-strasbg.fr/pub/cats/J/A+A/438/1163, ftp://cdsarc.u-strasbg.fr/pub/cats/J/A+A/440,403}% }; Kharchenko et al.~\\cite{starcat},~\\cite{clucat},~\\cite{newclu}, Paper I, II, III, respectively), we are able to reduce those uncertainties which are due to the inhomogeneity of the cluster parameters, and we can better estimate biases due to an incompleteness of the cluster sample. The \\clucat is originated from the All-Sky Compiled Catalogue of 2.5 million stars (\\ascc$\\negthickspace$% \\footnote{\\texttt{ftp://cdsarc.u-strasbg.fr/pub/cats/I/280A}% }; Kharchenko~\\cite{kha01}) including absolute proper motions in the Hipparcos system, $B$, $V$ magnitudes in the Johnson photometric system, and supplemented with spectral types and radial velocities if available. The \\ascc is complete up to about $V=11.5$~mag. We identified 520 of about 1700 known clusters (Paper~I) in the \\ascc and found 130 new open clusters (Paper~III). Therefore, the completeness of the cluster sample is mainly defined by the limiting magnitude of the \\ascc. For each cluster, membership was determined by use of spatial, kinematic, and photometric information (Paper~I), and a homogeneous set of structural, kinematic and evolutionary parameters was obtained by applying a uniform technique (Papers II and III). The sample was used to study the population of open clusters in the local Galactic disk by jointly analysing the spatial and kinematic distributions of clusters (Piskunov et al.~\\cite{clupop}, Paper IV).  In this paper we use the homogeneous data on structural parameters of open clusters from the COCD to study general correlations including cluster sizes as well as to analyse the spatial distribution of cluster members from the point of view of mass segregation. In Sec.~\\ref{sec:data} we briefly describe the data set and estimate the statistic properties of the cluster sample. The relations between cluster radius and its location in the Galaxy are discussed in Sec.~\\ref{sec:location}. The correlations of cluster size with age is considered in Sec.~\\ref{sec:age}. In Sec.~\\ref{sec:masseg} we examine the effect of mass segregation in open clusters. A summary is given in Sec.~\\ref{sec:concl}. ",
        "conclusions": "}  This study is based on the Catalogue of Open Cluster Data (\\clucat) and its Extension~1 described in Papers II and III. The \\clucat is derived from the \\ascc, a homogeneous all-sky catalogue with complete information on proper motions and $B,V$-photometry. So, all open clusters found in this catalogue can be treated in the same way to derive their astrophysical parameters. On the other hand, the price to be paid for this advantage is the bright completeness limit of \\ascc at about $V$~=~11.5. However, the biases resulting from a simply magnitude limited sample can be estimated, they have been discussed in the previous sections and have been taken care of in order not to influence the conclusions. Using samples of clusters from different sources with different photometry and/or different limiting magnitude may introduce biases in the results which cannot be estimated easily.  The whole sample from \\ascc consists of 641 open clusters. In papers II and III we determined membership in the clusters applying photometric as well as astrometric criteria. Apparent linear radii have been computed from individual distances and angular sizes of the clusters, based on members only. For the first time, the structural properties of the galactic open cluster system have been statistically analysed from an unbiased, homogeneous, and relatively large sample. A comparison of our cluster sizes with those given in Lyng\\aa{}~(\\cite{lyn87}) (about 500 clusters in common) shows that cluster radii from Lyng\\aa{}~ are in average lower by a factor of 2, and they fit rather the core than the corona.  Our large sample allowed us to investigate the dependence of the cluster size on the age of a cluster and on its location in the Galaxy. The clusters cover an age range between about 5~Myr to more than 1~Gyr. For younger clusters ($<$ 200 Myr) there is no significant correlation between linear size and Galactocentric distance. At an age corresponding to two revolutions around the Galactic centre we detect that the clusters are on average smaller ($\\overline{R_{cl}}=3.8\\pm0.2$~pc) inside the solar circle than outside ($\\overline{R_{cl}}=4.6\\pm0.3$~pc). According to a ($K-S$) test the probability that both subsamples are drawn from the same distribution is less than 4~\\%. This size dependence on Galactocentric radius lead to the conclusion that the inner Galactic disk is void with respect to older open clusters. No clusters older than the age of the Hyades should exist inside a Galactocentric radius of about 6~kpc. Perpendicular to the plane we note a systematic increase of cluster sizes with increasing $|Z^{\\prime}|$. This, however, turned out to be significant only for clusters older than $\\log t>8.35$, which already survived at least one revolution around the Galactic centre.  From these findings the following picture of the evolution of open clusters arises. Clusters in the wider Solar neighbourhood are formed within the thin disk, their initial size distribution does not show a significant correlation with the $R_{G}$-- and $|Z^{\\prime}|$- coordinates. The size distribution changes at ages corresponding to one revolution around the Galactic centre. At low $Z^{\\prime}$ we now note a relatively larger number of small clusters. This makes us conclude that close to the Galactic equator and inside the solar circle larger clusters are in danger to dissolve even during the first revolution around the Galactic centre. On the other hand, they have a higher chance to survive encounters and the impact of Galactic tidal forces, if their orbits are outside the Solar one and are inclined to the Galactic plane. Therefore, they reach higher ages at these locations. Finally, the apparent linear sizes of clusters and their cores are, on average, decreasing with time and this process is faster for the coronae than for the cores. Taking into account that our input catalogue is magnitude limited, this finding can be interpreted as a first hint for mass segregation.  In the majority of clusters of our sample clear evidence for mass segregation of stars with $m>1\\,m_\\odot$ has been established from the distribution of the radial mass gradient as a function of age. An apparent flattening of the radial mass gradient for clusters older than 50...100 Myr occurs due to stellar evolution when massive stars subsequently leave the main sequence, and, secondly, because we cannot observe the low-mass stars due to the bright limiting magnitude of the \\ascc. External gravitational shocks may also influence the mass distribution in clusters and can be partly responsible for a spread of the radial mass gradient at $\\log t > 8$. Nevertheless, a ``typical'' cluster older than about 100 Myr and within about 1~kpc from the Sun shows mass segregation.  The youngest clusters of our sample with ages less than 50 Myr show a large spread of the radial mass gradient: from clusters with a clear concentration of the most massive stars to the centres up to clusters with no or only a flat mass gradient. The different dynamical state of clusters of the same age possibly results from the different initial conditions and environments of the clusters."
    },
    "0606/astro-ph0606539_arXiv.txt": {
        "abstract": "We present a method for simulating the evolution of HII regions driven by point sources of ionizing radiation in magnetohydrodynamic media, implemented in the three-dimensional Athena MHD code. We compare simulations using our algorithm to analytic solutions and show that the method passes rigorous tests of accuracy and convergence. The tests reveal several conditions that an ionizing radiation-hydrodynamic code must satisfy to reproduce analytic solutions. As a demonstration of our new method, we present the first three-dimensional, global simulation of an HII region expanding into a magnetized gas. The simulation shows that magnetic fields suppress sweeping up of gas perpendicular to magnetic field lines, leading to small density contrasts and extremely weak shocks at the leading edge of the HII region's expanding shell. ",
        "introduction": "Observations show that giant molecular clouds (GMCs) in local group galaxies convert at most a few percent of their mass into stars per cloud crossing time \\citep{zuckerman74}, and that clouds are typically destroyed in a few crossing times \\citep{blitz06a}, well before they have converted a significant fraction of their mass into stars. Observations also strongly support the idea that GMCs are gravitationally bound \\citep{krumholz05c,blitz06a,krumholz06d}, so the fact that they survive for more than a crossing time yet do not collapse entirely into stars strongly suggests that internal feedback plays a dominant role in GMC evolution. HII regions driven by newly-formed massive stars are likely to be the dominant sources of energy injection and mass loss in clouds, and are therefore critical to GMC evolution \\citep{mckee97,williams97,matzner02}. \\citet{krumholz06d} show using semi-analytic models that feedback from HII regions can quantiatively reproduce the observed lifetime, star formation rate, and star formation efficiency of GMCs. However, these calculations rely on simple analytic solutions for the evolution of spherically-symmetric HII regions in non-magnetic gas. Understanding the detailed evolution of GMCs will require a considerably more sophisticated numerical approach, and for this reason three-dimensional simulation of the evolution of molecular clouds under the influence of internal sources of ionizing radiation is a critical problem in numerical astrophysics.  In this paper we have the dual purpose of presenting a new algorithm for simulation of HII regions in molecular clouds, and using this algorithm to explore potential computational problems that arise in general for simulations of this type. We also demonstrate our new algorithm in a simple application, the expansion of an HII region into a uniform neutral gas in which the magnetic pressure greatly exceeds the thermal pressure, as it does in molecular clouds. In the past year several authors have considered the problem of simulating HII regions, and presented both algorithms \\citep[e.g.][]{arthur06, mellema06} and results on the evolution of HII regions in turbulent hydrodynamic media \\citep{dale05, mellema05,maclow06}. For a recent review see \\citet{henney06}. Several authors have also presented methods for simulation of ionizing radiation-hydrodynamics (IRHD) in a cosmological context \\citep[e.g.][]{abel02, whalen06}, which differs from the problem in the context of present-day molecular clouds primarily in the amount of cooling to which the gas is subjected and the conditions of the gas before it is ionized. Finally, researchers studying the evolution of ultracompact HII regions planetary nebulae have presented algorithms and results for ionizing radiative transfer with hydrodynamics and magnetohydrodynamics in 2D and 3D under the simplifying assumption that the ionized gas has a perfectly sharp edge and is always in thermal and ionization equilibrium \\citep[e.g.][]{garciasegura96, garciasegura97, garciasegura00}.  Our algorithm improves on previous work in that it is the first to couple ionizing radiative transfer to magnetohydrodynamics (IRMHD) without imposing an assumption of thermal or ionization equilibrium. The inclusion of magnetic fields is important because observations indicate that the magnetic energy in molecular clouds is comparable to the kinetic and gravitational potential energies \\citep{crutcher99,crutcher05,heiles05}. Thus, dynamical expansion of ionized regions may be significantly altered by magnetic confinement, an effect we wish to explore. Indeed, we show that even in the simple case of expansion of an HII region into a uniform, magnetized medium, the magnetic field produces qualitatively new phenomena. The Alfven speeds of a few km s$^{-1}$ typically found in molecular clouds reduce the strength of shocks associated with expanding ionization fronts at early times, and at later times turn off the shocks entirely, greatly reducing the collection of gas and possibly thereby reducing the amount of triggered star formation. A non-equilibrium treatment of the ionization structure is important because we find that a non-equilibrium treatment of the thermal and ionization structure of the ionized gas leads to modest but significant effects on quantities such as the expansion rate of HII regions, effects that are lost in the assumption of perfect radiative equilibrium.  Before exploring these new phenomena, however, we point out that there has yet to be a detailed study of potential computational problems and constraints that arise in 3D simulations of IRHD and IRMHD with the strong cooling and large temperature contrasts that are expected for modern day (as opposed to primordial) interstellar chemistry. We have therefore performed a detailed comparison of simulations using our method to analytic solutions for computationally challenging problems, as a way of searching for potential difficulties that may arise in general IRHD methods. We discover several conditions that a simulation must satisfy to reproduce analytic results correctly. One must limit the amount by which the gas pressure is allowed to change between hydrodynamic updates, for a ray-tracing method one must periodically rotate the orientation of the rays, and one must either resolve the ionization front or restrict the rate of cooling at the front to suppress excess cooling due to numerical mixing. We show that failure to meet these conditions produces quantitatively incorrect results.  The remainder of this paper proceeds as follows. In \\S~\\ref{formulation} we describe the physical formulation of the problem that we adopt, including our approximations and assumptions. In \\S~\\ref{algorithm} we present our simulation algorithm. In \\S~\\ref{tests} we compare our code to analytic solutions, and thereby demonstrate the existence of conditions that numerical methods must satisfy in order to reproduce analytic solutions correctly. We use our method to simulate the evolution of HII regions in magnetized media in \\S~\\ref{MHDresults}, and finally we summarize and present conclusions in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have demonstrated an algorithm for computing the evolution of magnetized molecular gases subjected to internal sources of ionizing radiation, which is potentially applicable to molecular clouds. In testing our algorithm, we have discovered three conditions that are likely to apply to ionizing radiation hydrodynamic and magnetohydrodynamic codes in general. First, to achieve maximum accuracy the update time step must be limited so that the temperature in cells does not change by more than a factor of $f<10$ between hydrodynamic or magnetohydrodynamic updates. Larger values of $f$ produce small but significant errors in the expansion rates of D type ionization fronts. Second, when using a ray-tracing approach to compute the ionizing radiative transfer, one should rotate the orientation of the rays periodically to avoid a build-up of errors caused by the discretization of angles around the ionizing source. Failure to obey this conditions results in fronts that should be spherical developing aspherical features, and in more complex calculations this could potentially seed instabilities. Third, and most significantly, one must avoid overcooling caused by numerical smearing of the ionization front. This can be handled most easily by suppressing cooling in cells with mixed ionization fractions. Failure to obey this condition leads to an unphysical loss of energy from expanding HII regions that causes them to lag analytic solutions by tens of percent. A calculation that satisfies these three constraints can reproduce the analytic solution for the expansion of a D type ionization front to an accuracy of a percent for at least $\\sim 20$ ionized sound-crossing times.  Using our algorithm, we report the first three-dimensional simulations of the expansion of an HII region into a magnetized gas. We show that the presence of a magnetic field distorts the HII region, and greatly reduces the strength of the shock and the density contrast in directions perpendicular to the magnetic field. This leads to the formation of an HII region which is bounded by a dense shell of swept-up gas along the field, but not perpendicular to it. The absence of a dense shell over much of the solid angle means that, in the presence of strong, ordered magnetic fields, HII regions may not be able to collect and compress as much gas as one might expect from purely hydrodynamic estimates. This may reduce the efficiency of triggered star formation from HII regions."
    },
    "0606/astro-ph0606013_arXiv.txt": {
        "abstract": "{Motivated by recent observations of plateaus and minima in the radial abundance distributions of heavy elements in the Milky Way and some other spiral galaxies, we propose a dynamical mechanism for the formation of such features around corotation. Our numerical simulations show that the non-axisymmetric gravitational field of spiral density waves generates cyclone and anticylone gas flows in the vicinity of corotation. The anticyclones flatten the  pre-existing negative abundance gradients by exporting many more atoms of heavy elements outside corotation than importing inside it. This process is very efficient and forms plateaus of several kiloparsec in size around corotation after two revolution periods of a galaxy. The strength of anticyclones and, consequently, the sizes of plateaus depend on the pitch angle of spiral arms and are expected to increase along the Hubble sequence.} ",
        "introduction": "It has recently been recognized that density waves in the stellar component of spiral galaxies have a profound effect on the dynamics of stars, cold gas clouds, and dust in the vicinity of the corotation resonance. The non-axisymmetric gravitational field of spiral stellar density waves causes large changes ($\\sim 50$ per cent over the lifetime of a galaxy) in the angular momenta of individual stars and cold gas clouds around the corotation radius (Sellwood \\& Binney \\cite{Sellwood}). Considerable radial migrations associated with the angular momentum changes are expected to dilute the abundance gradients in the cold gas component of spiral galaxies (Sellwood \\& Preto \\cite{Sellwood2}).   The radial abundance distribution of heavy elements in log scale in at least some spiral galaxies (perhaps, including the Milky Way) cannot be described by a linear function with a negative slope. According to Zaritsky, Kennicutt, \\& Huchra (\\cite{Zaritsky}), the oxygen abundances in NGC~2997, NGC~3319, NGC~5033 and other spiral galaxies in their sample show a complex nonlinear behaviour -- plateaus and minima can be identified in the radial distribution of oxygen. Perhaps more convincing evidence for a complex radial distribution of heavy element abundances is found in the Milky Way. For instance, the radial abundance distributions of O, N, Mg, and other heavy elements derived by Daflon \\& Cunha (\\cite{Daflon}) from a sample of OB stars show a minimum near 8~kpc (although they have not accentuated the importance of this behaviour and approximated the radial abundance profiles by a linear function with a negative slope).  The existence of a plateau in the oxygen abundance distribution has also been reported by Andrievsky et al. (\\cite{Andrievsky}). Although a definite confirmation of plateaus or/and minima in the radial abundance profiles of heavy elements requires a larger sample of abundance tracers than has been used in the abovementioned studies, the existing evidence strongly suggests these features.   A simple multizone model of chemical enrichment in spiral galaxies has been recently proposed by Mishurov et al. (\\cite{Mishurov}) and Acharova et al. (\\cite{Acharova}). It explains the formation of minima and/or plateaus in the radial abundances of heavy elements near corotation by a selective action of star formation. The star formation rate around corotation is assumed to have a minimum (due to the lack of strong spiral shock waves) and consequently the heavy element production also has a minimum at the corotation radius. The assumed temporal migration of the corotation resonance can produce either plateaus or minima in the radial abundance distributions of heavy elements.  In this paper, we focus on a purely hydrodynamic explanation for a nonlinear radial distribution of heavy elements in spiral galaxies. We present the first numerical hydrodynamic simulations that self-consistently explain the formation of a plateau in the heavy element abundance distribution in the vicinity of corotation. We demonstrate the development of cyclones and anticyclones in the gas flow around corotation and study their influence on the radial abundance distribution of heavy elements in spiral galaxies. The existence of cyclones and anticyclones  has been observationally confirmed in at least two spiral galaxies (Fridman et al. \\cite{Fridman1,Fridman2}) and has been predicted in the laboratory experiments of rotating shallow water modelling (Nezlin \\cite{Nezlin}). The model equations are formulated in \\S~\\ref{model}. The numerical code is described in \\S~\\ref{code} and the initial conditions are given in \\S~\\ref{init}. The results of numerical simulations are presented in \\S~\\ref{results}. The possible implications for spiral galaxies of different Hubble types are discussed in \\S~\\ref{hubble}. The main results are summarized in \\S~\\ref{sum}. ",
        "conclusions": "\\label{sum} We have studied numerically the dynamics of warm gas and heavy elements in spiral galaxies. Two types of two-armed spiral galaxies are considered in which corotation is situated approximately in the middle of the spiral pattern (type~1) and at the very end of it (type~2), respectively. Type~2 spiral galaxies show little radial redistribution of heavy elements in the warm gas disk. Conversely, type~1 spirals  are distinguished by a substantial radial redistribution of heavy elements around corotation.  Strong cyclones and anticyclones around corotation are generated by the non-axisymmetric gravitational field of spiral stellar density waves in type~1 galaxies. The anticyclone flows transport gas from inside corotation to  the regions outside it and vice versa. If the radial abundances of heavy elements are characterized by a negative slope, the anticyclones bring many more atoms of heavy elements outside corotation than they import inside corotation. This results in a flattening of radial abundance profiles at the position of anticyclones after two revolution periods of a galaxy. On the other hand, the cyclone flows generate in-going and out-going streams of gas along the spiral arms. These streams meet at corotation, producing a contact discontinuity in the gas flow and associated step-like radial profiles of heavy element abundances at the position of cyclones. Nevertheless, the azimuthally averaged radial abundance distributions of heavy elements show {\\it  a well-defined plateau on both sides of corotation}, implying that anticyclones are more powerful in transporting the heavy elements than cyclones.   The sizes of plateaus around corotation in the azimuthally averaged radial abundance distributions of heavy elements are expected to increase along the Hubble sequence in spiral galaxies with equal number of arms. Our numerical simulations show that in two-armed Sa-Sb galaxies with a pitch angle of $10^\\circ$, the plateau has a maximum size of 2~kpc. In contrast, Sc galaxies with a pitch angle of $25^\\circ$ have the plateau that can become as large as 4.5~kpc. A growing efficiency of radial mixing of heavy elements (and an associated flattening of negative abundance gradients) along the Hubble sequence is related to the increasing radial separation between the spiral arms and a consequent increase in the strength of anticyclones.  A considerable portion of total gas mass in a spiral galaxy may be in the form of cold molecular hydrogen clouds. Our numerical simulations show that the efficiency of radial mixing grows as the gas temperature drops. A decrease in the restoring force of pressure gradients is responsible for this effect. However, our numerical hydrodynamics code is not appropriate for the modelling of molecular cloud dynamics, for which the sticky particle codes are well suited. Sellwood \\& Binney (\\cite{Sellwood}) and Sellwood \\& Preto (\\cite{Sellwood2}) have employed N-body simulations to study the dynamics of cold molecular clouds. They have reported the development of anticyclonic motions at corotation and predicted the flattening of any metallicity gradients within the disc.  In this paper, we neglect the effect of continuous star formation. It is a computationally difficult task to self-consistently include the production of heavy elements in the hydrodynamics code. A possible exception is oxygen which is released mostly by short-living massive stars, for which the instantaneous recycling approximation can be used (e.g. Acharova et al. \\cite{Acharova}). Our preliminary numerical simulations indicate that the on-going star formation in the spiral arms may change the shape of the plateau at corotation. More specifically, a mild minimum and maximum in the azimuthally averaged radial abundance profile of oxygen may develop at the inner and outer sides of corotation, respectively. Similar shapes in the radial abundance distribution of oxygen near corotation were reported by Acharova et al. (\\cite{Acharova}). The results of this study will  be presented in a follow-up paper."
    },
    "0606/hep-ph0606289_arXiv.txt": {
        "abstract": "A wide range of techniques have been developed to search for particle dark matter, including direct detection, indirect detection, and collider searches. The prospects for the detection of neutralino dark matter is quite promising for each of these three very different methods. Looking ahead to a time in which these techniques have successfully detected neutralino dark matter, we explore the ability of these observations to determine the parameters of supersymmetry. In particular, we focus on the ability of direct and indirect detection techniques to measure the parameters $\\mu$ and $m_A$. We find that $\\mu$ can be much more tightly constrained if astrophysical measurements are considered than by LHC data alone. In supersymmetric models within the $A$-funnel region of parameter space, we find that astrophysical measurements can determine $m_A$ to roughly $\\pm100$ GeV precision. ",
        "introduction": "\\label{}  If low scale supersymmetry exists in nature, it will likely be discovered in the next few years at the Large Hadron Collider (LHC), or perhaps even earlier at the Tevatron. In many supersymmetric scenarios, the LHC can identity the presence of the lightest neutralino in the form of missing energy in the cascade decays of squarks and/or gluinos. At approximately the same time, direct and indirect searches for dark matter will be reaching the level of sensitivity needed to discover neutralinos in a wide range of supersymmetric models.  Collider and astrophysical experiments tell us very different things about the nature of supersymmetry. The LHC is likely to reveal the approximate masses of a number of superpartners, including squarks, gluinos, the lightest neutralino, and in some cases sleptons and heavier neutralinos. Other properties of the supersymmetric model will remain unconstrained by the LHC, however. In particular, the composition of the lightest neutralino (the mixture of bino, wino, and higgsino components), and therefore its couplings, will be very difficult to deduce at a hadron collider. In contrast, the cross sections relevant to astrophysical dark matter experiments depend critically on the composition of the lightest neutralino.   \\begin{figure}[t] \\hspace{-1.5cm} \\includegraphics[width=2.2in,angle=-90]{direct.ps} \\includegraphics[width=2.2in,angle=-90]{nu4.ps} \\caption{Left: The relationship between the quantity  $|N_{11}|^2 |N_{13}|^2 \\tan^2 \\beta/m^4_A$ and the spin-independent neutralino-nucleon elastic scattering cross section. Right: The rate of neutrinos detected in a kilometer scale neutrino telescope, such as IceCube, from neutralino annihilations in the Sun as a function of the quantity $|N_{13}|^2-|N_{14}|^2$. A constraint 100 times more stringer than the current CDMS bound has been applied (in the right frame) in anticipation of increased sensitivity from direct detection experiments in the coming few years.} \\label{directneutrino} \\end{figure}   In Fig.~\\ref{directneutrino}, we show how the rates in astrophysical dark matter experiments depend on the composition of the lightest neutralino. In the left frame, we compare the scalar neutralino-nucleon elastic scattering cross section (which determines the rate in direct dark matter experiments) to the quantity:  $\\tan^2 \\beta |N_{11}|^2 |N_{13}|^2/m^4_A$. Here, $m_A$ is the mass of the CP-odd Higgs boson in the MSSM. The $|N|^2$'s are defined by the following: $\\chi^0_1 = N_{11}\\tilde{B}     +N_{12} \\tilde{W}^3 +N_{13}\\tilde{H}_1 +N_{14} \\tilde{H}_2$. The correlation shown in the left frame of Fig.~\\ref{directneutrino} is quite tight for models with large cross sections, which are produced through diagrams which exchange a heavy Higgs boson  coupling to $b$ and $s$ quarks (\\cite{carena}).  Neutralinos scattering with the Sun can become gravitationally bound and accumulate in the Sun's core, eventually annihilating to produce high-energy neutrinos. The rate of those neutrinos being observed at a kilometer-scale neutrino telescope such as IceCube or KM3 is shown in the right frame of Fig.~\\ref{directneutrino} as a function of the quantity $|N_{13}|^2-|N_{14}|^2$. This correlation comes from the $\\chi^0$-$\\chi^0$-$Z$ coupling which largely determines the capture rate in the Sun in nearly all observable models (\\cite{halzen}). ",
        "conclusions": ""
    },
    "0606/astro-ph0606543_arXiv.txt": {
        "abstract": "We present an \\xmm\\ observation of the massive edge-on Sb galaxy NGC~2613. We discover that this galaxy contains a deeply embedded active nucleus with a 0.3-10 keV luminosity of 3.3$\\times10^{40}{\\rm~ergs~s^{-1}}$ and a line-of-sight absorption column of $1.2 \\times 10^{23} {\\rm~cm^{-2}}$. Within the $25{\\rm~mag~arcsec^{-2}}$ optical B-band isophote of the galaxy, we detect an additional 4 sources with an accumulated luminosity of 4.3$\\times10^{39}{\\rm~ergs~s^{-1}}$. The bulk of the unresolved X-ray emission spatially follows the near-infrared (NIR) K-band surface brightness distribution;the luminosity ratio $L_X/L_K \\sim 8\\times10^{-4}$ is consistent with that inferred from galactic discrete sources. This X-ray-NIR association and the compatibility of the X-ray spectral fit with the expected spectrum of a population of discrete sources suggest that low-mass X-ray binaries (LMXBs) are the most likely emitters of the unresolved emission in the disk region. The remaining unresolved emission is primarily due to extraplanar hot gas. The luminosity of this gas is at least a factor of 10 less than that predicted by recent simulations of intergalactic gas accretion by such a massive galaxy with a circular rotation speed $V_c \\sim 304 {\\rm~km~s^{-1}}$ (Toft et al.~2002). Instead, we find that the extraplanar hot gas most likely represents discrete extensions away from the disk, including two ``bubble-like'' features on either side of the nucleus. These extensions appear to correlate with radio continuum emission and, energetically, can be easily explained by outflows from the galactic disk. ",
        "introduction": "X-ray observations of extraplanar hot gas ($T$ $\\gtrsim$ 10$^6$ K) around nearby edge-on disk galaxies are essential in the study of the galactic ecosystem in many aspects, particularly the disk-halo interaction. Such observations have helped establish the prevalence of galactic superwinds in starburst galaxies, e.g., NGC~253 (Strickland et al.~2000, 2002) and NGC 4666 (Dahlem, Weaver \\& Heckman~1998), among others. Extraplanar X-ray-emitting gas has also been detected unambiguously around several ``normal'' late-type galaxies with little evidence for nuclear starbursts:~NGC~891 (Sb; Bregman \\& Houck 1997), NGC 4631 (Scd; Wang et al.~2001), NGC 3556 (Sc; Wang, Chaves \\& Irwin 2003) and NGC 4634 (Scd; T$\\ddot{\\rm u}$llmann et al.~2006). In these galaxies (except for NGC~4634 which currently lacks direct evidence), extraplanar hot gas is clearly linked to outflows from recent massive star-forming regions in galactic disks. The global X-ray properties of extraplanar gas in these ``normal'' star-forming galaxies, when scaled with the star formation rate of the host galaxies, appear similar to those found in starburst galaxies (Strickland et al.~2004a, b; Wang 2005). Nevertheless, this needs to be confirmed by extended X-ray observations of ``normal'' star-forming galaxies.  On the other hand, current galaxy formation models also predict the existence of hot gaseous halos surrounding present-day disk galaxies, which arise from gravitational infall from the intergalactic medium (IGM; e.g., Toft et al.~2002 and references therein). The predicted extraplanar X-ray luminosity strongly depends on the mass of the host galaxy. X-ray observations thus have long been expected to detect such gaseous halos around nearby massive, typically earlier-type disk galaxies. However, there is so far little direct observational evidence for the presence of this kind of X-ray-emitting halo. Benson et al.~(2000) analyzed X-ray emission from the outer halos ($\\gtrsim 5'$) of primarily two early-type spirals NGC 2841 (Sb; $\\sim 15$ Mpc) and NGC 4594 (Sa; $\\sim 25$ Mpc), using {\\sl ROSAT} PSPC observations. No significant diffuse emission was detected, although the upper limits to the diffuse X-ray luminosities are consistent with the current predictions (Toft et al.~2002). Therefore, more dedicated searches for the X-ray signals of IGM accretion around disk galaxies are needed.  Here we present a study of an \\xmm ~observation toward NGC~2613, an edge-on Sb galaxy with ``normal'' star formation. We focus on probing the spatial and spectral properties of its large-scale X-ray emission. This galaxy (Table~\\ref{tab:N2613}) is a good candidate to probe the presence of hot gas, in the sense that: 1) it is very massive and thus expected to contain a large amount of hot gas; 2) its high inclination ($\\sim$79$^\\circ$) allows the possibility of detecting extraplanar emission, either from a halo of accreted gas or a large-scale outflow; 3) its moderately large distance (25.9 Mpc) places the galaxy and its $\\sim$50 kpc vicinity in the field of view (FOV) of a typical \\xmm ~observation, offering a good opportunity of studying the large-scale distribution of gas, and 4) it is known to show extraplanar features at other wavebands, specifically the radio continuum and \\HI (Chaves \\& Irwin 2001; Irwin \\& Chaves 2003) as well as earlier \\HI observations (Bottema 1989). ",
        "conclusions": "{\\label{sec:discussion}} \\subsection{The nature of the nuclear X-ray emission} {\\label{subsec:nucleus}} Our two-component spectral fit for the nucleus (\\S~\\ref{subsec:ps}; Table \\ref{tab:nucleus_fit}) indicates that the intrinsic neutral hydrogen column density is $\\sim$$1.2{\\times}10^{23}{\\rm~cm^{-2}}$. This is much higher than the beam-averaged \\HI column density of $\\sim$$2\\times10^{21}{\\rm~cm^{-2}}$ found by Chaves \\& Irwin (2001), but it is typical of values found for molecular circumnuclear disks.  Ott et al. (2001), for example, find a molecular column density of order $10^{23}{\\rm~cm^{-2}}$ for NGC~4945. Molecular data are not yet available for NGC~2613, but our results suggest that a substantial molecular component should be present in this galaxy.  For the nuclear component, the modelled intrinsic flux given in Table~\\ref{tab:nucleus_fit} leads to an intrinsic X-ray luminosity of $\\sim$$3.3{\\times}10^{40}{\\rm~ergs~s^{-1}}$ in the 0.3-10 keV range.  The photon index of this component is $\\sim$2, a typical value found in the X-ray spectra of AGNs (e.g. Pellegrini, Fabbiano \\& Kim 2003). No radio core was detected by Irwin, Saikia \\& English (2000), putting a 3$\\sigma$ upper limit of $4.5\\times10^{27}{{\\rm ~ergs~s^{-1}~Hz^{-1}}}$ on the radio spectral power at 1.425 GHz within the same 16$^{\\prime\\prime}$ region. Using the above luminosity over the 0.3-10 keV range, we derive an upper limit of $\\alpha\\,=\\,0.62$ on the energy spectral index ($S_\\nu\\,\\propto\\,\\nu^{-\\alpha}$) between the radio and X-ray bands.  Although the X-ray nucleus is heavily obscured, these values nevertheless suggest that the energy spectral index is likely flat or possibly rising at the low frequencies, a fact again consistent with the interpretation of the nuclear source as an AGN.  Thus, we conclude that the nuclear X-ray source represents an AGN in this galaxy, the first evidence that this is the case.  The non-nuclear component, characterized by the second power-law (PL2), shows a photon index of $\\sim1.7$ and an intrinsic 0.3-10 keV luminosity of $3.4{\\times}10^{39}{\\rm~ergs~s^{-1}}$. Irwin et al.~(2003) showed that the accumulated spectra of LMXBs in early-type galaxies can be uniformly described by a power-law model with a best-fit photon index of $1.56\\pm{0.02}$. By using a sample of nearby galaxies of various morphological types, Gilfanov (2004) studied the relation between the collective luminosity of LMXBs and the K-band luminosity, $L_K$, of the underlying stellar content. He found that, for LMXBs with luminosity higher than $10^{37}{\\rm~ergs~s^{-1}}$, their collective luminosity $L_X$ = ${\\rm(3.3-7.5)}{\\times}10^{39}{\\rm~ergs~s^{-1}}L_K/10^{11}L_{\\odot,K}$. Assuming that the spatial distribution of LMXBs follows that of the K-band star light (Jarrett et al.~2003), we estimate that the collective luminosity of LMXBs within the 16$^{\\prime\\prime}$ circle is ${\\sim}$${\\rm(1.8-4.2)}{\\times}10^{39}{\\rm~ergs~s^{-1}}$. Thus the non-nuclear component is consistent with the collective emission of unresolved LMXBs. We note that high-mass X-ray binaries (HMXBs) are expected to be present in star-forming disk galaxies and their composite spectral properties are somewhat similar to that of the LMXBs, thus the collective contribution of HMXBs may also be partly responsible for the non-nuclear component. We show below that in NGC~2613 the relative contribution of HMXBs is small as compared to that of LMXBs.  \\subsection{The collective X-ray emission of discrete sources} {\\label{subsec:cxs}} It is known that X-ray binaries, including LMXBs and HMXBs, dominate the X-ray source populations with luminosities $\\gtrsim 10^{35}{\\rm~ergs~s^{-1}}$ in galaxies. Owing to their distinct evolution time-scales, the numbers and thus the collective contributions of long-lived LMXBs and short-lived HMXBs to the X-ray emission of a galaxy are expected to be proportional to its stellar mass and star formation rate (SFR), respectively. Colbert et al.~(2004) analyzed {\\sl Chandra} observations of X-ray sources in a sample of nearby galaxies of various morphological types and SFRs. They found that the collective X-ray luminosity of point sources $L_{XP}$ is linearly correlated with the total stellar mass $M_{\\star}$ and the SFR of the host galaxy as \\begin{eqnarray} L_{XP}~(\\rm{ergs~s^{-1}})~=~(1.3\\pm{0.2})\\times 10^{29}~M_{\\star}~(\\rm{M_{\\odot}}) \\nonumber \\\\ + (0.7\\pm{0.2})\\times 10^{39}~\\rm{SFR}~(\\rm{M_{\\odot}~yr^{-1}}). \\label{eq:LXP} \\end{eqnarray}  We use this relation to assess the relative importance of LMXBs and HMXBs in contributing to the X-ray emission of NGC~2613.  The total stellar mass can be estimated from the K-band luminosity $L_K$ and the $B-V$ color index via (Bell \\& de Jong 2001) \\begin{equation} {\\rm log}(M_{\\star}/L_K) = -0.692 + 0.652 (B-V), \\label{eq:Mstar} \\end{equation} where $L_K$ is in units of the K-band Solar luminosity. The SFR can be estimated from the far-infrared (FIR) luminosity $L_{FIR}$ via (Kennicutt 1998) \\begin{equation} \\rm{SFR}~=~4.5\\times10^{-44}~L_{FIR}~(\\rm{ergs~s^{-1}}). \\label{eq:SFR} \\end{equation} $L_{FIR}$ is measured according to (Lonsdale, Helou \\& Good 1989) \\begin{equation} L_{FIR} = 3.1\\times10^{39}~D^2~(2.58~S_{60}+S_{100}), \\label{eq:LFIR} \\end{equation} where $D$ is the distance of the galaxy in units of Mpc, $S_{60}$ and $S_{100}$ are the flux densities in units of Jy at 60~${\\mu}$m and 100~${\\mu}$m, respectively. With the available photometric data for NGC~2613 (Table \\ref{tab:N2613}), we estimate that the total stellar mass is $2.1\\times10^{11}{\\rm~M_{\\odot}}$ and the SFR is $4.2{\\rm~M_{\\odot}~yr^{-1}}$. Based on Eq.~(\\ref{eq:LXP}), the contributions of LMXBs and HMXBs to the collective X-ray emission of discrete sources is $\\sim2.7\\times10^{40}{\\rm~ergs~s^{-1}}$ and $\\sim2.9\\times10^{39}{\\rm~ergs~s^{-1}}$, respectively, with the latter being about 10\\% of the former.  In the disk of NGC~2613, we find that the 0.5-2 keV unresolved emission is spatially correlated with the K-band star light.  Therefore, the normalization factor for the K-band profile (\\S~\\ref{subsubsec:spat_anal}) should represent the collective X-ray emissivity of the underlying old stellar population. Using the power-law model given by Irwin et al.~(2003) for the accumulated spectrum of LMXBs, we convert the observed 0.5-2 keV count rate into the intrinsic luminosity in the 0.3-10 keV band). The K-band flux density is also converted into intrinsic luminosity according to the 2MASS K-band photometry. The normalization factor, 3.0$\\times10^{-4}{\\rm~cts~s^{-1}~arcmin^{-2}}/({\\rm MJy~sr^{-1}})$, is then equivalent to an X-ray emissivity of $L_X = 4.2{\\times}10^{39}{\\rm~ergs~s^{-1}}L_K/(10^{11}L_{\\odot,K})$, or a luminosity ratio of $L_X/L_K \\sim 7.5\\times10^{-4}$. Gilfanova (2004) found that the collective X-ray luminosity of galactic LMXBs is related to the underlying K-band luminosity following $L_X = (3.3-7.5) {\\times}10^{39}{\\rm~ergs~s^{-1}}L_K/(10^{11}L_{\\odot,K})$, i.e., a luminosity ratio of $L_X/L_K \\sim (5.8-13.2)\\times10^{-4}$. Therefore the collective X-ray emissivity of unresolved discrete sources inferred for NGC~2613 is consistent with that of the galactic LMXBs, and we conclude that the collective X-ray emission of LMXBs dominates the soft emission of NGC~2613 in its disk region.  \\subsection{The origin of extraplanar gas} Presence of diffuse gas in NGC~2613 is evident by the soft X-ray excess over the K-band light. We consider two possible origins of the diffuse gas: 1) the continuously accreted IGM (Toft et al. 2002) and 2) the outflow from the galactic disk (Irwin \\& Chaves 2003).  \\subsubsection{An accreted gaseous halo?} Toft et al.~(2002) calculated global X-ray properties (e.g., luminosity, effective temperature and intensity distribution) of hot gaseous halos, based on their simulated galaxies. The predicted luminosity strongly depends on the circular speed of the host galaxy. The most massive galaxies in their simulations have circular speeds similar to that of NGC~2613 ($\\sim$ $300{\\rm~km~s^{-1}}$). The predicted 0.2-2 keV luminosity for such a galaxy is $\\sim8\\times10^{40}{\\rm~ergs~s^{-1}}$ (Fig.~3 in Toft et al.~2002). From our best-fit spectral models of the spectra of unresolved emission, we derive an intrinsic 0.2-2 keV luminosity of $\\sim$8$\\times10^{39}{\\rm~ergs~s^{-1}}$ for the sum of the thermal and power-law components, and $\\sim$6$\\times10^{39}{\\rm~ergs~s^{-1}}$ for the thermal components only. We note that the unresolved emission outside our spectral extraction region contributes little to the total luminosity. The simulated luminosity of gas emission by Toft et al.~(2002) is at least an order of magnitude higher than the observed value for NGC~2613. We therefore conclude that the simulations as presented by Toft et al.~(2002) substantially over-predict the X-ray emission from the cooling inflow of the IGM, if this is what is occurring in NGC~2613.  This over-prediction is related to the so-called over-cooling problem in current theories of galaxy formation. We speculate that the over-cooling problem is a result of an inappropriate treatment of stellar and/or AGN feedback. For example, the mechanical energy input from Type Ia supernovae (SNe) is typically not included in galaxy formation simulations, partly due to the difficulty in treating the astrophysics related to gaseous flows. Qualitatively, Type Ia SNe, which tend to occur in low-density hot environments, provide an especially effective mechanism for large-scale distributed heating, required to reduce the cooling of gas in galactic bulges and halos (Tang \\& Wang 2005). Massive stars in galactic disks may also serve as sources of mechanical energy that could produce outflows into halos and help slow down the cooling of the accrected gas. For example, with a star formation rate of $\\sim4.2~{\\rm M_{\\odot}~yr^{-1}}$ for stars between 0.1 and 100 ${\\rm M_{\\odot}}$, and assuming a Salpeter IMF and that stars with mass $> 8 {\\rm~M_{\\odot}}$ become core-collapse SNe, the rate of total energy release from the star-forming regions of NGC~2613 is $L_{SNII}\\sim1.0\\times10^{42} {\\rm~ergs~s^{-1}}$. Our spatial and spectral analyses suggest that the extraplanar gas is responsible for the thermal emission (\\S~\\ref{subsec:diffuse}) and has a total 0.3-10 keV luminosity of $\\lesssim5\\times10^{39}{\\rm~ergs~s^{-1}}$. Thus, SNe can provide enough energy to explain the X-ray emission of the extraplanar gas in NGC~2613.  \\begin{figure*}[!htb] \\vskip -1cm \\centerline{ \\epsfig{figure=n2613_x_r_c.eps,width=0.8\\textwidth,angle=0} } \\vskip -0.5cm \\caption{(a) VLA C+D configuration continuum contours overlaid on the same X-ray intensity image (grey scale) as contoured in Fig.~\\ref{fig:2dmap} and in (b) of this figure. The contour levels are 0.18, 0.27, 0.56, 0.84, 1.1, 1.7, 2.3, 3.2~mJy~beam$^{-1}$ and the beam is $22^{\\prime\\prime}{\\times}15^{\\prime\\prime}$ at a position angle of $-8\\fdg2$. A few X-ray extraplanar features are labelled (see text). (b) The same X-ray intensity contours as in Fig.~\\ref{fig:2dmap} overlaid on a greyscale image of the total intensity VLA C+D configuration \\HI map. The grey scale range (shown with a square root transfer function) is in units of $10^3$ Jy beam$^{-1}$ m s$^{-1}$ and the beam is $47^{\\prime\\prime}{\\times}32^{\\prime\\prime}$ at a position angle of $-8\\fdg2$. F1 and F2 refer to two \\HI extensions identified by Chaves \\& Irwin (2001). } \\label{fig:overlay_map} \\end{figure*}  \\subsubsection{Multiwavelength extraplanar features} {\\label{subsubsec:correlations}} In Fig.~\\ref{fig:overlay_map}a and b, we compare the X-ray emission with the radio continuum emission and \\HI total intensity emission, respectively. Of the two radio images, the radio continuum morphology more closely resembles the X-ray morphology in the sense that: a) the north bubble  has a radio continuum counterpart; b) the south extension also has a radio continuum counterpart; c) the south-west feature shows a small radio continuum protrusion; and d) the peaks of the large eastern extensions (north and south) also show radio emission.  The \\HI total intensity map shown here does not show all of the extended features identified by Chaves \\& Irwin (2001), but two of their features, F1 and F2 clearly extend above and below the galactic plane and are labelled in Fig.~\\ref{fig:overlay_map}b.  These two features might be related with the northern and southern arc of the eastern extensions seen in the X-ray.  It is not wise to read too much into these correlations, given the limited S/N of the maps. However, the relationship with the radio continuum is sufficiently strong that the X-ray emission in the extraplanar features, representing hot diffuse gas, is very likely associated with the radio continuum emission which represents predominantly the non-thermal component.  We further consider the energetics of a specific feature, namely the `north bubble', which is the only extraplanar feature that can be cleanly isolated from the ambient emission. Guided by Fig.~\\ref{fig:bubble_map}, we approximate the volume occupation of the bubble by a cylinder with 1$^\\prime$ in diameter and 0\\farcm8 in height, the center of which is 1\\farcm2 above the galactic center. Hence the volume of the bubble is $\\sim2.7\\times10^2{\\rm~kpc^3}$. We find a total 0.5-2 keV count rate of 2.4$\\times10^{-3}{\\rm~cts~s^{-1}}$ within the bubble. In the best-fit model to the spectra of unresolved emission, the high and low temperature components predict a 0.5-2 keV count rate of 7.5$\\times10^{-3}{\\rm~cts~s^{-1}}$ and 2.0$\\times10^{-3}{\\rm~cts~s^{-1}}$, respectively. Therefore the north bubble is unlikely to be due to the low temperature component alone. Instead, it could be dominated by the high temperature component. Taking an effective temperature of $\\sim$ 0.8 keV, we estimate the mean density of the bubble to be $\\sim\\eta^{-1/2}\\times10^{-3}{\\rm~cm^{-3}}$, where $\\eta$ is the filling factor of the hot gas inside the bubble. The total thermal energy of the bubble is $E_{th}\\simeq 3.6\\eta^{-1/2}\\times10^{55}{\\rm~ergs~s^{-1}}$, and the work done to steadily lift up the bubble against gravity is $E_g \\simeq 1.7\\eta^{-1/2}\\times10^{55}{\\rm~ergs~s^{-1}}$, given the gravitational potential introduced by the exponential disk of the galaxy (Irwin \\& Chaves 2003). Given the morphology and the position of the bubble, we speculate that it was produced near the nuclear region, either by a nuclear starburst or the AGN. If the bubble's total amount of thermal and gravitational energy is obtained from a starburst, it takes a time of $\\tau\\simeq(E_{th}+E_g)/(fL_{SNII})\\simeq1.7\\eta^{-1/2}\\times10^7{\\rm~yr}$ to form the present structure, where $f$ is a geometrical factor taken to be 0.1 to reflect a fractional star formation rate of the central 1 kpc in the disk. This timescale is typical for massive stars to become SN explosions. On the other hand, assuming the flux density of the AGN follows $S_\\nu\\,\\propto\\,\\nu^{-\\alpha}$ between the radio and X-ray bands, the total bolometric luminosity over this frequency range is $\\sim 4.5\\times10^{40} {\\rm~ergs~s^{-1}}$ with $\\alpha=0.62$ adopted (\\S~\\ref{subsec:nucleus}). It is uncertain what fraction of the AGN energy can be taken to energize the ambient gas, but we consider the AGN might also be capable of producing this feature. For example, the locations of the north bubble and south extension immediately on either side of the nucleus are reminiscent of extraplanar loops or lobes seen in nuclear outflow galaxies like NGC~3079 (e.g. Cecil et al. 2002) which is known to have an AGN.   \\begin{figure}[!htb] \\vskip -1.5cm \\centerline{ \\epsfig{figure=x_bubble.eps,width=0.5\\textwidth,angle=0} } \\vskip -2cm \\caption{Blow-up of the X-ray-emitting bubble to the north of the nucleus. The same X-ray intensity image as in Fig.~\\ref{fig:2dmap} is used. Contour levels are at 5, 5.2, 5.4, 5.6, 5.9, and 6.2 $\\times$10$^{-3}{\\rm~cts~s^{-1}~arcmin^{-2}}$. } \\label{fig:bubble_map} \\end{figure}   } We have analyzed an \\xmm~observation of the massive edge-on Sb galaxy NGC~2613. We find a deeply embedded AGN in this galaxy. The X-ray spectrum of this AGN can be characterized by a power-law model with a photon-index of $\\sim$2 and a 0.3-10 keV intrinsic luminosity of 3.3$\\times10^{40}{\\rm~ergs~s^{-1}}$. Linking the X-ray spectral properties of the AGN with the current upper limit at radio frequencies indicates a spectral flattening of the AGN at low frequencies.  The 0.5-2 keV unresolved X-ray emission is found to closely trace the near-IR emission in the disk region, and the X-ray to near-IR luminosity ratio is consistent with that inferred from galactic LMXBs. These two facts together indicate that the bulk of the unresolved emission is produced by the old stellar population of the galaxy, predominantly LMXBs.  A few extraplanar diffuse X-ray features are present in addition to the collective emission from discrete sources traced by the near-IR light. These features can be explained by the presence of hot gas, which can be spectrally characterized by a two-temperature plasma with $kT$ $\\sim$0.08 keV and $\\sim$0.8 keV. The total X-ray luminosity of hot gas is at least an order of magnitude lower than that predicted by current simulations of IGM accretion based on disk galaxy formation models. Thus the extraplanar features are very unlikely to result from IGM accretion.  Instead, morphologically most of these extraplanar features have extended radio counterparts, which are believed to arise from disk-related events. Also, energetically the extraplanar features can be generated by either supernova explosions or the AGN, the latter possibly related to the bubbles above and below the nucleus. Therefore, we conclude that the extraplanar features are most likely formed from outflows from the galactic disk.  Our observation suggests that a proper inclusion of galactic feedback is essential, not only to understanding galaxy formation, but also to its continued evolution.  NGC~2613 and galaxies like it provide nearby laboratories that may help to understand the over-cooling problem existing in current galaxy formation simulations.  For J. I., this work has been supported by the Natural Sciences and Engineering Research Council of Canada."
    },
    "0606/astro-ph0606069_arXiv.txt": {
        "abstract": "{We give a brief review of the known effects of a dynamical vacuum cosmological component, the dark energy, on the anisotropies of the cosmic microwave background (CMB). We distinguish between a ``classic\" class of observables, used so far to constrain the average of the dark energy abundance in the redshift interval in which it is relevant for acceleration, and a ``modern\" class, aiming at the measurement of its differential redshift behavior. \\\\ We show that the gravitationally lensed CMB belongs to the second class, as it can give a measure of the dark energy abundance at the time of equality with matter, occurring at about redshift 0.5. Indeed, the dark energy abundance at that epoch influences directly the lensing strength, which is injected at about the same time, if the source is the CMB. We illustrate this effect focusing on the curl (BB) component of CMB polarization, which is dominated by lensing on arcminute angular scales. An increasing dark energy abundance at the time of equality with matter, parameterized by a rising first order redshift derivative of its equation of state today, makes the BB power dropping with respect to a pure $\\Lambda$CDM cosmology, keeping the other cosmological parameters and primordial amplitude fixed. We briefly comment on the forthcoming probes which might measure the lensing power on CMB.}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606633_arXiv.txt": {
        "abstract": "In this Letter, we study a localized stellar overdensity in the constellation of Ursa Major, first identified in Sloan Digital Sky Survey (SDSS) data and subsequently followed up with Subaru imaging. Its color-magnitude diagram (CMD) shows a well-defined sub-giant branch, main sequence and turn-off, from which we estimate a distance of $\\sim 30$ kpc and a projected size of $\\sim 250 \\times 125$ pc. The CMD suggests a composite population with some range in metallicity and/or age. Based on its extent and stellar population, we argue that this is a previously unknown satellite galaxy of the Milky Way, hereby named Ursa Major II (UMa II) after its constellation.  Using SDSS data, we find an absolute magnitude of $M_V \\sim -3.8$, which would make it the faintest known satellite galaxy. UMa II's isophotes are irregular and distorted with evidence for multiple concentrations; this suggests that the satellite is in the process of disruption. ",
        "introduction": "Numerical simulations in the hierarchical cold dark matter paradigm of galaxy formation generally predict 1 to 2 orders of magnitude more satellite halos in the present day Local Group than the number of dwarf galaxies thus far observed \\citep[e.g.,][]{Mo99,Kl99,Be02}. Numerous solutions have been proposed for this ``missing satellite'' problem. For example, star formation may be inhibited in low-mass systems \\citep[e.g.,][]{Bu01,So02}, or the known satellites may represent a higher mass regime of the satellite initial mass function~\\citep[e.g.,][]{Str02,Kr04}.   However, it has become increasingly clear over the last two years that the census of Local Group satellites is seriously incomplete. Data from the Sloan Digital Sky Survey \\citep[SDSS;][]{Yo00} have revealed five new nearby dwarf spheroidals (dSphs) in quick succession: Andromeda IX~\\citep{Zu04}, Ursa Major~\\citep{Wi05a}, Andromeda X~\\citep{Zu06b}, Canes Venatici~\\citep{Zu06a} and Bo{\\\"o}tes~\\citep{Be06b}. All five galaxies were detected as stellar overdensities. The purpose of this Letter is to study another prominent stellar overdensity in SDSS Data Release 4 \\citep{Am06}. \\citet{Gr06} independently called attention to it, stating that it may be a ``new globular cluster or dwarf spheroidal''. Here we provide evidence from SDSS and subsequent deeper Subaru imaging for its interpretation as a dwarf spheroidal galaxy, the thirteenth around the Milky Way, with the proposed name Ursa Major II (UMa II).  \\begin{figure}[t] \\begin{center} \\epsscale{1.0} \\plotone{uma2_sdssfig.ps} \\caption{The UMa II Dwarf as seen by SDSS: {\\it Upper left:} Combined SDSS $g,r,i$ images of a $1.2^{\\circ} \\times 1.2^{\\circ}$ field centered on the overdensity (J2000 08:51:30 +63:07:48). $\\Delta \\alpha$ and $\\Delta \\delta$ are the relative offsets in right ascension and declination, measured in degrees of arc. The dashed lines indicate the two pointings observed with Subaru (see \\S 2). {\\it Middle left:} The spatial distribution of all blue objects ($g - i < 0.5$) classified as stars in the same area. {\\it Lower left:} Binned spatial density of all blue stellar objects, together with a dotted box that covers most of the object and a dotted annulus used to define the background. {\\it Upper right:} CMD of all stellar objects within the dotted box; note the clear main sequence turn-off and subgiant branch, along with hints of horizontal and red giant branches, even without removal of field contamination. {\\it Middle right:} Control CMD of field stars from the dotted annulus. {\\it Lower right:} A color-magnitude density plot (Hess diagram), showing the CMD of the box minus the control CMD, normalized to the number of stars in each CMD. All photometric data were corrected for Galactic foreground extinction using \\citet{schl98}. \\label{fig:uma_disc}} \\end{center} \\vskip -1.1cm \\end{figure} \\begin{figure} \\begin{center} \\epsscale{1.0} \\plotone{uma2_subarufig.ps} \\caption{The UMa II Dwarf as seen by Subaru: {\\it Upper left:} CMD of the central region of UMa II (see dashed box in the upper right panel of Figure~\\ref{fig:uma_disc}), constructed with Subaru $g,i$ data. The solid gray line graphically indicates the color-magnitude selection criteria used to construct the contour plots in the lower panels. The error bars on the left show the typical photometric errors at the $i-$band magnitude indicated. {\\it Upper right:} A color-magnitude density plot (Hess diagram), showing the CMD of the box minus a control-field CMD, normalized to the number of stars in each CMD. {\\it Lower left:} Isodensity contours of the stars selected from the Subaru data by the gray box in the upper left panel. The plotted contour levels are 1, 2, 3, 5, 7 and 9$\\sigma$ above the background level. $\\Delta \\alpha$ and $\\Delta \\delta$ are measured in degrees of arc.  {\\it Lower right:} Isodensity contours using SDSS data for comparison, with levels of 2, 3, 5, 7 and 9$\\sigma$ above the background plotted. Note that the three blobs appear in both panels.} \\label{fig:uma_disc2} \\end{center} \\vskip -1.2cm \\end{figure} \\begin{figure}[t] \\begin{center} \\epsscale{1.0} \\plotone{uma2_cmdsclust.ps} \\caption{ {\\it Upper left:} Composite CMD of the central region of UMa II (dashed box in the upper right panel of Figure~\\ref{fig:uma_disc}), with photometry of both SDSS and Subaru stars plotted (black circles and gray dots, respectively). Duplicate detections (i.e., detections of the same star in both sets of photometric data) have {\\em not} been removed. The error bars on the left show the typical photometric errors for each dataset at the $i-$band magnitude indicated. {\\it Upper right:} The same composite CMD, with all stars shown as gray dots, and Padova isochrones \\citep{Gi04} overplotted for (left to right) [Fe/H]$= -2.3$/12 Gyr, [Fe/H]$= -1.3$/12 Gyr and [Fe/H]$= -0.7$/10 Gyr, shifted to a distance modulus of 17.5.  {\\it Lower left:} Subaru $g-$band image of the apparent central cluster of UMa II. The image spans $3\\arcmin \\times 3\\arcmin$. The curved shadow to the left is scattered light from a nearby bright star.  {\\it Lower right:} Composite CMD of the central cluster region shown in the lower left panel, with SDSS and Subaru photometry plotted as gray and black dots, respectively. The three isochrones from the upper right panel are also overplotted; the middle isochrone ([Fe/H]$= -1.3$/12 Gyr) appears to be a reasonably good fit to the data, although even in this small region the main sequence is broader than might be expected from simple photometric errors (see upper left panel).} \\label{fig:uma_cmdscluster} \\end{center} \\vskip -0.5cm \\end{figure} \\begin{figure} \\begin{center} \\epsscale{0.95} \\plotone{uma2_complexa.ps} \\caption{The locations of UMa II and Complex A, together with the great circle of the Orphan Stream. The distance estimate to the Orphan Stream is comparable to that of UMa II, but Complex A is believed to lie much closer.  The gray scale shows the density of SDSS stars satisfying $g - r < 0.4$ and $20 < r < 22.5$. The inset image is a blow-up of the area immediately around UMa II, showing its long axis almost aligned with constant Galactic longitude. The column density contours for Complex A are taken from \\citet{Wa01}, while the great circle of the Orphan Stream is from \\citet{Be06c}.} \\label{fig:busy} \\end{center} \\vskip -1.2cm \\end{figure}  \\begin{deluxetable}{lc} \\tablecaption{Properties of the Ursa Major II Dwarf \\label{tbl:pars}} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter\\tablenotemark{a}} & {~~~ } } \\startdata Coordinates (J2000) & \\coords \\\\ Coordinates (Galactic) & $\\ell = 152.5^\\circ$, $b= 37.4^\\circ$ \\\\ Position Angle & $95^\\circ$\\\\ Ellipticity & $0.5$\\\\ Central Extinction, A$_{\\rm V}$ & $0\\fm29$\\\\ V$_{\\rm tot}$ & $14\\fm3\\pm0\\fm5$\\\\ (m$-$M)$_0$ & $17\\fm5\\pm0\\fm3$\\\\ M$_{\\rm tot,V}$ & $-3\\fm8\\pm0\\fm6$\\enddata \\tablenotetext{a}{Integrated magnitudes are corrected for the Galactic foreground reddening reported by \\citet{schl98}} \\label{tab:struct} \\end{deluxetable} ",
        "conclusions": "We have identified a new companion to the Milky Way galaxy in the constellation Ursa Major.  Based on its size, structure and stellar population, we argue that it is a new dwarf spheroidal galaxy and name it UMa II. It has a distance of $\\sim 30$ kpc and an absolute magnitude of $M_V \\sim -3.8$. Its color-magnitude diagram shows an upper main sequence, turn-off and sub-giant branch, as well as hints of red giant and horizontal branches.  UMa II has a bright central concentration, together with two further clumps. The irregular nature of the object suggests that it may have undergone disruption.  This is the fourth Milky Way dSph discovered by SDSS in little over a year. Together with the earlier discoveries of Ursa Major I, Canes Venatici and Bo{\\\"o}tes, this underscores how incomplete our current census actually is.  As SDSS covers only $\\sim 1/4$ of the celestial sphere, crude scaling arguments would suggest that there are tens of missing Milky Way dSphs.  If true, this would go some way toward resolving the missing satellite issue.  \\vskip-0.5cm"
    },
    "0606/astro-ph0606405_arXiv.txt": {
        "abstract": "{ We present new observations of 470 stars using the Fibre Large Array Multi-Element Spectrograph (FLAMES) instrument in fields centered on the clusters NGC\\,330 and NGC\\,346 in the Small Magellanic Cloud (SMC), and NGC\\,2004 and the N11 region in the Large Magellanic Cloud (LMC).  A further 14 stars were observed in the N11 and NGC\\,330 fields using the Ultraviolet and Visual Echelle Spectrograph (UVES) for a separate programme.  Spectral classifications and stellar radial velocities are given for each target, with careful attention to checks for binarity.  In particular, we have investigated previously unexplored regions around the central LH9/LH10 complex of N11, finding $\\sim$25 new O-type stars from our spectroscopy. We have observed a relatively large number of Be-type stars that display permitted Fe~{\\scriptsize II} emission lines.  These are primarily not in the cluster cores and appear to be associated with classical Be-type stars, rather than pre main-sequence objects.  The presence of the Fe~{\\scriptsize II} emission, as compared to the equivalent width of H$\\alpha$, is not obviously dependent on metallicity.  We have also explored the relative fraction of Be- to normal B-type stars in the field-regions near to NGC\\,330 and NGC\\,2004, finding no strong evidence of a trend with metallicity when compared to Galactic results. A consequence of service observations is that we have reasonable time-sampling in three of our FLAMES fields.  We find lower limits to the binary fraction of O- and early B-type stars of 23 to 36\\%.  One of our targets (NGC\\,346-013) is especially interesting with a massive, apparently hotter, less luminous secondary component. } ",
        "introduction": "As part of a European Southern Observatory (ESO) Large Programme we have completed a new spectroscopic survey of massive stars in fields centered on open clusters in the Large and Small Magellanic Clouds (LMC and SMC respectively) and the Galaxy.  The survey has employed the Fibre Large Array Multi-Element Spectrograph (FLAMES) instrument at the Very Large Telescope (VLT), that provides high-resolution ($R\\sim$20,000) multi-object spectroscopy over a 25' diameter field-of-view.  The scientific motivations for the survey, and the observational information for the three Galactic clusters (NGC\\,3293, NGC\\,4755, and NGC\\,6611), were presented by \\citet[][hereafter Paper~I]{mwpaper}.  In this paper we present the FLAMES observations in the Magellanic Clouds.  The material presented is largely a discussion of the spectral classifications and radial velocities of each star, and provides a consistent and thorough overview of what is a particularly large dataset.  In parallel to this catalogue, subsets of the sample are now being analysed by different groups.  The sources of photometry and astrometry used for target selection are given in Section~\\ref{targets}, followed by details of the observations in Section ~\\ref{obs}, and then a discussion of the observed sample.  Two FLAMES pointings were observed in each of the Clouds, centered on NGC\\,346 and NGC\\,330 in the SMC, and on NGC\\,2004 and the N11 region \\citep{hen56} in the LMC.  NGC\\,346 is a young cluster with an age in the range of 1 to 3$\\times$10$^6$ yrs \\citep{k89,wal00,jc03,mfast2}, that has clearly undergone prodigious star formation.  It is also the largest H~\\2 region in the SMC.  The best source of spectroscopic information in NGC\\,346 is the study by \\citet[][ hereafter MPG]{mpg}, who found as many O-type stars in the cluster as were known in the rest of the SMC at that time.  High-resolution optical spectra of five of the O-type stars from MPG were presented by \\citet[][together with AzV~220 that is also within the FLAMES field-of-view]{wal00}.  These were analysed by \\citet{jc03}, in conjunction with ultraviolet data, to derive physical parameters.  N11 is also a relatively young region and includes the OB associations LH9, LH10 and LH13 \\citep{lh}, that are of interest in the context of sequential star-formation, see \\citet{w92}, \\citet[][ hereafter P92]{p92}, and \\citet{b03}.  P92 illustrated how rich the region is in terms of massive stars, presenting observations of 43 O-type stars in LH9 and LH10, with three O3-type stars found in LH10.  These O3 stars were considered by \\citet{w02} in their extension of the MK classification scheme to include the new O2 subtype, with one of the stars from P92 reclassified as O2-type.  The FLAMES observations in N11 presented an opportunity to obtain good-quality spectroscopy of a large number of known O-type stars, whilst also exploring the spectral content of this highly-structured and dynamic region.  NGC\\,330 and NGC\\,2004 are older, more centrally condensed clusters. NGC\\,330 in particular has been the focus of much attention in recent years.  \\citet{f72} presented H$\\alpha$ spectroscopy of 18 stars in the cluster, and noted: `It is also an object of considerable importance in discussion of possible differences between stars of the same age in the SMC and in the Galaxy'.  The community has clearly taken his words to heart -- in the past 15 years there have been numerous studies in the cluster, most of which were concerned with the large population of Be-type stars therein, namely \\citet{grb92}, \\citet{l93}, \\citet{grb96}, \\citet{maz96}, \\citet{sk98}, \\citet{sk99}, \\citet{m99}, and \\citet{l03}.   The paper from these with the widest implications is that from \\citet{m99}, who compared the fraction of Be stars (relative to all B-type stars) in a total of 21 clusters in the SMC, LMC and the Galaxy.  The fraction of Be stars appears to increase with decreasing metallicity, although their study was limited to only one cluster (NGC\\,330) in the SMC.  This trend led \\citeauthor{m99} to advance the possibility of faster rotation rates at lower metallicities -- one of the key scientific motivations that prompted this FLAMES survey.  In contrast to NGC\\,330, with the exception of abundance analyses of a few B-type stars by \\citet{korn02,korn06}, relatively little was known about the spectroscopic content of NGC\\,2004 until recently.  A new survey by \\citet{mhf}, also with FLAMES, has observed part of the field population near NGC\\,2004 using the lower-resolution mode of the Giraffe spectrograph (and with a different field-centre to ours). \\citeauthor{mhf} concluded that the Be-fraction in their LMC field is not significantly different to that seen in the Galaxy.  The FLAMES observations for the current survey were obtained in service mode and so span a wide range of observational epochs, giving reasonable time-sampling for the detection of binaries.  There are surprisingly few multi-epoch, multi-object spectroscopic studies of stellar clusters in the literature in this respect, with one such study in 30 Doradus summarized by \\citet{bsm98}.  Placing a lower limit on the binary fraction in dense star-forming regions such as NGC\\,346 and N11, combined with stellar rotation rates, will help to provide useful constraints in the context of star formation and the initial mass function. ",
        "conclusions": "In Table~\\ref{overview} we give an overview of the entire sample in the FLAMES survey, incorporating the LMC and SMC observations reported here, with the Galactic data from Paper~I.  Analysis of many of these data is now well advanced by different groups.  Here, with the benefit of a broad view of the whole sample, we discuss some of the more general features of the survey.  \\subsection{H-R diagrams for each of the fields} In Figure~\\ref{hrds} we show Hertzsprung-Russell (H-R diagrams) for each of our LMC and SMC fields.  These have been compiled by employing various published calibrations -- clearly the detailed studies of different subsets of the survey will yield precise determinations of temperatures and luminosities, but we take the opportunity now to show the full extent of the LMC/SMC stars in the H-R diagram.  Objects listed as likely foreground objects are plotted as open circles, likely binaries as '+`, and emission-line objects as triangles.  For illustrative purposes the evolutionary tracks shown are from \\citet{s93} for the LMC targets and from \\citet{char93} for the SMC.  Temperatures were adopted on the basis of spectral type, luminosity and metallicity from \\citet[][O-type stars]{mfast2}, \\citet[][B-type supergiants]{clw06}, \\citet[][A-type superigants]{eh03}, and \\citet[][for other types]{sk82}.  Objects that luminosity classes of III or V were assigned temperatures from calibrations for dwarfs; more luminous stars (i.e. II, Ib, Iab and Ia) adopted temperatures from calibrations of supergiants.  With no metallicity-dependent temperature scale in the literature for early B-type supergiants, temperatures were taken from the Galactic results \\citep{clw06}, which were found to be in good agreement with those from analysis of individual stars in the SMC \\citep{tl04,tl05}.  Luminosities were calculated using intrinsic colours from \\citet{f70}, extrapolating or interpolating where required; bolometric corrections were calculated using the relations from \\citet{vacca} for the earliest types, and from \\citet{balbc} for the cooler stars; the ratio of total to selective extinction ($R$) in the LMC was taken as 3.1 \\citep[e.g.][]{h83} and as 2.7 in the SMC \\citep{b85}; distance moduli were taken as 18.5 to the LMC \\citep[e.g.][]{gibson} and 18.9 to the SMC \\citep{hhh}.  Figure~\\ref{hrds} highlights the predominantly less-massive populations observed in the two older clusters (i.e. NGC\\,2004 and NGC\\,330), particularly when compared to N11.  The N11 observations also sample a more luminous, more massive population than those in NGC\\,346.  This is, in part, influenced by the fact that some of the O-type stars observed by \\citet{wal00} were explicitly avoided in the FLAMES survey as state-of-the-art analyses have already been presented in the literature \\citep[e.g.][]{jc03}.  \\begin{figure*} \\begin{center} \\includegraphics[angle=180,width=12cm]{4988fig12.eps} \\vspace*{0.5cm}\\caption{H-R diagrams for the four FLAMES fields in the Magallanic Clouds. Open circles denote foreground stars, open triangles are emission-line objects, and likely binaries are indicated with crosses.  The evolutionary tracks are taken from \\citet{s93} for N11 and NGC\\,2004, and from \\citet{char93} for NGC\\,330 and NGC\\,346.}\\label{hrds} \\end{center} \\end{figure*}  \\begin{table*} \\caption[]{Overview of the distribution of spectral types of the FLAMES survey, including the Galactic clusters from Paper~I.}\\label{overview} \\begin{center} \\begin{tabular}{lp{1.5cm}p{1.5cm}p{1.5cm}p{1.5cm}p{1.5cm}} \\hline Field & O & Early-B & Late-B & AFG & Total \\\\ &   & (B0-3)  & (B5-9) &     &       \\\\ \\hline MW:\\o NGC\\,3293  & $-$ & 48 & 51 & 27 & 126 \\\\ MW:\\o NGC\\,4755  & $-$ & 54 & 44 & 10 & 108 \\\\ MW:\\o NGC\\,6611  & 13  & 28 & 12 & 32 & 85 \\\\ SMC: NGC\\,330     & 6 & 98 & 11 & 10 & 125 \\\\ SMC: NGC\\,346     & 19 & 84 & 2 & 11 & 116 \\\\ LMC: NGC\\,2004   & 4 ($+$ 1 WR) & 101 & 6 & 7 & 119 \\\\ LMC: N11              & 44 & 76 & $-$ & 4 & 124 \\\\ \\hline Total            & 87 & 489 & 126 & 101 & 803 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\subsection{Exploring new regions of N11} As mentioned earlier, part of our intention in the FLAMES observations was to explore some of the less dense, apparently star-forming regions in N11.  In addition to the discovery of the O2.5-type star to the north of LH10, the survey has revealed a large number of O-type stars in the regions surrounding LH9 and LH10.  In Figure~\\ref{n11_mess} we show the $V$-band WFI image used for target selection in N11, with the O- and B-type stars marked in different colours; LH9, LH10 and LH13 are also identified in the image.  Note the O-type stars to the south of LH9, newly discovered by this survey.  These include N11-020 [classified as O5 I(n)fp] on the northern edge of the N11F region (cf. Figure~\\ref{fchart_n11}).  Furthermore, N11-004 [O9.7 Ib] is the bright star in N11G, and N11-058 [O5.5 V((f))] is one of two visually bright stars in N11I, the other being N11-102 [B0.2~V].  Both N11G and N11I appear as small `bubbles' in near-IR images, presumably driven by the ionization and/or the stellar winds from these stars.  We have also observed N11-029 [OC9.7 Ib] in N11H, which appears radially smaller than N11G and N11I.  The densest regions in N11, i.e. N11B (LH10) and N11C (LH13) have been demonstrated by other authors \\citep{p92,w92,hm00,b03} to have rich, young populations of early-type stars.  The general consensus is that this region is a two-stage starburst, with the evolution of LH9 triggering star-formation around it.  For the first time the FLAMES survey has observed the regions to the south and west of LH9 -- whilst much smaller in size and stellar content, we find newly discovered O-type stars as the likely source of ionization and dynamic energy for the observed nebulae.  To place this region in context, the `bigger picture' is dramatically illustrated by the cover image of Edition \\#80 of the National Optical Astronomy Observatory/National Solar Observatory (NOAO/NSO) Newsletter.  This features an image of N11 from the Magellanic Clouds Emission Line Survey (credited to Drs.~S. Points, C. Smith, and M. Hanna).  In addition to N12, N13, and N14, the image includes the significantly extended shell of gas that constitutes N10 from \\citet{hen56}.    \\subsection{`Blue stragglers' in NGC\\,330}  Blue stragglers were reported in NGC\\,330 by \\citet{l93}, with \\citet{grb96} arguing that they might be a product of binary evolution. Blue stragglers are thought to arise from either mass-exchange or stellar mergers \\citep[e.g.][]{pm94}, with recent studies of globular clusters suggesting that both channels play a role \\citep{d04}. Six of our targets in the NGC\\,330 observations are late O-type stars, three of which are less than 5$'$ from the centre of the cluster: NGC\\,330-023 (see Section \\ref{330_023}), NGC\\,330-049, and NGC\\,330-123.  In particular, NGC\\,330-123 (R74-B18) is classified here as O9.5 V, cf.  B0~Ve from \\citet{l93} and O9~III/Ve from \\citet[][who were quoting unpublished types from Lennon]{grb96}.  Narrow (and weak) emission is seen in the core of the H$\\alpha$ Balmer line in the UVES spectrum of NGC\\,330-123, but is accompanied by [N~\\2] emission so the origins are likely nebular.  Indeed, \\citet{sk98} note that there is an elliptical region of diffuse nebular emission centred on this star, it would appear that NGC\\,330-123 is the ionization source of this nebulosity.  Interestingly, the radial velocity of NGC\\,330-123 is 177 km/s (with standard deviation, $\\sigma$ = 5), i.e. different from the systemic velocity of the cluster which is $\\sim$155 km/s \\citep[e.g.][]{l03}.  Binarity is a possible explanation of this difference, but additional spectra at other epochs (October 1992 and October 1995), though of lower resolution, are in good agreement with the FLAMES data. If this star is coeval with the cluster then its current radial velocity might indicate a previous, though relatively recent ejection event.  Its radial velocity however is similar to the three more distant O-type stars in the NGC\\,330 field: NGC\\,330-013, NGC\\,330-046 and NGC\\,330-052, which have $v_r =$ 176, 177 and 166~\\kms\\/ respectively.  Two of these, NGC\\,330-013 and NGC\\,330-052, are found to be helium rich \\citep{rmsmc}, whilst NGC\\,330-123 has also been reported as helium rich by \\citet[][note that this paper erroneously refers to R74-A01 as helium rich, when it should actually refer to R74-B18, i.e. NGC\\,330-123]{l93}.  All four stars could perhaps be considered as members of the general field population which, from the H-R diagram (Figure~\\ref{hrds}), can be seen to extend well beyond the notional cluster turn-off -- as represented by the $\\sim$20 stars with spectral types earlier than B1. We note that the majority of these stars have radial velocities rather different from the cluster radial velocity, suggesting that the probable number of true blue stragglers in this field belonging to NGC\\,330 is small.  Both NGC\\,330-023 and NGC\\,330-049 are late O-type stars with velocities consistent with the cluster \\citep[cf.][]{l03}, although the pecularities of NGC\\,330-023 have already been discussed in Section~\\ref{330_023}.  At respective radial distances of 2.3$'$ and 4.5$'$ they are not immediately proximate to the cluster core, but are perhaps outer members of the cluster and therefore could be true blue stragglers.  The situation for NGC\\,2004 is qualitatively similar to that of NGC\\,330.  The Wolf-Rayet and O2-3 stars stand out as potential blue stragglers but are more likely field stars, though the presence of an O2-3 star in the field is unusual in itself.  One further candidate blue straggler is NGC\\,2004-019, which is close to the core, with $r_{\\rm d} =$ 1.5$'$.  \\subsection{Incidence of Be-stars}\\label{discussbe}  The numbers of Be-type stars in each FLAMES pointing are summarised in Table \\ref{overview2}.  The observed field-population of B- and Be-type stars in our NGC\\,330 and NGC\\,2004 pointings should be relatively unaffected by specific selection effects (Section \\ref{seffects}), aside from the slightly different ($\\sim$0.5$^{\\rm m}$) faint cut-off of the observations.  The Be-stars in Table~\\ref{overview2} include some in the outer regions of the clusters themselves.  As an experiment, we considered the B-type stars with $r_d > {\\rm 2}'$ in NGC\\,330 and NGC\\,2004 to sample the field population away from the main clusters.  All stars in the range B0-3 were counted and then the relative percentage of those that are seen to be Be-type stars was found.  In NGC\\,2004 this ratio for the field population (i.e. Be / [Be$+$B]) was 16\\% (13/81), compared to 23\\% (18/77) in NGC\\,330.  It is difficult to quantify the uncertainties in these results.  Is there any effect in terms of absolute magnitude?  As a further experiment (and still with the $r_d > {\\rm 2}'$ condition) we considered the early B-type stars in our NGC\\,330 data with $V \\leq$~16.03.  Taking the difference of the distance moduli of the Clouds as 0.4$^{\\rm m}$, and allowing for the fact that the `typical' extinction toward the LMC is $\\sim$0.1$^{\\rm m}$ greater than the SMC \\citep[][]{m95}, this notionally imposes a similar cut-off as that in the NGC\\,2004 data.  The Be-fraction for the NGC\\,330 field population remains robust at 24\\% (11 Be/45 Be+B).  Our results are in reasonable agreement with those from \\citet{sk99} and, for NGC\\,2004, match those of \\citet{mhf}. Both \\citeauthor{sk99} and \\citeauthor{mhf} advanced their results as a lack of evidence for a strong dependence of the Be-fraction with metallicity when compared to the general result for the Milky Way of $\\sim$17-20\\% \\citep{zb}.  This is in contrast to the results for clusters from \\citet{m99}. The statistics in our NGC\\,346 and N11 data do not provide a meaningful test of the relative numbers of Be stars at different metallicities.  Aside from a likely mix of field and cluster populations in the NGC\\,346 pointing, the N11 region is even less suited given its complex star-formation history.  \\subsection{Binaries and the binary fraction}  With the time-sampling provided by the service-mode observations we are reasonably sensitive to detection of binarity in our SMC and LMC targets.  In Appendix~\\ref{mjd} we give full details of the observational epochs of the FLAMES spectra.  The observations in both N11 and NGC\\,2004 spanned a total of 57 days, and those in NGC\\,346 covered 84 days.  The time-coverage of the majority of the NGC\\,330 data is not as extensive as for the other fields, covering 10 days.  The new HR03 (\\lam4124) data offer some extra information in this regard (e.g. Section~\\ref{330_023}).  However, given the relatively poor signal-to-noise in the 2003 data it is difficult to compare measured velocities meaningfully -- from these comparisons a number of stars are flagged as having potentially variable velocities.  The FLAMES data are sufficient to derive periods for some of the newly discovered systems, the most interesting of which will receive a detailed treatment in a future study.  The incidence of binaries in summarised in Table~\\ref{overview2}.  We stress that stars considered as multiple in this discussion are those listed in Tables~\\ref{346}, \\ref{330}, \\ref{lh910}, and \\ref{2004} as `Binary' (be they SB1, SB2 or not specified).  Stars suggested to perhaps demonstrate $v_r$ variations are not considered further, pending follow-up.  As such, the percentages in Table~\\ref{overview} serve as strong, {\\it lower} limits on the binary fraction in our fields.  In young, dense clusters multiplicity seems a common, almost ubiquitous feature.  \\citet{gm01} found a significant binary fraction (79\\%) in the very young Galactic cluster NGC\\,6231, and \\citet{wbp99} and \\citet{p99} highlighted the high degree of multiplicity found in the massive members of the Orion Nebula, suggesting a different mode of star formation to that at lower masses (in which the binary fraction is lower).  The formation mechanism of massive stars is still a point of significant debate.  In their discussion of the competing scenarios of massive-star formation from accretion versus stellar mergers, \\citet{bz05} suggest that the multiple star fraction will be larger if merging dominates.  \\citeauthor{bz05} also suggest that mergers may be the dominant process in ultradense regions such as 30 Doradus; perhaps the high binary fraction of O- and early B-type stars in N11 is indicative of this mode of star-formation, remembering that 36\\% is a solid, lower limit obtained from a programme that was not optimised for detection of binaries.  By comparison, the low binary fraction in the NGC\\,330 targets is somewhat puzzling.  Although fewer binaries are generally found in the field population \\citep[e.g.][]{mgh98}, the NGC\\,330 fraction is significantly lower than that for NGC\\,2004, which similarly samples the field (see Figures~\\ref{fchart_330} and \\ref{fchart_2004}).  We speculate that, whilst the new HR03 data added an additional epoch to the time-sampling for the NGC\\,330 targets, it is still not as thorough as for the other fields.   \\begin{table*} \\caption[]{Overview of the number of Be-type stars and binaries in the LMC and SMC sample}\\label{overview2} \\begin{center} \\begin{tabular}{lp{1.2cm}p{1.2cm}p{1.2cm}p{1.2cm}p{1.2cm}} \\hline Field            & O ($+$ Oe) & Early-B & Be & Total & Binary \\\\ &   & (B0-3)  &    & (O$+$Early-B)   & Fraction \\\\ \\hline SMC: NGC\\,330    & 6 & 76 & 22 & 104 &  \\\\ {\\it Binary}     & $-$ & {\\it 3} & {\\it 1} & {\\it 4} & {\\it 4\\%} \\\\ &&&&&\\\\ SMC: NGC\\,346    & 19 & 59 & 25 & 103 & \\\\ {\\it Binary}     & {\\it 4} & {\\it 19} & {\\it 4} & {\\it 27} & {\\it 26\\%} \\\\ &&&&&\\\\ LMC: NGC\\,2004   & 4 & 83 & 18 & 105 & \\\\ {\\it Binary}     & {\\it 1} & {\\it 21} & {\\it 2} & 24 & {\\it 23\\%} \\\\ &&&&&\\\\ LMC: N11         & 44 & 68 & 8 & 120 & \\\\ {\\it Binary}     & {\\it 19} & {\\it 21} & {\\it 3} & {\\it 43} & {\\it 36\\%} \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "0606/astro-ph0606627_arXiv.txt": {
        "abstract": "We analyse observations, spanning 15 years, dedicated to the extreme emission-line object \\hen. Our photometric data indicate that the luminosity of the star undergoes marked variations with a peak-to-peak amplitude of 0.65~mag. These variations are recurrent, with a period of 16.093$\\pm$0.005~d. The spectrum of \\hen\\ is peculiar with many different lines (H\\,{\\sc{i}}, \\he, \\feii,...) showing P Cygni profiles. The line profiles are apparently changing in harmony with the photometry. The spectrum also contains \\oiii\\ lines that display a saddle profile topped by three peaks, with a maximum separation of about 600~\\kms. \\hen\\ is most likely an evolved luminous object suffering from mass ejection events and maybe belonging to a binary system. ",
        "introduction": "In the middle of the last century, dedicated \\ha\\ surveys have identified many emission-line stars. A large number of these stars were at first classified as Wolf-Rayet (WR) stars but, on second sight, it appears that their spectrum rarely presents the typical WR characteristics. Although these peculiar objects may represent keys for understanding stellar evolution, only a few of them, generally the brightest and least reddened ones, were analysed in detail.  \\hen, also known as WRAY\\,15-285, MR\\,14, SS73\\,10, or Ve\\,6-14 ($8^{\\rm h}48^{\\rm m}45.5^{\\rm s} -46 ^{\\circ}05'09''$, J2000), is such an emission-line object \\citep{vel}. It has only been poorly studied since its detection despite its peculiarities. The star actually displays a strong and broad \\ha\\ line which led to the classification of \\hen\\ as a Wolf-Rayet star \\citep{roberts,henize}, an extreme Be-like object \\citep{sand} and more recently as a B[e] star \\citep{dew}. \\citet{henize} even suggested that \\hen\\ could be a nova on the basis of its non-detection by \\citet{smi}. However, as the star is photometrically variable (see below), it could simply have been below Smith's detection limit at the time of her observation.  The aim of this paper is to present the results of a long-term observing campaign on this object. The paper is organized as follows: the observations are presented in Sect.\\,2, the photometric data are analysed in Sect.\\,3, and the spectral features of \\hen\\ are examined in Sect.\\,4. Finally, we discuss in Sect.\\,5 the nature of the star using the gathered evidence and we conclude in Sect.\\,6. ",
        "conclusions": "Our study shed new light on the poorly known peculiar system \\hen. In our data, this intriguing object displays recurrent photometric variations with a peak-to-peak amplitude of 0.65~mag and a period of 16.09$\\pm$0.01~days. Its spectrum presents many P~Cygni profiles (H\\,{\\sc{i}}, \\he, \\feii,...) and some forbidden lines like \\oiii. The Balmer line profiles vary along with the photometry. The \\oiii\\ profile is very peculiar and contains three sub-peaks: one is associated with the foreground nebula, but the other two, separated by $\\sim$600~\\kms, indicate the presence of an accretion disk or of expanding ejecta close to the star.  The actual nature of \\hen\\ is still difficult to ascertain with our limited sample of observations, but most existing evidence points towards a moderately bright and distant object ($d\\sim8-9$~kpc and $M_V\\sim-5.1$) having probably undergone a mass ejection event and maybe belonging to a binary system."
    },
    "0606/astro-ph0606761_arXiv.txt": {
        "abstract": "We present a survey of spiral arm extinction substructure referred to as feathers in 223 spiral galaxies using HST WFPC2 images.  The sample includes all galaxies in the RC3 catalog with $cz$ \\(<\\) 5000 km s\\(^{-1}\\), B\\(_{T}\\) \\(<\\)~15, {\\it i} \\(<\\)~60\\degr, and types Sa--Sd with well-exposed broadband WFPC2 images.  The detection frequency of delineated, periodic feathers in this sample is 20\\%\\ (45 of 223).  This work is consistent with Lynds (1970), who concluded that feathers are common in prototypical Sc galaxies; we find that feathers are equally common in Sb galaxies.  Sb--Sc galaxies without clear evidence for feathers either had poorer quality images, or flocculent or complex structure.  We did not find clearly defined feathers in any Scd--Sd galaxy.  The probability of detecting feathers was highest (83\\%) for spirals with well-defined primary dust lanes (the lanes which line the inner edge of an arm); well-defined primary dust lanes were only noted in Sab--Sc galaxies.  The detection frequency of feathers was similar in barred and unbarred spirals.  Consistent with earlier work, we find that neighboring feathers tend to have similar shapes and pitch angles.  Well-defined feathers often emerge from the primary dust lane as leading features before they curve to trailing; some are quite elongated, extending into the interarm and merging with other feathers.  OB associations are often found lining feathers, and many feathers transition to the stellar substructures known as spurs (Elmegreen 1980).  We find that feathers are coincident with interarm filaments strikingly revealed in Spitzer 8$\\mu$m images.  Comparison with CO 1-0 maps of NGC 0628 and NGC 5194 from BIMA SONG shows that feathers originate at the primary dust lane coincident with gas surface density peaks.  Contrary to the appearance at 8$\\mu$m, the CO maps show that gas surface density in feathers decreases rapidly with distance from the primary dust lane.  Also, we find that the spacing between feathers decreases with increasing gas surface density; consistent with formation via a gravitational instability. ",
        "introduction": "Spiral arms are not smooth, continuous features.  Typically an arm is composed of many substructures commonly referred to as feathers, spurs, and branches, which give the arm its patchy, divaricate appearance.  This substructure appears associated with much of the star formation in the arm.  \\subsection{Previous Observations of Spiral Arm Substructure}\\label{previous_obs}  Early reports of spiral arm substructure were based on examination of photographic plates such as those used for prints in \\textit{The Hubble Atlas of Galaxies} (Sandage 1961).  Weaver (1970) noted that the spiral arms of nearby galaxies appear clumpy, irregular, and mottled on small scales. In particular, he noted the presence of ``spurs\" (also referred to as branches or twigs), which appear to originate on the {\\it outside} of the arm, with larger pitch angles than the arm itself.  Spurs extend from the outside of the arm into the interarm, and are seen in photographic plates as {\\it stellar} features.  Weaver also noted the presence of {\\it dark} material concentrated along the {\\it inner} edges of spiral arms.  He remarked that the outside regions of arms appear to be made of material drawn out or ``brushed\" from the inner edges, and that this brushed-out structure has a pitch angle typically a factor of two larger than the arm itself.   At the same time, Lynds (1970) reported a detailed study of dark nebulae in 17 late-type spirals using photographic plates from Mt. Wilson and Palomar Observatories.  She commented on the well-known strong dust lanes along the inner edge of the arms, which she termed ``primary dust lanes\" (hereafter abbreviated PDLs), and noted the presence of thin dust lanes with large pitch angles cutting {\\it across} the luminous arms.  She called these features ``feathers\", and pointed out that these extinction features become mostly undetectable outside the luminous arm presumably owing to the absence of a sufficiently bright background.  She also emphasized that bright {\\HII} regions are typically near or embedded in dust lanes.  Later, Piddington (1973) noted that interarm stellar features called spurs by Weaver are often found associated with the extinction features cutting across luminous arms called feathers by Lynds, and suggested that the two are related.   Subsequently, Elmegreen (1980) studied seven spiral galaxies to investigate the properties of spurs, which as noted above are observable as stellar features.  Elmegreen was able to identify two to six well-delineated spurs in each of her galaxies, with lengths ranging from one to five kpc; her spurs are generally located in the outer parts of the luminous disk.  She observed that spurs are always located on the outside of spiral arms, have pitch angles equal to or greater than that of the originating arm, and commonly occur in pairs or groups with the spurs oriented roughly parallel to one another. Elmegreen also determined that on average spurs have pitch angles of roughly 50\\degr\\ with respect to spiral arms, which is comparable to the average pitch angles of the feathers measured by Lynds (1970).  Noting the similarity between the pitch angles of feathers and spurs, Elmegreen added further observational support to Piddington's (1973) suggestion that spurs and feathers have a common origin.   \\subsection{Theoretical Studies of Substructure}\\label{theories}  Motivated to explain the formation of observed spurs, branches, and feathers in spiral galaxies, Balbus (1988, hereafter B88) conducted a local gas dynamical stability analysis of a single-fluid polytropic flow through spiral arm potentials, following the linear evolution of self-gravitating perturbations.  In his analysis, the background gaseous surface density profile representing the arm has an arbitrary spatial form, and the differentially rotating, expanding background flow is consistent with this profile.  Balbus investigated all wavenumber directions in the plane of the disk and modeled the spiral arms as tightly wound, with no magnetic field.  B88 found that there are two preferred directions of growth in spiral arm flow: initial wavefronts roughly along the spiral arm, or perpendicular to it.  The rate of growth in both directions depends on the properties of the underlying flow.  In certain regimes, growth of instabilities leads to fragmentation parallel to the arms, observed as a thickening of the arms.  In other situations, growth of initially leading wavenumbers results in branches, feathers, or spurs.  For wavefronts initially perpendicular to the arm, as the gas moves into the interarm region, the flow expansion and shear (which increases downstream from the arm) shapes and stretches small-scale structure into the familiar large scale trailing `spur' shapes that are observed.   B88 also suggested that a two-dimensional lattice of small-scale structure can develop when the two dominant modes of growth intersect. In \\textit{The Hubble Atlas of Galaxies} (Sandage 1961), B88 found that some barred spiral galaxies, in particular the western arm of NGC 1300, seemed to exhibit a lattice structure of {\\HII} regions.  However, he found no examples of such structure in unbarred galaxies.  Kim and Ostriker (2002, hereafter KO) extended the local models of B88 by including the effects of non-linearity and magnetic fields.  KO conducted local, two-dimensional, time-dependent, magneto-hydrodynamic (MHD) simulations of self-gravitating, differentially rotating, razor thin disks of gas.  Their models followed the formation and fragmentation of ``gaseous spurs'' as the flow passes through spiral arm potentials. They found that local substructures are created via the magneto-Jeans instability (MJI); as the background flow passes through the spiral pattern it is shocked and compressed until it becomes Jeans unstable, at which point gravity, aided by magnetic forces, begins to create alternating compressed and rarefied regions along an arm.  The magnetic effects aid the formation of compressed, self-gravitating complexes because magnetic tension forces oppose the Coriolis forces that would otherwise stabilize the flow, helping to transfer angular momentum out of growing condensations.  KO found that the gaseous spurs fragment into clumps within which star formation could commence.  They suggest that these clumps could be the precursors of bright {\\HII} regions that jut from the outside of spiral arms inside corotation.  From their simulations, KO put together observable statistics of their gaseous spurs.  Most potentially comparable to observations is the spacing of these features along a spiral arm, which they find ranges from 2--5 times the local Jeans wavelength (at the spiral arm density peak) and corresponds to a spacing of approximately 750 pc on average.  KO also proposed that the shape and location of gaseous spurs within a spiral arm may potentially be used observationally to determine the spiral pattern speed of the arm.  Very recently, Kim \\& Ostriker (2006) have extended their thin disk simulations to three dimensions, also making comparison to two-dimensional ``thick disk'' models.  The results they find are overall consistent with the conclusions of KO, with the difference that spur spacings increase by a factor \\(\\sim2\\) due to the dilution of self gravity when the disk thickness increases.  Chakrabarti, Laughlin, and Shu (2003, hereafter CLS) studied the response of a thin, self-gravitating, singular isothermal gaseous disk to rigidly rotating spiral potentials, specifically focusing on the effects of ultraharmonic resonances by choosing parameters that minimize swing amplification. In simulations with low \\(Q_{g}\\), CLS found growth of spiral arm ``branches'' (which they define as trailing bifurcations of the main spiral arms). Long-term simulations with high \\(Q_{g}\\) exhibited the growth of stubby leading structures (referred to as ``spurs'' by CLS; this is however inconsistent with terminology of other authors).   Wada and Koda (2004, hereafter WK) performed two-dimensional, time dependent, global hydrodynamical simulations of a thin, isothermal, non-self-gravitating disk of gas in tightly and loosely wound, rigidly rotating spiral potentials.  WK's model rotation profiles included both differentially rotating and rigidly rotating cases. They found that the spiral shock front is stable when the gas is modeled with a flat rotation curve and unstable when modeled with a rising rotation curve.  In their models, they found the stability of the shock front is also dependent upon the pitch angle of the spiral arms: stable if \\(i\\leq10\\degr\\), unstable if \\(i\\geq10\\degr\\). In the unstable models of WK, strong shocks with arm-interarm density ratio \\(\\sim100\\) become unstable by rippling. WK attribute this ``wiggle'' instability to Kelvin-Helmholtz (K-H) modes involving the strong velocity shear behind the shock. Over time, the instabilities become non-linear forming what WK refer to as ``spurs'' in the interarm regions.  The spurs are quasi-regularly spaced, approximately 100-200 pc apart; the authors do not state, however, how spacing depends on model parameters.   Due to the shape of the rotation curve, the spurs formed in the WK simulations are curved near the arm in the opposite sense (i.e. trailing then leading) to the gaseous spurs produced in the KO simulations and the small-scale structure predicted in B88's analysis.  Dobbs and Bonnell (2006, hereafter DB) used three-dimensional SPH simulations to study the response of isothermal, non-self-gravitating gaseous disk models with varying temperatures to a four armed rigidly-rotating spiral potential.  DB find that the temperature of the disk has a crucial effect on the growth of spiral arm substructure.  In the lower temperature (T \\(<10^{3}\\) K) models of DB, the initially smooth arms become clumpy, and then the clumps are sheared into trailing features as they return to interarm regions.  The shapes of the interarm features found by DB are similar to those of KO, reflecting the flat rotation curve they adopt.  DB use somewhat nonstandard terminology in describing their results; they refer to the portions of interarm extinction features adjoining arms as ``spurs'', and the portions further downstream as ``feathers''.  In their T = 50 K model, which has arm-interarm contrast \\(\\sim50\\), the spacing of DB's spurs are \\(\\approx 700\\) pc;  DB also do not state, however, how spacing depends on model parameters.   \\subsection{Discussion of Substructure Nomenclature}  As is evident from the above summaries, the terms ``spurs'', ``feathers'', and ``branches'' have been used in many ways.  In this paper, we adopt the definitions from the initial, observational papers:  \\textit{feathers} -- thin dust lanes or extinction features that extend outward at a large angle from the \\textit{primary dust lane} (PDL) which lines the inner side of the arm, cutting across the outer bright part of the spiral arm (Lynds 1970).  \\textit{spurs} -- bright chains of OB associations and {\\HII} regions that jut at a large angle from the spiral arm into the interarm (Weaver 1970, Elmegreen 1980).  \\textit{branches} -- divarications of a spiral arm that lend to the overall spiral structure (Elmegreen 1980).  \\subsection{A New Survey of Spiral Arm Feathers} Feathers are of considerable interest because there are both observational (Lynds 1970, Piddington 1973, Elmegreen 1980, Scoville et al. 2001) and theoretical (B88, KO) reasons to associate them with a significant portion of star formation in spiral galaxies.  Also, they may provide information on basic physical conditions in spiral arms, such as the mean gaseous surface density and magnetic field strength (KO).  At least in principle, they may also be used to deduce the spiral pattern speed and details of the gas flow through spiral arms (KO, CLS).  Lastly, they are a striking characteristic of prototypical grand design spiral galaxies (e.g. M51, Beckwith 2005) and may therefore tell us something about the evolutionary and environmental aspects of spiral structure.  A complete description of spiral arms should include a characterization of the frequency and properties of such feathers, and a complete model of spiral structure should explain their dynamical origins.  The goals of this survey include: confirming that feathers are a common feature of spiral galaxies, evaluating their frequency and characteristics, and determining the types of spiral galaxies in which feathers occur (barred or unbarred, early or late-type, grand design or flocculent).  We aim to identify where feathers are located: in the inner or outer disk, on the inside of spiral arms or the outside?  In addition, we investigate the relation of the spacing of feathers to the gas surface density of a galaxy. ",
        "conclusions": "We find extinction feathers in nearly 20\\% of 233 spiral galaxies.  We show that feathers are most common in Sb--Sc galaxies; Sb--Sc galaxies in which we did not detect feathers either had poor quality images, or flocculent or complex structure.  Feathers are rare in Sa galaxies and undetected in Scd--Sd galaxies.  The presence of feathers is closely tied to the existence of a primary dust lane (PDL).  The probability of detecting feathers increases with PDL delineation; the highest being 83\\% for spirals with well-delineated PDLs, within which feathers are ubiquitous.  Characteristically, feathers: (1) are associated with bright star forming regions within spiral arms and interarm regions; (2) extend beyond the outer edge of spiral arms, sometimes far into interarm regions, and merging with the PDL of another arm; (3) transition or evolve into stellar spurs; and (4) often form lattice structures or multiple rows of feathers within a single spiral arm in barred galaxies.  Furthermore, we find that the spacing of feathers is related to the molecular surface density along spiral arms; (1) typically, the distance between feathers increases as the molecular surface density decreases and (2) the majority of feathers originate in regions of higher gas surface density.  The mean separation of feathers is \\(1.7 \\lambda_{Jeans}\\) and \\(10.4 \\lambda_{Jeans}\\) in NGC 0628 and NGC 5194; the value for NGC 0628 is likely an underestimate due to poor signal to noise and a lower limit of {\\HI} data. The above observable characteristics are consistent with models in which feathers are produced by local gravitational instabilities (i.e. Jeans or magneto-Jeans instability) in the gas."
    },
    "0606/astro-ph0606282_arXiv.txt": {
        "abstract": "The Galactic center (GC) provides a unique laboratory for a detailed examination of the interplay between massive star formation and the nuclear environment of our Galaxy. Here, we present an 100 ks {\\sl Chandra} ACIS observation of the Arches and Quintuplet star clusters. We also report on a complementary mapping of the dense molecular gas near the Arches cluster made with the Owens Valley Millimeter Array. We present a catalog of 244 point-like X-ray sources detected in the observation. Their number-flux relation indicates an over-population of relatively bright X-ray sources, which are apparently associated with the clusters. The sources in the core of the Arches and Quintuplet clusters are most likely extreme colliding wind massive star binaries. The diffuse X-ray emission from the core of the Arches cluster has a spectrum showing a 6.7-keV emission line and a surface intensity profile declining steeply with radius, indicating an origin in a cluster wind. In the outer regions near the Arches cluster, the overall diffuse X-ray enhancement demonstrates a bow shock morphology and is prominent in the Fe K$\\alpha$ 6.4-keV line emission with an equivalent width of $\\sim 1.4$ keV. Much of this enhancement may result from an ongoing collision between the cluster and the adjacent molecular cloud, which have a relative velocity $\\gtrsim 120 {\\rm~km^{-1}}$.  The older and less compact Quintuplet cluster contains much weaker X-ray sources and diffuse emission, probably originating from low-mass stellar objects as well as a cluster wind.  However, the overall population of these objects, constrained by the observed {\\sl total} diffuse X-ray luminosities, is substantially smaller than expected for both clusters, if they have normal Miller \\& Scalo initial mass functions. This deficiency of low-mass objects may be a manifestation of the unique star formation environment of the Galactic center, where high-velocity cloud-cloud and cloud-cluster collisions are frequent. ",
        "introduction": "\\label{s:intr}  Nuclear regions of galaxies are the breeding ground of high energy phenomena and processes, which are manifested observationally by active galactic nuclei (AGNs) and star-bursts. Such activities are believed to be important in both regulating galaxy evolution and generating thermal and chemical feedback into the intergalactic medium.  The best site for a detailed study of the activities and their complex interaction with the physically extreme environment in the nuclear regions of galaxies is our own nucleus, only $\\sim 8$ kpc away. We can observe the Galactic center (GC) with a spatial resolution and sensitivity that are factors $\\gtrsim 300$ and $ \\gtrsim 10^5$ better than those available for even nearby nuclear starburst galaxies (e.g., M82 and NGC 253) or AGNs (e.g., M81).  While the super-massive black hole at the dynamic center of the Galaxy is only weakly active at present \\citep{bag01}, much of the current high-energy activity in the GC is due to the presence of the three young massive stellar clusters in the central 50 pc: Arches (with age equal to $2-3 \\times 10^6$ yrs; core or half-mass radius 0.4 pc), Quintuplet ($3-6 \\times 10^6$ yrs; 1.0 pc), and the Central cluster ($3-7 \\times 10^6$ yrs; 0.5 pc) (Figer et al. 1999, 2002, 2004; Stolte et al. 2002; Genzel et al. 2003). These clusters are responsible for about half of the Lyman continuum flux emitted in the central several $10^2$ pc of the GC. Massive stars are also expected to release large amounts of mechanical energy into the GC region in form of fast stellar winds and supernovae, although the actual rate is highly uncertain. This mechanical energy input shapes the surrounding ISM. The present work focuses on the \\xs\\ and \\xq\\ clusters. The GC cluster is less massive and older. Its location in the circum-nuclear region makes its X-ray properties difficult to characterize and will not be dealt with here (however see \\citealt{nay05}).  Both the \\xs\\ and \\xq\\ clusters are known X-ray emitters. Discovered serendipitously at a large off-axis angle ($\\sim 7^\\prime$) in an initial 50 ks {\\sl Chandra} ACIS-I observation \\citep{yus02}, the X-ray emission arising from the \\xs\\ cluster was resolved into discrete and diffuse components. In a later 12 ks ACIS-I observation during a large-scale GC survey \\citep{wan02a}, the X-ray core of the cluster was further resolved into two separate components \\citep{wan03,law04}. The apparent diffuse X-ray component was speculated to arise from the so-called cluster wind \\citep{rag01,yus02}. Because of the high number density of massive stars, their stellar winds collide with each other and can be largely thermalized to form a plasma with an initial temperature of a few times $10^7$ K. The expanding of this plasma may then be considered as a wind from the entire cluster. However, the quality of the previous observations is not adequate for a quantitative test of this scenario. In particular, the diffuse X-ray spectrum shows a distinct 6.4-keV emission line from neutral or weakly ionized irons. The origin of this line is unknown. The X-ray emission from the \\xq\\ cluster is substantially weaker. The detection of a few discrete sources and possible diffuse emission in the region has been reported; but detailed spectral and timing information is not yet available \\citep{wan03,law04}  To further the study of these two clusters and their relationship to the environment, we have obtained an 100 ks \\chandra\\ ACIS-I observation. We have further carried out a complementary imaging study of the molecular gas in the immediate surroundings of the Arches cluster using the six-element Owens Valley Millimeter Array.  With these observations and other multi-wavelength data, we present an in-depth study of various point-like and diffuse X-ray sources in and around the clusters.  In this paper, we assume that the distance to the GC is 8 kpc (hence $1^\\prime =$ 2.5 pc) and quote statistical errors from our X-ray data analysis at the 90\\% confidence level, unless being pointed out otherwise. The solar abundance is in reference to \\citet{and89}; thus the number of iron relative to hydrogen is $4.68 \\times 10^{-5}$,  which is considerably greater than $2.69 \\times 10^{-5}$ in the so-called ISM abundance \\citep{wil02}, for example. ",
        "conclusions": "\\label{s:dis}  The above results show distinctly different X-ray properties between the \\xs\\ and \\xq\\ clusters. The \\xs\\ cluster contains three luminous point-like sources, all of which exhibit the strong 6.7-keV emission line, and two apparently diffuse components with either 6.4-keV or 6.7-keV line emission. The 6.4-keV line-emitting enhancement is strongly elongated, morphologically, tracing the east boundary of the CS cloud's southern extension. These characteristics are absent in the \\xq\\ cluster, in which we detect only weak X-ray sources, plus a very low surface brightness diffuse emission with a hard spectrum. There is also no evidence for any associated CS cloud. In the following, we discuss origins of these various X-ray components, possible causes of the distinct differences in the X-ray properties between the two clusters, and implications of our results.  \\subsection{Galactic Center Environment}\\label{ss:d_envi}  We attempt to  understand the \\xs\\ and \\xq\\ clusters in the context of the unique GC environment. The generally high gas density and pressure, strong gravitational tidal force, and large random and bulk motion  velocities in the GC affect both the formation and evolution of young stellar clusters \\citep{mor93}. Here we concentrate on the potential interplay between the molecular gas and the \\xs\\ cluster.  Is the ``$-$30 km s$^{-1}$ cloud'' and the \\xs\\ cluster physically associated? On one hand, because of their large velocity separation, the two systems would pass each other in only $\\sim 10^4$ yrs if the size of the cloud along the line of sight is comparable to that  projected in the sky and at the GC distance. On the other hand, the volume filling factor of dense molecular gas in the region is quite high ($\\gtrsim 0.3$; \\citealt{ser87}). In particular, the well-known Arched filaments all have negative velocities similar to that of the molecular gas; the photon-ionization modeling of these filaments suggests that they are physically in the vicinity of the \\xs\\ cluster \\citep{lan01a}. Thus the probability for a chance physical contact of a dense cloud with the cluster is not small. An independent argument for the association is an effective extinction deficit of $\\delta A_V \\approx 10$ over a region of $\\sim 15$\\as\\ from the cluster core, which can be interpreted as the displacement of the dusty gas by the cluster wind and/or the dust grain destruction by the UV radiation from the cluster \\citep{sto02}. The extinction is the largest towards the region just west of the cluster (\\citealt{sto02}; Note that East is to the right in their Figs.~3 and 8). Interestingly, this extinction deficit, corresponding to $\\delta N_H \\sim 3 \\times 10^{22} {\\rm~cm^{-2}}$, provides a natural explanation for the difference between our measured X-ray-absorbing column $N_H \\approx 8 \\times 10^{22} {\\rm~cm^{-2}}$ and the prediction from the total sight-line extinction $A_V = 24$ \\citep{sto02, boh78, sch98}. Furthermore, the interaction of the \\xs\\ cluster wind with the cloud may also explain the strong and distinct X-ray emission enhancement around the \\xs\\ cluster (\\S~\\ref{ss:d_line}). The far-IR spectroscopy further shows the presence of a component of dusty gas at a velocity of $-70 {\\rm~km~s^{-1}}$, unique at the location of the \\xs\\ cluster \\citep{cot05}. This component may represent shocked cloud gas, deflected toward us (e.g., in the lower left direction of Fig.~\\ref{f:ill}; see \\S~\\ref{ss:d_line} for further discussion). Therefore, we tentatively conclude that the ``$-$30 km s$^{-1}$ cloud'' and the cluster are undergoing a collision.  The collision of such clouds with the \\xs\\ cluster may have strongly affected its evolution. The absence of a natal cloud associated with the cluster at its velocity, for example, may be a consequence of the collision. The removal of this natal cloud from the cluster at an early time could have reduced the probability for low-mass stars to form. The cloud-cloud collision could also be responsible for the formation of the \\xs\\ cluster itself. The exceptionally high gas temperature and velocity dispersion in such a formation formation process could also result in a top-heavy initial mass function (IMF; see \\S~\\ref{ss:d_yso} for further discussion).  Our X-ray study further provides useful measurements about the GC environment. In addition to the $N_H$ measurement, we have also directly estimated the metal abundance (mainly iron) in the GC. Recent estimates based on near-IR spectroscopy of young and intermediate-age supergiants in GC \\citep[e.g.,][]{ram00} suggest an iron abundance that is consistent with being solar, i.e., similar to the abundance observed in the solar neighborhood. This result is {\\sl against} the general trend of an increasing metallicity with decreasing galacto-centric radius as observed in the disks of the Milky Way and nearby galaxies. Our X-ray measured iron abundance of $\\sim 1.8^{+0.8}_{-0.2}$ solar, based on the spectral analysis of the luminous colliding wind candidates in the \\xs\\ cluster, agrees with the trend. The thermal process involved in the X-ray emission is quite simple, and the ion fraction of the He-like Fe K$\\alpha$ emission is insensitive to the exact plasma temperature fitted. Furthermore, the iron abundance in the winds of the massive stars is not expected to be contaminated by their own nuclear synthesis in the deep cores of the stars. Therefore, we conclude that the iron abundance  in the ISM of the GC is super-solar.  \\subsection{Nature of Discrete X-ray Sources}\\label{ss:d_sou}  As shown in \\S~\\ref{ss:sou_pop}, our analysis confirms a relatively flat source NFR in the region of the \\xs\\ and \\xq\\ clusters, as indicated first in \\citet{mun06}. Our obtained power law slope ($\\alpha = 1.26_{-0.13}^{+0.14}$; 90\\% confidence) is flatter than those in the deep observations of Sgr B2 ($1.7\\pm0.2$) and Sgr A$^*$ ($1.4\\pm0.1$) as well as the $2^\\circ \\times 0\\fdg8$ shallow survey ($1.5\\pm0.1$). The implied over-population of relatively bright X-ray sources is clearly related to the presence of the two clusters.  The discrete X-ray sources in the core of the clusters are unlikely due to emission from individual normal massive stars or even binaries. The X-ray emission from such a star/binary can be characterized typically  by an optically-thin thermal plasma with a temperature of $\\sim 0.6$~keV and a luminosity following the empirical relation $\\frac{L_{X}}{L_{bol}} \\sim 10^{-7}$, where $L_{bol}$ is the bolometric luminosity. Thus the emission is too soft and faint to be observed from the GC. Even the Pistol star near the core of the \\xq\\ cluser (Fig.~\\ref{f:im_multi_bw}d) is not detected as an X-ray source. The star is a luminous blue variable with $L_{bol} \\gtrsim 10^{6.6} L_\\odot$ and has an extinction of $A_{K} \\approx 3.2$, corresponding to $N_{H} \\approx 5.1\\times10^{22} {\\rm~cm^{-2}}$. Assuming the {\\sl MEKAL} thermal plasma with a temperature of 0.6 keV, we estimate that the 3$\\sigma$ upper limit to the 0.3-8 keV luminosity is 3 $\\times10^{33} {\\rm~ergs~s^{-1}}$, consistent with $L_{x}/L_{bol} \\sim 10^{-7}$.  Most likely, the luminous X-ray sources associated with the \\xs\\ cluster represent colliding stellar winds in massive star close binaries. The characteristic shock temperature of a colliding wind is \\begin{equation} T \\simeq (3\\times10^{7} {\\rm~K}) v_{w,3}^{2}, \\end{equation} where $v_{w,3}$ is the relative colliding wind velocity in units of $10^3 {\\rm~km~s^{-1}}$. Well-known examples of such systems are WR11 (kT$\\approx 4.3$ keV, $L_X \\sim 8 \\times 10^{33} {\\rm~ergs~s^{-1}}$; \\citealt{sch04}) and WR140 (kT$\\approx 3$ keV, $L_X \\sim 2 \\times 10^{33} {\\rm~ergs~s^{-1}}$; \\citealt{zhe00}). Clearly, the expected temperatures are similar to the measured values for the sources in the \\xs\\ cluster, although their luminosities seem to be substantially higher than those confirmed colliding wind systems, which all have $L_X < 1 \\times 10^{34} {\\rm~ergs~s^{-1}}$ \\citep[e.g.,][]{osk05}. The unusually high X-ray luminosities of the colliding wind systems may be related to the compactness of the \\xs\\ cluster, in which very close binaries may form dynamically.  In contrast, the X-ray sources in the \\xq\\ cluster are probably typical colliding wind systems. They all have  individual $L_X$ in the range of $(0.2-3) \\times 10^{33} {\\rm~ergs~s^{-1}}$  as well as the hard X-ray spectra with the 6.7-keV emission line, as expected.  While only relatively luminous X-ray sources are detected individually, sources below our detection limit are hidden in the ``diffuse'' emission. Indeed, the diffuse emission in the cores of the \\xs\\ and \\xq\\ clusters shows a general correlation with their stellar distributions (Fig.~\\ref{f:rbp}). Thus, relatively faint colliding wind binaries could significantly contribute to the emission. But the bulk of the  diffuse X-ray emission in outer regions of the clusters may have different origins for several reasons. First, the emission extends much further away from the cluster cores than the stellar light distributions (Fig.~\\ref{f:rbp}). Second, the spectrum of the diffuse emission is harder than that of the discrete sources. Third, the emission in the outer region of the \\xs\\ cluster mainly exhibits the 6.4-keV line, inconsistent with the the colliding wind interpretation.  \\begin{figure*}[!htb] \\centerline{ \\epsfig{figure=f15.eps,width=1\\textwidth,angle=0} } \\caption{\\small Radial ACIS-I 1-9 keV intensity profiles ({\\sl crosses} with $1\\sigma$ error bars) around the \\xs\\ (a) and \\xq\\ (b) clusters, compared with the respective NICMOS F205W stellar light distributions (connected {\\sl triangles}). The cluster wind predictions are shown approximately as the solid line from 3-D simulations for the ``standard'' stellar wind mass-loss rates of the two clusters \\citep{roc05}. \\label{f:rbp}} \\end{figure*}  \\subsection{Cluster Winds}\\label{ss:d_cw}  In addition to colliding winds in individual massive star binaries, the collision among stellar winds collectively becomes important in a compact cluster such as the \\xs. The collision results in the thermalization of the stellar winds and their subsequent merging into a so-called cluster wind. Various 1-D models and 3-D hydrodynamic simulations have been carried out on cluster winds \\citep{rag01,ste03,roc05}. Within the uncertainties of such model parameters as overall stellar wind velocities and mass loss rates, simulated cluster winds are shown to explain the luminosities of diffuse X-ray emission from several star clusters \\citep[e.g.,][]{ste03,roc05}; but little detailed comparison has yet been performed.  Fig.~\\ref{f:rbp} compares the radial diffuse X-ray intensity profiles from the 3-D hydro-dynamical simulations, carried out specifically for the \\xs\\ and \\xq\\ clusters, approximately accounting for the discrete positions of massive stars and their individual stellar wind properties \\citep{roc05}. For the \\xq, the cluster wind could account for $\\sim 1/4-1/3$ of the observed diffuse X-ray emission. For the \\xs, which is much more compact, the simulated profile gives a reasonably good match to our measured distribution of the diffuse X-ray intensity within $\\sim 10$\\as, but is too steep to explain the emission at larger radii. The flattening of the observed intensity distribution in the radius range of $\\sim 10^{\\prime\\prime}$ to $\\lesssim 15^{\\prime\\prime}$) may arise from the reverse shock heating and confinement of the wind. At larger radii, the overall diffuse X-ray enhancement demonstrates a bow shock morphology and is prominent in the Fe K$\\alpha$ 6.4-keV line emission (\\S~\\ref{sss:dif_a}), inconsistent with the expectation for the cluster wind interpretation (see below). Therefore, the cluster wind may be important in the core, but not in the outer region of the \\xs\\ cluster.  The complexity of the diffuse X-ray emission from the \\xs\\ cluster probably reflects its interaction with the CS cloud. Both the morphology of the diffuse X-ray emission, particularly the elongation of the 6.7-keV line emission, and the extinction deficit distribution indicate that the motion of the cluster relative to the cloud is from East to West in the sky.  Because of their supersonic relative motion, a bow-shock is expected to form around the cluster. Fig.~\\ref{f:ill} illustrates this  simple-minded scenario for the interaction, although the true situation is certainly more complicated.  Following \\citet{bur88}, we can estimate from the ram-pressure balance the characteristic radius of the reverse shock in the cluster wind as \\begin{equation} r_s = (0.7 {\\rm~pc}) \\dot{M}_{w,-4}^{1/2} v_{w,3}^{1/2} v_{r,2}^{-1} n_{a,2}^{-1/2}, \\label{e_r_s} \\end{equation} where $\\dot{M}_{w,-4}$ is the mass-loss rate of the cluster wind (in units of  $ 10^{-4} M_\\odot$), $v_{r,2}$ is the relative velocity between the cluster and the cloud ($10^2 {\\rm~km~s^{-1}}$), and $n_{a,2}$ is the gas density in the colliding cloud ($10^2 {\\rm~cm^{-3}}$). Because the contact discontinuity has a scale $l_c \\sim 1.5 r_s$, we can estimate the volume of shocked wind materials as $V \\sim \\frac{4\\pi}{3}(1-1/1.5^3) l_c^3$. Assuming that this volume corresponds to the 6.7-keV line plume, which has a radius $l_c \\sim 0.6$ pc (15\\as) and we can infer $n_e \\sim 5 {\\rm~cm^{-3}}$ from the integrated emission measure of the central plume, $IEM \\sim 16 {\\rm~cm^{-6}~pc^3}$ (the {\\sl MEKAL} fit in Table~\\ref{spec_d}). The ram-pressure balance also gives the density of the shocked ambient gas  $n_a \\sim \\frac{n_e}{4}(\\frac{v_w}{v_{r}})^2$ $\\sim (1.3 \\times 10^2  {\\rm~cm^{-3}}) v_{w,3}^2 v_{r,2}^{-2}$. This, together with Eq.~\\ref{e_r_s}, gives $\\dot{M} \\sim (1 \\times 10^{-4} M_\\odot) v_{w,3}$.  The above inferred $n_a$ and $\\dot{M}$ values depend  on $v_w$, which may be quite uncertain. In particular, the near-IR spectroscopic estimate of stellar winds may have significantly underestimated $v_w$ as possible low-emissivity winds in the line profiles were not taken into account \\citep{cot96}, i.e., the wind terminal velocity could be considerably higher than $1 \\times 10^3 {\\rm~km~s^{-1}}$. Nevertheless, the above inferred mass-loss rate still appears substantially smaller  (by a factor of up to $\\sim 10$) than the current estimates based on radio continuum estimates \\citep[e.g.,][]{lan05}. Such estimates may be very uncertain (e.g., \\citealt{roc05}), particularly for binaries with strong wind-wind interaction. The relatively small $n_a$ value is consistent with the weak CS emission from the ambient gas, probably representing the inter-clump medium of the colliding cloud.  While the shocked cluster wind should be constantly flowing out from the bow shock at a velocity comparable to the sound velocity $c_s \\sim (8 \\times 10^2 {\\rm~km~s^{-1}}) v_{w,3}$, we can also estimate the ionization time scale as $\\tau \\sim n_e l_c/c_s \\sim (1 \\times 10^{11} {\\rm~cm^{-3}~s}) v_{w,3}^{-1}$, which is much too large to explain the 6.4-keV line emission with an {\\sl NEI} plasma, but is consistent with that inferred from the spectrum of the central plume (\\S~\\ref{sss:dif_a}). Therefore, the observed size and shape of the 6.7-keV line plume (Fig.~\\ref{f:im_line}), at least qualitatively, match the predictions of this simple bow shock interpretation, within the uncertainties of the relevant parameters.  \\subsection{Origin of the 6.4-keV line emission}\\label{ss:d_line}  The above discussion indicates that the 6.4-keV line emission associated with the \\xs\\ cluster is unlikely due to an {\\sl NEI} process. We thus consider the possible origin of the line emission as the filling of iron K-shell vacancies produced by either ionizing radiation with photon energies $> 7.1$ keV or collision with low-energy cosmic-ray electrons \\citep[LECRe;][]{val00}. The fluorescent line emission and Thompson continuum scattering seem to give a reasonable good explanation for those most prominent 6.4-keV enhancements associated with well-known giant molecular clouds such as Sgr B2 and Sgr C in the GC \\citep{koy96a, cra02, rev04}. This explanation requires the presence of a luminous  X-ray source with a spectrum consistent with the observed power law continuum with a photon index of $\\Gamma \\approx 1.8$. Because such a source is currently not present in the GC, the observed emission is proposed to be the reflection of past Sgr A*, with an X-ray luminosity of $\\gtrsim 10^{39} {\\rm~ergs~s^{-1}}$, about a few hundreds years ago.  However, the fluorescence interpretation has difficulties in accounting for the 6.4-keV line emission regions closer to Sgr A*. A comparison of the CS emission and the diffuse 6.4-keV line intensity does not show a peak-to-peak correlation, which should be expected because the gas traced by the CS emission is expected to be optically thin to the iron ionizing radiation \\citep{wan03}. As shown in \\S~\\ref{ss:cs_dist}, the detailed correlation is also absent in the \\xs\\ CS cloud. This difficulty may be avoided, if the CS emission does not trace well the actual gas distribution (e.g., due to the destruction of the molecule by the strong UV radiation from the \\xs\\ cluster). Even in this case, however, the gas column density of the cloud cannot be much greater than $\\delta N_H \\sim 10^{22} {\\rm~cm^{-2}}$, constrained by both the X-ray absorption and the near-IR extinction distribution (\\S~\\ref{ss:d_envi}). Following \\citet{sun98}, we estimate the required X-ray luminosity of Sgr A* to produce the detected 6.4-keV line intensity of the \\xs\\ (Fig.~\\ref{f:cs_x}) as \\begin{equation} L_X = (4 \\times 10^{39}{\\rm~ergs~s^{-1}}) (d/27 {\\rm~pc})^2 (\\delta N_H/10^{22} {\\rm~cm^{-2}})^{-1}, \\end{equation} where we have assumed the iron abundance to be 2 $\\times$ solar and have scaled the distance ($d$) between the cloud and Sgr A* to be their projected separation in the sky, corresponding a light travel time of only about 90 years. Of course, the actual distance is likely to be greater, and the required $L_X$ should then be higher. This common interpretation of the 6.4-keV line enhancement and those associated with Sgr B2 and Sgr C, though difficult to rule out completely, would not explain the apparent position coincidence between  the cloud and the cluster.  Alternatively, one may consider the \\xs\\ cluster as the origin of the hard X-rays. But this possibility can be easily dismissed because of the absence of the 6.7-keV line (which is strong in both the point-like sources and in the cluster core) in the 6.4-keV line enhancement. Furthermore, the observed X-ray luminosity of the cluster is more than a factor of $10^2$ short of what is required for the fluorescence interpretation.  A more plausible scenario for the \\xs\\ 6.4-keV line enhancement is the LECRe-induced Fe K-shell vacancy filling \\citep{val00}. In this scenario, the continuum is due to the bremsstrahlung radiation of the LECRe. The expected power-law photon index of the continuum is 1.3-1.4 over the range of 1-10 keV, consistent with our measured value of the SE extension (Table~\\ref{spec_d}). The LECRe may be produced in strong shocks that are present within the \\xs\\ cluster and in both the forward bow shock and the reverse-shock in the cluster wind (see the discussion above). For example, \\citet{byk00} have shown that a shock of velocity $\\gtrsim 10^2 {\\rm~km~s^{-1}}$ into a molecular cloud, accompanied by magneto-hydrodynamic turbulence, can provide a spatially inhomogeneous distribution of nonthermal LECRe. \\citet{yus03} have further presented observational evidence for nonthermal diffuse radio emission from the \\xs\\ cluster and have suggested that colliding wind shocks may generate the responsible relativistic particles. The diffuse X-ray enhancement has a bow-shock morphology and is presumably linked to the site of particle acceleration. But, because of  particle diffusion and gas flow, one  does not expect a peak-to-peak correlation of the X-ray emission with the CS emission from the colliding cloud. Following \\citet{yus02a}, we estimate the LECRe energy density required to produce the observed 6.4-keV line intensity. If the {\\sl shocked} gas density is $\\sim 10^{3} {\\rm~cm^{-3}}$, the required energy density is then $\\sim 6 \\times 10^3 {\\rm~eV~cm^{-3}}$, substantially greater than the value $0.2 {\\rm~eV~cm^{-3}}$ from averaging over the Galactic ridge \\citep{val00}. But the implied pressure inside the bow shock can still be balanced by the high ram-pressure ($\\sim 2 \\times 10^{-8} v_{r,2}^2 n_{a,2} {\\rm~dyn~cm^{-2}}$) of the collision between the cluster wind and the CS cloud. In short, the bow shock provides a plausible interpretation of the distinct spatial and spectral properties of the diffuse X-ray emission around the cluster and its physical relationship to the CS cloud.  \\begin{figure}[!thb] \\centerline{ \\epsfig{figure=f16.eps,width=0.4\\textwidth} } \\caption{\\small  An illustration of the proposed cluster-cloud collision scenario for the \\xs. The shocked cloud gas is partly traced by the CS and 6.4-keV lines (Fig.~\\ref{f:cs_x}), whereas the shocked cluster wind plasma near the cluster is by the 6.7-keV line (Fig.~\\ref{f:im_line}). } \\label{f:ill} \\end{figure}  Finally, we consider the possibility that the 6.4-keV line enhancement represents the reprocessed X-rays from numerous discrete and faint sources embedded around the \\xs\\ cluster. A natural candidate for such sources might be low-mass pre-main sequence young stellar objects (YSOs). But they are in general not known to emit strong 6.4-keV line emission. In the Orion nebula, for example, the line emission is detected from only a few YSOs and with EWs less than 300 keV. Therefore, YSOs are probably not a significant contributor to the 6.4-keV line enhancement.  \\subsection{YSO population and stellar IMF}\\label{ss:d_yso}  The overall luminosity of the diffuse X-ray emission provides a fundamental limit to the population of YSOs and hence the IMF of the \\xs\\ and \\xq\\ clusters. YSOs in the mass range of $(0.3 - 3) M_\\odot$ typically have large $L_X/L_{bol}$ ratios and hard X-ray spectra. Most importantly, they can be numerous, as shown in the {\\sl Chandra} Orion Ultra-deep Project \\citep{fei05}. Though with a {\\sl mean} 2-8 keV luminosity of only $\\sim 1.2 \\times 10^{30} {\\rm~ergs~s^{-1}}$ per star, YSOs collectively account for about 75\\% of the luminosity of the Orion nebula,  the IMF of which is consistent with the standard Miller \\& Scalo (1979; MS hereafter), based on the work by \\citet{hil97}. If the clusters in the GC have a similar IMF, YSOs should then be equally important.  At the GC, typical YSOs cannot be detected individually, but can be constrained collectively in {\\sl Chandra} observations. Based on the diffuse X-ray intensity observed around Sgr A*, \\citet{nay05} find that the population of YSOs in the GC cluster is extremely small. They conclude that it cannot be a remnant of a massive star cluster, originated at several tens of parsecs away from Sgr A* and then dynamically spiralled in, and is thus most likely formed {\\sl in situ} in a self-gravitating circum-nuclear disk and with a top heavy IMF. While star formation around a super-massive black hole represents an extreme, it is clearly important to examine the IMF of the \\xs\\ and \\xq\\ clusters in the general environment of the GC.  We find a similar deficiency of YSOs in \\xs\\ and \\xq\\ clusters, which places important constraints on their IMF. Existing near-IR studies have provided estimates on the present-day MF of stars with masses greater than a few solar masses in the core of the \\xs\\ cluster; MF measurements in outer regions are difficult because of severe confusion with field stars. Fig.~\\ref{f:imf} shows the MF within $r < 0.4$ pc of the \\xs\\ cluster \\citep{sto05}, compared with various predictions of YSOs. The standard MS IMF, normalized with the number of stars in the mass range of $M > 60$ M$_\\odot$, predicts at least 2$\\times10^{5}$ YSOs. Using our measurements of the diffuse X-ray emission, we can directly get an upper limit to the population of YSOs over the entire cluster ($r < 2.5$ pc). We assume that the mean X-ray luminosity of the YSOs in the \\xs\\ cluster to be the same as that in the Orion nebula, because of their similar ages. As is shown above, much of the diffuse X-ray emission, though difficult to quantify, likely has other origins (e.g., cluster winds) to account for the prominent iron K$\\alpha$ lines. Therefore, the use of the total 2-8 keV diffuse X-ray luminosity of $2 \\times 10^{34} {\\rm~ergs~s^{-1}}$ (Table~\\ref{spec_d}) gives the upper limit as $2 \\times 10^{4}$ YSOs, which is a factor of 10 smaller than the above prediction from the MS IMF. The actual discrepancy should be substantially larger. We have neglected the mass loss in the stellar evolution, which is important for the massive stars. Considering the mass loss, the number of stars in the above initial mass range, hence the normalization of the IMF, would be greater. The number of cluster stars in the region of $r= 0.4-2.5$ pc is also not included, though difficult to fully quantify; for example, there are 77 stars with $M>40 M_{\\odot}$ in the radius $r< $ 0.6 pc \\citep{fig99b}, compared to about 48 in the same mass range of the MF obtained by \\citet{sto05} for $r< $ 0.4 pc.  \\begin{figure}[htb] \\centerline{ \\epsfig{figure=f17.eps,width=0.5\\textwidth} } \\caption{\\small Present-day MF as obtained by \\citet{sto05} in the $r < 0.4$~pc core of the \\xs\\ cluster, compared with the power law ($\\propto M^\\Gamma$, where $\\Gamma=-0.86$; dashed curve) fitted in the 6-60 M$_\\odot$ range \\citep{sto05}) and the MS half-Gaussian (dot-dashed line), normalized to the number of stars in the mass range of $> 60$ M$_\\odot$. The X-ray-inferred upper limit to the number of YSOs (0.3-3 M$_\\odot$) in the entire cluster is marked as the bar with the arrow. The MF at M $\\lesssim$ 6 M$_\\odot$ may be significantly contaminated by field stars \\citep{sto05}.} \\label{f:imf} \\end{figure}  An extrapolation of a power law fit in the mass range of 6-60 M$_\\odot$, as obtained by \\citet{sto05}, is consistent with the X-ray-inferred number of YSOs (Fig.~\\ref{f:imf}). But, \\citet{sto05} shows that the MF steepens with radius. This steepening is expected as a result of the dynamic mass segregation of stars in the cluster core \\citep{kim02,por02}. Furthermore, mergers among stars may also be important in flattening the MF toward the cluster core. Therefore, the MF of the entire cluster, including the region in $r = 0.4 - 2.5$ pc, is likely to be steeper. If this is the case, a turnover in the MF (e.g., at $\\sim 6$ M$_{\\odot}$, as indicated in the study of \\citet{sto05}, may indeed be required to explain the X-ray-inferred upper limit on the overall YSO population in the \\xs\\ cluster.  Similarly, we can constrain the YSO population in the \\xq\\ cluster. There are 30 stars with masses larger than 20 $M_{\\odot}$ within $r = 25$\\as\\ (1 pc) of the \\xq\\ cluster \\citep{fig99a}. Assuming the MS IMF would predict a total number of YSOs to be at least $2 \\times10^{4}$. These YSOs would have a 2-8 keV luminosity of $\\sim 1 \\times10^{34} {\\rm~ergs~s^{-1}}$, accounting for the weak dependence of the mean X-ray luminosity on the cluster age (a factor of 1.6; \\citealt{pre05}). This predicted value is a factor of 5 greater than our measured total diffuse X-ray luminosity of $2 \\times10^{33} {\\rm~ergs~s^{-1}}$ within $\\sim 1^\\prime$ of the \\xs2\\ cluster (\\S~\\ref{sss:dif_a})."
    },
    "0606/nucl-th0606063_arXiv.txt": {
        "abstract": "Using the relativistic impulse approximation with the Love-Franey \\textsl{NN} scattering amplitude developed by Murdock and Horowitz, we investigate the low-energy ($100$ MeV$\\leq E_{\\mathrm{kin}}\\leq 400$ MeV) behavior of the nucleon Dirac optical potential, the Schr\\\"{o}dinger-equivalent potential, and the nuclear symmetry potential in isospin asymmetric nuclear matter. We find that the nuclear symmetry potential at fixed baryon density decreases with increasing nucleon energy. In particular, the nuclear symmetry potential at saturation density changes from positive to negative values at nucleon kinetic energy of about $200$ MeV. Furthermore,the obtained energy and density dependence of the nuclear symmetry potential is consistent with those of the isospin- and momentum-dependent MDI interaction with $x=0$, which has been found to describe reasonably both the isospin diffusion data from heavy-ion collisions and the empirical neutron-skin thickness of $^{208} $Pb. ",
        "introduction": "The energy dependence of the nuclear symmetry potential, i.e., the isovector part of nucleon mean-field potential in asymmetric nuclear matter, has recently attracted much attention \\cite% {91bomb,97ulrych,das03,li04a,li04b,chen04,rizzo04,fuchs04,mazy04,chen05,li05xmed,05behera,baran05,samma05,zuo05,fuchs05,fuchs05prc,rizzo05,mazy06}% . Its knowledge together with that of the density dependence of the nuclear symmetry energy are important for understanding not only the structure of radioactive nuclei and the reaction dynamics induced by rare isotopes but also many critical issues in astrophysics \\cite{ireview98,ibook}. Various microscopic and phenomenological models, such as the relativistic Dirac-Brueckner-Hartree-Fock (DBHF) \\cite% {97ulrych,fuchs04,mazy04,samma05,fuchs05,fuchs05prc,mazy06} and the non-relativistic Brueckner-Hartree-Fock (BHF) \\cite{91bomb,zuo05} approach, the relativistic mean-field theory based on nucleon-meson interactions \\cite% {baran05}, and the non-relativistic mean-field theory based on Skyrme-like interactions \\cite{das03,05behera}, have been used to study the nuclear symmetry potential, but the predicted results were found to vary widely. While most models predict a decreasing nuclear symmetry potential with increasing nucleon momentum, albeit at different rates, a few nuclear effective interactions used in some of the models lead to an opposite conclusion.  Using the relativistic Dirac optical potential obtained from the relativistic impulse approximation (RIA) \\cite% {mcneil83prc,mcneil83,shepard83,arnold79,clark83,miller83,horowitz85,murdock87,tjon85,ott88,jin93,toki01} with the empirical nucleon-nucleon (\\textsl{NN}) amplitude calculated by McNeil, Ray, and Wallace (MRW) \\cite{mcneil83prc,mcneil83,shepard83}, which works well for elastic nucleon-nucleus scattering at medium and high energies (above $500$ MeV), three of present authors \\cite{chen05ria} have recently studied the high-energy behavior of the nuclear symmetry potential in asymmetric nuclear matter. It was found that for nucleons at high energies, the symmetry potential at fixed baryon density is essentially constant and is slightly negative below nuclear density of about $\\rho =0.22$ fm$^{-3}$ but increases almost linearly to positive values at high densities.  A nice feature of RIA is that it permits very little phenomenological freedom in deriving the nucleon Dirac optical potential in nuclear matter. The basic ingredients in this method are the free \\textsl{NN} invariant scattering amplitude and the nuclear scalar and vector densities in nuclear matter. This is in contrast to the relativistic DBHF approach, where different approximation schemes and methods have been introduced for determining the Lorentz and isovector structure of the nucleon self-energy \\cite{97ulrych,fuchs04,mazy04,samma05,fuchs05,fuchs05prc,mazy06}. However, the original RIA of MRW failed to describe spin observables at laboratory energies lower than $500$ MeV\\cite{ray85}, and its predicted oscillations in the analyzing power in proton-\\textrm{Pb} scattering at large angles were also in sharp disagreement with experimental data \\cite{drake85}. These shortcomings are largely due to the implicit dynamical assumptions about the relativistic \\textsl{NN} interaction in the form of the Lorentz covariance \\cite{adams84} and the somewhat awkward behavior under the interchange of two particles \\cite{horowitz85} as well as the omitted medium modification due to the Pauli blocking effect. To solve these theoretical limitations at lower energies, Murdock and Horowitz (MH) \\cite{horowitz85,murdock87} extended the original RIA to take into account following three improvements: i) an explicit exchange contribution was introduced by fitting to the relativistic \\textsl{NN} scattering amplitude; ii) a pseudovector coupling rather than a pseudoscalar coupling was used for the pion; iii) medium modification from the Pauli blocking was included. With these improvements, the RIA with free \\textsl{NN} scattering amplitude was then able to reproduce successfully measured analyzing power and spin rotation function for all considered closed shell nuclei in proton scattering near $200$ MeV. Particularly, the medium modification due to the Pauli blocking effect was found to be essential in describing the spin rotation function for $^{208}$% Pb at proton energy of $290 $MeV \\cite{murdock87}.  Extending our previous work by using the generalized RIA of MH and the nuclear scalar and vector densities from the relativistic mean-field theory, we study in the present paper the low-energy ($100$ MeV$\\leq E_{\\mathrm{kin}% }\\leq 400$ MeV) behavior of the nucleon Dirac optical potential, the Schr% \\\"{o}dinger-equivalent potential, and the nuclear symmetry potential in isospin asymmetric nuclear matter. We find that for low energy nucleons the nuclear symmetry potential at fixed nuclear density decreases with increasing nucleon energy. In particular, the nuclear symmetry potential at saturation density changes from positive to negative values at nucleon kinetic energy around $200$ MeV. The resulting energy and density dependence of the nuclear symmetry potential is further found to be consistent with the isospin- and momentum-dependent MDI interaction with $x=0$ \\cite% {das03,chen05}, which has been constrained by the isospin diffusion data in heavy-ion collisions and the empirical neutron-skin thickness of $^{208}$Pb \\cite{chen05,li05xmed,steiner05nskin,chen05nskin}. Our results thus provide an important consistency check for the energy dependence of the nuclear symmetry potential in asymmetric nuclear matter.  The paper is organized as follows. In Section \\ref{optical}, we briefly review the generalized relativistic impulse approximation for the nuclear optical potential. Results on the relativistic nuclear optical potential, the resulting Schr\\\"{o}dinger-equivalent potential, and the nuclear symmetry potential in asymmetric nuclear matter are presented in Section \\ref{results}% . A summary is given in Section \\ref{summary}. ",
        "conclusions": "\\label{summary}  Based on the generalized relativistic impulse approximation of MH and the scalar and vector densities from the relativistic mean-field theory with the parameter set HA, we have studied the low-energy behavior of the nuclear symmetry potential in asymmetric nuclear matter. In the relativistic impulse approximation of MH, the low energy behavior of the Dirac optical potential has been significantly improved by including the pseudovector coupling for pion, the exchange contribution, and medium modification due to the Pauli blocking effect. We find that compared with results from the original RIA of MRW, the generalized RIA of MH gives essentially identical real parts of the scalar and vector amplitudes for both proton-proton and neutron-proton scattering but significantly reduced strength in their imaginary parts at low energies $E_{\\mathrm{kin}}\\leq 300$ MeV. These improvements in the RIA of MH modify the real scalar and vector Dirac optical potentials at lower energies and make the resulting energy dependence of the Schr\\\"{o}% dinger-equivalent potential and nuclear symmetry potential more reasonable.  At saturation density, the nuclear symmetry potential is found to change from positive to negative values at nucleon kinetic energy of about $200$ MeV. This is a very interesting result as it implies that the proton (neutron) feels an attractive (repulsive) symmetry potential at lower energies but repulsive (attractive) symmetry potential at higher energies in asymmetric nuclear matter. Adding also the repulsive Coulomb potential, a high energy proton in asymmetric nuclear matter thus feels a very stronger repulsive potential. This behavior of the nuclear symmetry potential can be studied in intermediate and high energy heavy-ion collisions that are induced by radioactive nuclei, e.g., by measuring two-nucleon correlation functions \\cite{chen-nn} and light cluster production \\cite{chen-cluster} in these collisions.  Comparing the energy and density dependence of the nuclear symmetry potential from the RIA of MH with that from the MDI interaction indicates that results from the MDI interaction with $x=0$ are in good agreement with those from the RIA of MH. For baryon density less than $0.25$ fm$^{-3}$ and nucleon energy less than $400$ MeV as considered in the present work, the nuclear symmetry potential from the MDI interaction with $x=0$ lies approximately between the two results from the RIA of MH with isospin-dependent and isospin-independent Pauli blocking corrections. This provides a strong evidence for the validity of the MDI interaction with $x=0$ in describing both the isospin diffusion data in intermediate energy heavy ion collisions and the neutron skin thickness data for $^{208}$Pb.  The results presented in present work thus provide an important consistency check for the energy/momentum dependence of the nuclear symmetry potential in asymmetric nuclear matter, particularly the momentum dependent MDI interaction with $x=0$, which is an essential input to the isospin-dependent transport model \\cite{li04b,baran05,li05xmed} in studying heavy-ion collisions induced by radioactive nuclei at intermediate energies. They are also useful in future studies that extend the Lorentz-covariant transport model \\cite{ko87,mosel92} to include explicitly the isospin degrees of freedom."
    },
    "0606/astro-ph0606231_arXiv.txt": {
        "abstract": "We report in this letter our analysis of a large sample of photospheric vector magnetic field measurements. Our sample consists of 17200 vector magnetograms obtained from January 1997 to August 2004 by Huairou Solar Observing Station of the Chinese National Astronomical Observatory. Two physical quantities, $\\alpha$ and current helicity, are calculated and their signs and amplitudes are studied in a search for solar cycle variations. Different from other studies of the same type, we calculate these quantities for weak ($100G<|B_z|<500G$) and strong ($|B_z|>1000G$) fields separately. For weak fields, we find that the signs of both $\\alpha$ and current helicity are consistent with the established hemispheric rule during most years of the solar cycle and their magnitudes show a rough tendency of decreasing with the development of solar cycle. Analysis of strong fields gives an interesting result: Both $\\alpha$ and current helicity present a sign opposite to that of weak fields. Implications of these observations on dynamo theory and helicity production are also briefly discussed. ",
        "introduction": "Magnetic helicity is a physical quantity that measures the topological complexity of a magnetic field, such as the degree of linkage and/or twistedness in the field (Moffatt 1985, Berger \\& Field 1984). It has been shown that its total amount is approximately conserved in the Sun even when there is an energy release during fast magnetic reconnection (Berger 1984). This conservation of total magnetic helicity is considered to play an important role in the dynamical processes in the Sun. For example, by considering helicity conservation in the mean-field dynamo, theories have predicted that solar dynamo would produce opposite helicity signs in the mean field and in the fluctuations (Blackman \\& Field 2000, see Ossendrijver 2003 for a review). It has also been considered that magnetic helicity and its conservation may play an important role in CME dynamics (Low 2001, Demoulin et al. 2002) where accumulation of total magnetic helicity in the respective northern and southern hemispheres leads to a natural magnetic energy storage for CME eruptions (Zhang \\& Low 2005, Zhang, Flyer \\& Low 2006).   A direct measurement of magnetic helicity and hence a direct test of above theories by observations are still out of our reach because so far the photosphere is still the only layer that we can measure vector magnetic fields with reasonable temporal and spatial resolutions. However, by calculating derived physical quantities, such as $\\alpha$ and current helicity, from observed photospheric vector magnetograms we do get a glimpse of properties of magnetic helicity in the Sun. For example, from photospheric magnetic field measurements we learn that magnetic fields emerging from the solar convection zone to the photosphere are already significantly twisted (Kurokawa 1987, Leka et al. 1996) and statistically these fields possess a positive helicity sign in the southern hemisphere and a negative helicity sign in the northern hemisphere (Pevtsov et al. 1995, Bao \\& Zhang 1998). These observations thus provide us implications on how magnetic helicity might be produced in the convection zone (Berger \\& Ruzmaikin 2000) and how magnetic helicity conservation might have played a role in balancing the twist and writhe helicity in an originally untwisted flux rope (Longcope et al. 1998).   In this letter, we intend to use photospheric vector magnetic field measurements to find further observational indications of helicity production and conservation. Different from other works of the same type, we separate studied fields into two parts: strong magnetic fields and weak magnetic fields. We organize our paper as follows: In \\S 2, we describe our observation and data reduction. In \\S 3, we present our analysis and discussions. We conclude the letter with a brief summary in \\S 4. ",
        "conclusions": "\\subsection{Comparison with previous studies}  Before we proceed to present our results it is useful to check our data reduction of this dataset with previous results obtained by other instruments and datasets. We select a subsample of our dataset, containing observations made between 1997 July to 2000 September, in order to compare with Pevtsov et al. (2001) where $\\alpha_{best}$ and current helicity were also calculated for the same period of time. The difference is that their magnetograms were obtained by the Haleakala Stokes Polarimeter (HSP) at Mees Solar Observatory.   Figure 1 presents the latitudinal profile of $\\alpha_{best}$ for the 391 active regions observed by Huairou magnetograph during this period of time. Each point presents the average value of $\\alpha_{best}$ when multiple magnetograms of the same active region were obtained. Note in producing this figure we did not separate the weak and strong fields but instead use all data points with $|B_z|>100G$ and $|B_x, B_y|>200G$, in order to make a reasonable comparison with Pevtsov et al. (2001). The green line shows the least-square best-fit linear function of these $\\alpha_{best}$ values. The similarity between our figure and Figure 1 of Pevtsov et al. (2001) indicates a good consistence between the two datasets.   Out of our 391 active regions during this period, 58.9\\% of 214 active regions in the northern hemisphere have $\\alpha_{best}<0$ and 67.2\\% of 117 active regions in the southern hemisphere have $\\alpha_{best}>0$. These numbers are consistent with the numbers of 62.9\\% and 69.9\\% for the northern and southern hemispheres respectively in Pevtsov et al. (2001). Our data shows no tendency of hemispheric rule by current helicity. 44.4\\% of 214 active regions in the northern hemisphere have $h_c<0$ and 45.8\\% of 117 active regions in the southern hemisphere have $h_c>0$. Note in Pevtsov et al. (2001) a much weaker tendency is also found with numbers of 50\\% and 57.5\\% for their $h_c$ values in the northern and southern hemispheres respectively. They contribute this difference to Faraday rotation. But we suggest the difference is largely (although possibly not all) because of a physical point which we will return to address below.   Averages of $\\alpha_{best}$ for active regions observed in each 10 degrees of solar latitudes are also plotted in Figure 1, presented as red square symbols. The large error bars of these averages remind us that our established hemispherical rule is of a statistical result. Individual active regions may present large deviations from the mean values. This is also true for other statistical results that we will present below.   \\subsection{Helicity observation of weak fields}   Figure 2 presents our result of solar cycle variations of $\\alpha$ (top panel) and current helicity (middle panel) for weak fields ($100G<|B_z|<500G$). Each point in these plots is a weighted average of $\\langle\\alpha_z\\rangle$ or current helicity for active regions observed during one year. For active regions in the southern hemisphere the weight is set to $1$ and for active regions in the northern hemisphere the weight is set to $-1$. The weighted averages then indicate the magnitudes of $\\alpha$ or current helicity averaged over the global surface during a whole year, assuming the northern and southern hemispheres have opposite helicity signs. We see that both averaged $\\alpha$ and current helicity have positive signs except for the Year 2004. This tells us that both $\\alpha$ and current helicity for weak fields obey the established hemispheric rule during most years of the solar cycle. The averaged $\\alpha$ and current helicity for Year 2004 are negative, which indicates the usual hemispheric rule is not followed in this year. This is consistent with Hagino \\& Sakurai (2005) where they also found a violation of the usual hemispheric rule during solar minimums.   Figure 2 also presents a rough tendency of a decrease of $\\alpha$ and current helicity with the development of solar cycle. We notice in Berger \\& Ruzmaikin (2000) the helicity production rate by differential rotation in solar interior is calculated and their calculation also shows a similar decrease of magnitudes of the rate of helicity transported into the northern and southern hemisphere respectively. This can be seen from the bottom panel of Figure 2 where the helicity transportation rate into the southern hemisphere by the m=0 mode is replotted, with data taken from Berger \\& Ruzmaikin (2000). This interesting consistence seems to suggest that differential rotation is the source of helicity production in solar interior although we are not able to make a conclusion because we do not know whether the $\\alpha$ effect will also produce the same tendency or not.   As pointed out by the careful referee, the calculated transferred helicity ends at zero during solar minimums whereas our observation as well as Hagino \\& Sakurai (2005) show the helicity goes to the opposite sign during solar minimums. We intend to explain this as a result of trans-equatorial reconnection (Pevtsov 2000) which has consumed the helicity of the dominate sign in each hemisphere, a point interesting of itself but is out of the scope of current letter.   Another interesting implication of Figure 2 is that, whereas we usually consider helicity variation as a function of latitude as presented in Figure 1, another possibility is that the helicity variation is more associated with solar cycle dependence and the known latitude dependence is just a derived relation from this solar cycle dependence of helicity and the Butterfly diagram.    \\subsection{Helicity observation of strong fields}   For strong magnetic fields ($|B_z|>1000G$), calculation of weighted averages of $\\alpha$ and current helicity presents an interesting result, shown in Figure 3. All averaged $\\alpha$ and current helicity are negative, which means they do not follow the usual hemispheric rule. This also means that strong fields have a helicity sign opposite to that of weak fields.   As we have mentioned earlier, if we interpret our observed weak fields in active regions as the representatives of the general weak fields distributed over the global Sun, then we may use them to represent the large-scale field. Our strong fields may be used to represent the small-scale fluctuations compared to the large-scale of the global Sun. Then under this interpretation our observation seems to be consistent with the theory that solar dynamo would produce opposite helicity signs in the mean field and in the fluctuations.   It is also interesting to notice that in Berger \\& Ruzmaikin (2000) the higher modes helicity, such as the m=5 mode replotted in Figure 3, also has a sign opposite to that of the m=0 mode. Again, if we interpret their low-degree (such as m=0) mode field corresponds to our weak field because both of them represent a more uniformly-distributed field over the global Sun and their high-degree (such as m=5) mode field corresponds to our strong field because both of them are sporadically appeared on the surface, then their calculation and our observation show a consistence again.   The observation that strong fields have a helicity sign opposite to that of weak fields may help us understand why $\\alpha_{best}$ usually shows a better hemispheric rule than current helicity if both quantities are calculated from vector magnetograms of the whole field (Pevtsov et al. 2001). We interpret it as follows. When we calculate $\\alpha_{best}$ of the whole field, each data point is given an equal weight. This results in the calculated $\\alpha_{best}$ presenting the sign of weak fields, whose number of data points dominates over that of strong fields. But when we calculate the current helicity of the whole field, defined as $h_c=B_z\\cdot(\\nabla\\times{\\bf B})_z = \\alpha B_z^2$, we have attributed a weight of $B_z^2$ to each data point. This then results in a nearly cancellation of current helicity between the weak and strong fields because weak and strong fields happen to have opposite helicity signs and the former has a larger number of data points but smaller $B_z^2$ values for each data point whereas the latter has a smaller number of data points but each data point has a larger $B_z^2$ value.   It has been suggested that Faraday rotation contributes to the difference between $\\alpha_{best}$ and current helicity. We suggest the main reason is the opposite helicity signs between weak and strong fields. J. T. Su \\& H. Q. Zhang (2006, in preparation) recently did a calculation and it shows that whereas Faraday rotation may rotate the transverse fields to 20 - 30 degrees, the resultant $\\alpha$ values are less influenced, with changes of $\\alpha$ values all less than a few percentages. Another comment is that if Faraday rotation is the reason of the difference we should not see the difference in the dataset obtained by spectrograph-type magnetographs where the effect of Faraday rotation can be taken care of by inversion methods. But the difference is observed in Pevtsov et al. (2001) where HSP data are used. We have recently checked several active regions observed by ASP/HAO. Similar feature of opposite helicity signs between weak and strong fields is found, although not in every region examined. Also kindly pointed out by the referee, similar tendency of opposite helicity signs is also indicated in a decaying active region observed by ASP (Figure 4 of Pevtsov \\& Canfield 1999).   Finally we point out another consistence of our observation with previous study. By applying a known reconstruction technique to MDI data Pevtsov and Latushko (2000) calculated the current helicity of the global Sun. They found that the usual hemispheric rule is followed for regions above 40 degrees of solar latitudes whereas the rule is surprisingly not obvious for regions within 40 degrees of solar latitudes. With our observation, we now can interpret it as follows. In high latitudes magnetic fields are dominated by weak fields with their signs following the usual hemispherical rule, whereas in low latitudes strong fields with an opposite helicity sign present to result in a reduction to the usual hemispherical rule."
    },
    "0606/astro-ph0606188_arXiv.txt": {
        "abstract": "{We observed the $^{12}$CO(3$\\rightarrow$2) emission of the emission-line regions Hubble\\,I, Hubble\\,V, Hubble\\,X, Holmberg~18, and the stellar emission-line object S28 in NGC6822 with the ESO Atacama Pathfinder Experiment (APEX) 12m telescope as part of its science verification. The very low system temperature of $130-180$~K enabled us to achieve detections in 4 single pointings and in a high spatial resolution $70''\\times 70''$ map of Hubble\\, V. We compare the spectra with H{\\sc i} observations, obtained with the Australia Telescope Compact Array, of the same regions. In combination with previous multi-line CO observations, we perform a preliminary investigation of the physical conditions in Hubble\\, V using a simple LTE model. We estimate the mass of the Hubble~V region and the H$_2/I_{\\rm CO(3\\rightarrow 2)}$ conversion factor. Also, we show that Hubble~V is located very near the line-width versus size relation traced by the Milky Way and LMC molecular clouds. ",
        "introduction": "We present $^{12}$CO(3$\\rightarrow$2) observations of the emission-line regions Hubble\\,I, Hubble\\,V, Hubble\\,X, Holmberg~18, and the stellar emission-line object S28 \\citep{kd} in NGC6822 with the ESO Atacama Pathfinder Experiment (APEX) 12m telescope\\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut f\\\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}. NGC6822 is a Local Group dwarf irregular galaxy, of type IB(s)m. Its study started with the landmark paper of \\cite{h25}, who identified it as a galaxy in its own right, external to the Milky Way. The most recent estimate of its distance, based on observations of 116 Cepheid variables, places it at $466 \\pm 20$~kpc \\citep{p04}. \\cite{ppr05} find a metallicity $12 + \\log($O/H$) = 8.37 \\pm 0.09$ for Hubble V and $12 + \\log($O/H$) = 8.19 \\pm 0.16$ for Hubble X, corresponding to roughly half the solar metallicity. Star formation proceeded at an almost constant rate up to the present, except for the central bar region, where star-formation increased by a factor of $3-4$ during the last 600~Myr \\citep{w01}. Its proximity allows us to study the different components and phases of its interstellar medium on scales of order 10$-$100 parsec.  The detection of the compact molecular clouds associated with Hubble\\,V was first reported by \\cite{w94}. Later on, the emission-line regions Hubble\\,I, Hubble\\,V, and Holmberg~18, and the stellar emission-line object S28 were observed in $^{12}$CO(1$\\rightarrow$0) emission by \\cite{i97}. Moreover, for Hubble\\,V, the brightest HII region in NGC6822, detections of $^{12}$CO(2$\\rightarrow$1), $^{12}$CO(3$\\rightarrow$2), $^{12}$CO(4$\\rightarrow$3), and $^{13}$CO(1$\\rightarrow$0) have been reported \\citep{i03}. $^{12}$CO(1$\\rightarrow$0) and $^{12}$CO(2$\\rightarrow$1) observations centered on Hubble\\,X, on the other hand, did not yield a detection \\citep{i03}. These results will serve as a comparison for the APEX data presented here.  In this Letter, we combine the $^{12}$CO(3$\\rightarrow$2) line intensities measured with APEX with line intensities of other $^{12}$CO and $^{13}$CO transitions, taken from the literature, to constrain the physical conditions of the molecular interstellar medium of NGC6822 using simple LTE models. We also investigate the spatial distribution of the $^{12}$CO(3$\\rightarrow$2) emission and how it correlates with previous high resolution HI observations.  \\begin{figure*} \\vspace*{9cm} \\special{hscale=35 vscale=35 hsize=540 vsize=560 hoffset=-20 voffset=45 angle=0 psfile=\"5411fig1a.ps\"} \\special{hscale=35 vscale=35 hsize=540 vsize=560 hoffset=235 voffset=45 angle=0 psfile=\"5411fig1b.ps\"} \\special{hscale=35 vscale=35 hsize=540 vsize=520 hoffset=-20 voffset=-70 angle=0 psfile=\"5411fig1c.ps\"} \\special{hscale=35 vscale=35 hsize=540 vsize=720 hoffset=235 voffset=-70 angle=0 psfile=\"5411fig1d.ps\"} \\caption{CO(3$\\rightarrow$2) spectra of the star-forming regions Hubble\\rm{I}, Hubble\\rm{X}, and Holmberg~18 and the stellar emission-line object S28 in NGC6822, overplotted with the best fitting Gaussian. The black curve indicates the H{\\sc i} emission as derived by \\cite{deblok06}. Evidently, the molecular gas associated with Hubble~I and Holmberg~18 has the same velocity as the neutral gas. The velocity of the molecular gas associated with the stellar object S28 differs from that of the H{\\sc i}. Hubble~X was not detected. \\label{HIV}} \\end{figure*} ",
        "conclusions": "\\begin{table*} \\caption{CO(3$\\rightarrow$2) properties of the targeted regions in NGC6822:~the peak main-beam temperature, $T_{\\rm mb}$, the velocity of the line with respect to the Local Standard of Rest, the line FWHM, and the integrated line intensity, $I_{\\rm CO}$ . \\label{tabres}} \\begin{center} \\begin{tabular}{|c|cccccc|} \\hline name & RA & DEC & $T_{\\rm mb}$ (K) & LSR velocity (km/s) & FWHM (km/s) & $I_{\\rm CO}$ (K~km~s$^{-1}$) \\\\ \\cline{1-7} Hubble\\rm{I} & 19 44 31.64 & -14 42 01.2& $0.07 \\pm 0.02$ &$-66.3 \\pm 0.4$ & $6.1 \\pm 1.1$ & $0.49 \\pm 0.11$ \\\\ Hubble\\rm{V} & 19 44 52.80 & -14 43 11.0 & $0.61 \\pm 0.02$ &$-41.3 \\pm 0.1$ & $6.0 \\pm 0.2$ & $3.89 \\pm 0.14$ \\\\ Hubble\\rm{V} & 19 44 52.80 & -14 43 11.0 & deconvolved &$-41.3 \\pm 0.4$ & $6.0 \\pm 0.2$ & $6.65 \\pm 0.59$ \\\\ Hubble\\rm{X} & 19 45 05.20 & -14 43 13.0 & & & & $< 0.30$ \\\\ Holmberg~18 & 19 44 48.93 & -14 52 38.0 & $0.11 \\pm 0.02$ &$-43.6 \\pm 0.2$ & $2.5 \\pm 0.4$ & $0.29 \\pm 0.07$ \\\\ KD82\\_S28 & 19 44 57.79 & -14 47 51.5 & $0.12 \\pm 0.02$ &$-57.5 \\pm 0.3$ & $5.2 \\pm 0.7$ & $0.67 \\pm 0.11$ \\\\ \\cline{1-7} \\end{tabular} \\end{center} \\end{table*}  We fitted Gaussians to the detected emission lines in order to estimate the peak intensity, $T_{\\rm mb}$ (K), and the integrated intensity, $I_{\\rm CO} = \\int T_{\\rm mb}(v)\\,dv$ (K~km~s$^{-1}$), of the $^{12}$CO(3$\\rightarrow$2) emission lines of the observed star-forming regions (see Table \\ref{tabres}). We used the best fitting Gaussian and the 1$\\sigma$ noise on the spectrum, estimated from the spectral region between $-100$ and 0~km/s and excluding the emission line, to generate 1000 new noisy spectra. These were analysed the same way as the original spectrum, allowing us to estimate the 1$\\sigma$ errors on these quantities. For the non-detected star-forming region Hubble\\rm{X}, we give a 3$\\sigma$ upper limit over a velocity width of 6~km/s.  \\subsection{Physical conditions in Hubble{\\rm V}}  \\begin{figure} \\vspace*{8cm} \\special{hscale=50 vscale=50 hsize=540 vsize=520 hoffset=-25 voffset=-37 angle=0 psfile=\"5411fig2.ps\"} \\caption{$^{12}{\\rm CO}(3\\rightarrow2)$ map of Hubble\\rm{V}. Each panel shows the brightness temperature $T_{\\rm mb}$ as a function of velocity with respect to the Local Standard of Rest (LSR). Some spectra still show some residual variations even after subtracting off a double sinusoidal baseline. Nearest neighbor panels overlap by 8$''$, i.e. about half of the APEX beam width at this frequency, so that next nearest neighbor panels are roughly independent. The CO source associated with Hubble{\\rm V} is clearly resolved in this map and is extended towards the north-east. \\label{HV}} \\end{figure}  We made a preliminary assessment of the physical conditions in the CO cloud associated with Hubble{\\rm V} assuming local thermodynamical equilibrium (LTE). In that case, there is one excitation temperature, $T_{\\rm ex}$ responsible for populating the energy levels of the $^{12}{\\rm CO}$ and $^{13}$CO isotopomers. This need of course not be the case in reality, with the higher-$J$ lines not being thermalized due to their larger Einstein $A$ coefficients. In the following, we will assume the isotopic ratio $X={^{12}{\\rm CO}}/{^{13}{\\rm CO}}=60$ \\citep{lp93,saz02}, in which case the optical depths of $^{12}{\\rm CO}$, denoted by $\\tau_{12}$, and $^{13}{\\rm CO}$, denoted by $\\tau_{13}$, obey the relation $\\tau_{12} = X\\,\\tau_{13}$. The calculated line intensities are coupled to the observed quantities by the beam filling factor $f_{\\rm b}$, which we assume to be the same for the $^{12}{\\rm CO}$ and $^{13}{\\rm CO}$ emission. Furthermore, we will assume that both the $^{12}{\\rm CO}$ and the $^{13}{\\rm CO}$ emission arises in the same region so that $\\Omega_{\\rm source}$, the solid angle spanned on the sky by the CO source, is the same for both isotopomers. This is to keep this preliminary modeling as simple as possible since there is no physical reason why $f_{\\rm b}$ and $\\Omega_{\\rm source}$ should be the same for all transitions. We can use this source solid angle to correct the observed emission line brightness temperatures for beam dilution using the relation $T'_{\\rm mb} = T_{\\rm mb} (\\Omega_{\\rm source} + \\Omega_{\\rm beam})/\\Omega_{\\rm source}$, with $\\Omega_{\\rm beam}$ the beam solid angle. The main-beam brightness temperature of an observed transition can be written as \\begin{equation} T'_{{\\rm mb},i} = f_{\\rm b}\\,\\left( 1 - e^{-\\tau_i} \\right) \\frac{h \\nu_i}{k} \\left( \\frac{1}{e^{h \\nu_i/kT_{\\rm ex}}-1} - \\frac{1}{e^{h \\nu_i/kT_{\\rm cmb}}-1} \\right), \\end{equation} with $\\nu_i$ the frequency of the transition, $\\tau_i$ its optical depth, and $T_{\\rm cmb} = 2.725$~K the background radiation temperature.  \\begin{table} \\caption{Comparison of the LTE model with the measured intensities of the $^{12}$CO and $^{13}$CO transitions. All intensities have been corrected for beam dilution using $\\Omega_{\\rm source} = 209\\,{\\rm arcsec}^2$. \\label{tablte}} \\begin{center} \\begin{tabular}{|ccc|} \\hline transition & measurement (K) & model (K) \\\\ \\cline{1-3} $^{12}$CO(1$\\rightarrow$0) & $1.72 \\pm 0.11$ & 1.73 \\\\ $^{12}$CO(2$\\rightarrow$1) & $1.69 \\pm 0.08$ & 1.66 \\\\ $^{12}$CO(3$\\rightarrow$2) & $1.56 \\pm 0.05$ & 1.57 \\\\ $^{12}$CO(4$\\rightarrow$3) & $1.50 \\pm 0.09$ & 1.48 \\\\ $^{13}$CO(1$\\rightarrow$0) & $0.11 \\pm 0.07$ & 0.11 \\\\ \\cline{1-3} \\end{tabular} \\end{center} \\end{table}  \\begin{figure} \\begin{center} \\includegraphics[width=7cm,angle=0]{5411fig3.eps} \\caption{Spatial distribution of the CO(3$\\rightarrow$2) emission of Hubble\\rm{V} at 18$''$ resolution (contours) plotted over an H$\\alpha$ image (greyscale). The square indicates the region mapped by our APEX observations. The contour values are 1 (which corresponds to 3$\\sigma$), 1.5, 2, 2.5, 3.0, and 3.5~K~km~s$^{-1}$. The CO emission is not centered on the H$\\alpha$ emission. The positions of the CO clouds MC1 and MC2, identified by \\cite{w94}, are indicated in the figure. \\label{spatial}} \\end{center} \\end{figure} Using a non-linear minimisation routine, we simultaneously fitted a Gaussian model for the spatial distribution of the CO emission of Hubble{\\rm V}, convolved with the APEX beam, which constrains the source solid angle $\\Omega_{\\rm source}$ and the parameters $T_{\\rm ex}$, $f_{\\rm b}$, and $\\tau_{12}$ to the $^{12}$CO(3$\\rightarrow$2) map, presented in fig. \\ref{HV}, and to the $^{12}$CO(3$\\rightarrow$2) brightness temperature measured by us, and the $^{12}$CO(1$\\rightarrow$0), $^{12}$CO(2$\\rightarrow$1), $^{12}$CO(4$\\rightarrow$3), and $^{13}$CO(1$\\rightarrow$0) brightness temperatures presented in \\cite{i03}.  We then used the best fitting values for $\\Omega_{\\rm source}$, $T_{\\rm ex}$, $f_{\\rm b}$, and $\\tau_{12}$ to generate 10000 mock data sets with added Gaussian noise on each of the observed quantities, using the measured 1$\\sigma$ uncertainties on the measured quantities as estimates for the dispersions of each of the noise distributions. These mock data sets were analysed the same way as the original set, allowing us to estimate the 1$\\sigma$ errors on the derived quantities. This way, we find that the parameter values $\\Omega_{\\rm source} = 209 \\pm 50\\,{\\rm arcsec}^2$, $T_{\\rm ex} = 49 \\pm 27\\,{\\rm K}$, $\\tau_{12} = 3.7 \\pm 2.3$, and $f_{\\rm b}= 0.04 \\pm 0.03$ provide the best fit to the whole data-set. There is a large degree of degeneracy between the parameters of this model, e.g. between $T_{\\rm ex}$ and $f_{\\rm b}$. This is reflected by the very large errorbars on these quantities. Still, the minimisation routine converges to the same solution independent of the starting point of the minimisation which proves that the minimum of the $\\chi^2$ is well defined. Moreover, this temperature estimate agrees with the dust temperature derived from the ratio of the $60 \\mu$ and $100 \\mu$ IRAS flux densities, $f_\\nu(60)=7.89$~Jy and $f_\\nu(100)=11.81$~Jy, of Hubble{\\rm V}, $T_{\\rm dust} \\approx 40$~K. This estimate was derived assuming a single temperature component and a $\\lambda^{-1}$ emissivity law. Given the apparently rather high temperature of this CO emission cloud, observations of higher-$J$ transitions, e.g. with FLASH or the future SIS heterodyne receivers, are required for a more precise assessment of its physical properties using more sophisticated LVG models, taking into account non-LTE effects. Also, some of the published high-$J$ transition temperatures, such as $^{12}$CO(4$\\rightarrow$3) value of \\cite{i03}, may be affected by the small area that was mapped. If some of the emission was missed, this may lead to an underestimation of the brightness temperature of these transitions.  Using the $^{12}$CO(3$\\rightarrow$2) FWHM line-width, $\\Delta V$, in km~s$^{-1}$, and the radius of the emission region, $R$, in parsec, we can also estimate the virial mass of Hubble~V as $M_{\\rm vir} \\approx 190 R (\\Delta V)^2$, in solar masses \\citep{m88}. For $\\Delta V = 6.0$~km~s$^{-1}$, $R = \\sqrt{\\Omega_{\\rm source}}/2.355 =6.1'' = 13.9$~pc, we find $M_{\\rm vir} \\approx 9.5 \\times 10^4\\,M_\\odot$. Using the relation $M_{\\rm dust} = 1.27 f_\\nu(100) D^2 (\\exp(144\\,K/T)-1)\\,M_\\odot$ \\citep{blg02} for the dust mass, with $D$ the distance in Mpc, we find $M_{\\rm dust} \\approx 130\\,M_\\odot$. Using the metallicity-dependent gas versus dust mass relation (eqn. (6) in \\cite{blg02}), this yields $M_{\\rm gas} \\approx 9 \\times 10^4\\,M_\\odot$, in good agreement with the virial mass. This is much less than the estimate of the total gas mass $M_{\\rm gas} = 10 \\pm 5 \\times 10^5\\,M_\\odot$ of \\cite{i03}. This is most likely because we are measuring quantities (radius, velocity dispersion) that pertain only to the CO emission region : any mass distribution outside this region, which was taken into account by \\cite{i03}, would not have a very large effect on the virial mass estimate. Hubble~V is also very near the line-width versus size relation traced by the Milky Way and LMC molecular clouds \\citep{h95,m90}. We estimate the $^{12}$CO(3$\\rightarrow$2) to H$_2$ conversion at $5.4-5.8 \\times 10^{20}$ cm$^{-2}$ (K\\,km\\,s$^{-1}$)$^{-1}$, depending on whether the deconvolved or the observed brightness is used.  \\subsection{Comparison with previous observations}  In Figures~\\ref{HIV} and \\ref{HV}, we plot the (unscaled) HI spectra derived by \\citet{deblok06}, of the same regions as our observations on top of the $^{12}$CO(3$\\rightarrow$2) spectra. One can see that apart from S28, both emission lines are coincidental although the HI emission is systematically broader than the $^{12}$CO(3$\\rightarrow$2) emission. In Figure~\\ref{spatial}, we plot the spatial distribution of the $^{12}$CO(3$\\rightarrow$2) emission of Hubble\\rm{V} at 18$''$ resolution. We find a different morphology than the one derived by \\citet{i03}, however it is in accordance with their [C\\rm{II}] emission, found in the same paper. A possible cause might be the higher system temperature of 2460K during their observation, producing more noise which might cause a shift in the spatial distribution. The main emission peak in our map corresponds to the molecular cloud MC2, first detected by \\cite{w94}; the eastward extension corresponds only very roughly to MC1.  With our observations we prove that APEX is very suited for deriving spatially extended, high signal-to-noise maps of emission-line regions in Local Group dwarf galaxies, where one can achieve a spatial resolution of a few tens of parsecs."
    },
    "0606/astro-ph0606141_arXiv.txt": {
        "abstract": "In the inner regions of an accretion disk around a black hole, relativistic protons can  interact with ambient matter to produce electrons, positrons and $\\gamma$-rays.  The resultant steady state electron and positron  particle distributions are self-consistently computed taking into account Coulomb and Compton cooling, $e^-e^+$ pair production (due to $\\gamma-\\gamma$ annihilation) and pair annihilation. While earlier works used the diffusion approximation to obtain the particle distributions, here  we solve a more general integro-differential equation that correctly takes into account the large change in particle energy that occur when the leptons Compton scatter off hard X-rays.  Thus this formalism can also be applied to the hard state of black hole systems, where the dominant ambient photons are hard X-rays.  The corresponding photon energy spectrum is calculated and compared with broadband data of black hole binaries in different spectral states.  The results indicate that the $\\gamma$-ray spectra ($E > 0.8$ MeV) of both the soft and hard spectral states and the entire hard X-ray/$\\gamma$-ray spectrum of the ultra-soft state, could be due to $p-p$ interactions.  These results are consistent with the hypothesis that there always exists in these systems a $\\gamma$-ray spectral component due to $p-p$ interactions which can contribute between $0.5$ to $10$\\% of the total bolometric luminosty. The model predicts that {\\it GLAST} would be able to detect black hole binaries and provide evidence for the presence of non-thermal protons which in turn would give insight into the energy dissipation process and jet formation in these systems. ",
        "introduction": "Black hole X-ray binaries are generally observed to be in two distinct states.    These states, which are named hard and soft, differ in their luminosity and spectral shapes.  In the hard state, which is in general less luminous, the spectrum of the system can be described as a hard power-law with a spectral index $\\Gamma \\sim 1.7$ and a cutoff around $100$ keV. This spectrum can be modeled  as thermal Comptonization of soft photons by a plasma having temperature $T \\sim 50$  keV \\citep[e.g.][]{Gier97}. In contrast, the spectrum during the soft state  consists of a blackbody-like component with $kT \\sim 1$ keV, which typically dominates the luminosity and is generally considered to be thermal emission from an optically thick accretion disk. Apart from this soft (or disk)  component, a hard X-ray power-law tail,  with  photon index $\\Gamma \\sim 2.5$ and no detectable cutoff up to  $\\sim 8$ MeV \\citep{mcc02}, is also observed.   There have been several interpretations of the high energy ($E > 200$ keV) emission from black hole systems. It may arise from a photon starved inner most region of a disk which cools due to bremsstrahlung self-Comptonization \\citep{Mel93} or as  emission due to $\\pi^0$ decay, which are created by  proton-proton interaction in a hot proton gas, $T > 10^{11}$ K \\citep{Kol79,Jou94}. While these possibilities maybe viable, they are based on assumptions of the geometry and physical properties of the system, which are not directly and independently verifiable, like the presence of a very hot proton gas or a photon starved region. Another interpretation is that the spectra arises due to Comptonization  of photons by the bulk motion of matter falling into the black hole \\citep{Lau99}.  However, \\cite{Nie06} have argued   that the non-detection of spectral breaks at $E < 500$ keV is contrary to this model's prediction.  A detailed radiative model, which has been used to fit good quality broad band data of black hole systems, is the hybrid model which is inscribed in a spectral fitting code called EQPAIR \\citep{cop98,Gier99}.  In the framework of this model, this high energy component arises from   a plasma  consisting of both thermal and non-thermal electrons  that Comptonize external soft photons. Detailed spectral modeling of the observed soft state spectra demands the coexistence  of thermal and non-thermal electrons in the emission region and hence the steady state non-thermal electron distribution is computed by assuming that there is an injection of non-thermal  particles into the system where they  cool by Comptonization and Coulomb interactions.    The origin of these non-thermal particles is uncertain. A possible site may be a corona on top of a cold disk which is heated by magnetic field reconnections \\citep{har93, pout99}, but the details of the process are largely unknown. The acceleration process has to be highly efficient to produce non-thermal electrons in an environment where electrons cool rapidly by inverse Comptonization. If this acceleration process is  mass-independent then protons are also expected  to be accelerated to relativistic energies, for example by scattering off magnetic \"kinks\" in  a Keplerian accretion disk \\citep{sub96}. Some of these non-thermal protons may escape from the system and contribute to the jet formation \\citep{sub99}. This is particularly interesting since there are some evidence that the X-ray producing region may be same as the base of the extended jet \\citep{Mar05}.  Hence the detection of these non-thermal protons will provide valuable clues to the nature of black hole systems.  Non-thermal protons would interact with the ambient thermal protons and produce electron-positron  pairs, which would Comptonize photons to high energies. These high energy photons would produce further pairs by $\\gamma-\\gamma$ interaction and a pair cascade would ensue. Pair cascades initiated by the injection of pairs (or equivalently high energy non-thermal electrons)  have been extensively examined \\citep[e.g.][]{zdzlight,sven}. The effect of  $p-p$ interactions and  the resultant spectra have also been computed and studied \\citep[e.g.][]{sss,E80,E83,maha,sera,zdz86}.  These works, in general, did not consider the presence of copious photons or have assumed that the ambient photons are in the UV range, a scenario relevant to AGN and under-luminous black hole systems.  However, for black hole binaries, the system is dominated by either soft or hard X-rays depending on the spectral state. One of the important radiative interaction that would occur in such an environment is the inverse Compton scattering of X-rays by pairs with Lorentz factor $\\gamma \\approx 200$ that are produced by the $p-p$ interaction.  The standard methods to compute the inverse Compton spectra, assume that the interaction in the rest frame of the electron takes place in the non-relativistic Thompson limit, which is true only if electron Lorentz factor  times the photon energy $\\gamma \\epsilon$ $ \\ll m_ec^2$.  This assumption is violated if the ambient photon energy is $ \\le m_ec^2/\\gamma \\approx 2$ keV.  Thus, when \\citet{bhat03} considered $p-p$ interaction in the presence of blackbody  photons having temperature $T \\approx 1$ keV, they used  a general formalism to describe the inverse Compton process given by \\citet{blu70}. They found that for such a situation, which is relevant to the soft state of black hole binaries, the effect of non-thermal protons is to produce a broad feature around $1-50$ MeV. Using the observed OSSE data for GRS 1915+105, they could constrain the fraction of non-thermal protons in the system to be $< 5\\%$.  Although \\citet{bhat03}, computed the change in energy of the photon (and hence the change in energy of the lepton) appropriately in  the Klein-Nishina regime,  they used, for simplicity, a diffusion equation to describe the kinetic evolution of the pairs, which intrinsically assumes that the change in energy of the particle per scattering is small. This assumption breaks down when the ambient photon energy is in X-rays, especially when there are a copious amount of ambient hard X-rays.  In this work, we extend the formalism developed by \\citet{bhat03}, by solving an integro-differential equation for the pair kinetic evolution which correctly takes into account the large energy changes that a lepton undergoes upon scattering with an X-ray photon. This not only allows for a more accurate estimation of the emergent spectra for systems in the soft state, but enables the scheme to be applied to the hard state also.  In \\S $2$ we describe the model and the assumptions made to compute the steady state non-thermal  electron/positron distributions and the resultant photon spectra. In \\S $3$ some general results of the computation are presented along with comparison with observations of black hole systems Cyg X-1  and GRS 1915+105 in different spectral states.  The main results of the work are summarized and discussed in \\S $4$. ",
        "conclusions": "A self-consistent scheme, which computes the electron/positron and radiation energy distribution inside a thermal plasma having non-thermal protons, is developed. The non-thermal protons interact via $p-p$ collisions to produce electron/positron pairs and $\\gamma$-rays.  These high energy pairs cool by inverse Comptonization of ambient photons. The effect of subsequent pair production due to photon-photon interactions and annihilation of positrons with ambient electrons, is taken into account. Unlike previous works, the scheme does not assume that the energy change of the leptons during inverse Compton scattering is small. Hence it can be applied to black hole binary systems, where the dominant ambient photons are X-rays.  Comparison of the computed spectra with the broad band X-ray/$\\gamma$-ray spectra of black hole binaries in different spectral states reveal:  \\noindent $\\bullet$ The hard X-ray spectra (3--800 keV) of  the soft and Very High state (VHS), are too steep to be explained as emission due to $p-p$ interactions. However, the observed $\\gamma$-ray spectra (0.8--8 MeV), especially for the soft state,  could be due to such interactions.  In this interpretation, an unknown acceleration process energises both electron and protons, producing their non-thermal distributions. While the electron non-thermal distribution gives rise to the hard X-ray emission, the protons are the origin of the $\\gamma$-rays. The powers going into accelerating protons and electrons are within a factor of two.  \\noindent $\\bullet$ For the ultra-soft state, the observed hard X-ray spectra ( E $> 5$ keV) could be entirely due to $p-p$ interactions. This interpretation gives a natural explanation for the similar hard X-ray spectral slope ($\\Gamma \\approx 2$) observed whenever a black hole is in the ultra-soft state.  \\noindent $\\bullet$ For the hard state, the observed $\\gamma$-ray spectrum (0.5--8 MeV) can be explained as emission due to  $p-p$ interactions alone. The predicted steep spectral shape in this energy range, is not sensitive to the model parameters, and matches well with the observations.  \\noindent $\\bullet$ For both the soft and hard states, the model predicts that for reasonable parameters, {\\it GLAST} should be able to detect black hole binaries. This in contrast to the situation when only non-thermal electrons are present in the system, where very low compactness and large maximum Lorentz factor of the electrons have to be postulated, in order for {\\it GLAST} to make a similar detection.  These results are consistent with the hypothesis that there always exists a spectral component due to a non-thermal proton distribution in black hole binaries.  This component peaks at $\\gamma$-rays and can contribute between $0.5$ to $10$\\% of the bolometric luminosity. In the ultra-soft state, this component is visible for $E > 5$ keV. During the soft and hard states, the emission is detectable only when $E > 0.8$ MeV, since  other spectral  components dominate at lower energies. This hypothesis can be verified using future observations by {\\it GLAST}.  The scheme does not include scattering of  high energy photons  with thermal electrons. Although the Klein-Nishina cross section decreases with photon energy, such scatterings can be important  especially when the Thompson optical depth is significantly greater than unity, a case more pertaining to the soft state. In fact, the inability of the present model to explain the hard X-ray emission during the soft state, may be an artifact of this assumption. If that is true, then it may be possible that only protons are accelerated in black hole binaries and not electrons.  This is theoretically appealing given the complexities of  accelerating electrons to high energies in a region where  inverse Compton cooling is efficient. However, several sophisticated modification of the present scheme have to be undertaken before this speculation can be tested.  Evidence of non-thermal protons in black hole binaries, would shed light on the energy dissipation process that occur in the  inner regions of the accretion disk. These energetic protons could also be the origin of the outflows/jets that are observed in many of these systems. Such insights may finally lead to a comprehensive physical picture of these enigmatic sources."
    },
    "0606/astro-ph0606694_arXiv.txt": {
        "abstract": "A  magnetic helicity integral is proposed  which can be applied to domains which are not magnetically closed,  i.e.~have a non-vanishing normal component of the magnetic field on the boundary. In contrast to the relative helicity integral, which was previously suggested for magnetically open domains, it does not rely on a reference field and thus avoids all problems related to the choice of a particular reference field. Instead it uses a gauge condition on the vector potential, which corresponds to a particular topologically unique closure of the magnetic field in the external space. The integral has additional elegant properties and is easy to compute numerically in practice.  For magnetically closed domains it reduces to the classical helicity integral. ",
        "introduction": "Magnetic helicity  is an important quantity in describing the structure and evolution of magnetic fields in many fields of physics, in particular in plasma physics and astrophysics. It was introduced to Plasma Physics  by H.K.~Moffatt in \\cite{Moffatt69} and was originally defined as an integral over a magnetically closed volume, i.e.~a volume for which  the normal component of the magnetic field (${\\bf B}$) vanishes on the boundary $\\partial V$: \\begin{equation} H({\\bf B}) := \\int_{V} {\\bf A} \\!\\cdot\\! {\\bf B} \\ dV \\ ; \\quad B_{n}= {\\bf B}\\!\\cdot\\! {\\bf n}\\vert_{\\partial V}=0 . \\label{totalhelicity} \\end{equation} Here ${\\bf A}$ is the vector potential, $\\nabla \\times {\\bf A} = {\\bf B}$, of the magnetic field. The integral measures - roughly speaking -  the Gaussian linkage of magnetic flux within $V$. More precisely, it is the asymptotic linking number of pairs of field lines averaged over the volume  \\cite{Arnold86}. It is an  important property  of this integral that we can derive an equation of continuity for the helicity density, which uses only the homogenous Maxwell's equations (here ${\\bf E}= - \\partial_t {\\bf A}- \\nabla \\phi$ is the electric field): \\begin{equation} \\partial_{t} ({\\bf A} \\!\\cdot\\! {\\bf B}) + \\nabla \\!\\cdot\\! \\left( \\phi {\\bf B} + {\\bf E} \\times {\\bf A} \\right) = -2 \\ {\\bf E} \\!\\cdot\\! {\\bf B}  \\  . \\label{helevolu} \\end{equation} It can be shown that the integral is a topological invariant, i.e.~it does not change under a deformation of the field within $V$, as given for instance by the motion of a magnetic field embedded  in an ideal plasma, satisfying (${\\bf v}$ is the plasma velocity) \\begin{equation} {\\bf E} + \\mathbf{v}\\times\\mathbf{B}  =  0   \\label{idealevolu} \\ . \\end{equation} Under such a condition,  (\\ref{helevolu}) becomes \\begin{equation} \\partial_{t} ({\\bf A} \\!\\cdot\\! {\\bf B}) + \\nabla \\!\\cdot\\! \\left( (\\phi -{\\bf v}\\!\\cdot\\! {\\bf A}) {\\bf B} + {\\bf v}  {\\bf A}\\!\\cdot\\!{\\bf B} \\right)  =   0   \\ , \\label{idealevoluhel} \\end{equation} so that integrating over a volume with ${\\bf v}\\!\\cdot\\! {\\bf n}=0$ on the boundary (or more generally a comoving volume) results in \\begin{eqnarray} \\frac{d}{dt} \\int_{V} {\\bf A} \\!\\cdot\\! {\\bf B} \\ dV & = & \\int_{V} \\partial_{t} ({\\bf A} \\!\\cdot\\! {\\bf B}) + \\nabla  \\!\\cdot\\!  ({\\bf v} \\  {\\bf A}\\!\\cdot\\!{\\bf B} ) \\ dV  \\nonumber \\\\ & =&  - \\int_{\\partial V}\\!\\!\\! (\\phi -{\\bf v}\\!\\cdot\\! {\\bf A}) {\\bf B}\\!\\cdot\\!{\\bf n}  \\, da  =  0. \\label{changehel2} \\end{eqnarray} Moreover, the total helicity is often an approximate invariant for non-ideal plasmas \\cite{Berger84}, and is therefore a valuable tool in determining the evolution of many technical and natural plasmas. One of the earliest results was the prediction of the relaxed state of a Reversed-Field Pinch \\cite{Taylor:RelaxationReconnection}, but there are many more applications (see  \\cite{HelicitySpaceLab} for an overview).  However, the boundary condition $B_{n}=0$ on the integral (\\ref{totalhelicity}), which is necessary to ensure gauge invariance, means that it can not  be applied to cases where the magnetic field crosses the boundary. Typical examples are the vacuum vessels of laboratory plasmas where an external magnetic field crosses the boundaries, or the atmospheres of stars or planets, where the studied volume is usually bounded by the surface of the body, through which the magnetic field emerges.  In such cases it was previously necessary to resort to the calculation of the relative helicity, i.e.~the helicity was calculated with respect to a reference field ${\\bf B}_{\\rm ref}$ satisfying the same boundary conditions. One can prove \\cite{FinnAntonsen,Berger1984} that for an arbitrary closure of the magnetic field outside $V$, denoted by ${\\bf B}_{\\rm ext}$, the relative helicity \\begin{eqnarray} H({\\bf B} \\vert {\\bf B}_{\\rm ref}) & = & H({\\bf B}+{\\bf B}_{\\rm ext})-H({\\bf B}_{\\rm ref}+{\\bf B}_{\\rm ext}) \\\\ & = & \\int ({\\bf A} + {\\bf A}_{\\rm ref}) \\cdot ({\\bf B} -{\\bf B}_{\\rm ref}) \\ dV \\ , \\end{eqnarray} is actually independent of the external closure of the field.  The reference field is in most cases choosen to be  a potential field (see e.g.~\\cite{BergerRuzmaikin:HelicityProduction}) since a potential field is easy to compute and  physically distinguished as the lowest energy state compatible with the boundary conditions. The introduction of a reference field, however, not only complicates the calculation of magnetic helicity, but also complicates  its already difficult  interpretation.  For instance, the question arises as to whether a change of relative helicity in a volume has a physical meaning, or whether it is only due to our particular choice of reference field.  In this contribution it is proposed to replace  the reference field by a more general boundary condition on the vector potential and it is shown that this leads to a well defined quantity. ",
        "conclusions": "In this letter it was shown how the total helicity integral can naturally be generalized to allow for magnetic fields which are not closed within the domain, i.e.~which have a non-vanishing normal component on the boundary. The construction does not require an explicit reference field as the relative helicity integral does, which was previously used in this situation. Instead we have a gauge condition for ${\\bf A}$ on the boundary which corresponds to closing the domain with a topologically unique field. This field is an external complement with zero helicity density to the field in the given domain. The new integral  has all desirable properties, i.e.~it is gauge invariant, topologically invariant, and it reduces to the total helicity whenever the latter is well defined. Moreover, it shows the proper additivity with respect  to  fields  and complementary volumes. This facilitates not only many calculations of helicity, but also its interpretation."
    },
    "0606/gr-qc0606025_arXiv.txt": {
        "abstract": "We use covariant and first-order formalism techniques to study the properties of general relativistic cosmology in three dimensions. The covariant approach provides an irreducible decomposition of the relativistic equations, which allows for a mathematically compact and physically transparent description of the 3-dimensional spacetimes. Using this information we review the features of homogeneous and isotropic 3-d cosmologies, provide a number of new solutions and study gauge invariant perturbations around them. The first-order formalism is then used to provide a detailed study of the most general 3-d spacetimes containing perfect-fluid matter. Assuming the material content to be dust with comoving spatial 2-velocities, we find the general solution of the Einstein equations with non-zero (and zero) cosmological constant and generalise known solutions of Kriele and the 3-d counterparts of the Szekeres solutions. In the case of a non-comoving dust fluid we find the general solution in the case of one non-zero fluid velocity component. We consider the asymptotic behaviour of the families of 3-d cosmologies with rotation and shear and analyse their singular structure. We also provide the general solution for cosmologies with one spacelike Killing vector, find solutions for cosmologies containing scalar fields and identify all the PP-wave 2+1 spacetimes. ",
        "introduction": "General relativity in three spacetime dimensions is known to possess a number of special simplifying features: there are no gravitational waves, no black holes in the absence of a negative cosmological constant, the Weyl curvature is identically zero, and the weak-field limit of the theory does not correspond to Newtonian gravity in two space dimensions~\\cite{GAK}-\\cite% {teit}. The theory is therefore considerably 'smaller' than general relativity spacetimes with four (or more) dimensions, and the strong-energy condition that creates geodesic focussing does not depend on the density of the material sources. These simplifying features mean that considerable progress can be made in the search for the general cosmological solution of the three-dimensional Einstein equations. In an $(N+1)$-dimensional spacetime the number of independently arbitrary $N$-dimensional functions of the space coordinates that are needed to specify the Cauchy data for the general cosmological problem on a spacelike hypersurface in vacuum is $% (N+1)(N-2)$; in the presence of a general (non-comoving) perfect fluid it is $N^{2}-1;$ and for a comoving perfect fluid it is $N^{2}-N-1$~\\cite{BBL}. Thus, in the $N=2$ case, we see that the number reduces to zero for the vacuum solution (reflecting the absence of free gravitational fields in vacuum), reduces to one arbitrary spatial function in the comoving perfect-fluid case ,and to three arbitrary spatial functions for a perfect fluid.  In this paper we will set up the general cosmological problem in three-dimensional spacetimes and find the general solution of the field equations in the case of comoving pressure-free matter, with and without a cosmological constant, $\\Lambda $. We go on to find solutions for the case of non-comoving dust and classify the singularities and asymptotic behaviours that arise in both cases with and without a cosmological constant. The relative tractability of the general cosmological problem in 2+1 dimensions allows us to go some way towards finding a general solution of the Einstein equations and we are able to isolate those features which prevent a full solution being found. In particular, we are able to find and classify the solutions for dust containing one of the (two possible) non-zero spatial 2-velocity components.  There have been several past investigations of the structure of 2+1 dimensional general relativity and studies of the properties of particular solutions with high symmetry (see~\\cite{GAK}-\\cite{D}, and~\\cite{Clema}-\\cite% {Ida}). Important motivations for these studies were provided by the astrophysical interest in the possible observational signatures of cosmic strings and domain walls in the universe \\cite{vil}-\\cite{his}. Higher-order curvature contributions were discussed in~\\cite{BBL}, together with the special features of the Newtonian-relativistic correspondence in general relativity and related theories, while the study of quantum gravity is reviewed in~\\cite{carl}. Cosmological solutions and singularities were discussed in~\\cite{BBL} and~\\cite{collas,gar1}, while gravitational collapse of spherically symmetric dust clouds have been considered in~\\cite{Kriele}-% \\cite{gutti}.  The outline of this paper is as follows. In section 2 we define the 3-d Einstein equations and our notations. Section 3 introduces the 2+1 covariant formalism and the general kinematics of 3-d spacetimes, identifying the special features that arise from the lower dimensions and from the vanishing of the Weyl curvature. These include the key role of the isotropic pressure as the sole contributor to the gravitational mass of the system and the fact that vorticity never increases with time. In section 4 we give a number of new cosmological solutions, review the characteristics of the homogeneous and isotropic models, including those that are singularity-free, and provide the generalisation of the G\\\"{o}del universe to 3 dimensions. We also consider linear perturbations around the 3-d analogues of the `dust'-dominated FRW models and find them to be (neutrally) stable. In section 5 we employ Witten's first-order formalism~\\cite{witten} to formulate the equations governing the most general 3-d cosmological spacetime metric containing perfect-fluid matter. Then, in section 6 we specialise the matter source to pressure-free dust with non-zero $\\Lambda $ and comoving 2-velocities and find the general solution of the field equations. These fall into three classes, one of which generalises the solution of Kriele~\\cite{Kriele} to $\\Lambda \\neq 0$, while another is the generalisation of the Szekeres metric with nonzero $\\Lambda $~to 2+1 dimensions \\cite{szek, BSS}. Section 7 considers the most general dust cosmologies with non-comoving velocities and finds various new classes of solutions. We study the asymptotic behaviour of these solutions and analyse in detail the structure of their spacelike and timelike singularities. Also, by means of a number of examples, we illustrate the wide range of possible behaviours in the presence of vorticity and shear. The same section also introduces a transformation that generates exact solutions with nonzero cosmological constant from those with vanishing $\\Lambda $. Finally, in section 8 we look at the case of a pure scalar field, provide the general solution of Einstein's equations with one spacelike Killing vector, and identify all the 2+1 PP-wave spacetimes. Our results are summarised and discussed in section 9. ",
        "conclusions": "Whereas the general cosmological solutions of the 3+1 dimensional Einstein equations are intractably complicated and likely dominated by non-integrability, the structure of the theory in 2+1 offers the possibility of making considerable progress towards finding the general solution in several interesting situations. This fact, together with our current perception that quantum field theory fits more naturally in three rather than four dimensions, has motivated the study of Einstein's theory in 3-dimensional spacetimes.  In this article we employed covariant and first-order formalism techniques to study the properties of general relativity in three dimensions. The covariant approach provided an irreducible decomposition of the relativistic equations and allowed for a mathematically compact and physically transparent description of their properties. Using this information we reviewed the kinematical, dynamical and geometrical features of 3-dimensional spacetimes and identified the special features that distinguish them from the standard 3+1 models. These include the key role of the isotropic pressure as the sole contributor to the gravitational mass of the system and the fact that vorticity never increases with time. We also reviewed the 3-d analogues of the spatially homogeneous and isotropic FRW models and investigated their stability against linear perturbations. We found that, unlike their conventional counterparts, dust-dominated 3-d homogeneous and isotropic spacetimes are stable under shear and vorticity distortions and (neutrally) stable against disturbances in their density distribution. The latter reflects the vanishing of the total gravitational mass in 3-dimensional dust models, which ensures the absence of linear Jeans-type instabilities. In addition to isotropic spacetimes, we also looked at 3-dimensional anisotropic models providing Kasner-like solutions for the case of pressure-free matter and generalising G\\\"{o}del's universe to three dimensions. The covariant formalism allowed us to carry out these analyses by a study of the kinematic variables characterising the expansion of the universe. The absence of both electric and magnetic Weyl curvature components in three dimensions considerably simplifies the analysis. We then specialised further to the case of a pressureless matter source. In addition to being physically realistic, this assumption produces a significant further simplification of the cosmological field equations in three-dimensional spacetimes. We were able to find the general cosmological solutions of the theory in the case where the matter was comoving. No symmetry assumptions were made. We then considered the fully general pressureless fluid system with non-comoving velocities. We were able to solve the system in the case where one spatial velocity component was zero whilst the other was non-zero. This allowed us to carry out an asymptotic study, close to and far from singularities, of an inhomogeneous cosmology with rotation, expansion and shear. All the singularities arising in these solutions were classified using the different criteria of strength introduced by Krolak and Tipler. We were able to provide a simple transformation which generalised all the solutions we found with vanishing cosmological constant into new solutions with non-zero cosmological constant. Finally, we considered scalar-field metric with one Killing vector and found all the PP-wave solutions in 2+1 dimensional universes.  These investigations suggestion a number of problems for further study. Exact solutions in the cases with non-zero isotropic and anisotropic pressure remain to be investigated. In the case of zero pressure, we have analysed the problem of the general solution of the three-dimensional Einstein equations into a well-defined system of partial differential equations. We have solved for the case with comoving velocities and a single non-comoving velocity but the problem remains to find the general solution of the equations when both non-comoving fluid velocities are present.  \\subsection*{Acknowledgement}  D. Shaw is supported by the PPARC."
    },
    "0606/astro-ph0606062_arXiv.txt": {
        "abstract": "Recent work on globular cluster systems in dwarf galaxies outside the Local Group is reviewed.  Recent large imaging surveys with the {\\it Hubble Space Telescope} and follow-up spectroscopy with 8-m class telescopes now allow us to compare the properties of massive star clusters in a wide range of galaxy types and environments.  This body of work provides important constraints for theories of galaxy and star cluster formation and evolution. ",
        "introduction": "\\label{sec:intro}  Studies of globular clusters (GCs) in dwarf galaxies provide very important insights into galaxy formation, the formation and evolution of GCs, and the relationship between GCs and nuclei. Comparisons of the properties of star clusters in different types of galaxies can test the theories of galaxy formation.  In hierarchical scenarios of galaxy formation dwarf-size galaxies form first and then merge into larger systems.  If star cluster formation coincided with galaxy formation, then a significant fraction of the star clusters in massive galaxies should have been formed in dwarfs.  In this case the star clusters in dwarf galaxies in dense environments should be at least as old and metal-poor as the oldest star clusters in giant galaxies. However, recently evidence has mounted that stellar populations in surviving low mass galaxies are younger than in giant ellipticals\\cite{treu05}.  In this ``downsizing'' view the dwarf galaxies formed after the giants or at least had their star formation rates suppressed at early times. A signature of downsizing would be that the star clusters in dwarfs are younger than those in giant galaxies.  Another question that star clusters can help answer is the relationship between dwarf irregular (dI) and dwarf elliptical (dE) galaxies.  All dwarf galaxies must have formed with substantial gas fractions like today's dI galaxies.  However, in massive local galaxies clusters the majority of the dwarfs are gas-free, smooth-isophote dEs.  The differences may be due to environment or dIs may get transformed into dEs by gas stripping, supernovae winds, or galaxy interactions.  A comparison of the star clusters in the two types of dwarfs provides insight into the processes that shaped these galaxies and into why some dEs form nuclei.  In addition, the shape of the initial mass function of star clusters and how it evolves is not well understood.  There are still debates about whether the form of initial mass function is a single or broken power-law (resulting in a log-normal distribution in magnitudes) and about the effects of various destruction processes \\cite{baum98}\\cite{fz01}\\cite{vesp01}.  By comparing the present-day mass functions in dwarf galaxies with those in giant galaxies it may be to disentangle the destructive processes and therefore determine the shape of the initial star cluster mass function.  This paper reviews the properties of star cluster systems in dwarf galaxies outside of the Local Group.  Large imaging surveys with the {\\it Hubble Space Telescope} are now starting to provide us with statistically significant samples of GCs in dEs and dIs in different environments.  Follow-up spectroscopy with 8-m class telescopes are now providing complementary results on the ages, metallicities, abundance ratios, and kinematics of GCs and nuclei in dwarf galaxies. ",
        "conclusions": "\\label{sec:conc}  Substantial progress in understanding the GCSs of dwarf galaxies has been made in recent years due to large imaging surveys in different environments with {\\it HST}, new spectroscopic work using 8-m class telescope, and the inclusion of globular clusters in cosmological galaxy formation models.    More work is still needed on photometry and spectroscopy of GCs in dIs in order to improve the comparisons with the results on dEs.  Also, there is much to be learned from the kinematics of GCs that could not be discussed here.  GCs will continue to be a fundamental tool for understanding the formation of dwarf galaxies and testing theories of galaxy formation in general.\\\\  \\noindent This work was supported by the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., on behalf of the international Gemini partnership of Argentina, Australia, Brazil, Canada, Chile, the United Kingdom, and the United States of America."
    },
    "0606/astro-ph0606581_arXiv.txt": {
        "abstract": "Bearing in mind the application to the theory of core-collapse supernovae, we performed a global linear analysis on the stability of spherically symmetric accretion flows through a standing shock wave onto a proto neutron star. As unperturbed flows, we adopted the spherically symmetric steady solutions obtained with realistic equation of state and formulae for neutrino reaction rates taken into account. These solutions are characterized by the mass accretion rate and neutrino luminosity. Then we solved the equations for linear perturbations numerically, and obtained the eigen frequencies and eigen functions. We found (1) the flows are stable for all modes if the neutrino luminosity is lower than a certain value, e.g. $\\sim 1\\times 10^{52}$ ergs/s for $\\dot{M}=1.0M_{\\odot}/{\\rm s}$. (2) For larger luminosities, the non-radial instabilities are induced, probably via the advection-acoustic cycles. Interestingly, the modes with $\\ell=2$ and $3$ become unstable at first for relatively low neutrino luminosities, e.g. $\\sim 2-3\\times 10^{52}$ ergs/s for the same accretion rate, whereas the $\\ell=1$ mode is the most unstable for higher luminosities, $\\sim 3-7\\times 10^{52}$ ergs/s. These are all oscillatory modes. (3) For still larger luminosities, $\\gtrsim 7\\times 10^{52}$ ergs/s for $\\dot{M}=1.0M_{\\odot}/{\\rm s}$, non-oscillatory modes, both radial and non-radial, become unstable. These non-radial modes were identified as convection. The growth rates of the convective modes have a peak at $\\ell=5-11$, depending on the luminosity. We confirmed the results obtained by numerical simulations that the instabilities induced by the advection-acoustic cycles are more important than the convection for lower neutrino luminosities. Furthermore, we investigated the effects of the inner boundary on the stabilities and found that they are not negligible though the existence of the instabilities is not changed qualitatively for a variety of conditions. ",
        "introduction": "The explosion mechanism of core-collapse supernova is still remaining to be revealed. Several observations suggest that the supernova explosion is asymmetric in general. For example, the HST images of supernova 1987A, the most thoroughly investigated supernova so far, shows clear global asymmetry of the ejecta. The radiations scattered by the supernova ejecta commonly have one percent level of linear polarization, the fact naturally explained by the asymmetry of the scattering surface~\\citep{leo00, wan02}. Furthermore, it is now well known that young pulsars have a large peculiar velocity \\citep{lyn94}, which is supposed to be gained by a nascent neutron star at its birth owing to the asymmetry of dynamics. On the theoretical side, the possible roles of the asymmetry in the explosion mechanism are being explored from various points of view at the moment~\\citep{kot06}.  As a result of many multi-dimensional studies, it is now believed that the dynamics of supernova core is intrinsically asymmetric even though extrinsic factors like the rotation or magnetic field of progenitor is absent. The asymmetry is thought to be produced by some hydrodynamical instabilities, the most well known of which is convection. A couple of recent numerical simulations demonstrate, on the other hand, that another instability plays an important role in the accretion flows through the standing shock onto the proto neutron star \\citep{blo03,sch04,blo06,ohn06a}. The remarkable feature of this instability is the dominance of the $\\ell=1$ mode, where $\\ell$ stands for the azimuthal index of the spherical harmonic functions, $Y_{\\ell}^{m}$. The so-called standing accretion shock instability or SASI is thus expected to be a natural explanation for the large kick velocities of young pulsars.  The mechanism of SASI is still controversial. One of the promising mechanisms is the amplified cycle of the inward advections of vortex and entropy fluctuations and outward propagations of acoustic perturbations, which was originally discussed in the context of black hole accretions \\citep{fog00,fog01,fog02}. The mechanism was also known in the engineering field (see \\citet{how75} and references therein). Most recently, however, \\citet{blo06} claimed, based on their 2D numerical simulations with a simplified cooling taken into account, that the instability is induced by the purely acoustic cycles. On the other hand, \\citet{ohn06a} obtained the results that appear to support the advection-acoustic cycle in their 2D numerical models with more realistic heating and cooling processes taken into consideration. Obviously, more detailed linear analyses for appropriate background models and boundary conditions are needed to clarify the nature of the instabilities and that is the main purpose of this paper, in which we accomplished global linear analyses for the spherically symmetric steady accretion flows that are supposed to represent the post-bounce supernova core reasonably well. It should be emphasized that SASI is a non-local instability.  As mentioned above, the convection has been numerically explored by many researchers for the past decades. If the neutrino luminosity is larger than a certain value, there appears a so-called heating region in the post-shock flows, where the net neutrino heating occurs~\\citep{yam06}. In this region, the entropy gradient is negative and the flow is convectively unstable if the conventional criterion for convection is applied. Recently, \\citet{fog06a} demonstrated by linear analysis that the advection tends to suppress the convection. In fact, the condition for the instability is tightened and the growth rate is lowered in general. In this paper, SASI was not discussed. We think, however, the SASI and convection should be treated simultaneously and on an equal basis. It is also noted that it is important to employ a realistic equation of state and neutrino heating/cooling rates, since not only the evolution of perturbations but also the unperturbed background structure is dictated by them. In fact, the stability strongly depends on the latter. For example, the radial distributions of sound velocity and inflow velocity are crucially important for the advection-acoustic cycle. The growth rate and stability criterion for convection, on the other hand, depend not only on the entropy gradient but also on the advection time of the flow as discussed by \\citet{fog06a}. Moreover, the stability of accretion flows through a standing shock against radial perturbations is mainly determined by the acceleration or deceleration of the flow just behind the shock wave \\citep{nak92,nak93,nak94}. It is, hence, important to employ appropriate unperturbed background flows and boundary conditions. Some of the discrepancy mentioned above may be attributed to the difference of them.  In this paper, we accomplished global linear analyses of spherically symmetric steady accretion flows through a standing accretion shock onto a proto neutron star, using a realistic equation of state and neutrino-heating/cooling rates, and investigated the stability of the flows against both radial and non-radial perturbations. In so doing, we treated the convective instabilities and SASI on an equal basis. In the next section, we summarize the formulations, describing briefly the background steady solutions and then giving basic equations for the global linear analysis. The main results of the numerical calculations are presented in section~\\ref{eigen}. We discuss in section~\\ref{mec} the mechanisms of instabilities found in the previous section. In section~\\ref{dep}, we consider the dependence of the results on the assumptions we make for the inner boundary conditions and the perturbation of neutrino temperature. The final section is devoted to discussion. ",
        "conclusions": "\\label{disc}  In this paper, we investigated the stability of the spherically symmetric accretion flow through the standing shock wave onto the proto neutron star, which is supposed to approximate the post-bounce situation in the core-collapse supernova. We performed systematically the global linear-stability analyses for both radial and non-radial perturbations. We found that the flow is stable for all modes if the neutrino luminosity is lower than $\\sim 1\\times 10^{52}$ ergs/s in our models with $\\dot{M}=1.0M_{\\odot}/{\\rm s}$. For larger luminosities, the non-radial instabilities are induced, probably via the advection-acoustic cycles. The modes with $\\ell=2$ and $3$ become  unstable at first for relatively low neutrino luminosities, e.g. $\\sim 2-3\\times 10^{52}$ ergs/s for the same accretion rate, whereas the $\\ell=1$ mode is the most unstable for higher luminosities, $\\sim 3-7\\times 10^{52}$ ergs/s. These are all oscillatory modes. For still higher luminosities, $\\gtrsim 7\\times 10^{52}$ ergs/s, the non-oscillatory modes become unstable, among which we identified the non-radial ones as convection. In this way, we discussed the convective instability in the presence of advection on the same basis. We found that the growth rates of the convective modes change gradually with $\\ell$ and have a peak at $\\ell=5-11$, depending on the luminosity. The luminosity, at which each convective mode becomes unstable, decreases with $\\ell$ up to $\\ell=6$ and then increases for the same accretion rate. We confirmed the results obtained by numerical simulations that the instabilities induced by the advection-acoustic cycles become more important than convections for lower neutrino luminosities. We modified the criterion for the convection in the presence of advection previously discussed by \\citet{fog06a}. Furthermore, we investigated the sensitivity of the results to the inner boundary conditions and found that the nature of the instabilities is not changed qualitatively.  In most of realistic simulations, the neutrino luminosity seems to be much smaller than the critical value (see e.g. \\citep{jan05,sum05}). Hence, from the results of this paper, the over-stabilizing modes induced probably by the advection-acoustic cycles are more important than convections in causing the global anisotropy as observed in the ejecta of core-collapse supernovae \\citep{leo00,wan02}. The instability is normally saturated at some amplitudes owing to the non-linear coupling of various modes~\\citep{ohn06a}. Although the instability of the standing shock wave, whatever the cause, is helpful for the shock revival, recent multi-dimensional numerical simulations \\citep{jan05} demonstrated that these instabilities alone could not induce an explosion. Then there must be some other agents to further boost the shock revival. Very recently \\citet{bur06} proposed a new mechanism, in which the dissipation of out-going acoustic waves that are excited by g-mode oscillations in the proto neutron star is the agent of the shock revival. In this context, the possible coupling between the instabilities of the accretion flows and the oscillations of the proto neutron star are interesting~\\citep{yos06,ohn06b}.  In this paper, we considered only the spherically symmetric background flows. Considering the fact that massive stars are rotating in general, we are naturally concerned with the stability of rotational accretion flows. In the previous paper~\\citep{yam05}, we showed that the rapid rotation changes steady accretion flows in such a way that the critical luminosity is decreased and, as a result, the shock revival is facilitated. In so doing, we found that some of the models have a negative entropy gradient and can be convectively unstable. >From the results of this paper, we also expect some oscillatory modes will become unstable and will be more important than convection for low neutrino luminosities. Magnetic fields may also play an important role for the stability of accretion flows. Then, the so-called magneto-rotational instability (MRI) should be discussed in the same framework. In the case of the spherically symmetric background considered in this paper, each eigen value is degenerate with respect to the index, $m$. In the presence of rotation and/or magnetic field, the degeneracy will be removed at least partially. Furthermore, how the growth rates themselves are modified by these effects is an interesting issue. We are currently undertaking the task~\\citep{yam06b} and the results will be published elsewhere."
    },
    "0606/astro-ph0606548_arXiv.txt": {
        "abstract": "The temperature anisotropies and polarization of the cosmic microwave background (CMB) radiation provide a window back to the physics of the early universe. They encode the nature of the initial fluctuations and so can reveal much about the physical mechanism that led to their generation. In this contribution we review what we have learnt so far about early-universe physics from CMB observations, and what we hope to learn with a new generation of high-sensitivity, polarization-capable instruments.  Key-words: cosmic microwave background --- early universe ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606254_arXiv.txt": {
        "abstract": "We use axisymmetric magnetohydrodynamic simulations  to investigate  the  spinning-down of magnetars rotating in the propeller regime and moving supersonically through the interstellar medium. The simulations indicate that magnetars spin-down rapidly due to this interaction,  faster than for the case of a non-moving star. From many simulation runs we have derived an approximate scaling laws for the angular momentum loss rate, $\\dot{L} \\propto-~\\eta_m^{0.3}\\mu^{0.6}\\rho^{0.8}{\\cal M}^{-0.4} \\Omega_*^{1.5}$, where $\\rho$ is the  density of the interstellar medium, ${\\cal M}$ is Mach number, $\\mu$ is the star's magnetic moment, $\\Omega_*$ is its angular velocity, and $\\eta_m$ is magnetic diffusivity. A magnetar with a surface magnetic field of $10^{13} - 10^{15}$ G is found to spin-down to a period $P > 10^5-10^6$ s in  $\\sim 10^4 - 10^5$ years. There is however  uncertainty about the value of the magnetic diffusivity so that the time-scale may be longer. We discuss this model in respect of Soft Gamma Repeaters (SGRs) and the isolated neutron star candidate RXJ1856.5-3754.  \\begin{keywords}  neutron stars --- magnetars --- magnetic field --- MHD \\end{keywords} ",
        "introduction": "Some neutron stars referred to as ``magnetars'' have unusually large magnetic fields, $B\\sim 10^{13}-10^{15} ~{\\rm G}$ (Duncan \\& Thompson 1992; Thompson \\& Duncan 1995). Possible candidates for magnetars include anomalous X-ray pulsars and soft gamma-ray repeaters (SGRs) (Kulkarni \\& Frail 1993; Kouveliotou et al. 1994; Hurley et al. 1999). These objects are associated with  supernovae remnants and hence are relatively young (Vasist \\& Gotthelf 1997; Kouveliotou et al. 1998). Only a few candidates for magnetars have been found so far. The estimated birthrate of magnetars is $\\sim  10\\% $ of ordinary pulsars (Kulkarni \\& Frail 1993; Kouveliotou et al. 1994, 1999) so that there might be many more magnetars which are presently invisible. Their ``visibility\" depends on a number of factors. One important factor is the rate of the star's spin-down. If magnetars spin-down rapidly to very long periods, then one will not detect spin modulated variability during flares. Consequently the identification of the variability with a rotating neutron star would be more difficult.  \\begin{figure*} \\includegraphics[scale=.8]{f1.eps} \\caption{Matter flow around a strongly magnetized star rotating in the propeller regime and propagating through the interstellar medium with Mach numbers ${\\cal M}=1$ (top panel) and ${\\cal M}=3$ (bottom panel). Other parameters correspond to the main case. The background represents the logarithm of density and the length of the arrows is proportional to the poloidal velocity. The solid lines are magnetic field lines. The dashed solid line shows the Alfv\\'en surface. Distances are measured in units of the Bondi radius. Time corresponds to $50$ rotation periods of of the star.} \\label{Figure 1} \\end{figure*}    During the pulsar stage of evolution, magnetars spin down much more rapidly than ordinary pulsars. Consequently, they pass through their pulsar stage much faster, in $\\sim 10^4$ years (Thompson \\& Duncan 1995). When the light cylinder radius becomes larger than magnetospheric radius $r_m$, the relativistic wind is suppressed by the inflowing matter (Shvartsman 1970) and the star enters the propeller regime where the spin-down is due to the interaction of the star's rotating field with the interstellar medium (ISM) (Davidson \\& Ostriker 1973; Illarionov \\& Sunyaev 1975). Magnetars with velocities $v > 100-200$ km/s  interact directly with the ISM. That is, the magnetospheric radius is larger than gravitational capture radius (Harding \\& Leventhal 1992; Rutledge 2001; Toropina et al. 2001). In the propeller regime the rapidly rotating magnetosphere interacts strongly with the supersonically inflowing ISM.   The spin-down rate of supersonically moving  magnetars in the propeller regime has been estimated earlier, but different authors have obtained rather different results. For example, Rutledge (2001) estimates a spin-down time of $\\sim 4 \\times 10^9$ yr for a neutron star with a surface magnetic field $B = 10^{15}~{\\rm G}$ and velocity $v = 300~{\\rm km/s}$. On the other hand, Mori and Ruderman (2003) estimate that a magnetar spins-down to periods greater than $10^4$ s within $\\sim 5 \\times 10^5$ years. Mori and Ruderman put forward this model as an explanation of the isolated neutron star (INS) candidate RX J1856.5-3754 which does not show variability.  The propeller  stage of evolution has been investigated in non-magnetar cases both theoretically (e.g., Illarionov \\& Sunyaev 1975; Davies, Fabian \\& Pringle 1979; Davies \\& Pringle 1981; Lovelace, Romanova \\& Bisnovatyi-Kogan 1999; Ikhsanov 2002; Rappaport, Fregeau, \\& Spruit, 2004) and with MHD simulations (Wang \\& Robertson 1985; Romanova et al. 2003; Romanova et al. 2004; Romanova et al. 2005; Ustyugova et al. 2006). However, only limited theoretical work has been done for magnetar-type propellers, which propagate through the ISM supersonically (Rutledge 2001; Mori \\& Ruderman 2003). No MHD simulations of this flow regime have been done previously.    This work presents the first numerical  simulations of the  interaction of supersonic, fast rotating magnetars with the ISM. Earlier, we investigated supersonic propagation of {\\it non-rotating} strongly magnetized stars through the ISM (Toropina et al. 2001). The present simulations are analogous to those in Toropina et al. (2001). However, here the star rotates in the propeller regime. The main objective of this work is to determine the dependences of spin-down rate of the star on the different variables and to estimate corresponding time-scale of the spin-down.   Sections 2 and 3  describe the physical situation and simulation model. In \\S 4, we discuss the results  of our  simulations. In \\S 5, we apply our results to magnetars and  in \\S 6 we discuss possible magnetar candidates. Conclusions are given in Section 7.     \\begin{figure*} \\includegraphics[scale=.7]{f2.eps} \\caption{Temporal variation of the total angular momentum flux from the star obtained by integrating over a surface surrounding the star (solid line) and the flux through  a surface across the magnetotail at $z=0.6$ (dotted line).} \\label{Figure 2} \\end{figure*}  \\begin{figure*} \\includegraphics[scale=.8]{f3.eps} \\caption{Distribution of the angular momentum fluxes in the magnetotail for our  main case. The color background shows the specific angular momentum carried by the magnetic field (top panel) and by the matter (bottom panel). The solid lines are magnetic field lines.} \\label{Figure 3} \\end{figure*}  \\begin{figure*} \\includegraphics[scale=.8]{f4.eps} \\caption{Dependence of the angular momentum flux on different parameters, (a) the Mach number ${\\cal M}$, (b) the angular velocity of the star $\\omega_*\\equiv \\Omega_*/\\Omega_K$, (c) the magnetic moment $\\mu$, and (d) the magnetic diffusivity $\\tilde{\\eta}_m$.} \\label{Figure 4} \\end{figure*} ",
        "conclusions": "Using axisymmetric MHD simulations we have studied the supersonic propagation through the ISM of magnetars in the propeller stage. We have done many simulation runs for the purpose of determining the angular momentum loss rate of the star due to the interaction of its magnetosphere with the shocked ISM. We conclude, that the interaction may be highly effective in spinning-down  magnetars. A star with magnetic field $B \\sim 10^{13}-10^{15} G$ is expected to spin-down in $\\Delta T \\sim 10^4-10^5 {\\rm years}$. This time may be longer if the ISM material does not efficiently interact with the external regions of the magnetar's magnetosphere. Therefore, after relatively short stages of pulsar and propeller activity, a magnetar becomes a very slowly rotating object, with  a period $P > 10^5-10^6 s$, which is much longer than the periods expected for ordinary pulsars. This may be a reason why the number of soft gamma repeaters, which are candidate magnetars, is so small. We should note however, that the rate of spinning-down depends on the  magnetic diffusivity which is not known. At lower diffusivity the rate of spinning-down will be lower. The INS candidate RX J1856.5-3754 may be an example of a slowly rotating magnetar. However,  this model does not explain the $H_{\\alpha}$ nebulae. An ordinary misaligned pulsar explains the different features more easily,  excluding the fact that no periodic fluctuations were observed from this object."
    },
    "0606/astro-ph0606312_arXiv.txt": {
        "abstract": "Prompt optical emission from the $\\gamma$-ray burst of GRB 041219A has been reported by Vestrand et al. There was a fast rise of optical emission simultaneous with the dominant $\\gamma$-ray pulse, and a tight correlation with the prompt $\\gamma$-ray emission has been displayed. These indicate that the prompt optical emission and $\\gamma$-ray emission would naturally have a common origin. We propose that this optical component can be modeled by considering the Comptonization of $\\gamma$-ray photons by an electron cloud. As a result of this mechanism, the arrival time of the optical photons is delayed compared with that of the $\\gamma$-rays. We restrict that the lag time to be shorter than $10$ s, within which the prompt optical emission is considered to vary simultaneously with the prompt $\\gamma$-ray emission. Taking the observations of GRB 041219A into account, we derive the number density of the surrounding electron cloud required for Comptonization. The redshift of GRB 041219A is predicted to be $z\\lesssim 0.073$ as well. ",
        "introduction": "GRB 041219A was detected by both the IBIS detector on the {\\rm International Gamma-Ray Astrophysics Laboratory} satellite \\citep{got04} and the Swift Burst Alert Telescope (BAT) \\citep{bar04}. It was an unusually bright and long burst. The fluence of $1.55\\times 10^{-4}\\,ergs\\,cm^{-2}$ measured by BAT in the 15-350\\,keV band  would put it among the top few per cent of the 1637 $\\gamma$-ray burst (GRB) events listed in the comprehensive fourth BATSE (Burst and Transient Source Experiment) catalog \\citep{pac99}. The duration of the prompt $\\gamma$-ray emission was approximately $520\\,s$, making it one of the longest bursts ever detected \\citep{ves05}.  So far, prompt optical and infrared emission that occurs when the main burst is still in progress has been detected from a few GRBs. These include GRB 990123 \\citep{ake99}, GRB 041219A \\citep{ves05, bla05}, GRB 050401 \\citep{ryk05}, GRB 060111B (Klotz et al. 2006), and GRB 050904 \\citep{Boe06, Wei06}. The data obtained by RAPTOR (Rapid Telescopes for Optical Response) shows that the prompt optical light curve of GRB 041219A could be well fitted with a constant prompt optical-to-$\\gamma$-ray flux ratio $F_{opt}/F_{\\gamma}=1.2\\times 10^{-5}$ \\citep{ves05}. This strongly suggests that a direct correlation of both the time-varying spectral shape and the flux magnitude exists between the prompt optical emission and the $\\gamma$-ray emission.  Prompt long-wavelength radiation accompanying prompt $\\gamma$-ray emission has been widely discussed by many authors in the pre-afterglow era \\citep{kat94, sch94, wei97, tav96,zha05} and in the afterglow era \\citep{Sai99, mes99, wu06, van00, fan04, fw04, Bel05}. Recently, simultaneous variation of the optical-infrared emission with the prompt $\\gamma$-rays from GRB 041219A has been discussed by \\cite{fan05} using a neutron-rich internal shock model. Alternatively, here we argue that Compton attenuation of the $\\gamma$-ray photons with a power-law spectrum by intervening electron clouds could give birth to the prompt optical emission. The $\\gamma$-ray photons are assumed to be produced by the central engine (e.g., standard fireball model), while an electron cloud with the extremely high number density required could likely be ejected by the progenitor of the GRB, for example, as ejecta from the associated supernova in an earlier phase of the explosion \\citep{mac99}. When the incident $\\gamma$-ray photons travel through the electron cloud, some of them will be reprocessed into optical photons through the Compton attenuation mechanism. The model is addressed in detail in \\S\\, 2, and we give the discussion and conclusions in \\S\\, 3. ",
        "conclusions": "We propose that the prompt optical and infrared emission from GRB 041219A can be modeled by the saturated Comptonization of its $\\gamma$-ray emission. The current evidence shows that the optical emission of GRB 041219A is rather distinct from that of other GRBs, suggesting a possible environmental influence \\citep{ves05}. We argue that this special event could have been surrounded by a very high density electron cloud. If the sub-$10\\,s$ lag time prompt optical emission can be considered as varying simultaneously with the prompt $\\gamma$-ray emission, by fitting the observations we find that the required electron number density to degrade the $\\gamma$-rays into the optical band is $\\sim 7\\times 10^{17}\\,cm^{-3}$, and this dense scattering region is expected to cover $\\sim 87\\%$ of the total area. Although the required value of the electron number density is much higher than that of ordinary interstellar matter, it is lower than the electron number density of the Sun's outer convection layer \\citep{bah88}.  We argue that the electron cloud could be a result of, for example, either the ejecta of the outer layer of the GRB progenitor ejected in an earlier phase of the explosion \\citep{mac99} or the capture of a star by the GRB progenitor \\citep{wan01}. If the scattering electron cloud is produced by the former mechanism, its total mass can be estimated as $M\\sim 4\\pi R_{\\ast}^2Ln_{e}m_{p}\\approx 2 \\times10^{-6}\\,M_{\\odot}$, where $m_{p}$ is the proton mass and $R_*$ is the radius of the GRB progenitor star, with $R_*\\sim10R_{\\odot}$ adopted here. Suppose that the ejected outer layer of the progenitor has a typical length of $10^5\\,cm$ and density of $10^{-3}\\,g\\,cm^{-3}$; when it expands into a size of $10^8\\,cm$, it could create an electron cloud as required by the model. As a result, we predict that an upper limit on the redshift of GRB 041219A of $z\\sim 0.073$ if the isotropic energy of this GRB is less than that of GRB 990123.  Further more, many photons radiated in the UV/soft X-ray band could be inferred. Nevertheless, the XRT and UVOT instruments on board \\emph{Swift} did not autonomously slew to the burst, since automated slewing is not yet enabled \\citep{fen04}; it is interesting to note that the \\emph{Rossi X-Ray Timing Explorer} All Sky Monitor  obtained some soft X-ray data for the initial 120\\,s of GRB 041219A \\citep{mcb06}. The results presented here highlight the need for continued broadband observations of $\\gamma$-ray burst and the afterglow.  \\noindent{\\bf Acknowledgement:} We thank the anonymous referee for valuable comments and suggestions that lead to an overall improvement of this study. We are thankful to Y.-F. Huang, S.-N. Zhang and K.-S. Cheng for both helpful comments and useful discussions. Thanks also to H.-X. Yin and W. Qiao for useful discussions. This research was supported by the National Natural Science Foundation of China (grants 10573021, 10273011, and 10433010), and by the Special Funds for Major State Basic Research Projects."
    },
    "0606/astro-ph0606638_arXiv.txt": {
        "abstract": "We compare simulations of the \\lya forest performed with two different hydrodynamical codes,  \\gad and \\textsc{Enzo}. A comparison  of the dark matter power spectrum for simulations run with identical  initial conditions show differences of  1-3\\% at the  scales relevant for quantitative studies of the  \\lya forest. This allows a meaningful comparison  of the effect of the  different implementations of the hydrodynamic part of the two codes. Using the same  cooling and heating algorithm in both codes   the differences in the  temperature and the density probability distribution function are of the order of 10 \\%. The differences are comparable  to the effects  of boxsize and resolution on these statistics. When self-converged results for each code are taken into account the differences in the flux power spectrum -- the statistics most widely used  for estimating  the matter power spectrum and cosmological parameters from \\lya forest  data --  are  about 5\\%. This is again comparable to the effects of boxsize  and resolution. Numerical uncertainties  due to a particular implementation  of solving the hydrodynamic or gravitational equations appear therefore to contribute only moderately to the error budget in estimates of  the flux power  spectrum from numerical simulations. We further find that the differences  in the flux power spectrum for \\enzo simulations run with and without adaptive mesh refinement are also of order 5\\% or smaller. The latter require 10 times less CPU time making the CPU  time requirement similar to that of a version of \\gad that is optimised  for \\lya forest simulations. ",
        "introduction": "There is  now a well established  paradigm for the origin  of the \\lya forest,  the  ubiquitous  absorption  lines due  to neutral hydrogen in the  spectra of  high-redshift quasars.  The absorption blue-wards of 1216 $\\AA$ is predominantly due to density fluctuations in the intervening warm  ($\\sim 10^4$ K) photoionized inter-galactic medium (IGM)  on scales larger than  the Jeans  length  of the  gas  (see \\cite{Rauch_1998} for a review).  Numerical simulations  were instrumental in  establishing the new paradigm in the 1990s (\\cite{Cen_1994}, \\cite{Zhang_1995}, \\cite{Hernquist_1996}, \\cite{Theuns_1998a}, \\cite{Zhang_1997}). The  \\lya forest and the associated  metal absorption probe the thermal and ionization history of the IGM  as well as the interplay of galaxies and the IGM from which they are formed. More recently  the \\lya forest has also  been  established  as a means   of  quantitative measurement of the underlying matter distribution and  thus a variety of cosmological parameters ({\\it e.g.} \\cite{Croft_1998}, \\cite{Croft_2002a}, \\cite{Viel_2004}, \\cite{Seljak_2005a}, \\cite{Viel_2006b}, \\cite{Seljak_2006}; Viel, Haehnelt \\& Lewis (2006)).   Numerical simulations thereby play a crucial role in  inferring the linear matter power spectrum and other derived quantities from the \\lya forest data.  With increasing sample   sizes statistical errors of measurements of the flux  distributions  have reached the percent level and the  error budget is  dominated  by systematic uncertainties  (Viel, Haehnelt \\& Springel (2004), \\cite{McDonald_2005a} \\nocite{McDonald_2005b}).  Uncertainties due to numerical simulations  contribute significantly to the error budget and  the accuracy with  which  the  flux distribution  for   given input  physics  and cosmological parameters can be simulated  has become important.  Most studies so far have used convergence tests to  assess uncertainties due to the numerical simulations and direct comparisons of cosmological hydrodynamical simulation performed with different codes have been rare. The  differences between hydro-dynamical  simulations of galaxy clusters with a wide range of different codes/methods have been studied in the Santa Barbara cluster project (Frenk et al. 1999)\\nocite{Frenk_1999}.   Recently \\cite{O'Shea_2005} performed a  comparison  between the grid based adaptive mesh  refinement (AMR) code \\textsc{Enzo}\\footnote{http://cosmos.ucsd.edu/enzo/} and the smoothed particle hydrodynamics (SPH) code \\textsc{Gadget-2}\\footnote{http://www.mpa-garching.mpg.de/gadget/}. However, little  has   been  done  in this  respect  for hydrodynamical simulations of the \\lya  forest data (see \\cite{Theuns_1998b} for a notable exception of a comparison between two SPH codes) Some comparisons of hydrodynamical simulations with approximate simulations of the \\lya forest data have been  performed   by McDonald  et al.  (2005),\\cite{Zhan_2005}  and Viel, Haehnelt \\& Springel  (2006). We  present here a comparison of hydrodynamical simulations of the \\lya  forest with \\enzo and \\gad which concentrates on the statistical properties of the flux distribution. \\\\ We are therefore mostly interested in  properties of the moderate to low over-density  gas  which is  responsible  for \\lya forest absorption.  Of particular interest is the probability distribution  of the gas density, the temperature, the resulting flux distribution and the  flux power spectrum.  A major difference between grid-based and SPH codes is their treatment of shocks and their effects on the temperature distribution. We will also examine these differences.\\\\  \\indent The plan of  the  paper is as follows.  In \\S\\ref{Codes} we describe  the \\enzo code and the \\gad code and the different  ways  in  which  the codes  solve  the gravitational  and hydrodynamics  equations. In \\S\\ref{Sims}  we describe  the simulation set used in the comparisons. In \\S\\ref{comparison} we investigate how physical properties  of both codes  compare,  in  particular the  gas distribution and the 1-D flux power spectrum. Finally in \\S\\ref{time} we will  look at the performance of each code in terms of CPU time consumption. ",
        "conclusions": "We have performed a detailed comparison of \\lya forest simulations with \\gad, a \\textsc{TreePM-SPH} code, and \\enzo a Eulerian AMR code in order to asses the numerical uncertainties due to a particular numerical implementation of solving the hydrodynamical equations. The codes are similar with respect to the way in which they compute the gravitational forces at large scales but differ in the way they calculate gravitational forces on small scales; the codes use a Tree-PM and PM N-body algorithm, respectively. Their main differences lie, however,  in the way in which they solve the gas  hydrodynamics. \\gad discretises mass using SPH methods while \\enzo  discretises space using adaptive meshes. The main results are as follows.   \\begin{itemize}  \\item{The differences in the dark matter power spectrum between simulations with \\enzo and \\gad on scales relevant for measurements of the matter power spectrum from \\lya forest data are about 2\\% for an appropriate choice of box size and resolution.}   \\item{The temperature density relation of simulations with \\enzo and \\gad differ very little. The PDF of the volume weighted temperature differ by $\\sim 10$ \\% probably mainly due to differences in the PDF of the gas density which are of the same order and at least partially caused by a slight  mismatch in resolution.}  \\item{The PDF  of the flux distribution of simulations with \\enzo and \\gad agree very well. Typical differences are $\\sim 5-10$\\% probably again mainly due to a slight mismatch of the resolution.}  \\item{The differences of the flux power  spectrum of simulations with \\enzo and \\gad on scales relevant for measurements of the matter power spectrum from \\lya forest data are about 5\\% for an appropriate choice of box size and resolution and simulations which fully resolve the Jeans mass. For simulations of lower resolution but larger boxsize the difference increase up to  $\\sim 10$\\%. Note that the differences are scale and redshift dependent. }  \\end{itemize}   \\begin{figure}  \\psfig{figure=timing.ps,width=0.48\\textwidth} \\caption{\\label{times} A comparison of the CPU time required for the simulation to run through a fixed redshift interval. Simulation parameters are as described in the text and annotated on the plot. The thin lines are for a SUN distributed memory cluster and the thick lines are for a SGI shared memory (SM) machine. Only relative values for the same architecture should be considered. Note the reversal in relative speed between \\gad simulations with star formation  and \\enzo static grid simulation between the two architectures. The thick dashed line shows a linear scaling of CPU time with the number of particles for  comparison. }  \\end{figure}    Overall the \\lya forest simulations with \\enzo and \\gad agree astonishingly well. The choice of method for solving the hydrodynamical simulations appears to affect the gas distribution and its thermal state very little. It is also reassuring that two different implementations for solving the gravitational equations agree well.  The corresponding uncertainties should contribute to the overall  error budget of measurements of the matter power spectrum from \\lya forest data at the level of 3\\%. The total error in current measurement is significantly larger  and they should thus not be important. The main numerical uncertainties are  instead due to a lack of sufficient dynamic range which typically makes correction of 5\\% for boxsize and resolution necessary. This will obviously improve as computational resources  become more powerful. In practical terms memory requirements of simulations with \\enzo without AMR and \\gad are similar. \\enzo without AMR offers the highest speed but requires somewhat larger corrections. Our results suggest that if sufficient computational resources are available and sufficient care is employed the accuracy of numerical simulations should  not yet be a limiting factor in improving the accuracy of measurements of the matter power spectrum from \\lya forest data."
    },
    "0606/astro-ph0606197_arXiv.txt": {
        "abstract": "Dark matter or modifications of the Newtonian inverse-square law in the solar system are studied with accurate planetary astrometric data. From extra-perihelion precession and possible changes in the third Kepler's law, we get an upper limit on the local dark matter density, $\\rho_\\mathrm{DM} \\ls 3 {\\times} 10^{-16}~\\mathrm{kg/m}^3$ at the 2-$\\sigma$ confidence level. Variations in the $1/r^2$ behavior are considered in the form of either a possible Yukawa-like interaction or a modification of gravity of MOND type. Up to scales of $10^{11}$~m, scale-dependent deviations in the gravitational acceleration are really small. We examined the MOND interpolating function $\\mu$ in the regime of strong gravity. Gradually varying $\\mu$ suggested by fits of rotation curves are excluded, whereas the standard form $\\mu(x) = x/(1+x^2)^{1/2}$ is still compatible with data. In combination with constraints from galactic rotation curves and theoretical considerations on the external field effect, the absence of any significant deviation from inverse square attraction in the solar system makes the range of acceptable interpolating functions significantly narrow. Future radio ranging observations of outer planets with an accuracy of few tenths of a meter could either give positive evidence of dark matter or disprove modifications of gravity. ",
        "introduction": "Gravitational inverse-square law and its relativistic generalization have passed significant tests on very different length- and time-scales. Precision tests from laboratory and from measurements in the solar system and binary pulsars provide a quite impressive body of evidence, considering the extrapolation from the empirical basis \\citep{ade+al03,wil06}. First incongruences seem to show up only on galactic scales with the observed discrepancy between the Newtonian dynamical mass and the directly observable luminous mass and they are still in order for even larger gravitational systems. Two obvious explanations have been proposed: either large quantities of unseen `dark' matter (DM) dominate the dynamics of large systems \\citep{zwi33} or gravity is not described by Newtonian theory on every scale \\citep{fin63}. Dark matter is all the general theory of relativity needs to overcome apparent shortcomings and provides a coherent picture for gravitational phenomena from the laboratory to the cosmological context. The paradigm of cold DM when complemented with a positive cosmological constant (the so called $\\Lambda$CDM scenario) is successful in explaining the whole range of galactic and extra-galactic body of evidence, from flat rotation curves in spiral galaxies to large scale structure formation and evolution \\citep{pea99}.  The $\\Lambda$CDM paradigm could be regarded as the definitive picture apart from that the presumed existence of DM relies so long only on its putative global gravitational effect, whereas direct detection by any independent mean is still lacking. This makes room to alternative proposals based on modifications of Newtonian gravity. In general, such proposals do not extend the inverse-square law to a regime in which it has never before been tested and they do not introduce any exotic component. Proposals are very different from each other. Some of them can make gravity stronger on scales of galaxies and explain flat rotation curves without dark matter \\citep{mil83}; others realize a mechanism for the cosmic acceleration without dark energy, for example as a result of gravity leaking on scales comparable to the horizon \\citep{dva+al00}. Two main alternative proposals have been discussed. In the first one, the gravitational potential deviates from the usual form at large distances. A classical example is the inclusion of a Yukawa-like term in the gravitational potential. This is strictly related to more fundamental theories where these additional contributions appear as the static limit of interactions due to the exchange of virtual massive bosons \\citep{ade+al03}. According to the second main choice, Newton's law fails when the gravitational acceleration is small rather than when the distance is large. The prototype and still one of the most empirically successful alternative to DM is Milgrom's modified Newtonian dynamics (MOND) \\citep{mil83,sa+mc02}. With some basis in sensible physics, MOND can provide an efficient description of the phenomenology on scales ranging from dwarf spheroidal galaxies to cluster of galaxies but its cosmological extension is still in the childhood \\citep{bek04}.  High precision solar system tests could provide model independent constraints on possible modifications of Newtonian gravity. The solar system is the larger one with very well known mass distribution and can offer tight confirmations of Newtonian gravity and general relativity. Any deviation emerging from classical tests would give unique information either on dark matter and its supposed existence or on the nature of the deviation from the inverse-square law. Several authors have discussed this possibility. \\cite{tal+al88} derived limits from the analysis of various planetary astrometric data set on the variation in the $1/r^2$ behavior of gravity. Experimental bounds on non luminous matter in solar orbit were derived either by considering the third Kepler's law \\citep{and+al89,and+al95} or by studying its effect upon perihelion precession \\citep{gr+so96}. The influence of a tidal field due to Galactic dark matter on the motion of the planets and satellites in the solar system was further investigated by \\cite{bra+al92} and \\cite{kl+so93}. The orbital motion of solar-system planets has been determined with higher and higher accuracy \\citep{pit05b} and recent data allow to put interesting limits on very subtle effects, such as that of a non null cosmological constant \\citep{je+se06,se+je06}. In this paper, we discuss what state-of-art ephemerides tell us about non Newtonian or DM features. In section~\\ref{basi}, we review standard expectations about Galactic dark matter at the solar circle and discuss some standard frameworks for deviations from the inverse square-law, i.e. a Yukawa-like fifth force and the MOND formalism. Observational constraints from perihelion precessions and changes in the third Keplerian law are discussed in section \\ref{peri} and \\ref{mean}, respectively. Section \\ref{conc} is devoted to some final considerations. ",
        "conclusions": "\\label{conc}  Debate between dark matter and departures from inverse-square law is still open. Considering both theoretical and observational aspects, dark matter seems to be slightly preferred. If on a galactic scale the two hypotheses match, on the cosmological side only DM can give a consistent framework. This might shortly change with the steady improvements in relativistic generalization of the MONDian paradigm. So, in our opinion, it is of interest to examine results on a very different scale, that of the solar system. Solar system data have been confirming predictions from the general theory of relativity without any need for dark matter and it is usually assumed that deviations can show up only on a larger scale. In this paper, we have explored what we can learn from orbital motion of major planets in the solar system. Results are still non-conclusive but nevertheless interesting. Best constraints come from perihelion precession of Earth and Mars, with similar results from modifications of the third Kepler's law. The upper bound on the local dark matter density, $\\rho_\\mathrm{DM} \\ls 3 {\\times} 10^{-16}~\\mathrm{kg/m}^3$, falls short to estimates from Galactic dynamics by six orders of magnitude.  Deviations of the gravitational acceleration from $1/r^2$ are really negligible in the inner regions. A Yukawa-like fifth force is strongly constrained on the scale of $\\sim 1$~AU. For a scale-length $\\lambda_\\mathrm{Y} \\sim 10^{11}~\\mathrm{m}$, a Yukawa-like modification can contribute to the total gravitational action for less then one part on $10^{11}$. Similar limits could be achieved by precise measurements on the proof masses carried on board of the LISA Pathfinder satellite~\\footnote{http://www.rssd.esa.int/index.php?project=LISAPATHFINDER} (Speake, private communication). In fact, instantaneous measurements of the drag-free test-mass acceleration during the transfer orbit towards the first Sun-Earth Lagrange point could in principle test the inverse square law on a scale length of $\\sim 1$~AU (Speake, private communication). Results on a similar scale-length could be obtained through a detailed analysis of binary pulsars. The periastron shift, the gravitational redsfhift/second-order Doppler shift parameter and the rate of change of orbital period are sensitive to scalar-tensor gravity and to any other deviation from the general theory of relativity \\citep{wil06}. Dipole gravitational radiation associated with violations of the equivalence principle in its strong version could cause an additional form of gravitational damping and a significant change of the orbital period could occur, in particular for a binary pulsar system with objects of very dissimilar mass \\citep{wil06}. A massive graviton associated with a Yukawa-like fifth force could also affect the speed of propagation of gravitational waves and induce radiation effects at the reach of future gravitational wave detectors \\citep{wil06}.    A large class of MOND interpolating function is excluded by data in the regime of strong gravity. The onset of the asymptotic $1/r$ acceleration should occur quite sharply at the edge of the solar system, excluding the more gradually varying $\\mu(x)$ suggested by fits of rotation curves. On the other hand, the standard MOND interpolating function $\\mu(x) = x/(1+x^2)^{1/2}$ is still in place. Studies on planetary orbits could be complemented with independent observations in the solar system. Mild or even strong MOND behavior might become evident near saddle points of the total gravitational potential, where MONDian phenomena might be put at the reach of measurements by spacecraft equipped with sensitive accelerometers \\citep{be+ma06}. As a matter of fact, fits to galactic rotation curves, theoretical considerations on the external field effects and solar system data could determine the shape of the interpolating function with a good accuracy on a pretty large intermediate range between the deep Newtonian and MONDian asymptotic behaviors.  Future experiments performing radio ranging observations of outer planets could greatly improve our knowledge about gravity in the regime of large accelerations. The presence of dark matter could be detected with a viable accuracy of few tenths of a meter on the measurements of the orbits of Neptune or Pluto, whereas an uncertainty as large as hundreds of meters would be enough to disprove some pretty popular MOND interpolating functions.  In order to become really competitive with general relativity and the $\\Lambda$CDM paradigm, MOND should be predictive on the whole range of observed systems from solar system to the cosmic microwave background radiation. On a galactic scale, effects of DM or MOND are pretty similar and very difficult to distinguish each other but there might be some detectable differences on a smaller scale. In fact, the local value of DM at the solar circle is pretty much fixed by Galactic dynamics whereas the MOND behavior in the regime of strong accelerations probed locally is not univocal on a theoretical and observational basis. Nevertheless, only a very small class of interpolating free functions would give the same perturbation on the orbits of outer planets as that from local DM. Matching the expectations from DM with future radio ranging observations would be an important, nearly conclusive confirmation of its existence. On the other hand, deviations at a different order of magnitude, as those expected for a large variety of MOND interpolating functions, would be a strong indication of departure from the inverse-square gravitational law."
    },
    "0606/astro-ph0606368_arXiv.txt": {
        "abstract": "\\emph{HST} ACS images reveal blue cores in four E+A, or post-starburst, galaxies. Follow-up spectroscopy shows that these cores have LINER spectra. The existence of LINERs, consistent with those in many elliptical galaxies, is yet one more piece of evidence that these post-merger, post-starburst, bulge-dominated galaxies will evolve into normal ellipticals. More interestingly, if LINERs are powered by low-luminosity AGN, their presence in these E+As suggests that any rapid growth phase of the central black hole ended in rough concert with the cessation of star formation. This result emphasizes the importance of E+As for exploring how the evolution of black holes and AGN may be tied to that of galactic bulges. ",
        "introduction": "The strong correlation between black hole mass and galaxy bulge velocity dispersion \\cite[$M_\\bullet-\\sigma_B$;][]{Ferrarese00,Gebhardt00} suggests one of two things. Either early-type galaxies do not arise from dissipative mergers or there is a connection --- perhaps causal --- between the small-scale physics of black hole (BH) growth and the large-scale physics that organizes the host galaxy morphology, kinematics, and stellar populations during and after the merger. Given the observational evidence that gas-rich mergers do occur and that at least some produce pressure-supported, bulge-dominated remnants, it remains to understand how the processes that drive the evolution of the smallest and largest galactic scales are related.  A key to resolving this question is to identify galaxies undergoing large-scale transitions via mergers and to consider the properties of their cores. On-going mergers are too complicated to provide clear answers, while their likely remnants are too far removed from the merger event. E+A galaxies, which have post-starburst spectra, frequent tidal features, and the kinematic and morphological signatures of early type galaxies \\citep{Zabludoff96,Norton01,Chang01,Yang04}, are true transitional objects and thus plausible test cases. In this {\\it Letter}, we explore whether there is evidence for a central BH/AGN in nearby E+As, and, if so, whether the core is consistent with those of early-type galaxies and evolving in concert with the galaxy as a whole. ",
        "conclusions": "We identify four E+A galaxies with blue cores, which are revealed by our follow-up spectroscopy to have LINER spectra.  The existence of LINERs, similar to those in elliptical galaxies, is more evidence that E+A galaxies, with their post-starburst spectra, post-merger, gas-poor, bulge-dominated morphologies, and pressure-supported kinematics \\citep[e.g.,][]{Norton01,Chang01,Yang04}, evolve into normal early-types.  More interestingly, if LINERs are low-luminosity AGN, their presence in E+As suggests that any rapid growth phase of the central AGN has ended in rough concert with the star formation and therefore that the evolution of the black hole is tied to that of the galactic bulge.  What is not clear from our work is whether the coupling between AGN and bulge evolution is causal, as is suggested by some theoretical models incorporating AGN feedback \\citep{Springel05a,Springel05b}. The study of a large sample of E+As, including an investigation of any correlation between AGN strength and the time elapsed since the starburst, could provide a test of the AGN-feedback hypothesis, as those models predict that the black hole accretion rate peaks shortly after the starburst and declines quickly as the merger remnant ages."
    },
    "0606/astro-ph0606642_arXiv.txt": {
        "abstract": "In this paper we analyze previously published spectra with high signal-to-noise ratios of E and S0 galaxies in the rich cluster CL1358+62 at $z=0.33$, and introduce techniques for fitting stellar population models to the data. These data and methods will be used further in a larger study of the evolution of absorption line strengths in intermediate redshift clusters. Here we focus on the 19 elliptical and lenticular galaxies with an homogeneous set of eight blue Lick/IDS indices. These early-type galaxies follow very narrow line strength-line width relations using Balmer and metal lines, indicating a high degree of uniformity in their formation and enrichment histories. We explore these histories using recently published, six-parameter stellar population models \\citep{thomas,thomas2}, and describe a novel approach for fitting these models {\\it differentially\\/}, such that the largest sources of systematic error are avoided. The results of the model fitting are accurate {\\it relative\\/} measures of the stellar population parameters, with typical formal errors of $\\simlt 0.1$ dex. The best-fit models yield a mean $\\chi^2$ of 1.2 per degree of freedom, indicating that the models provide good descriptions of the underlying stellar populations. We find: (1) no significant differences between the best-fit stellar population parameters of Es and S0s at fixed velocity dispersion; (2) the stellar populations of the Es and S0s are uniformly old, consistent with results previously published using the fundamental plane; (3) a significant correlation of [Z/H] with galaxy velocity dispersion, in a manner consistent with the observed $B-V$ colors of the galaxies, and indicating that dust is not a significant contributor to the colors of early-type galaxies; (4) a possible, modest anti-correlation of [$\\alpha$/Fe] with velocity dispersion, with $< 2\\sigma$ significance, and discrepant with the correlation inferred from data on nearby galaxies at the $<3\\sigma$ level; and (5) a significant anti-correlation of [$\\alpha$/N] with galaxy velocity dispersion, which we interpret as a correlation of nitrogen enhancement with mean metallicity. Neither [$\\alpha$/C], nor [$\\alpha$/Ca] shows significant variation. While the differences between our conclusions and the current view of stellar populations may point to serious deficiencies, our deduced correlation of mean metallicity with velocity dispersion does reproduce the observed colors of the galaxies, as well as the slope of the local Mg-$\\sigma$ relation. Our tests indicate that the inferred population trends do describe real galaxies quite well, and matching our results with published data on nearby galaxies, we infer that the discrepancy stems largely from the historical treatment of broadening corrections to the narrow indices. The data also strongly indicate that secondary nitrogen is an important component in the chemistry of elliptical and lenticular galaxies. Taken together, these results reduce early-type galaxies in clusters to a family with one-parameter, velocity dispersion, greatly simplifying scenarios for their formation and evolution. More specifically, our data conclusively show that cluster S0s did not form their stars at significantly later epochs than cluster ellipticals of the same mass, and the presence of secondary nitrogen indicates that both Es and S0s formed from self-enriching progenitors, presumably with extended star-formation histories. ",
        "introduction": "Ever since the discovery that early-type galaxies follow scaling relations \\citep{mm1957}, it was clear they had profound implications for cosmology as well as for galactic structure, formation, and evolution \\cite[e.g.][and subsequent literature]{minkowski1962,terlevich}. Perhaps the most well-studied of these scaling relations is the color-magnitude relation \\citep{baum}, and its interpretation as a correlation of metal abundance with galaxy luminosity \\citep{rood1969} has survived to the present-day. Over the past several decades, numerous other scaling relations have been discovered and interpreted, such as $L$-$\\sigma$ \\citep{fj76}, $r_e$-$I_e$-$\\sigma$ \\citep[the fundamental plane;][]{dress7s,dd87}, Mg-$\\sigma$ \\citep{terlevich}, and other line strength-line width relations, such as H$\\delta_A$- or H$\\gamma_A$-$\\sigma$ \\citep[e.g.][]{kuntschner2000,kelson01}.  Several important breakthroughs have helped to constrain the nature of early-type galaxy scaling relations. The clearest results have come from explicit tests for changes in the ages of stellar population with redshift along the sequence of early-types. By measuring the slopes of the color-magnitude relations in clusters to redshifts of unity \\cite{stanford} concluded that the relation originates largely from systematic variations in metallicity with galaxy mass \\citep[also see][]{blakeslee2006}. Likewise, the fundamental plane of early-type galaxies has a slope that does not appear to evolve significantly with redshift to $z=0.33$ \\citep{kelsonc}, further evidence that cluster E/S0s are a family well-described by uniform ages and a mass-metallicity relation. At redshifts of $z=0.8-0.9$ there are hints that the slope of the fundamental plane may evolve modestly \\citep{jorgensen2006}, though this result appears to be inconsistent with the slope of the color-magnitude relation as measured by \\cite{blakeslee2006} and it remains to be seen how these data will be reconciled.  There have been many efforts to observe the high-redshift line strength-line width relations \\citep{ziegler,kelson01,jorgensen0152,barr2005,moran2005}. However, significant constraints have not been forthcoming because of the difficulty of obtaining sufficient signal-to-noise ratios.  The situation is much better at low redshifts. Detailed analysis of line strengths in low-redshift E/S0s have been very revealing since \\cite{jesus} (and others) broke the age-metallicity degeneracy \\citep{worthey}. Using line strengths many authors have shown that massive galaxies in clusters are uniformly old \\citep[e.g.][]{jorgensen1997,jorgensen1999,trager2000a,kuntschner2000}, but some ambiguities do remain. For example it has long been understood that the Mg$b$-$\\sigma$ \\citep{terlevich} and Fe-$\\sigma$ \\citep{jesus} relations are not consistent with the hypothesis that they both arise solely from a correlation between the mean metallicity of the stellar populations and galaxy mass. The discrepancy has been interpreted using stellar populations models in which ``$\\alpha$-enhancement, or [$\\alpha$/Fe], is allowed to vary systematically with velocity dispersion \\citep[see, e.g.][for details and many important references]{jorgensen1999,trager2000a,trager2000,worthey2003,thomas2005}. While \\cite{trager2000a} reminded readers that this is a misnomer, because the $\\alpha$ elements are not enhanced --- Fe is simply under-abundant. We will, however, continue to use the term ``$\\alpha$-enhancement'' to be consistent with past literature on the subject. After several decades of analysis, improved techniques of observation, and despite the immense amount of progress in stellar population models \\citep{faber1972,worthey,trager2000a,schiavon1,thomas}, the modeling of absorption lines remains ambiguous and contradictory \\citep[e.g.,][and others]{worthey,trager2000,worthey2003,schiavon3,thomas-ca}. Some have even questioned the ``$\\alpha$-enhancements'' altogether \\citep[]{proctor2004a}.  Over the long-term we hope that the physics of stellar atmospheres will be understood with sufficient detail to allow for the creation of synthetic stellar spectra for stars, over the full range of stellar masses and phases of stellar evolution, that comprise the observed spectra of galaxies. Presently, however, numerous uncertainties in the formation of molecular lines, and incomplete line lists prevent one from generating completely synthetic spectra for stellar populations. As a result the features in the spectra of stars and galaxies cannot be fully modeled. Despite major progress in generating high resolution spectral energy distributions of simple stellar populations \\citep{vaz2a,bc2003}, we cannot fully solve for the age(s) and elemental abundances of populations in a galaxy by directly fitting synthetic spectra.  \\begin{figure*} \\centerline{\\epsscale{0.7} \\plotone{fig1.eps}} \\caption{Color-magnitude diagram of the cluster in $R$ and $V_{606}-I_{814}$, taken from \\cite{kelsonb}. All 194 confirmed cluster members in the HST imaging are shown. The red, blue, and green symbols represent those galaxies observed by \\cite{kelsonb} for use in their study of the fundamental plane. Ellipticals are shown using red, filled circles. The S0 galaxies are shown using blue open circles. Spiral galaxies are shown in green. The light gray points were not observed in the high-resolution study of \\cite{kelsonb}. The red, blue, and green points are shown using two sizes. The larger ones mark those galaxies in the homogeneous sample (for which all eight of the indices used in the modeling are available; see \\S \\ref{sec:fitting}). The E and S0 (red and blue) galaxies in the figure are discussed in this paper while the full fundamental plane sample of CL1358+62 will be discussed in a subsequent paper. \\label{fig:selection}} \\end{figure*}  Without the ability to directly model the spectra of galaxies, one must measure and model spectral indices. One widely used system is the Lick/IDS system, created by \\cite{burstein84} \\citep[and subsequently revised by][]{trager}. To this day, the measurement and modeling of these line strengths remain the most useful means of assessing the bulk properties of the stellar populations in passively evolving galaxies.  Over the last several years several groups have expanded our ability to model these indices with more parameters than age, metallicity, and $\\alpha$-enhancement. The models of \\cite{trager2000a},\\cite{thomas}, and \\cite{schiavon2005} allow one to probe specific elemental abundances, such as nitrogen, carbon, and calcium. Such models only make predictions for passively evolving stellar populations. Fortunately models of galaxy evolution involving star-formation are not needed in our study of massive cluster galaxies through intermediate redshifts, because such galaxies have been shown to be passively evolving \\citep[e.g.,][]{kelson97,kelsonc,wuyts}. At higher redshifts, on-going star-formation may become increasingly important \\citep[e.g.][]{juneau}, and more sophisticated models may then be required.  In our survey of galaxies in rich clusters at intermediate redshifts, the sample of galaxies with the highest signal-to-noise spectroscopy over the widest range of galaxy luminosities is that from our fundamental plane survey of CL1358+62 at $z=0.33$. Because of the depth and quality of that sample, we adopt it as the reference sample to which the other clusters in our survey will subsequently be compared \\citep{evolpaper}.  Here we study the absorption lines of the E/S0 galaxies in that sample, using the models of \\cite{thomas} and \\cite{thomas2} to explore not only ages and metallicities, but additional stellar population parameters, namely the relative abundances of nitrogen, carbon, and calcium, along with the mean $\\alpha$-enhancement. Because we are keenly interested in potential variations in the abundance ratios, and because the state-of-the art high-resolution models \\citep{bc2003} do not reproduce, in detail, the strengths of many features in the spectra of real galaxies \\citep[e.g. CN, Ca4227, Mg, H$\\delta$, H$\\gamma$;][]{gallazzi2005}, we prefer to derive stellar population parameters from an analysis of absorption line indices. The \\cite{bc2003} SEDs are employed for other purposes and these are discussed below in the context of deriving broadening corrections to our data.  Many of the details regarding the processing of the data, and subsequent corrections, are discussed elsewhere \\citep{kelsonb,datapaper}. However several of the key points are discussed below in \\S \\ref{sec:data} because of their crucial role in the analysis. We then describe our comparison of the absorption line strengths of the E/S0 galaxies to the models of \\cite{thomas} and \\cite{thomas2}, in which we perform non-linear least-squares fits to each galaxy's set of line strengths \\cite[see, also][]{proctor2004b}, {\\it but only for relative differences in the stellar population parameters\\/}. This allows us to explicitly avoid potentially large systematic uncertainties in the direct comparison of the data and models. This methodology is used in the rest of the survey and so is described in some detail. The resulting relative ages and patterns of fitted chemical abundances are described in \\S \\ref{sec:fits} and then discussed further in \\S \\ref{sec:summary}. Our conclusions are summarized in \\S \\ref{sec:conclusions}. While our findings do not depend on the cosmology, we use the cosmological parameters $H0=72$ km/s/Mpc, $\\Omega_M=0.27$, and $\\Omega_\\Lambda = 0.73$, when such parameters are required (e.g. for corrections to the line strengths for any variation in the metric sizes of the apertures from which the spectra were obtained). ",
        "conclusions": "\\label{sec:conclusions}  We have measured absorption line strengths for galaxies in the cluster CL1358+62 using the spectra published in \\cite{kelsonb}. A homogeneous population of early-type galaxies has been analyzed, with the full selection of galaxies in \\cite{kelsonb} to be discussed in a subsequent paper. The largest source of systematic errors that plague the analysis of absorption lines has been eliminated by matching the stellar populations models \\citep{thomas,thomas2} to the mean observed line strengths of the most massive early-type galaxies in CL1358+62. Furthermore, not only does a recalibration of the models allow for accurate differential measurements of the stellar population parameters, but the topology of the $\\chi^2$ minimum can be used to accurately estimate formal errors.  Using only those E/S0 galaxies which have accurate measurements of eight blue Lick/IDS indices (H$\\delta_A$, H$\\gamma_A$, CN$_2$, Ca4227, G4300, Fe4383, Fe4531, and C4668), we fit for the {\\it relative} measures of age and chemical abundances. The fitting has resulted in homogeneous sets of stellar population parameters relative to the reference point used in resetting the model zero-point.  The key results can be summarized as follows:  \\noindent (1) The populations of the E and S0 galaxies are statistically identical in their age and abundance patterns, down to the magnitude limit of $R=21$ mag, consistent with what was found using the fundamental plane \\citep{kelsonc}. Beyond the magnitude limit, selection biases become important.  \\noindent (2) The observed scatter in relative ages is 0.06 dex, where the scatter expected from measurement errors is 0.09 dex. The scatter in ages is consistent with that inferred from the color-magnitude relation \\citep{vdcm33}, and from the fundamental plane \\citep{kelsonc}.  \\noindent (3) We find a tight correlation between relative metallicity and velocity dispersion, with a slope of $0.86\\pm 0.17$ dex/dex. The scatter about this relation is 0.06 dex, where the scatter expected from the formal errors is 0.05 dex. The observed $B-V$ colors are also consistent with the inferred ages and metallicities.  \\noindent (4) The scatter in relative $\\alpha$-enhancement is only 30\\% larger than expected from the formal errors alone. When the nitrogen abundance ratio is allowed to vary, we find a mild anti-correlation of $\\Delta$[$\\alpha$/Fe] with $\\log \\sigma$, but with a significance less than $<2\\sigma$. While we find little evidence for a significant correlation (or variation) of [$\\alpha$/Fe] with velocity dispersion, the slope of the local Mg-$\\sigma$ relation can be reproduced by our $\\Delta$[Z/H]-$\\sigma$ correlation. \\cite{thomas2005} argued that Mg$b$-$\\sigma$ arises mostly from the correlation of metallicity with velocity dispersion. Our results go beyond that and suggest that no variation of $\\alpha$/Fe is required. The discrepancy between our conclusions and previously published work on this topic arises from our steeper [Z/H]-$\\sigma$ relation. This discrepancy may be due to the treatment of the broadening corrections, in particular for the narrow Lick iron indices, though more detailed analysis is required to verify this hypothesis. Previous work \\cite[e.g.][]{trager2000} indicated that cluster E/S0s form a three-parameter family of objects, with [Z/H] and $\\alpha$/Fe correlated with velocity dispersion. Historically, models with variable ratios of Type I and Type II supernovae and subsequent galactic winds have been invoked to explain such a complex family of objects. Elimination of the correlation of $\\alpha$-enhancement with velocity dispersion would greatly simplify models of E/S0 formation.  \\noindent (5) We find that the enhancement of nitrogen is strongly correlated with velocity dispersion, while carbon abundance appears to show no variation. This N-$\\sigma$ correlation has not been observed before. This new correlation originates either in deficiencies in the stellar population models, or implies new physics to be incorporated into our understanding of E/S0 formation and evolution. The most likely explanation is one in which the nitrogen abundance ratios do not arise from mechanisms specifically tied to $\\sigma$, but that the nitrogen-enhancement is correlated with metallicity. In other words, the nitrogen is secondary in origin, and the progenitors of E/S0 galaxies experienced significantly extended star-formation.  These data form the only sample of early-type galaxies in which these particular abundance patterns have been found, so far. In all other respects the sample shows the uniform age distribution and a steep correlation of mean metallicity with velocity dispersion inferred from other diagnostics.  While models which include variable abundance ratios \\citep[e.g.][]{trager2000,thomas,thomas2} are relatively immature, their development has been encouraging, and the data are already well-fit by such models. The blue Lick indices appear to be powerful diagnostics of stellar population ages and provide excellent, self-consistent fits to the metal lines, even if the resulting abundance ratios run counter to much of the previous literature. Despite these discrepancies, the observed line strength-line width relations of nearby galaxies, published by \\cite{nelan2005}, are recovered, with some notable exceptions. Those relations that match well do bolster our faith in these findings. However, nearly all of the iron line strength-line width relations are not matched well, and the published data on nearby galaxies have had large multiplicative Doppler corrections applied. We suggest that perhaps past treatment of these corrections may be in error. That the $(B-V)$ colors are consistent with our stellar population parameters not only reinforces our conclusions, but also indicates that dust is not an important component in the optical SEDs of early-type galaxies.  The results on enhancements (over Fe) of $\\alpha$, nitrogen, and carbon have important implications. While it remains to be seen whether the carbon- and nitrogen-sensitive indices are being modeled correctly, the nitrogen enhancement-velocity dispersion correlation agrees very well with the correlation of [N/Z] with [Z/H] expected to arise from secondary production of nitrogen by AGB stars. Taken together, our differential stellar population parameters suggest that these massive early-type galaxies form a family with only a single parameter: galaxy velocity dispersion (or possible mass).  The uniform ages and simple trends of metallicity with velocity dispersion appear to contradict the observation that fewer S0s exist at higher redshift \\citep{morph2,postman2005}. We conclude that the timescales and mechanisms with which the populations in S0s form are statistically identical to that of ellipticals of the same mass. Early-type galaxies, now shown to have significant self-enrichment and extended star-formation histories, must be reconciled with hierarchical formation \\citep[e.g.][]{kauffman1993} in a $\\Lambda$CDM universe. The uniform ages and tight mass-metallicity relation may not be compatible with significant merging and histories of extended star-formation.  In a subsequent paper, we will employ our machinery for measuring these age and abundance patterns in a range of clusters from $z=0$ to $z=0.83$. With comparisons of the properties of cluster galaxies over a range of redshifts, we will test the universality of some of the conclusions presented here, quantify changes in the stellar population parameters with cluster redshift \\citep[see, e.g.,][]{kelson01,jorgensen0152,barr2005}, and attempt to reconcile the above results with a coherent picture of the formation and evolution of galaxies in rich clusters."
    },
    "0606/astro-ph0606532_arXiv.txt": {
        "abstract": "We present the far-infrared (IR) maps of a bipolar planetary nebula (PN), NGC 650, at 24, 70, and $160\\micron$ taken with the Multiband Imaging Photometer for Spitzer (MIPS) on-board the Spitzer Space Telescope. While the two-peak emission structure seen in all MIPS bands suggests the presence of a near edge-on dusty torus, the distinct emission structure between the $24\\micron$ map and the 70/$160\\micron$ maps indicates the presence of two distinct emission components in the central torus. Based on the spatial correlation of these two far-IR emission components with respect to various optical line emission, we conclude that the $24\\micron$ emission is largely due to the [\\ion{O}{4}] line at $25.9\\micron$ arising from highly ionized regions behind the ionization front, whereas the 70 and $160\\micron$ emission is due to dust continuum arising from low-temperature dust in the remnant asymptotic giant branch (AGB) wind shell. The far-IR nebula structure also suggests that the enhancement of mass loss at the end of the AGB phase has occurred isotropically, but has ensued only in the equatorial directions while ceasing in the polar directions. The present data also show evidence for the prolate spheroidal distribution of matter in this bipolar PN. The AGB mass loss history reconstructed in this PN is thus consistent with what has been previously proposed based on the past optical and mid-IR imaging surveys of the post-AGB shells. ",
        "introduction": "NGC 650 (PK 130-10\\fdg1, M 76, Little Dumbbell Nebula) is a large ($\\sim300\\arcsec$; \\citealt{balick92}) bipolar planetary nebula (PN) of the ``late butterfly'' type \\citep{balick87}. The nebula structure in the optical consists of the bright rectangular core of $95\\arcsec \\times 40\\arcsec$ (the long side perpendicular to the bipolar axis) and a pair of fainter lobes extending $\\sim90\\arcsec$ and $\\sim150\\arcsec$ from the central star that is attached to the long side of the rectangular core. The central core was long suspected to be a nearly edge-on torus, and the most recent kinematical study has eloquently demonstrated that the core is an inclined torus and the lobes are blown-bubbles expanding into the polar directions \\citep{bryce96}.  Under the framework of the widely-accepted general interacting stellar wind (GISW) model (e.g., \\citealt{balick87}), the formation of bipolar PNs is understood as a two-step process. First, the progenitor star has to lose its envelope material via mass loss (${\\dot M} \\sim 10^{-6}$ M$_{\\odot}$, $v \\sim 10$ km s$^{-1}$) during the asymptotic giant branch (AGB) phase of the evolution. This AGB wind leads up to the so-called superwind (${\\dot M} \\sim 10^{-4}$ M$_{\\odot}$, $v \\sim 20$ km s$^{-1}$) at the end of the AGB phase prior to the exhaustion of the envelope material. During the superwind phase, mass loss is expected to cause concentration of the ejected matter into the equatorial plane. The AGB mass loss therefore results in the equatorially-enhanced circumstellar shell.  Then, a hot and tenuous fast wind (${\\dot M} \\sim 10^{-9}$ M$_{\\odot}$, $v \\sim 10^{3}$ km s$^{-1}$) begins to emanate from the progenitor just prior to the beginning of the PN phase when the circumstellar material starts to ionize. This fast wind pushes into the surrounding equatorially-concentrated envelope, and the flow of the wind is thus channeled into the polar directions. The preferential wind flow towards the polar directions creates the typical bipolar lobes as wind-blown bubbles. The GISW scheme is thus capable of producing various PN shapes (both bipolar and spheroidal shells) depending primarily on the degree of the equatorial enhancement in the surrounding AGB shell.  Hence, in order for the GISW scheme to work, the equatorially-enhanced AGB shell must be present when a fast wind is initiated. The GISW model itself, however, does not address how to generate the equatorial enhancement in the AGB wind shell. The presence of such circumstellar shells with a built-in equatorial density enhancement has been confirmed by imaging surveys of proto-PNs in thermal IR dust emission from the innermost torus (e.g., \\citealt{skinner94,meixner97,dayal98,meixner99,ueta01}) and dust-scattered star light (e.g., \\citealt{ueta00}). Based on the results from the largest mid-IR and optical imaging surveys to date \\citep{meixner99,ueta00}, a so-called ``layered shell model'' has been proposed to explain {\\sl both} the bipolar and elliptical proto-PN morphologies \\citep{ueta02,ueta03}. In this model, the proto-PN shells are thought to consist of three generic layers, each of which possesses a specific structure reflecting the geometry of mass loss at the time. The outer spherically symmetric layer represents isotropic mass loss during the early AGB phase, while the inner toroidal layer embodies equatorially-enhanced mass loss during the superwind epoch at the end of the AGB phase. Between these layers there is an intermediate spheroidal layer, which results from a gradual transformation of the mass loss geometry from isotropic to equatorially-enhanced over the course of the AGB phase.  The unique strength of this layered shell model is its versatility. Radiative transfer calculations using this model have successfully reproduced both bipolar and elliptical proto-PN morphologies by simply varying the degree of the equatorial density enhancement (i.e., optical depth) of the shell \\citep{meixner02,ueta03}. The implication of the layered shell model is that there is no intrinsic distinction between the elliptical and bipolar proto-PN shells other than the degree of the equatorial density enhancement. Bipolar PNs would emerge from equatorially-enhanced proto-PNs and elliptical PNs would descend from proto-PNs of rather isotropic density distribution. The presence of the spheroidal intermediate layer is the key to account for the elongation along the polar axis in proto-PN shells (e.g., \\citealt{ueta01,meixner02,ueta03,meixner04}).  It is clear that the equatorially-enhanced superwind {\\sl initiates} the subsequent aspherical structure development that eventually results in the observed complex PN structures. However, it is still not yet evident how the shell structure development ensues in the AGB and post-AGB winds. How does the spheroidal density distribution arise prior to the equatorial enhancement at the end of the AGB phase? Is the post-AGB mass loss significant in shell shaping? Does a fast wind simply ``snow-plow'' the surrounding material, tracing the pre-existing toroidal density distribution, or contribute to an additional equatorial enhancement? These are only a few questions concerning the shell structure development in the circumstellar shells of evolved stars.  Further to understand the role of mass loss in the shell structure development during the AGB phase and beyond, the remnant AGB shells need to be investigated in PNs via sensitive far-IR observations of thermal dust emission, since such remnant AGB shells are very much dispersed and cold. In the following, we present the results of Spitzer Space Telescope far-IR mapping observations of a bipolar PN, NGC 650, to probe the history of AGB and post-AGB mass loss in this object. Below, we describe the data set and reduction procedure (\\S 2),  discuss the results (\\S 3), and summarize conclusions (\\S 4). ",
        "conclusions": ""
    },
    "0606/astro-ph0606704_arXiv.txt": {
        "abstract": "Quasars are powered by accretion onto supermassive black holes, but the problem of the duty cycle related to the episodic activity of the black holes remains open as one of the major questions of cosmological evolution of quasars. In this Letter, we obtain quasar duty cycles based on analyses of a large sample composed of 10,979 quasars with redshifts $z\\le2.1$ from the Sloan Digital Sky Survey (SDSS) Data Release Three. We estimate masses of quasar black holes and obtain their mass function (MF) of the present sample.  We then get the duty cycle $\\bar{\\delta}(z)=10^{-3}\\sim 1$ based on the So\\l tan's argument, implying that black holes are undergoing multiple episodic activity. We find that the duty cycle has a strong evolution. By comparison, we show that evolution of the duty cycle follows the history of cosmic star formation rate (SFR) density in the Universe, providing intriguing evidence for a natural connection between star formation and triggering of black hole activity. Feedback on star formation from black hole activity is briefly discussed. ",
        "introduction": "Supermassive black holes are relics of quasars  in the Universe (So\\l tan 1982; Rees 1984, 1990). Evolution of quasars is led by switching on and off accretion onto the black holes. During their entire evolution, how many times and how many black holes are triggered? what is the triggering mechanism? and why do quasars switch off?  The duty cycle, defined as the fraction of active black holes to their total number is a key to tackling the above problems (Richstone et al. 1998; Martini 2004). A popular method to get the duty cycle invokes the continuity equation and the MF of quasar black holes. With an assumption that quasar black holes have the same Eddington ratio, their MF can be obtained from the luminosity function and finally the duty cycle can be found from the continuity equation (Small \\& Blandford 1990; Marconi et al. 2004). This is a convenient way to discuss evolution of the black holes, but the degeneracy of the Eddington ratio and the duty cycle still holds. Actually, not only are the Eddington ratios {\\em not} constant for different quasars at different epochs, but they appear to be quite scattered (Vestergaard 2004). % The duty cycle is poorly understood as a statistical parameter tracing the evolution of quasar populations.  In recent years, there has been much progress in estimating black hole masses both in nearby galaxies and distant quasars. The empirical relation of reverberation mapping allows us to conveniently obtain the black hole masses from a large sample and directly get their MF. Thus it becomes realistic to get new clues to understand the evolution of quasars. Fortunately, by invoking the MF, we can decouple the degeneracy of the duty cycle and the Eddington ratio to get the duty cycle.  In this Letter, we use available SDSS data to directly get the MF of the black holes so as to discuss the duty cycle based on the MF and find that there is a strong cosmological evolution. Our calculations assume a cosmology with the Hubble constant $H_0=70~{\\rm Mpc^{-1}~km~s^{-1}}$, $\\Omega_{\\rm m}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "The evolution of quasars is jointly controlled by the triggering mechanism and accretion. The duty cycle is a key parameter to unveil the evolution of quasars. The results of the present paper show a very strong cosmological evolution of quasar's duty cycle. The triggering history represented by $\\bar{\\delta}(z)$ is quite similar to the evolution of cosmic SFR density. This indicates that star formation may be the direct mechanism to trigger the activity of the black holes.  The duty cycle can be roughly justified from the galaxy luminosity function. According to the luminosity function of galaxies at $1.8\\le z\\le 2.0$ (Dahlen et al. 2005), the galaxy number density is $n_{\\rm G}\\approx 826$Gpc$^{-3}$ for galaxies brighter than $R-$band magnitude $M_{\\rm R}=-24$. This corresponds to the number density of galaxies with black hole mass larger than $10^9\\sunm$ converted from $\\log \\left(M_{\\rm BH}/\\sunm\\right)=-0.5M_{\\rm R}-3$ (McLure \\& Dunlop 2001). Number density of quasars brighter than $M_i=-28$ (corresponding to a black hole with mass $>10^9\\sunm$ if it is accreting at the Eddington limit) is $n_{\\rm Q}\\approx 175$ Gpc$^{-3}$ based on the quasar luminosity function (Richards et al. 2006). We thus estimate a duty cycle of $\\sim 0.18$, which is roughly consistent with the present results (see Figure 2a).  The SFR density rises with redshift out to $z=1.5$ and appears to be roughly flat between $z\\approx 1.5$ and $z\\approx 3.0$. This tendency could be explained by strong feedback from activity of black holes to their host galaxies (Silk 2005; Di Matteo et al. 2005; Croton et al. 2006).  With a balance between star formation and feedback in $z\\sim 1.5-3.0$, the violent star formation is then suppressed. However, the SFR density is going to decrease with time due to a shortage of gas and the BH duty cycle follows this trend. To further confirm this, future work will focus on the dependence of the duty cycle on the black hole masses. It could show the feedback-dependence on the BH growth itself. Additionally the total accretion time (net lifetime) of black holes will be then obtained."
    },
    "0606/hep-ph0606007_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ Upper limits on neutrino masses from cosmology have been reported recently to reach the impressive sub-eV level, which is competitive with future terrestrial neutrino experiments.  In this brief review of the latest limits from cosmology we point out some of the caveats that should be borne in mind when interpreting the significance of these limits. ",
        "introduction": "The latest results from the WMAP satellite \\cite{spergel} confirm the success of the $\\Lambda$CDM model, where $\\sim 75\\;\\%$ of the mass-energy density is in the form of dark energy, with matter, most of it in the form of cold dark matter (CDM)  making up the remaining 25 \\% .   Neutrinos with masses on the eV scale or below will be a hot component of the dark matter and will free-stream out of overdensities and thus wipe out small-scale structures.  This fact makes it possible to use observations of the clustering of matter in the universe to put upper bounds on the neutrino masses.  An excellent review of the subject is that of Lesgourgues and Pastor \\cite{lesgourgues}. With the improved quality of cosmological data sets seen in recent years, the upper limits have improved, and some quite impressive claims have been made in the recent literature. We will in the following summarize the latest upper bounds and point out some of the potential systematic uncertainties that need to be clarified in the future. ",
        "conclusions": "Current cosmological observations provide strong upper limits on the sum of the neutrino masses.  However, when assessing the significance of these limits one should bear in mind that several assumptions are involved in deriving these limits, both cosmological and astrophysical. It is an important task for further research to clarify how sensitive the results are to these assumptions, and how well they are justified."
    },
    "0606/astro-ph0606674.txt": {
        "abstract": "There has recently been growing evidence for the existence of neutron stars possessing magnetic fields with strengths that exceed the quantum critical field strength of $4.4 \\times 10^{13}$ G, at which the cyclotron energy equals the electron rest mass.  Such evidence has been provided by new discoveries of radio pulsars having very high spin-down rates and by observations of bursting gamma-ray sources termed magnetars.  This article will discuss the exotic physics of this high-field regime, where a new array of processes becomes possible and even dominant, and where familiar processes acquire unusual properties.  We review the physical processes that are important in neutron star interiors and magnetospheres, including the behavior of free particles, atoms, molecules, plasma and condensed matter in strong magnetic fields, photon propagation in magnetized plasmas, free-particle radiative processes, the physics of neutron star interiors, and field evolution and decay mechanisms.  Application of such processes in astrophysical source models, including rotation-powered pulsars, soft gamma-ray repeaters, anomalous X-ray pulsars and accreting X-ray pulsars will also be discussed. Throughout this review, we will highlight the observational signatures of high magnetic field processes, as well as the theoretical issues that remain to be understood. ",
        "introduction": "\\label{sec:intro}  Since their theoretical conception by Baade \\& Zwicky (1934) neutron stars have been fascinating celestial objects, both for study of their exotic environments and for their important place in stellar evolution.  Among the first signals to be detected from neutron stars came from radio pulsars (Hewish et al. 1968), spinning many times per second, distinguishing themselves from the background of interplanetary scintillation signals by their extremely regular pulsations.  Pulsars were also soon discovered to be spinning down, their periods increasing also very regularly.  The rotating magnetic-dipole model (Pacini 1967, Gold 1968, Ostriker \\& Gunn 1969), in which the pulsar loses rotational energy through magnetic dipole radiation, was dramatically confirmed with the discovery that the spin-down power predicted for the Crab pulsar was a nearly perfect energetic match with the radiation of its synchrotron nebula.  The rotating dipole model also accounts for the observed rate of spin-down of the Crab and other pulsars, with required surface magnetic fields in the range of $10^{11} - 10^{13}$ Gauss for the first detected pulsars.  This range has since significantly broadened, first with the discovery of a class of pulsars having periods of several milliseconds (Backer et al. 1982), believed to have been spun-up by accretion torques of a binary companion (Alpar et al. 1982), and much lower surface magnetic fields in the range of $10^8 - 10^{10}$ Gauss.  Recent surveys have also discovered pulsars with very high period derivatives (e.g. Morris et al. 2002, McLaughlin et al.~2003) that imply surface fields up to around $10^{14}$ Gauss.  Another class of neutron stars was discovered at X-ray and $\\gamma$-ray energies and may possess even stronger surface magnetic fields.  Such stars are now referred to as magnetars % NEW: (Duncan \\& Thompson 1992), since they most probably derive their power from their magnetic fields rather than from spin-down energy loss (see Woods \\& Thompson 2005). Within the magnetar class there are two types of sources that were originally thought to be very different objects, although they are now believed to be closely related.  The Anomalous X-Ray Pulsars (AXPs) were discovered as pulsating X-ray sources in the early 1980s and were thought at first to be an unusual type of accreting X-ray pulsar, from which the name is derived.  The AXPs have periods in a relatively narrow range of 5 - 11 s, are observed to be spinning down with large period derivatives (Vasisht \\& Gotthelf 1997) and have no detectable companions or accretion disks that would be required to support the accretion hypothesis.  Interpretation of their period derivatives as magnetic dipole spin down imply magnetic fields in the range $10^{14} - 10^{15}$ Gauss.  But such high fields were not widely accepted initially, since their detected X-ray luminosities of around $10^{35}\\,\\rm erg\\,s^{-1}$ exceed their spin-down luminosities by several orders of magnitude.  It was only by connection to another subclass of magnetar, the Soft Gamma-Ray Repeaters (SGRs), that the extremely high magnetic fields of AXPs were adopted.  SGRs were discovered as transient sources that undergo repeated soft $\\gamma$-ray bursts, usually in widely separated episodes.  They undergo both repeated smaller bursting of subsecond duration as well as much more luminous superflares, lasting hundreds of seconds, which so far have not repeated in any single source but may be repeating on much longer timescales.  It was not until some twenty years after their discovery that their quiescent X-ray pulsations were detected and also very high period derivatives (Kouveliotou et al. 1998), both with a range of values very similar to those of AXPs.  Recently, bursts resembling the smaller bursts of SGRs were seen from several AXPs (Kaspi et al. 2003), making it likely that SGRs and AXPs are two variations of the same type of object % NEW - AKH (Thompson \\& Duncan 1993), very strongly magnetized, isolated neutron stars possibly powered by magnetic field decay.  The periods and period derivatives of the various types of isolated pulsars are shown in the $P$-$\\dot P$ diagram of Figure \\ref{fig:PPdot}. Assuming that the spindown torque is due to magnetic dipole radiation, two quantities can be defined from the measured $P$ and $\\dot P$ for each pulsar: (1) Characteristic age $P/(2\\dot P)$: From $\\dot\\Omega\\propto -\\Omega^3$ (where $\\Omega=2\\pi/P$), the age of the pulsar is found to be $T=(P/2\\dot P)[1-(P/P_i)^2]$, where $P_i$ is the pulsar's initial spin period. (2) Surface dipole magnetic field: \\be B_s  = \\left( {\\frac{{3Ic^3 P\\dot P}}{{2\\pi ^2 R^6 }}} \\right)^{1/2} \\simeq 2 \\times 10^{12}{\\rm G}\\,(P \\dot P_{15})^{1/2}, \\label{eq:B0}\\ee where $\\dot P_{15} \\equiv \\dot P/(10^{-15}\\,\\rm s\\,s^{-1})$ (and $P$ is in units of second), and and $I$ ($\\simeq 10^{45}$~g~cm$^2$ and $R$ ($\\simeq 10^6$~cm) are the neutron star moment of inertia and radius. There are presently around 1600 spin-powered radio pulsars known, with periods from $1.5$ ms - 8 s (Manchester 2004)\\footnote{ see http://www.atnf.csiro.au/research/pulsar/psrcat/}. Some fraction of these pulse at other wavelengths, including about 10 in $\\gamma$-rays and about 30 in X-rays.  The magnetars, eight AXPs and four SGRs\\footnote{ see http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html}, occupy the upper right-hand corner of the diagram and curiously overlap somewhat with the region occupied by the high-field radio pulsars.  However, the two types of objects display very different observational behavior.  The high-field radio pulsars have very weak or non-detectable X-ray emission and do not burst (e.g. Kaspi \\& McLaughlin 2005), while the magnetars have no detectable radio pulsations, with the exception of the recent detection of radio pulsations in the transient AXP XTE J1810-197 (Camilo et al. 2006).  The intrinsic property that actually distinguishes magnetars from radio pulsars is presently not understood.  \\begin{figure} \\includegraphics[width=15cm]{PPdot.eps} \\caption{Plot of period vs. period derivative for the presently known rotation-powered pulsars and magnetars. Lines of constant characteristic age, $P/2\\dot P$, and surface dipole field (see Eqn [\\ref{eq:B0}]) are superposed.} \\label{fig:PPdot} \\end{figure}  A third class of strongly magnetized neutron stars are the accreting X-ray pulsars (Parmar 1994).  These sources are members of binary systems with high-mass companions that either have strong stellar winds or overflow their Roche lobes, transferring material to the neutron star.  Inside an Alfven surface where the pressure of the accretion flow equals the neutron star magnetic field pressure, the accreting material is funneled along the magnetic field to the poles. The heated accretion flow then radiates from hot spots that rotate with the star, thereby producing pulsations.  However, since accretion torques dominate the period derivative evolution, the surface magnetic fields of X-ray pulsars cannot be determined using the rotating dipole model, as they are for the rotation-powered pulsars and magnetars. Instead, magnetic fields of these objects have been measured from the energies of the cyclotron resonant scattering features (CRSFs) that appear in their spectra (see Orlandini \\& dal Fiume 2001, for review), since the fundamental occurs at the electron cyclotron energy, $E_{ce} = \\hbar eB/(m_ec)=11.58\\,(B/10^{12}\\,{\\rm G})$~keV. Table 1 lists the X-ray pulsars, the energies of CRSFs that have been detected in their spectra and the inferred magnetic field strengths.  In most cases, the measured magnetic fields range from $1-5 \\times 10^{12}$ Gauss, with the highest fields being around $4 \\times 10^{12}$Gauss.  The formation of such line features will be discussed in section \\ref{sec:XRPs}.  \\begin{table} \\caption{X-Ray Pulsars with Cyclotron Resonant Scattering Features$^a$} \\begin{center} \\begin{tabular}{llc} \\hline Source & Energy (keV)$^b$ & B ($10^{12}$ Gauss)$^c$\\\\ \\hline 4U 0115+63 & 12 (4) & 1.0\\\\ 4U 1907+09 & 18 (1)& 1.6\\\\ 4U 1538-52 & 20 & 1.7 \\\\ Vela X-1 & 25 (1) & 2.2 \\\\ V 0332+53 & 27 (2) & 2.3 \\\\ Cep X-4 & 28 & 2.4 \\\\ Cen X-3 & 28.5 & 2.5 \\\\ X Per & 29 & 2.5 \\\\ XTE J1946+274 & 36 & 3.1 \\\\ MX0656-072 & 36 & 3.1 \\\\ 4U 1626-67 & 37 & 3.2 \\\\ GX 301-2 & 37 & 3.2 \\\\ Her X-1 & 41 & 3.5 \\\\ A 0535+26 &  47$^d$ (2) & 4.1 \\\\ \\hline \\end{tabular} \\end{center} {\\small $^a$ Data is from Heindl et al. (2004).  $^b$ Numbers in parentheses are the number of cyclotron harmonics detected.  $^c$ Magnetic field strength for viewing angle along field direction.}  $^d$ Kretschmar et al. (2005) \\end{table}  Clearly, there is a broad range of astrophysical sources in which the magnetic fields approach and exceed the quantum critical field strength, $B_Q \\equiv m_e^2c^3/(e\\hbar)=4.414 \\times 10^{13}$~G, at which the cyclotron energy equals the electron rest mass.  In this regime, the magnetic field profoundly affects physical processes and introduces additional processes that do not take place in field-free environments.  From a physics point of view, strongly magnetized neutron stars provide the only environment in which to measure and test these effects.  From an astrophysics point of view, it is necessary to investigate the physics of strong magnetic fields in order to effectively model and understand the nature of the sources. This article will attempt to address both points of view by providing a review of the basic physical processes important in neutron star interiors and magnetospheres, as well as a review of the source models in which these processes play a critical role.  We begin with a description of a free electron in a magnetic field in both classical and quantum regimes.  Subsequent sections then discuss and review current understanding of the behavior of matter, atoms, molecules and plasma, in strong magnetic fields, photon propagation in magnetized plasmas, free-particle radiative processes, the physics of neutron star interiors, and field evolution and decay mechanisms. Then we will review models for magnetized atmospheres, non-thermal radiation from rotation-powered pulsars, burst and quiescent radiation from SGRs and AXPs, and emission from accreting X-ray pulsars.  Other useful books and reviews on strongly magnetized neutron stars include Meszaros (1992), Duncan (2000), Lai (2001) and Harding (2003).  A complimentary work concentrating on stellar magnetism in general is the book by Mestel (1999). ",
        "conclusions": "In this article we have reviewed the physics that applies in extremely strong magnetic fields, the properties of strongly magnetized neutron stars, atmospheres and magnetospheres, and the present status of our understanding and models of the astrophysical sources that are their manifestation.  While the use of non-relativistic quantum mechanics is adequate (even for $B>B_Q$) for describing bound states or any processes where the electron stays in the ground Landau level, one must use relativistic quantum mechanics (QED) in computing the rates and cross sections for many free particle and photon processes. Nearly all of the cross sections for the first and second-order free particle processes in strong magnetic fields, such as cyclotron radiation and absorption, one photon and two-photon pair creation and annihilation and Compton scattering, have been calculated and studied. Only a few of the higher-order processes, such as Bremsstrahlung and photon splitting, have been investigated since they become increasingly complex.  The unusual and interesting properties of photon propagation in strong magnetic fields has been a topic of intensive study in recent years, especially after the discovery of magnetars.  Among all the processes that have been investigated, vacuum polarization has been found to be of particular importance.  A variety of astrophysical sources, that include rotation-powered and accretion-powered neutron stars and the magnetars, SGRs and AXPs, are believed to be very strongly magnetized neutron stars.  The processes that operate only in such strong fields may be fundamental to the functioning of these sources.  For rotation-powered pulsars, one-photon pair creation is thought to be the primary process attenuating high-energy photons, and the created pairs may be critical in the production of the observed radio pulsations.  We do not yet fully understand though how and why the radio emission seems to turn off well before the pulsar spin-down has ceased.  Theoretical models show that pair production decreases as the pulsar ages, but does it turn off suddenly, or are there some threshold properties (e.g. multiplicity or energy spectrum) for radio emission?  In the case of the magnetars, power levels far exceeding their spin-down power requires fast magnetic field decay on timescales only possible through processes such as ambipolar diffusion, that become effective in fields above $10^{14}$ G.  We do not yet understand, however, why these sources are radio quiet.  Is pair production suppressed by the strong field or are the collective plasma processes disrupted? Theoretical studies have so far not been able to find a convincing mechanism for the suppression of the radio emission in magnetars. Even more puzzling is the growing number of magnetars and radio pulsars having very similar spin properties, and therefore similar implied surface magnetic field strengths.  Aside from their radio emission properties, radio pulsars and magnetars also have very different X-ray emission levels, transient emission and glitching behavior.  What are the hidden characteristics of these sources that distinguish them?  Also of interest are the emerging population of ``dim'' isolated neutron stars with apparently pure thermal emission. The spectral features detected in some of these sources are exciting, but their identifications remain unclear. The nature of these sources is unknown. Could they be descendants of magnetars?  To find answers to these questions, new ideas and more theoretical work are surely needed.  But new and more sensitive instruments are on the horizon that will also provide some clues as well.  The ALFA pulsar survey (Cordes et al. 2005) began operation last year at Arecibo, one of the world's most powerful radio telescopes, and this survey is expected to discover at least 1000 new radio pulsars, nearly doubling the current number.  Among the newly discovered pulsars will be many more radio pulsars with magnetar field strengths.  The Gamma-Ray Large Area Telescope (GLAST) (McEnery et al. 2004), due to launch in 2007, will have a point source detection threshold about 20 times below that of EGRET and will discover hundreds of new $\\gamma$-ray pulsars.  GLAST will also have sensitivity up to 200 GeV and will be able to make very sensitive measurements of spectral high-energy cutoffs.  Third-generation ground-based Air Cherenkov detectors, such as H.E.S.S.  (Hinton et al. 2004) in Namibia, have begun operation.  They are sensitive to $\\gamma$-rays in the range 50 GeV to 50 TeV and may discover or put important constraints on pulsar and magnetar spectra and their nebulae.  Further into the future are planned X-ray telescopes, such as Constellation X, XEUS, and X-ray polarimeters, such as AXP, POGO and ACT.  Polarimeters in particular will be extremely important in looking for some of the signatures of very strong magnetic fields that have been discussed in this article, such as the vacuum polarization resonance and photon splitting cutoffs.   \\ack We thank Matthew Baring for comments on the manuscript. This work has been supported in part by NSF grant AST 0307252 (DL)."
    },
    "0606/astro-ph0606710_arXiv.txt": {
        "abstract": "The colour and luminosity distributions of red galaxies in the cluster Abell\\,1185 ($z=0.0325$) were studied down to $M^*+8$ in the $B$, $V$ and $R$ bands. The colour--magnitude (hereafter CM) relation is linear without evidence for a significant bending down to absolute magnitudes which are seldom probed in literature ($M_R=-12.5$ mag). The CM relation is thin ($\\pm0.04$ mag) and its thickness is quite independent from the magnitude.  The luminosity function of red galaxies in Abell\\,1185 is adequately described by a Schechter function, with a characteristic magnitude and a faint end slope that also well describe the LF of red  galaxies in other clusters. There is no passband dependency of the LF shape other than an obvious $M^*$ shift due to the colour of the considered population. Finally, we conclude that, based on colours and luminosity, red galaxies form an homogeneous population over four decades in stellar mass, providing a second evidence against faint red galaxies being a recent cluster population. ",
        "introduction": "One of the major problems facing models of galaxy formation is why cold dark matter models predict a larger number of low-mass galaxies than observed. Today, extremely faint ($M \\gg M^*+5$ mag) galaxies are still a poorly studied population because they are difficult to find and, once found, the measurement of their properties is not trivial since the measured properties are usually strongly affected by selection effects. Thus, understanding the properties of the lowest luminosity galaxies, which are presumably also very low mass, might shed light on some of the unsolved questions related to the production of galaxies in low-mass halos.  The bulk of the observational work on low-luminosity galaxies is limited to clusters environment. The pioneering work by Visvanathan \\& Sandage (1977) presents the colour-magnitude relation over an eight magnitude range, although the very large majority of the data is relevant to the brightest four magnitudes. Secker et al. (1997) reports that the colour-magnitude relation (also known as the red sequence) is linear in the Coma cluster over a seven magnitude range, i.e. down to $M_R=-15.0$ mag. Conselice et al. (2002) confirms that galaxies in the first four magnitudes of the Perseus clusters obey to the usual colour--magnitude relation (e.g. Bower, Lucey \\& Ellis 1992), but faint candidate members of the cluster tend to depart from the colour of the red sequence, being bluer or redder. The scatter around the red sequence is small ($\\sigma \\approx 0.07$ mag) down to $M_R\\sim-17$ mag, but rises to $\\sigma \\approx 0.5$ mag at $M_R\\sim-14$ mag. Instead, Evans et al. (1990) found that the fainter Fornax galaxies become redder, not bluer. Therefore, it is unclear what is the shape of the red sequence at faint magnitudes and if it is universal or if it changes from cluster to cluster.  In other environments, information is even scarcer. In the review article by Mateo (1998) on dwarf galaxies in the local group, the colour-magnitude relation includes about 30 galaxies in the range $-22<V<-11$ mag. Blanton et al. (2005) derived luminosity and colour distributions of faint galaxies in the local universe for a flux and surface brightness limited sample. However, it is unclear how much their results are affected by the fact that they assume no environment dependence and a uniform spatial distribution of the studied sample, and later on, the distributions were claimed environmental-dependent and the sample clustered (see Andreon, Punzi \\& Grado 2005 for a discussion).  \\begin{figure*} \\psfig{figure=ima_field.ps,width=12truecm,clip=} \\caption[h]{(Compressed) image of the studied field in the $R$ band, with red galaxies (members and background) marked with an ellipse. North is up and east is to the left. The field of view is $42\\times28$ arcmin$^2$.} \\end{figure*}  In this paper, we report results about a volume complete sample of galaxies in a nearby cluster of galaxy, Abell\\,1185. Our data are deep enough to reveal extremely faint objects, allowing us to investigate the properties of red galaxies over a nine magnitude range.  Abell\\,1185, a cluster of richness one class (Abell\\,1858), has a redshift $cz = 9800$ km s$^{-1}$ and a velocity dispersion of $\\sigma_v=740\\pm60$ km s$^{-1}$ (Mahdavi et al. 1996), and an X-ray luminosity of  $L_X$(0.5-3 keV) $= 1.6 \\ 10^{43}$ $h_{50}^{-1}$ ergs s$^{-1}$ (Jones \\& Forman 1984).  For the Abell\\,1185 cluster, we adopt a distance modulus of 35.7 mag ($H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $z=0.0325$). ",
        "conclusions": "We have studied the colour of red galaxies of the Abell\\,1185 cluster at $z=0.0325$ down to $M^*+8$ in the $B$, $V$ and $R$ bands. The colour--magnitude relation is linear without evidence for a significant bending down to absolute magnitudes seldom probed in literature ($M_R=-12.5$ mag). It is also thin ($\\pm0.04$ mag) and its thickness is not a strong function of magnitude. The luminosity function of red galaxies in Abell\\,1185 is adequately described by a Schechter function, with characteristic magnitude and faint end slope that also well describe the LF of red  galaxies in other clusters. There is no passband dependency of the LF shape,  other than an obvious $M^*$ shift due to the colour of the considered population. Red galaxies form an homogeneous population, over four decades in stellar mass, for what concerns colours and luminosity. Homogeneity in colour implies an old star age or a synchronization between the history of star formation of the various galaxies, down to $M^*+8$."
    },
    "0606/gr-qc0606082_arXiv.txt": {
        "abstract": "Non-minimal actions with matter represented by a scalar field coupled to gravity are considered in the context of a homogeneous and isotropic universe. The coupling is of the form $-\\f{1}{2}\\xi \\phi^2 R$. The possibility of successful inflation is investigated taking into account features of loop cosmology. For that end a conformal transformation is performed. That brings the theory into the standard minimally coupled form (Einstein frame) with some effective field and its potential. Both analytical and numerical estimates show that a negative coupling constant is preferable for successful inflation. Moreover, provided fixed initial conditions, larger $|\\xi|$ leads to a greater number of {\\em e}-folds. The latter is obtained for a reasonable range of initial conditions and the coupling parameter and indicates a possibility for successful inflation. ",
        "introduction": "The simplest form of a non-minimally coupled action is given by \\be \\label{Action} S[g_{ab}, \\phi]=\\int{d^4 x \\sqrt{-g} \\left( \\frac{f(\\phi)}{2 \\kappa} R - \\frac{1}{2}g^{ab}\\nabla_a \\phi \\nabla_b \\phi - U(\\phi) \\right)} \\ee where $g \\equiv \\det(g_{ab})$, $\\kappa = 8 \\pi G$, $R$ is the scalar curvature of the metric and $U(\\phi)$ is the self-interacting potential of the matter scalar field $\\phi$. The curvature term is coupled to matter through the coupling function $f(\\phi)$ and is manifestly Lorentz invariant. There are some obvious restrictions on the coupling function. First of all, as we shall see, it enters the symplectic structure for the gravitational (connection) variables. Thus it must not vanish anywhere. Moreover, one can view this function as defining an effective gravitational constant $\\kappa$, in the sense \\be \\label{kappaEff} \\kappa_{\\rm eff} = \\kappa/f(\\phi) \\ee Hence we require $f(\\phi)$ to be always positive. Note also that the scalar field is expected to change sign, e.g. during the reheating epoch. It is then natural to restrict $f(\\phi)$ to be an even function of $\\phi$. Finally, the coupling function must satisfy the limit $f(\\phi) \\rightarrow 1$ when $\\phi \\rightarrow 0$, so that the standard normalization is recovered in the absence of a scalar field.  Given the restrictions above, the simplest form of the coupling function is \\be \\label{CFcn} f(\\phi)=1-\\sigma \\phi^2.\\ee (One can think of this expression as the Taylor expansion up to the quadratic order. We are interested in the case of a weak coupling, so the latter should be a good approximation. Note also that coupling of higher degree would require a dimensionful coupling constant, while $\\xi$ is unitless.) In fact, this is the form of coupling function that is normally considered in the literature (see, for example, \\cite{Carroll}). We will refer to $\\sigma$ as the {\\em coupling strength} (not to be confused with the {\\em coupling constant} $\\xi\\equiv \\f{\\sigma}{\\kappa}$ of \\cite{Conf_NMC}). Note that zero $\\sigma$ corresponds to the case of minimal coupling. If $\\sigma>0$, the coupling function can, in principle, become zero and even negative if the value of the scalar field exceeds the critical value \\be \\label{phi_crit}\\phi_{\\rm crit}:=\\f{1}{\\sqrt{|\\sigma|}} \\ee We have deliberately written the absolute value of the coupling strength because, as we shall see later on, it is also of importance in the case of negative coupling.  In this paper, we restrict our attention to a flat homogeneous and isotropic model, described by the FRW-metric \\be ds^2=-dt^2+a(t)^2 d\\vec{r}^{\\,2} \\ee Using a conformal transformation \\be \\label{CT} \\t g_{ab} := f(\\phi) g_{ab}\\ee the action (\\ref{Action}) can be recast in the form \\cite{Conf_NMC} \\be \\label{action} S[\\t g_{ab},\\phi]=\\int{dt \\left[ \\f{3}{\\kappa}(a^2 \\ddot a + a \\dot a^2)+a^3\\left(\\f{1}{2}F(\\phi)^2\\dot\\phi^2-\\t V(\\phi)\\right)\\right]} \\ee where we have introduced the effective potential \\be \\label{Veff} \\t V(\\phi):=\\f{U(\\phi)}{f(\\phi)^2} \\ee and \\be F(\\phi)^2:=\\f{1-\\sigma \\phi^2 (1-\\f{6 \\sigma}{\\kappa})}{f(\\phi)^2} \\approx \\f{1}{f(\\phi)}. \\ee The last approximation holds as long as we are interested in weak coupling such that $\\sigma \\sim (10^{-3}\\div 10^{-2}) \\kappa$; then $\\left( 1- \\f{6 \\sigma}{\\kappa}\\right)$ can be set to 1 in this order of magnitude. Note that in the Hamiltonian formulation, there is a similar canonical transformation under which $F(\\phi)^2:=\\f{1}{f(\\phi)}$ exactly \\cite{nonMin}. Consequently, the relation between the scalar fields (\\ref{phiRel}), derived below, also becomes exact, which considerably simplifies the analysis even if $\\f{6 \\sigma}{\\kappa} \\neq 1$. The last step, that will bring the kinetic term into its canonical form (Einstein frame), is to redefine the field \\bq \\label{phiDef} \\t \\phi &:=&\\int{d\\phi F(\\phi)} \\approx \\int{\\f{d\\phi}{\\sqrt{1-\\sigma \\phi^2}}} \\nonumber\\\\ &=&\\f{1}{\\sqrt{\\sigma}}\\left\\{% \\begin{array}{ll} \\sin^{-1}(\\sqrt{\\sigma\\phi}), & \\hbox{if}  \\quad \\sigma > 0\\\\ \\sinh^{-1}(\\sqrt{|\\sigma|}\\phi), & \\hbox{if} \\quad \\sigma < 0 \\\\\\end{array}% \\right. \\eq Note that for a positive $\\sigma$ the quantity $\\sqrt{\\sigma}\\phi$ must be less than unity, for $f(\\phi)$ must not vanish and, in fact, be positive. From (\\ref{phiDef}) we get \\be \\label{phiRel} \\sqrt{|\\sigma|}\\phi \\approx\\left\\{% \\begin{array}{ll} \\sin(\\sqrt{\\sigma}\\t\\phi), & \\sigma > 0\\\\ \\sinh(\\sqrt{|\\sigma|}\\t\\phi), & \\sigma < 0 \\\\\\end{array}% \\right. \\ee  Aiming at a loop quantization, we now have all the ingredients to proceed to the Hamiltonian formulation of the theory. For that end we introduce the phase space variables: \\be |\\t p|:=a^2, \\quad \\t c:=\\gamma \\dot a, \\quad \\t \\phi \\quad {\\rm and} \\quad \\t \\pi:=p^{3/2}\\dot{\\t \\phi} \\ee here $\\t p$ and $\\t c$ are the gauge-invariant triad and connection components respectively, $\\gamma$ is the Barbero-Immirzi parameter \\cite{AshVarReell,Immirzi}, and $\\t \\pi$ is the field momentum. Generally, the sign of $\\t p$ determines the orientation of the triad. From now on we assume that $\\t p>0$. Again, the tilded variables stand for quantities written in the Einstein frame. For notational convenience, however, we will omit tildes until transforming back to the Jordan frame. In terms of newly defined variables the action (\\ref{action}) takes the form \\be \\label{newAction} S=\\int{dt \\left[ \\frac {3}{\\kappa \\gamma} p \\dot{c} + \\pi \\dot{\\phi} - \\H \\right]}. \\ee with the Hamiltonian \\be \\label{Ham} \\H = -\\f{3}{\\kappa \\gamma^2} \\sqrt{p} c^2 + \\f{\\pi^2}{2 p^{3/2}} + p^{3/2} V(\\phi) \\ee From (\\ref{newAction}) we see that the pairs $\\{p,c\\}$ and $\\{\\phi,\\pi\\}$ are indeed canonically conjugate variables with the Poisson brackets \\be \\label{PB} \\{c,p\\}=\\frac{\\kappa \\gamma}{3}, \\quad \\{\\phi,\\pi\\}=1\\ee Before proceeding to inflation, let us comment on a phenomenological relation between the Einstein and Jordan frames. More specifically, eventually we are interested in the number of e-folds during the inflationary epoch \\be \\t N:=\\ln\\left(\\f{\\t a_e}{\\t a_b}\\right)\\ee where the subscripts stand for the initial (before inflation) and final (after inflation) values of the scale factor $\\t a$. Suppose we have solved the equations of motion in the Einstein frame and obtained $\\t a_b$ and $\\t a_e$. Then it is clear from (\\ref{CT}) that the scale factors in the two frames are related by \\be \\label{aRel} \\t a = \\sqrt{f(\\phi)} a\\ee Therefore the Jordanian number of e-folds is \\be N \\equiv \\ln\\left(\\f{a_e}{a_b}\\right) = \\ln\\left(\\f{\\sqrt{f(\\phi_b)}\\t a_e}{\\sqrt{f(\\phi_e)}\\t a_b}\\right)= \\t N + \\f{1}{2}\\ln\\left(\\f{f(\\phi_b)}{f(\\phi_e)}\\right) \\ee At the end of inflation, the scalar field is essentially zero, while at the beginning of inflation it has its maximum value $\\phi_{\\rm max}$. Thus, using (\\ref{CFcn}), we can explicitly write \\be \\label{Nrel} N=\\t N + \\f{1}{2} \\ln(1-\\sigma \\phi_{\\rm max}^2)\\ee In principle, the beginning of inflation may occur when the scalar field is close to its critical value, such that $|\\sigma \\phi_{\\rm max}| \\approx 1$. In this case, a negative coupling $\\sigma<0$ would yield a negligible correction to the righthand side of (\\ref{Nrel}). A positive coupling, on the contrary, might give a significant contribution resulting in a substantial difference between the Jordanian and Einsteinian number of e-folds. ",
        "conclusions": "In this section, we study numerical solutions to the differential equations (\\ref{HamEqSc}). Again we will be mostly interested in the maximum value of the inflaton for the allowed range of initial conditions and the parameter $\\sigma$.  Let us first discuss the initial field and the scale factor. As was shown in the previous section, the dependence on the former is rather trivial: the initial value $\\phi_0$ acts as constant of integration and should be merely added to the change in the inflaton during the climbing phase.  The dependence of $\\phi_{\\rm max}$ on the initial scale factor is more complicated, but the value of $p_i$ itself is quite restricted. It appears that there is a natural choice of $p_i$, associated with the smallest eigenvalue of the area operator \\cite{ABL}, related to $\\mu_0$. In fact, as one can see from (\\ref{phinode}), the effective growth of the scalar $\\Delta \\phi_{\\rm eff}$ depends on the limits of integration, hence scale factor, only through the ratio $q\\equiv \\f{p}{p_*}$. The lower limit is given by $q_i=\\f{1}{2 j}$, while the upper limit is not fixed a-priory. It has a physical meaning of the marginal value that separates the effective and classical behavior of the spectrum of the geometrical density operator. For the estimate in the previous section we used $q=1$ as the upper limit. A closer look at the graph of $D(q)$ (see Fig. \\ref{fig_Dq}) shows that there is a maximum around $q\\approx 1$. In other words, the behavior of the spectrum is not yet classical. It would be more reasonable to stop the integration (which is meant to be over the effective domain) at a somewhat greater value of $q$, where $D(q)$ is essentially close to unity. Recall that $q_0=1$ was giving $\\Delta \\phi_{\\rm eff}=0.14\\mpl$. At the same time, analysis of Eq. (\\ref{phi_eff}), for $\\Delta \\phi_{\\rm eff}$ as a function of the upper limit of integration $q_0$, shows that the effective growth of the inflaton sensitively depends on $q_0$, whereas the classical formula (\\ref{apprTrans}) indicates a much weaker dependence upon $q_0$. For instance, taking $q_0=2$, one would get $\\Delta \\phi_{\\rm eff}=0.3\\mpl$, while $q_0=4$ yields $\\Delta \\phi_{\\rm eff}=0.5\\mpl$. This implies that if one increases $q_0$, one should expect a somewhat greater value of total $\\Delta \\phi \\equiv \\Delta \\phi_{\\rm eff} + \\Delta \\phi_{\\rm cl}$. Nevertheless, for the sake of estimate, we will stick with $q_0=1$ and compare the numerical results with analytical formulas.  \\begin{figure} \\begin{center} \\includegraphics[width=7cm,height=7cm, angle=270]{D2.eps} \\end{center} \\vskip-0.5cm \\caption{The spectrum of the geometrical density operator $D(q)$ introduced in Eq. (\\ref{defD})} \\label{fig_Dq} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm,height=5.0cm]{sigma.ps} \\end{center} \\vskip-0.5cm \\caption{The change of the inflaton during the classical phase as function of its initial momentum for different values of $\\sigma$ (measured in the units of $\\lp^2$).} \\label{phiSc_num} \\end{figure}  Summing up, as the effective growth of the inflaton is not very substantial and does not depend on the initial field momentum $\\pi_0$ or coupling strength $\\sigma$, the quantity of the most interest is the classical part of $\\Delta \\phi$. The latter is determined by $\\pi_0$ and $\\sigma$.  The dependence of $\\Delta \\phi_{\\rm cl}$ on the initial momentum $p_0$ for five different values if $\\sigma$ is displayed in Fig. \\ref{phiSc_num}. The curves correspond (top to bottom) to $\\sigma/\\lp^2=-0.1,-0.05,0,0.05$ and 0.1 and are plotted for the allowed range of initial momenta: $\\f{\\pi_0}{\\lp}<5 j = 25$. The maximum inflaton value grows with initial momentum and qualitatively resembles the Lambert function. It is steepest for small momenta and flattens out when $\\pi_0$ becomes large. Quantitatively, for $\\sigma=0$ there is a very good agreement with the analytical result (\\ref{apprTrans}) and the graph of Fig. \\ref{fig2}. The curves, corresponding to opposite values of the coupling strength, are situated symmetrically around the minimal curve, which justifies the linear (in $\\sigma$) approximation used in (\\ref{linearSigma}). Furthermore, direct calculation shows that the $\\sigma$-corrections to the central curve are very close to those predicted by formula (\\ref{sigmaCorrections}) for both small and large initial momenta.  Classical climbing brings the scalar field fairly close to the required values determined by $\\t N = 60$ and (\\ref{Nquad}). For instance, $\\phi_{\\rm req}=2.6 \\mpl$ for $\\sigma=-0.05 \\lp^2$ and $\\phi_{\\rm req}=2.5 \\mpl$ for $\\sigma=-0.10 \\lp^2$). However, $\\Delta\\phi_{\\rm cl}$ on its own is not enough for sufficient inflation. At the same time, the number of e-folds is very sensitive to deviations in $\\phi_{\\rm max}$ near the required values. In other words, even small (positive) variations of $\\phi_i$ and/or $\\Delta \\phi_{\\rm eff}$, such as of order of $0.1 \\mpl$, would lead to sufficient $\\t N$ and should be seriously taken into account.  We see from Fig. \\ref{phiSc_num} that a negative coupling constant indeed gives a greater change of the inflaton, and more negative values are preferable for successful inflation. This fact can be also understood from the Jordanian perspective. The non-minimally coupled action (\\ref{Action}) can be thought of as a standard one with a field-dependent gravitational constant (\\ref{kappaEff}). Clearly, a negative $\\sigma$ results in `weaker' gravitational coupling, and it is not surprising that the inflaton would reach to a greater value (yielding a greater $N$), than for $\\sigma=0$. Moreover, if $\\sigma<0$, the second term of the righthand side of the relation between the number of e-folds in the Einstein and Jordan frames (\\ref{Nrel}) is of order one and can be neglected. Thus $\\t N \\approx N$ and is greater than in the case of minimal coupling. Similar considerations work for a positive coupling as well and imply a smaller number of e-folds.  We should now clarify the seeming discrepancy with the result of Futamase and Maeda \\cite{Conf_NMC}. As we have already mentioned, in that paper, the authors argued that, in order to allow successful inflation, the coupling constant $\\xi$ had to be fine-tuned within a very narrow interval: $\\xi < 10^{-12}$ for $\\xi>0$ and $|\\xi|<10^{-3}$ for $\\xi<0$. These restrictions on $\\xi$ arose from the heuristic argument based on Linde's chaotic inflationary scenario. The values of $\\phi_i$ are assumed to be randomly distributed from zero up to Planckian energy densities $V(\\phi) \\sim \\mpl^4$. Such potential is attained at $\\phi \\sim 10^6 \\mpl$. The latter must necessarily be less than the critical value $\\phi_{\\rm crit}$ (\\ref{phi_crit}), which implies $\\sigma \\sim 10^{-12} \\lp^2$ and $\\xi \\sim 10^{-14}$.  The above conditions can be relaxed in the framework of LQC. As the initial values of the scalar field appear as vacuum fluctuations, that are amplified during the `climbing' phase, one just needs to restrict the coupling constant so that $\\phi_{\\rm crit}$ merely exceeds $\\phi \\sim 3 \\mpl$ - the maximum inflaton value we are interested in. That yields $\\sigma \\sim 10^{-1} \\lp^2$ and $\\xi \\sim 10^{-3}$. The former is exactly the maximum value of the parameter $\\sigma$ we have considered in the paper.  To summarize, the main characteristic of inflation, the number of e-folds, depends on the initial matter fluctuations and the coupling parameter $\\sigma$. The coupling strength is bounded to be less than $0.1 \\lp^2$ by the condition $\\phi_{\\rm max}<\\phi_{\\rm crit} \\equiv 1/\\sqrt{|\\sigma|}$ for both positive and negative couplings. At the same time, the most negative $\\sigma$ would work best for successful inflation. The restriction on the initial field momentum $\\pi_i$ appears as a requirement for the matter density to remain subcritical and implies $|\\pi_i|<5 j$. With these restrictions, the value of the inflaton will increase by approximately $2.0-2.6 \\mpl$ during the `climbing' (effective and classical) phase. Together with the initial vacuum fluctuations of the scalar field of order of several tenths of $\\mpl$ this would lead to $\\t N \\geq 60$, i.e. sufficient inflation. It should be noted that one cannot get $\\t N$ much greater than 60, which indicates that observations may be sensitive to the mechanism discussed here.  {\\bf Acknowledgements:} We thank P.Singh and K.Vandersloot for valuable discussions."
    },
    "0606/astro-ph0606526_arXiv.txt": {
        "abstract": "First using a sample of 32 GRBs with known redshift (Guidorzi et al. 2005) and then a sample of 551 BATSE GRBs with derived pseudo-redshift (Guidorzi 2005), the time variability/peak luminosity correlation ($V$ vs. $L$), originally found by Reichart et al. (2001) using a sample of 18 GRBs, was tested. For both samples the correlation is still found but less relevant due to a much higher spread of the data. Assuming a straight line in the $\\log{L}$--$\\log{V}$ plane ($\\log{L} = m \\log{V} + b$), as done by Reichart et al., both Guidorzi et al. and Guidorzi found that the line slope for both samples is much lower than that derived by Reichart et al.: $m = 1.3_{-0.4}^{+0.8}$ (Guidorzi et al. 2005), $m = 0.85\\pm 0.02$ (Guidorzi 2005), $m = 3.3^{+1.1}_{-0.9}$ (Reichart et al. 2001). Reichart \\& Nysewander (2005) discuss our results and attribute the different slope to the fact we do not take into account in the fit the variance of the sample (also called slop), and demonstrate that, using the method presented by Reichart (2001), the expanded data set of Guidorzi et al. (2005) in $\\log{L}$--$\\log{V}$ plane is still well described by a line with slope $m = 3.4^{+0.9}_{-0.6}$. Here we compare the results of two methods accounting for the slop of the sample, the method implemented by Reichart (2001) and that by D'Agostini (2005). We demonstrate that the method used by Reichart et al. (2001) to estimate the straight line slope, provides an inconsistent estimate of the parameter when the sample variance is comparable with the interval of values covered by the GRB variability. We also demonstrate that, using the D'Agostini method, the slope of the $\\log{L}$--$\\log{V}$ correlation is still consistent with that derived by us earlier and inconsistent with that derived by Reichart \\& Nysewander (2005). Finally we discuss the implications on the interpretations proposed for the $V-L$ correlation and show that our results are in agreement with the peak energy/variability correlation found by Lloyd-Ronning \\& Ramirez-Ruiz (2002) and the peak energy/peak luminosity correlation (Yonetoku et al. 2004; Ghirlanda et al. 2005). ",
        "introduction": "\\label{s:intro}  Most of our knowledge about the Gamma-Ray Burst (GRB) phenomenon is derived from their spectra and light curve profiles, but it is recognised that other observational probes (e.g., polarisation of the gamma-rays) would give key  information to the solution of the GRB enigma. Among these probes, it is recognised the importance of the erratic time variability of the GRB time profiles. For example, in the GRB internal shock model, not very variable radiation is expected to be produced at radii lower than the photosphere radius in which the shocks remain optically thin to pairs \\citep{Kobayashi02}, while highly variable radiation is expected to be produced in shocks above this radius \\citep{Meszaros02}. Also in the sub-jet model by \\citet{Ioka01}, time variability is expected and its amplitude related to the viewing angle of the burst.  A key objective of the study of GRBs is to establish whether GRBs can be reliably used as standard candles and to determine the optimal relationship between observed and intrinsic properties. The recent discovery of GRBs with bright afterglows at redshifts $z>6$ highlights their power as probes of the high redshift Universe \\citep{Haislip06,Kawai06}, but spectroscopically-confirmed redshifts are only a fraction. In contrast, the characteristics of gamma-ray light curves, which are available for all GRBs, offer a potentially independent estimate of luminosity distance for statistically-significant samples.  An important problem, however, in deriving the intrinsic correlations is the interpretation of scatter in correlations, which may be produced by measurement methods, construction of samples with properties measured by satellites with differing response functions, small samples and different statistical analysis methods and intrinsic physical differences in the GRB population \\citep{Nava06}. As sample sizes slowly increase, addressing these issues remains critical for the correct inference of intrinsic GRB properties and thus their use as cosmological probes.  In this paper, we concentrate on a long-standing empirical relation that initially suggested a possible Cepheid-like correlation between gamma-ray variability and peak luminosity of GRBs \\citep{Reichart01,Fenimore00}.  \\citet{Reichart01} (hereafter R01), using a robust measure $V$ of the GRB variability, for a sample of 13 GRBs with known redshift, found that in the GRB rest frame this measure  is correlated with the GRB peak luminosity $L$. In the $\\log{L}$--$\\log{V}$ plane the correlation was modelled with a linear function $\\log L = m \\log V  + b$ with the slope of the straight line $m = 3.3^{+1.1}_{-0.9}$ and a sample variance along $V$ of $\\sigma_{\\log{V}}=0.18$, both parameters being obtained with the method described by \\citet{Reichart01b} (hereafter called Reichart method). This method was proposed to fit data sets affected by a sample variance in addition to the statistical variance (called \"intrinsic variance\") of each data point. Recently, first \\citet{Guidorzi05a} (hereafter GFM05) and then \\citet{Guidorzi05b} (hereafter G05) tested this correlation first using an extended sample of 32 GRBs with known redshift (GFM05), and then with 551 GRBs detected by {\\em CGRO}/BATSE \\citep{Paciesas99} for which a pseudo-redshift was derived by G05 exploiting the spectral lag-luminosity correlation \\citep{Norris00, Norris02,Band04}. In both cases, the correlation was confirmed and in the $\\log V$--$\\log L$ plane the slope of the straight line was found much lower than that derived by R01 ($m = 1.3_{-0.4}^{+0.8}$ derived by GFM05; $m = 0.85\\pm 0.02$ derived by G05). It was also found that, with the sample variance neglected, the straight line did not provide a good description of the data ($\\chi^2/{\\rm dof} = 1167/30$ and $\\chi^2/{\\rm dof} = 4238/549$ for the samples considered by GFM05 and G05, respectively). Neglecting the sample variance was correctly questioned by \\citet{Reichart05} (hereafter RN05), who however show that, also using the extended data set of GFM05, the $m$ slope is given by $m = 3.4_{-0.6}^{+0.9}$, still in perfect agreement with the original value found by R01 and in strong disagreement with the value found by GFM05. They attribute this disagreement to the fact that GFM05 do not include among the parameters of the fit the sample variance, which they estimate to be $\\sigma_{\\log{V}}=0.20\\pm0.04$. In an unrefereed note, \\citet{Reichart05_comm} takes issue with the most recent paper by G05, who confirms the results previously found by GFM05.  In this paper, we discuss the Reichart method and compare the results with those obtained following the treatment by \\citet{Dagostini05} (hereafter called ``the D'Agostini method''), which deals with the same problem of fitting data points affected by extrinsic scatter in addition the intrinsic uncertainties along both axes. We show with numerical simulations that the Reichart method provides an inconsistent estimate of the $m$ slope in the specific case of the $\\log V$--$\\log L$ data set, we discuss the likely reason for this inconsistency. We show that our original results are substantially confirmed by the D'Agostini method even taking into account the sample variance. Finally we discuss the subsequent implications for the inferred physics of GRB central engines and their relativistic outflows. The usefulness of the so-called variability/peak luminosity correlation is discussed in the broader context of recently discovered correlations between other observed and derived properties of GRBs.    \\section[]{The Reichart method} \\label{s:meth} The Reichart method has as starting point the well known Bayes theorem that, for inference of physical parameters, is widely discussed in several text books (see, e.g., D'Agostini 2003). It  states that the conditional probability that a set of parameters \\mbox{\\boldmath$\\theta$}$=\\theta_1,\\theta_2,...\\theta_n$ takes a particular value $\\mbox{\\boldmath$\\theta$}^0$, given a data set $D$, whose values depend on \\mbox{\\boldmath$\\theta$} (for instance, N independent observations of a quantity $X$), is given by: \\begin{equation} p(\\mbox{\\boldmath$\\theta$} | D I) = \\frac{p(\\mbox{\\boldmath$\\theta$} | I) \\ p(D | \\mbox{\\boldmath$\\theta$} I)}{p(D | I)} \\end{equation} where $I$ is the available prior information, $p(\\mbox{\\boldmath$\\theta$} |I)$ is the probability distribution of \\mbox{\\boldmath$\\theta$} on the basis of the information $I$, and $p(D |\\mbox{\\boldmath$\\theta$} I)$ is the conditional probability of getting the measured data set given the value $\\mbox{\\boldmath$\\theta$}^0$ and the information $I$. The probability distribution  $p(\\mbox{\\boldmath$\\theta$} |I)$ is called prior probability, while the probability distribution $p(D |\\mbox{\\boldmath$\\theta$} I)$  is called likelihood function. The probability $p(D | I)$ is introduced as  normalization factor. In the case that no prior information is known, $p(\\mbox{\\boldmath$\\theta$}|I)$ is generally assumed to be a uniform distribution and the range of the possible values of \\mbox{\\boldmath$\\theta$} are those logically allowed. In this case the posterior probability and the likelihood function are equivalent.  Reichart (2001) concentrates on the derivation of the prior probability $p(\\mbox{\\boldmath$\\theta$}|I)$ in the special case in which for two of the $N$ parameters, $x$ and $y$ (in our specific case $x = \\log V$ and $y = \\log L$), we have a set of $N$ pairs of measurements, from which it appears that the two are correlated with $y = y_c(x; \\theta_m)$, were $\\theta_m$, with $m = 1, 2$, \\ldots, $M$, are M intermediate parameters that describe the curve $y_c(x;\\theta_m)$ (in our case $y = m x + b$, $M = 2$). Correctly the scatter of the parameter values $x$ and $y$ around  the curve is assumed to be due partly to the measurement errors (intrinsic scatter) and partly due to weaker dependences of either parameters $x$ or $y$ on other, yet unmeasured, and even yet unknown, parameters (extrinsic scatter or sample scatter). Both the scatters for $x$ and $y$, are assumed by Reichart to be normally distributed and uncorrelated, with unknown standard deviations $\\sigma_x$ and $\\sigma_y$ of the extrinsic scatter of $x$ and  $y$, respectively. The conditional probability of the  values of the parameters $\\theta_m$, $\\sigma_x$, $\\sigma_y$ given the measured set of data points and their uncertainties is thus derived under simplifying assumptions (see Section 2.2.2 of Reichart 2001), among which that the curve $y = y_c(x; \\theta_m)$ can be approximated by a straight line ($y \\approx y_{t,n} + s_{t,n}(x-x_{t,n})$). The result is the following (eq.~43 of Reichart 2001): \\begin{eqnarray} \\lefteqn{ p(\\theta_m, \\sigma_x, \\sigma_y | x_n, y_n, \\sigma_{x,n}, \\sigma_{y,n}) \\approx \\prod_{n=1}^N \\sqrt{1+s_{t,n}^2} \\ G_n \\big[ y_n, y_{t,n} +{}} \\nonumber\\\\ & & {}+s_{t,n}(x_n - x_{t,n}), \\sqrt{\\sigma_y^2 + \\sigma_{y,n}^2 + s_{t,n}^2 (\\sigma_x^2 + \\sigma_{x,n}^2)} \\big] \\label{e:reichart} \\end{eqnarray} where $\\sigma_{x,n}^2$ and $\\sigma_{y,n}^2$ are the intrinsic variances of the $N$ pairs of data points $x_n$ and $y_n$ ($n = 1, 2$, \\ldots, $N$), respectively; $(x_{t,n}, y_{t,n})$ is the point on the curve $y=y_c(x; \\theta_m)$ which maximises the two-dimensional Gaussian $G_n(x, x_n, \\sqrt{\\sigma_x^2+\\sigma_{x,n}^2}) \\times G_n(y, y_n, \\sqrt{\\sigma_y^2+\\sigma_{y,n}^2})$.  In order to simplify the derivation of the prior probability in this special case, eq.~\\ref{e:reichart} is assumed to be a likelihood function and the maximum likelihood method is applied to constrain the values of the intermediate parameters $\\theta_m$, $\\sigma_x$ and $\\sigma_y$ and the uncertainty in the values of $\\theta_m$.    \\subsection{Application of the Reichart method to the Luminosity--Variability correlation} \\label{s:appl}  We now apply the Reichart method in the specific case of the variability--luminosity correlation test, showing that we are capable of reproducing the results obtained by R01 and RN05. In the case of this test, $N$ pairs of measured values  of $V_n$ and $L_n$ (one per each burst) are available, and, in the eq.~\\ref{e:reichart}, $x_n=\\log{(V_n/\\bar{V})}$, $y_n=\\log{(L_n/\\bar{L})}$ ($\\bar{V}$ and $\\bar{L}$ being the correspondent median values), $\\sigma_{x,n} = \\sigma_{\\log{V_n}}$ and $\\sigma_{y,n}=\\sigma_{\\log{L_n}}$. Having modelled the correlation by a straight line: \\begin{equation} \\displaystyle y(x) \\ = \\ m\\,x \\ + \\ q \\label{eq:PL} \\end{equation} from eq.~\\ref{e:reichart}, it is possible to derive the log--likelihood function: \\begin{eqnarray} \\displaystyle \\log{p(m,q,\\sigma_x,\\sigma_y|\\{x_i,y_i,\\sigma_{x,i},\\sigma_{y,i}\\})} =~~~~~~~~~~~~~~~~~~~~~~~\\nonumber\\\\ ~~~~\\frac{1}{2}\\,\\sum_{i=1}^N \\Big[\\log{\\Big(\\frac{1+m^2}{2 \\pi (\\sigma_y^2 + m^2\\,\\sigma_{x}^2 + \\sigma_{y,i}^2 + m^2\\,\\sigma_{x,i}^2)}\\Big)}\\quad +\\nonumber\\\\ ~~~~-\\quad\\frac{(y_i - m\\,x_i - q)^2}{\\sigma_y^2 + m^2\\,\\sigma_{x}^2 + \\sigma_{y,i}^2 + m^2\\,\\sigma_{x,i}^2}\\Big] \\label{eq:prior} \\end{eqnarray} in which, without loss of generality, it is possible to assume either $\\sigma_x=0$, or, alternatively, $\\sigma_y=0$ (both appear only in the term $\\sigma_y^2+m^2\\,\\sigma_x^2$).  Assuming $\\sigma_x=0$ in eq.~\\ref{eq:prior}, with the maximum likelihood method the free parameters $m$, $q$, and $\\sigma_y$ can be derived. Alternatively, if one assumes $\\sigma_y=0$, the value for $\\sigma_x$ can then be determined from $\\sigma_x=\\sigma_y/m$, where $\\sigma_y$ is the value obtained in the previous case under the assumption of null $\\sigma_x$.  For our extended sample of 32 GRBs with known redshift (GFM05), we find that at 90\\% confidence level, the best-fitting parameter values are the following: $m= 3.8_{-1.1}^{+2.8}$, $q=0.07_{-0.33}^{+0.32}$, $\\sigma_y = \\sigma_{\\log L} = 0.93_{-0.29}^{+0.77}$ (alternatively, $\\sigma_x = \\sigma_{\\log V} \\approx 0.24$) (solid lines in Fig.~\\ref{f:z}). As can be seen, all these values are in excellent agreement with those derived by RN05 (they report the 1$\\sigma$ uncertainty). Similarly, for the sample of 551 BATSE GRBs with pseudo-redshift (G05), we find $m=3.5_{-0.4}^{+0.6}$, $q=0.15_{-0.9}^{+0.9}$, $\\sigma_y=1.21_{-0.15}^{+0.19}$, or $\\sigma_x \\approx 0.35$ (solid lines in Fig.~\\ref{f:pz}). All these results are also summarised in Table~\\ref{tab:fit_results}. \\begin{figure} \\includegraphics[width=8.5cm]{guidorzi_fit_VL_f1.eps} \\caption{Variability vs. Peak Luminosity for the 32 GRBs with known redshift considered by GFM05. Also shown are the best fit  curves and 1$\\sigma$ regions. {\\it Solid lines}: best fit results with the Reichart method. {\\it Dashed lines}: results obtained by GFM05. {\\it Dashed-dotted lines}: best fit results with the D'Agostini method.} \\label{f:z} \\end{figure} \\begin{figure} \\includegraphics[width=8.5cm]{guidorzi_fit_VL_f2.eps} \\caption{Variability vs. Peak Luminosity for the 551 BATSE GRBs with pseudo-redshift considered by G05. Also shown are the best fit  curves and 1$\\sigma$ regions. {\\it Solid lines}: best fit results with the Reichart method. {\\it Dashed lines}: results obtained by G05. {\\it Dashed-dotted lines}: best fit results with the D'Agostini method. } \\label{f:pz} \\end{figure} \\begin{table*} \\centering \\caption{Best-fitting parameters obtained with different methods for the GRB samples used by GFM05 and G05. The confidence intervals are at 90\\%.} \\label{tab:fit_results} \\begin{tabular}{llllll} \\hline GRB Set       & Method & $m$ & $q$ & $\\sigma_y$     & $\\sigma_x$    \\\\ &        &     &     & ($\\sigma_x=0$) & ($\\sigma_y=0$)\\\\ \\hline 32 (from GFM05) & Reichart   & $3.8_{-1.1}^{+2.8}$ & $0.07_{-0.33}^{+0.32}$ & $0.93_{-0.29}^{+0.77}$ & $\\sim0.24$\\\\ 32 (from GFM05) & D'Agostini & $1.7_{-0.4}^{+0.4}$ & $0.07_{-0.19}^{+0.18}$ & $0.58_{-0.12}^{+0.15}$ & $\\sim0.34$\\\\ 32 (from GFM05) & {\\tt fitexy}  & $1.9\\pm0.1^{{\\rm (a)}}$ & $0.14\\pm0.02^{{\\rm (a)}}$ & -- & --\\\\ 551 (from G05)   & Reichart   & $3.5_{-0.4}^{+0.6}$ & $0.15_{-0.9}^{+0.9}$ & $1.21_{-0.15}^{+0.19}$ & $\\sim0.35$\\\\ 551 (from G05)   & D'Agostini & $0.88_{-0.13}^{+0.12}$ & $0.01_{-0.03}^{+0.03}$ & $0.65_{-0.04}^{+0.04}$ & $\\sim0.74$\\\\ 551 (from G05)   & {\\tt fitexy}  & $1.37\\pm0.02^{{\\rm (a)}}$ & $0.09\\pm0.01^{{\\rm (a)}}$ & -- & --\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\rm (a)}$] $1\\sigma$ confidence interval. \\end{list} \\end{table*}   \\section[]{An unbiased method} \\label{s:dago_meth} \\citet{Dagostini05} addressed the same problem described in the previous section, i.e. how to perform a linear fit between two data sets with errors on both axes and with an extra variance. Similarly to the Reichart method, the D'Agostini method is based on the parametric inference typical of the Bayesian approach. However, the resulting log-likelihood derived by \\citet{Dagostini05} (see eq. 35 and 52 therein) differs from that by Reichart, reported in \\ref{e:reichart}, for just one term: in the D'Agostini likelihood function the term $(1+m^2)$ is just 1. For a detailed description of the D'Agostini method and of its derivation we address the reader to the original paper. We will demonstrate that, unlike the Reichart method, the likelihood function derived by D'Agostini provides unbiased estimates of the unknown parameters.  According to D'Agostini (2005), the log-likelihood function for the case here considered, is considered is given by \\begin{eqnarray} \\displaystyle \\log{p(m,q,\\sigma_x,\\sigma_y|\\{x_i,y_i,\\sigma_{x,i},\\sigma_{y,i}\\})} =~~~~~~~~~~~~~~~~~~~~~~~\\nonumber\\\\ ~~~~\\frac{1}{2}\\,\\sum_{i=1}^N \\Big[\\log{\\Big(\\frac{1}{2 \\pi (\\sigma_y^2 + m^2\\,\\sigma_{x}^2 + \\sigma_{y,i}^2 + m^2\\,\\sigma_{x,i}^2)}\\Big)}\\quad +\\nonumber\\\\ ~~~~-\\quad\\frac{(y_i - m\\,x_i - q)^2}{\\sigma_y^2 + m^2\\,\\sigma_{x}^2 + \\sigma_{y,i}^2 + m^2\\,\\sigma_{x,i}^2}\\Big] \\label{eq:prior_dago} \\end{eqnarray} Using this equation, in the case of the 32 GRBs of GFM05, at 90\\% confidence level we find the results reported in Table~\\ref{tab:fit_results}. In particular we find a slope $m= 1.7\\pm0.4$ against a value $m= 3.8_{-1.1}^{+2.8}$ found with the Reichart method. Likewise, using the sample of 551 GRBs with pseudo-redshift of G05, we find the best fit parameter values reported in Table~\\ref{tab:fit_results}, $m=0.88_{-0.13}^{+0.12}$. The best-fitting power-laws obtained with the D'Agostini method and the Reichart method are compared in Fig.s~\\ref{f:z} and \\ref{f:pz} for the two data sets, respectively.  In Table~\\ref{tab:fit_results} we report the also the best-fitting parameters obtained using the Least square fit in the case of data in the case of data affected by errors on both axes (``{\\tt fitexy}'' tool, \\citet{Press93}), but not with no extra variance. We report these results to examine how the best-fitting parameters are affected when the sample variance is taken into account.  As can be seen from Table~\\ref{tab:fit_results}, the Reichart method and D'Agostini method give different best fit parameter values, especially the value of the slope $m$. The Reichart method yields steeper slopes even that those found with the D'Agostini method, whose results are consistent with those originally obtained by GFM05 and G05 ($m = 1.3_{-0.4}^{+0.8}$, GFM05; $m = 0.85\\pm 0.02$, G05), even though the fit, in the latter case, was found unacceptable (very high $\\chi_r^2$). Also with the '{\\tt fitexy}' algorithm \\citep{Press93}, the slopes obtained are much shallower than the correspondent obtained with the Reichart method. ",
        "conclusions": "\\label{s:conc} We applied both Reichart and D'Agostini methods to the samples of 32 GRBs with known redshift and 551 BATSE GRBs with pseudo-redshift considered by GFM05 and G05, respectively. The goal was to estimate the slope $m$ as well as the scatter of the power law describing the correlation between variability and peak luminosity originally presented by \\citet{Reichart01}. Both methods account for an extra variance in addition to the intrinsic affecting each single point. From simulations, we found that when the sample variance $\\sigma_x^2$ is comparable with the total scatter along the same axis, $\\sigma_{x,t}^2$, the Reichart method tends to overestimate $m$, while the D'Agostini still estimates it correctly. When the sample variance is negligible with respect to the total variance, the two methods give consistent results. In the specific case of the $V-L$ correlation, the two variances are comparable: in the case of the 32 GRBs of GFM05, it is $\\sigma_{x,t}=0.32$, $\\sigma_{x}=0.24$ (Reichart) or $\\sigma_{x}=0.34$ (D'Agostini). This explains the discrepancy between the two methods. We showed that the D'Agostini method is a reliable estimator of $m$ in this regime, whereas the Reichart is no more. In particular, the best-fitting value for $m$ obtained with the D'Agostini method is $1.7\\pm0.4$ and $0.88_{-0.13}^{+0.12}$ at 90\\% confidence level for the 32 GRBs of GFM05 and for the 551 GRBs of G05, respectively. These values are significantly smaller than those obtained with the Reichart method, which are consistent with previous results of R01 and RN05. These results hold as far as the definition of variability given by R01 is assumed.  Alternatively, from other definitions of variability based on different kinds of filter used to smooth the light curves, it seems to be possible to obtain a range of values for $m$ from $\\sim$1 to $\\sim$3 \\citep{Fenimore00,Li06}.  One of the possible implications of a smaller value of $m$ than originally found by R01 on the interpretation of the $V-L$ correlation is that, in the jet emission scenario, we would expect a stronger dependence of the Lorentz factor of the expanding shells on the jet opening angle \\citep{Kobayashi02}.  Finally, other more recent and tighter correlations, such as the Amati, Ghirlanda, Liang \\& Zhang relations appear to be more reliable luminosity estimators than the $V-L$ one. In particular, our results of values of $m$ in the interval 1--2 obtained with the D'Agostini method, appear to be consistent with two independent relations: $E_{\\rm p}\\propto V^{\\delta}$ ($0.4<\\delta<1.15$) \\citep{Lloyd02b}, and the equivalent version of the Amati relation with the isotropic-equivalent peak luminosity instead of the isotropic-equivalent total released energy, $E_{\\rm p}\\propto L^{0.5}$. By combining the two, one expects $m=\\delta/0.5$, in agreement with the results reported here as well as those reported by GFM05 and G05.  Finally, the increasing number of GRBs with spectroscopic redshift detected by {\\em Swift} \\citep{Gehrels04} will help extend the range in $V$ and better constrain the power-law fit of the $V-L$ correlation, with the benefit of a homogeneous data set of light curves all detected with the Burst Alert Telescope (BAT) onboard {\\em Swift}. A thorough test of the $V-L$ correlation with BAT data is in progress (Rizzuto et al., in preparation)."
    },
    "0606/astro-ph0606656_arXiv.txt": {
        "abstract": "{ We present an {\\sl XMM-Newton} observation of the eclipsing binary Algol which contains an X-ray dark B8V primary and an X-ray bright K2IV secondary. The observation covered the optical secondary eclipse and captured an X-ray flare that was eclipsed by the B star. The EPIC and RGS spectra of Algol in its quiescent state are described by a two-temperature plasma model. The cool component has a temperature around 6.4$\\times$10$^{6}$ K while that of the hot component ranges from 2 to 4.0$\\times$10$^{7}$ K. Coronal abundances of C, N, O, Ne, Mg, Si and Fe were obtained for each component for both the quiescent and the flare phases, with generally upper limits for S and Ar, and C, N, and O for the hot component. F-tests show that the abundances need not to be different between the cool and the hot component and between the quiescent and the flare phase with the exception of Fe. Whereas the Fe abundance of the cool component remains constant at $\\sim$0.14, the hot component shows an Fe abundance of $\\sim$0.28, which increases to $\\sim$0.44 during the flare. This increase is expected from the chromospheric evaporation model. The absorbing column density $N_H$ of the quiescent emission is 2.5$\\times10^{20}$ cm$^{-2}$, while that of the flare-only emission is significantly lower and consistent with the column density of the interstellar medium. This observation substantiates earlier suggestions of the presence of X-ray absorbing material in the Algol system. ",
        "introduction": "Algol ($\\beta$ Per) is a nearby eclipsing binary with an early main-sequence (B8V) primary and a cool subgiant (K2IV) secondary. This system is only 28.46 pc away (Hipparcos parallax measurements, ESA 1997) and has a period of about 2.87 d and an orbit inclination of $81^{\\circ}$. The radii of the two companion stars ($R_{B}$ and $R_K$) are 3.0 and 3.4 $R_{\\odot}$, respectively, and their separation is 14.14 $R_{\\odot}$ (Hill et al. 1971; Richards 1993). It is a semi-detached system with the K star filling its Roche lobe. The mass transfer occurs in the form of gas streams from the K star to the B star and ends up in an annulus around the B star. (Richards 1993; Richards et al. 1995; Richards 2001; Richards 2004; Retter et al. 2005). Algol was first detected in X-rays by {\\sl SAS} 3 (Schnopper et al. 1976) and confirmed as a strong X-ray emitter by Harnden et al. (1977). It is generally accepted that the K star has an active corona and accounts for most of the X-ray flux of the system, whereas the B star is X-ray dark (White et al. 1980). This was proved for the first time by Chung et al. (2004) by  the detection of Doppler shifts of the spectra caused by orbital motion of the K star. Given an active corona, flares appear typically every day with a duration of several to many hours.  Algol shows elemental abundances which differ from those of the sun (Antunes et al. 1994; Schmitt \\& Ness 2004). Antunes et al. (1994) quote abundances of Fe, O, Mg, Si, S, Ar and Ca which are lower than the solar photospheric values by a factor of 2-3, and N to be less than 0.1. With the exception of N these results have been generally confirmed by {\\sl Chandra} observations (Schmitt \\& Ness 2004). Schmitt \\& Ness (2002) studied a sample of late type stars including Algol and they found an enhancement of N. Drake (2003) also suggested that the N abundance is enhanced by a factor of 3, and C is depleted  by a factor of 10 (both relative to HR1009, whose C/N abundance is consistent with that of the solar photosphere). These C and N abundances would lead to the conclusion that the K star has lost at least half of its initial mass.  That flares have different elemental abundances has been concluded from previous missions, such as {\\sl GINGA} (Stern et al. 1992), {\\sl ROSAT} (Ottmann \\& Schmitt 1996), {\\sl BeppoSAX} (Favata \\& Schmitt 1999), {\\sl Chandra} {Nordon \\& Behar 2007} and {\\sl XMM-Newton} {Nordon \\& Behar 2008}. They suggest that the elemental abundances of Algol's flaring region resemble more the photospheric than the quiescent coronal abundances, which might be a consequence of chromospheric evaporation (G\\\"{u}del et al. 1999; Favata \\& Micela 2003 and references therein). At the beginning of the event, fresh chromospheric material is evaporated in the flaring loop(s), and enhances the corona elemental abundances; once the material has been transported to the flaring corona structure, the fractionation mechanism, which is responsible for the lower abundance, would start operating, bringing the corona abundance back to its quiescent value.  The absorbing column density ($N_H$) inferred  from X-ray observations is usually higher than the column density of the interstellar medium (ISM) between Algol and the observer. Welsh et al. (1990) gave an upper limit of the ISM column density of 2.5$\\times10^{18}$ cm$^{-2}$ using the ISM Na{\\small I} D line absorption of the B8 primary. Stern et al. (1995) got a similar result from the He/H ratio using EUVE observations. The $N_H$ derived from the {\\sl ROSAT} PSPC observation of Algol is about 5$\\times10^{19}$ cm$^{-2}$ (Ottmann 1994). Favata \\& Schmitt (1999), using {\\sl BeppoSAX}, reported that $N_H$ increased up to $>10^{21}$ cm$^{-2}$ during the flare rise and decreased during the decay, which may possibly be associated with moving, cool material in the line of sight, e.g. a major coronal mass ejection.   In this paper we present a time resolved spectroscopy of Algol, using {\\sl XMM- Newton} European Photon Imaging Camera (EPIC) and Reflection Grating Spectrometer (RGS) observations of Algol. Compared with the instruments used in earlier studies the EPIC cameras are the first ones to have a combination of moderate spectral resolution, wide energy coverage, and high sensitivity, while RGS provides us with unparalleled sensitivity for high resolution spectroscopy in the soft X-ray band. These advantages permit us to perform  a more detailed diagnosis of the plasma properties and their evolution as far as the corona of the K type secondary is concerned. The {\\sl XMM-Newton} observation covered the secondary eclipse and an X-ray flare was detected. Schmitt et al. (2003) presented a detailed study of the geometry of the flare, and we analyze the X-ray spectra of Algol in both the quiescent and flaring states. We can therefore study the mechanism of the X-ray emission of Algol in both the quiescent and the flaring states. In $\\S$ 2, we describe the observation and data reduction. The spectral properties of the overall and flare-only emissions are given in $\\S$ 3. The proposed local cool absorbing gas structures are discussed in $\\S$ 4 and a summary is provided in $\\S$ 5. ",
        "conclusions": "We have analyzed an {\\sl XMM-Newton} observation of the eclipsing binary Algol, in which the interval of the secondary optical minimum and some preceding section in the orbit were covered. During the eclipse of the X-ray bright K star a flare occurred, so that we can study both the quiescent corona and an eclipsed flare. We joined the data of all four X-ray instruments on board of {\\sl XMM-Newton}, which give spectral data at high resolution with the RGSs from 0.3 - 2 keV and medium resolution data with the EPIC cameras from 0.3 - 10 keV.  Satisfactory fits have been obtained over the entire energy band using the two-temperature VMEKAL model for the overall spectrum. We cannot rule out that the elemental abundances of the spectrum's low temperature and high temperature components are identical, but the abundance of Fe is clearly different at the 4$\\sigma$ level for the quiescent phase and significantly more than 5$\\sigma$  for the flare phase, indicating that the high temperature  and low temperature plasma components are not well mixed. We also observed a significant Fe abundance increase in the hot component by a factor of $\\sim$ 1.6 during the flare, which supports the idea of chromospheric evaporation.  The fits to the $N_H$ column density reveal values around 2.5$\\times10^{20}$ cm$^{-2}$, for the quiescent corona which is far in excess of the interstellar column density. On the other hand, the fits to the flare-only emission are consistent with no absorption exceeding the interstellar value. We propose that the line of sight column density across Algol is not uniform. It is sufficiently high towards the polar regions, and reduced towards the equatorial regions.   \\normalem"
    },
    "0606/astro-ph0606183_arXiv.txt": {
        "abstract": "{A cosmological model with total density close to critical (and flat geometry), dominated by dark matter and dark energy of unknown nature, and consistent with the basic predictions of the inflationary scenario is a very good fit to a variety of cosmological probes: the anisotropy of the CMB, the large scale distribution of matter, the luminosity distance of high-redshift type Ia supernovae and so on. These high-quality data have established a new standard of precision in the determination of cosmological parameters.}  \\FullConference{CMB and Physics of the Early universe\\\\ 20-22 April 2006\\\\ Ischia, Italy}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "Cosmology has developed into a fully mature science. The parameters of the big bang model are now known with great accuracy, and the constraints are expected to get tighter in the future. Inflation has not been falsified, and its main predictions are strikingly consistent with observations. The results obtained using completely different cosmological probes are in remarkable agreement among themselves, as well as with theoretical predictions.  Nonetheless, many fundamental questions are still open \\cite{open questions}. The pace of experimental and theoretical progress, however, does not seem to be close to a halt."
    },
    "0606/astro-ph0606460_arXiv.txt": {
        "abstract": "We present a comparison of barstrength $Q_b$ and circumnuclear dust morphology for 75 galaxies in order to investigate how bars affect the centers of galaxies.  We trace the circumnuclear dust morphology and amount of dust structure with structure maps generated from visible-wavelength HST data, finding that tightly wound nuclear dust spirals are primarily found in weakly barred galaxies.  While strongly barred galaxies sometimes exhibit grand design structure within the central 10\\% of $D_{25}$, this structure rarely extends to within $\\sim 10$~pc of the galaxy nucleus.  In some galaxies, these spiral arms terminate at a circumnuclear starburst ring.  Galaxies with circumnuclear rings are generally more strongly barred than galaxies lacking rings.  Within these rings, the dust structure is fairly smooth and usually in the form of a loosely wound spiral.  These data demonstrate that multiple nuclear morphologies are possible in the most strongly barred galaxies: chaotic central dust structure inconsistent with a coherent nuclear spiral, a grand design spiral that loses coherence before reaching the nucleus, or a grand design spiral that ends in a circumnuclear ring.  These observations may indicate that not all strong bars are equally efficient at fueling material to the centers of their host galaxies.  Finally, we investigate the longstanding hypothesis that SB(s) galaxies have weak bars and SB(r) galaxies have strong bars, finding the opposite to be the case: namely, SB(r) galaxies are less strongly barred and have less dust structure than SB(s) galaxies.  In general, more strongly barred galaxies tend to have higher nuclear dust contrast. ",
        "introduction": "\\label{sec:intro} Most spiral galaxies have large-scale bars, and any complete picture of galaxies must include the impact of bars on their evolution.  Bars are important because they are an obvious mechanism for funneling gas and dust to the centers of galaxies: gas clouds orbiting in the disk lose angular momentum as they encounter the bar, thereby sinking towards the galaxy's center \\citep{binney87, athanassoula92, piner95, regan97, maciejewski02}.  Bars are therefore expected to play a major role in circumnuclear star formation, bulge growth, and the fueling of the central, supermassive black hole.  Bar-driven radial gas inflow should drive large quantities of gas and dust into the central regions of galaxies, thereby increasing the central gas concentration of barred galaxies relative to unbarred galaxies.  Observations by \\citet{sakamoto99} and \\citet{sheth05} support this view, finding higher concentrations of molecular gas in the central kiloparsec of barred galaxies than in unbarred galaxies. As higher gas density is empirically correlated with higher star formation rates \\citep[e.g.,][]{kennicutt98}, there should be a correlation between bar-driven inflow and central star formation.  Observational evidence for enhanced nuclear or circumnuclear star formation in barred galaxies provides broad support for this picture \\citep{ho97c, maoz01, knapen06}.  It has long been speculated that this same radial gas inflow could provide active galactic nuclei (AGN) with fuel \\citep{simkin80, schwarz81}.  There is so far no evidence, though, that bars (large-scale or nuclear) are the primary fuel source for AGN: not all active galaxies are barred, not all barred galaxies are active, and, in fact, well-matched samples of active and inactive galaxies have the same bar fraction \\citep{ho97c, mulchaey97a, laine02}.  \\citeauthor{ho97c} speculate that this is because radially-transported gas is prevented from reaching the nucleus.  If bars are indeed fueling galaxies' circumnuclear regions, then this should be reflected in the morphology of the cold circumnuclear interstellar medium (ISM).  \\citet{martini03a} establish a nuclear dust morphology classification system with four spiral classes (grand design, tightly wound, loosely wound, and chaotic spiral) and two non-spiral classes (chaotic and no structure) to quantify differences between various galaxy types.  \\citet{martini03b} find no differences in the nuclear dust structure of active and inactive galaxies.  This study does find evidence that nuclear grand design spirals are primarily found in barred galaxies and that these grand design spirals connect to large scale dust lanes in barred galaxies.  Likewise, \\citeauthor{martini03b} primarily find tightly wound nuclear spirals in unbarred galaxies.  All of these studies of the role of bars in fueling central star formation or black holes simply compare ``barred'' versus ``unbarred'' galaxies.  This discretization glosses over, and over simplifies, the long-known continuum of barstrengths \\citep{devaucouleurs59}. One early way of quantifying the strength of a bar is the deprojected bar ellipticity $\\epsilon_b$ \\citep{martin95}.  The disadvantage of using the bar ellipticity is that defining the bar---i.e., its exact dimensions---is a somewhat subjective process \\citep{buta01}. \\citet{abraham00} introduce the parameter $f_{\\mbox{\\scriptsize bar}}$, which is a function of the bar axis ratio, and therefore susceptible to the same systematics as $\\epsilon_b$.  \\citeauthor{buta01}, by expanding on the method of \\citet{combes81}, circumvent this problem by using a force ratio, which they call $Q_b$.  Assuming that the light traces the underlying mass distribution, they use near-infrared images to calculate a map of a galaxy's potential.  From these potential maps, and with some assumptions about the characteristic scale-height, they then calculate the tangential force $F_T$ and mean (axisymmetric) radial force $\\langle F_R\\rangle$.  They define the ratio map $Q_T$ as \\begin{equation}\\label{eqn:Qt} Q_T(i,j) = \\frac{F_T(i,j)}{\\langle F_R(i,j)\\rangle}. \\end{equation} This ratio map has the property that in each of four quadrants $Q_T$ reaches a local extremum.  Letting $Q_{T}^{\\mbox{\\scriptsize max},\\,k}$ be the absolute value of such an extremum in the $k$th quadrant, the barstrength $Q_b$ is defined as \\begin{equation}\\label{eqn:Qb} Q_b = \\frac{1}{4} \\sum\\limits^{4}_{k=1}  Q_{T}^{\\mbox{\\scriptsize max},\\,k}. \\end{equation} The cited uncertainties on $Q_b$ are a measure of how much $Q_{T}^{\\mbox{\\scriptsize max}}$ varies by quadrant.  From a comparison with the three barstrength classes defined by \\citeauthor{devaucouleurs59} (unbarred SA, weakly barred SAB, and strongly barred SB), \\citeauthor{buta01} find a reasonable correlation between the RC3 bar classification and $Q_b$, although with substantial scatter. In particular, SA galaxies have $Q_b \\lesssim 0.1$, SAB galaxies have $0.05 \\lesssim Q_b \\lesssim 0.2$, and SB galaxies have $Q_b \\gtrsim 0.15$.  The most strongly barred galaxies have $Q_b \\approx 0.6$ \\citep{buta01}. While we adopt $Q_b$ as the best available measure of barstrength, we note that \\citet{athanassoula92} shows that both the quadrupole moment (strength) and the pattern speed play an important role in how effectively the bar drives mass towards the center of the galaxy. Pattern speed, however, is much more difficult to measure than parameters that can be calculated from photometric data.   \\citet{martini04} used $Q_b$ and the nuclear classifications of \\citet{martini03a} to compare the barstrengths and nuclear dust morphologies of 48 galaxies with archival Hubble Space Telescope (HST) data.  He found that grand design nuclear spirals are primarily found in strongly barred galaxies, while axisymmetric tightly wound nuclear spirals are exclusively found in galaxies with $Q_b < 0.1$.  With a larger sample and reconsideration of the circumnuclear dust morphology classification at the smallest scales, we examine how nuclear dust structure varies with barstrength.  We identify 75 galaxies, described in \\S\\ref{sec:data}, with a measured barstrength $Q_b$ from the literature.  To analyze the dust morphology, we create ``structure maps'' from archival HST data using the technique of \\citet{pogge02} as discussed in \\S\\ref{sec:smap}.  We then classify the galaxies according to a refined version of the classification scheme proposed in \\citet{martini03a}, as described in \\S\\ref{sec:nuc}.  The main difference between the classification scheme used here and that of \\cite{martini03a} is a stronger focus on the central-most regions of the galaxy.  We discuss our results in \\S\\ref{sec:results}. ",
        "conclusions": "\\label{sec:conc} We present a study of the circumnuclear dust morphology for a sample of 75 galaxies with archival HST data and measured barstrength $Q_b$.  This $Q_b$ is a measure of the maximal force ratio due to the presence of a bar in a galaxy, and thus is arguably superior to the rudimentary bar axis ratio \\citep{buta01}.  We use the structure map technique of \\citet{pogge02} to enhance the visibility of the galaxies' dust content and to classify the circumnuclear dust morphology within the central 5\\% of $D_{25}$, according to a refined version of the classification scheme proposed by \\citet{martini03a}.  We also introduce the structure map rms $\\sigma_{\\mbox{\\scriptsize sm}}$ within the central regions of the galaxy as a quantitative measure of the amount of nuclear dust structure.  A comparison of the morphological classifications and measured barstrengths reveals that tightly wound nuclear dust spirals (with pitch angles less than 10\\dg) are preferentially found in galaxies with lower $Q_b$.  No other nuclear class exhibits a significant correlation with barstrength.  Previous observations found that grand design nuclear dust spirals are hosted exclusively by strongly barred galaxies.  While we do see grand design structure in many barred galaxies, this grand design structure does not extend all the way into the unresolved nucleus ($\\sim 10$~pc), although it is often present at larger scales.  Earlier studies did not strictly require that the grand design structure extend into the nucleus.  Taking this into account, we identify two distinct types of circumnuclear grand design spirals.  Small grand design (SGD) spirals are nuclear dust spirals in which the two symmetric spiral arms are coherent from 1\\% of $D_{25}$ (typically a few hundred parsecs) to the central tens of parsecs of the galaxy (where tracing the structure becomes resolution-limited).  Large grand design (LGD) nuclear spirals, on the other hand, show two prominent symmetric arms within 10\\% of $D_{25}$ (typically on the scale of a few kiloparsecs); these arms do not necessarily extend to the center of the galaxy.  In fact, these two types of grand design structure are nearly disjoint: the nuclear spiral arms in only two of the twenty LGD galaxies in our sample extend into the unresolved center of the galaxy (i.e., also displayed SGD structure).  The LGD spirals are found in systematically more strongly barred host galaxies.  This strongly confirms previous indications from much smaller samples and demonstrates that the dust lanes along the leading edges of large-scale spirals do not generally extend all the way into the nuclear region, but instead lose coherence at the scale of several hundred parsecs.  While SGD spirals are not found in galaxies with a significantly different barstrength than typical galaxies, they are found in galaxies with significantly less dust structure in the central regions than LGD galaxies.  The reduced dust structure may reflect a requirement for the formation of SGD morphology, or simply a requirement for its detection.  The LGD spiral arms may not maintain coherence to the nucleus because mass inflow due to the presence of a large-scale bar prompts star formation, which can disrupt a nuclear grand design spiral.  In addition, forty percent of the LGD spirals do not extend into the nuclear region because there is a circumnuclear starburst ring.  We find that all of the galaxies with circumnuclear starburst rings have LGD structure and are more strongly barred than other galaxies.  Within the rings, three fourths of the galaxies with circumnuclear rings have coherent loosely wound spirals within the ring, while the others have less coherent chaotic spirals.  We also find that among SB galaxies, SB(s) galaxies have more dust structure and are more strongly barred than SB(r) galaxies.  This is partially at odds with the prevailing view in the literature, which is that SB(r) galaxies should be more strongly barred, although there is consensus that SB(r) galaxies have less central dust \\citep[e.g.][]{kormendy04}.  \\citet{sanders80} suggest that differences in the bar pattern speed may also explain the large-scale morphological differences between SB(s) and SB(r) galaxies; as our results indicate that SB(s) galaxies are more strongly barred, pattern speed may be the more relevant parameter.  Overall, there is agreement in the literature that more strongly barred galaxies should---and do---have more dust and gas in their centers \\citep{kormendy04}.  We find that for the most strongly barred galaxies, there are several possible morphologies the circumnuclear dust can take. This may indicate that not all bars of a given strength $Q_b$ funnel material toward the centers of galaxies with equal efficiency, potentially due to the effects of pattern speed on bar efficiency, or simply the fact that $Q_b$ is a one-parameter description of the bar. In the most strongly barred galaxies, there can be an LGD spiral whose arms do not extend to the galaxy nucleus but instead lose coherence. Some LGD spirals end in circumnuclear starburst rings.  In the absence of these structures, however, the nuclear dust in the most strongly barred galaxies tends to be fairly chaotic, potentially hosting star formation.  It would be interesting to investigate how circumnuclear dust morphology and dust structure $\\sigma_{\\mbox{\\scriptsize sm}}$ varies with nuclear star formation rates in strongly barred galaxies."
    },
    "0606/astro-ph0606306_arXiv.txt": {
        "abstract": "{We study the cosmological evolution of the universe when quintessence is modeled within supergravity, supersymmetry is broken in a hidden sector, and we also include observable matter in a third independent sector. We find that the presence of hidden sector supersymmetry breaking leads to modifications of the quintessence potential. We focus on the coupling of the SUGRA quintessence model to the MSSM and investigate two possibilities.  First one can preserve the form of the SUGRA potential provided the hidden sector dynamics is tuned. The currently available limits on the violations of the equivalence principle imply a universal bound on the vacuum expectation value of the quintessence field now, $\\kappa ^{1/2}Q\\ll 1$. On the other hand, the hidden sector fields may be stabilised leading to a minimum of the quintessence potential where the quintessence field acquires a mass of the order of the gravitino mass, large enough to circumvent possible gravitational problems. However, the cosmological evolution of the quintessence field is affected by the presence of the minimum of the potential. The quintessence field settles down at the bottom of the potential very early in the history of the universe. Both at the background and the perturbation levels, the subsequent effect of the quintessence field is undistinguishable from a pure cosmological constant.}  \\begin{document} ",
        "introduction": "There is a host of observational evidence in favor of the existence of a non-zero vacuum energy density of the Universe driving the acceleration of the expansion of the Universe~\\cite{LSS,IA,CMB}. The simplest explanation for this new era in the history of the Universe is the presence of a cosmological constant of extraordinarily small value, some 120 orders of magnitude lower than the Planck scale. Such a small value is particularly difficult to accommodate when dynamical effects such as the Quantum Chromo-Dynamics (QCD) and electroweak phase transitions or even Grand Unified Theory (GUT) scale physics are taken into account. This has prompted the possibility of using extra dimensional models such as self--tuning scenarios~\\cite{kachru} or brane induced gravity models~\\cite{deffayet}. Unfortunately these alternatives have drawbacks such as hidden fine--tunings~\\cite{Lalak}. Within four dimensional physics, there is an experimental way of discovering whether the vacuum energy is a true constant of nature or the result of more complicated dynamical effects. Indeed very active experimental programs are dedicated to the analysis of the so-called equation of state of the dark energy sector (the ratio between the pressure and the energy density). If the equation of state differs from $-1$ (and is greater than $-1$, otherwise see for instance Ref.~\\cite{MSU}), then a plausible candidate for dark energy is quintessence~\\cite{RP,quint,PB}, \\ie the dynamics of a scalar field rolling down a runaway potential. Of course, quintessence only accounts for the small and non-vanishing vacuum energy, it has nothing to say about the large cancellation of the overall cosmological constant.  \\par  One of the most stringent requirements imposed on quintessence models is the existence of attractors, \\ie long time stable solutions of the equations of motion~\\cite{quint}. Indeed the presence of an attractor implies an insensitivity to initial conditions of the quintessence field for the vacuum energy now.  For a large class of potentials, attractors leading to vacuum energy dominance exist provided their large field behavior is of the inverse power law type. Such potentials are known under the name of Ratra--Peebles potentials~\\cite{RP}. In these cases, the quintessence field reaches an attractor, only to leave it when dominating the energy content of the universe. This happens when the field is of the order of the Planck scale.  \\par  This has drastic consequences on quintessence model building. Indeed, it requires a natural framework within which Planck scale physics is taking into account. Supergravity is a promising field theoretical arena where both particle physics and Planck scale physics can be described~\\cite{Nilles}. Models of quintessence in supergravity have been constructed~\\cite{BM1,BM2,BMR1,BMR2} leading to interesting phenomenological consequences such as low values of the equation of state. In particular, the simplest model of quintessence in supergravity, often dubbed the SUGRA model in the literature, leads to the following potential \\begin{equation} \\label{sugra} V_{\\rm quint}(Q)={\\rm e}^{\\kappa Q^2/2+\\kappa \\xi ^2}\\frac{M^{4+\\alpha }}{Q^{\\alpha}}\\, , \\end{equation} with $\\kappa \\equiv 8\\pi/\\mpl ^2$ and $M^{4+\\alpha }=\\lambda ^2\\xi ^4 m_{\\rm c}^{\\alpha }2^{\\alpha/2}$ and where, in this equation, $Q$ is canonically normalized. The quantity $\\alpha $ is a free positive index and $\\lambda $ is a dimensionless coupling constant and, in order to avoid any fine-tuning, we will always consider that $\\lambda \\sim 1$. $m_{\\rm c}$ is the cut-off scale of the effective theory used in order to derive the SUGRA potential. Typically $m_{\\rm c}$ can be thought as the GUT scale but we will see that the Planck scale is also possible (and, sometimes, necessary). Finally, $\\xi $ is a vacuum expectation value (vev) of some other field present in the quintessence sector, see below for more details. As a specific example, $\\xi$ can be realized as a Fayet-Iloupoulous term arising from the Green--Schwarz anomaly cancellation mechanism~\\cite{BM1,BM2}. The main feature of the above potential is that supergravity corrections have been exponentiated and appear in the prefactor. Phenomenologically, this potential has the nice feature that the equation of state $\\omega \\equiv p_Q/\\rho _Q$ can be closer to $-1$ than with the Ratra--Peebles potential when the field approaches its present value $\\kappa^{1/2} Q_{\\rm now}\\approx 1$. Moreover, a small value for $M$ can be avoided. Indeed, since the vev of the quintessence field is now of the order of the Planck mass, requiring that the quintessence energy density be of the order of the critical energy density $\\rho_{\\rm cri}\\sim 10^{-122}\\mpl ^4$ implies that \\begin{equation} \\frac{M}{\\mpl}\\sim 10^{-122/(4+\\alpha )}\\, , \\end{equation} and, therefore, $M$ can be a large scale (by particle physics standard) for very reasonable values of the index $\\alpha $. For instance, it is above the TeV scale for $\\alpha \\gta 4$. This mechanism is reminiscent of a ``see-saw'' mechanism where a very small scale (the cosmological constant scale) is explained in terms of a large one (the scale $M$) and a very large one (the Planck scale $m_{_{\\rm Pl}}$).  Moreover, the scale $\\xi $ can have acceptable values. From the expression of the scale $M$, one obtains \\begin{equation} \\label{xisugra} \\frac{\\xi }{\\mpl }\\sim \\left(\\frac{\\rho _{\\rm cri}}{\\mpl }\\right)^{1/4} \\left(\\frac{\\mpl}{m_{\\rm c}}\\right)^{\\alpha /4}\\, , \\end{equation} and for $\\alpha \\gta 11$ and a cut-off $m_{\\rm c}$ of the order of the GUT scale, $\\xi $ is above the TeV scale\\cite{BM1}. However, considering $m_{\\rm c}<m_{_{\\rm Pl}}$ can also be viewed as problematic since, as already mentioned, the vev of the quintessence field tends to be of the order of the Planck mass. In this case, it is difficult to control the shape of the K\\\"ahler potential, see Eq.~(\\ref{quintsector}). Facing this issue, a natural choice could be $m_{\\rm c}=m_{_{\\rm Pl}}$. Then the above equation indicates that $\\xi $ needs to be fine-tuned as the same level as the cosmological constant even if the model remains different since the equation of state is redshift dependent. Moreover, in this situation, the choice of the K\\\"ahler potential strongly influences the shape of the scalar potential~\\cite{BM2} since the SUGRA corrections are of order one. Hence, it is no longer possible to see the K\\\"ahler potential of Eq.~(\\ref{quintsector}) as a Taylor expansion but it should rather be considered as a specific choice. Another route is to argue that, in these circumstances, the quintessence energy density can easily remain less than $m_{\\rm c}^4$, indicating that the theory can still be meaningful if the cut-off is interpreted as the cut-off to the energy density scales and not to the vev's of the fields. In any case, it is clear that there are some fine-tuning problems at this level even if it is arguably less acute than in the cosmological constant case.  \\par  However, it is clear that the quintessence sector cannot be considered as disconnected from the particle physics standard model (or its extensions) and should be embedded in a more general structure. In Ref.~\\cite{BMpart}, we have investigated the coupling between the quintessence sector and a hidden sector where supersymmetry is broken. A general formalism to calculate the corresponding implications was presented. On very general grounds, it was shown that, as a consequence of this coupling, the shape of the quintessence potential is changed and that the particle masses become dependent on the quintessence field which implies the presence of a fifth force, a violation of the equivalence principle and possibly, depending on the complexity of the model, a variation of the gauge couplings and of the proton to electron mass ratio.  \\par  The main goal of the present article is to apply this general formalism to a concrete case, namely the SUGRA one, where the quintessence potential is described by Eq.~(\\ref{sugra}). We find that the coupling between the quintessence, observable and hidden sectors has tremendous consequences on the dynamics of the quintessence field. The first one is that the shape of the quintessence potential is modified.  When the hidden sector fields have been stabilized, the potential acquires a minimum located at a very small value of the field (in comparison with the Planck mass). As a result, the model becomes equivalent to a pure cosmological constant as the quintessence field settles at the minimum of a potential before Big Bang Nucleosynthesis (BBN). The mass of the field is also changed and becomes equal to the gravitino mass. The above conclusion is also true at the perturbative level as no growing mode is present despite the smallness of the Jeans length.  On top of this, the energy scales in the quintessence sector have to be fine--tuned at a highest level than the cosmological constant itself. This makes the whole scenario very unappealing.  On the other hand, preserving the shape of the SUGRA potential requires a fine-tuning of the hidden sector dynamics. If this is done, one may wonder whether the model is still phenomenologically acceptable. We show that, in this case, the scenario tends to be ruled out by local tests of gravity rather than by cosmological considerations.  \\par  Our results have been obtained using the SUGRA model~\\cite{BM1} of quintessence. Effectively the only crucial ingredient is the fact that the potential in the quintessence sector reduces to the Ratra--Peebles form for small values of the quintessence field.  Within this framework, \\ie a quintessence sector with a Ratra--Peebles potential at small values of the quintessence field coupled to a hidden sector gravitationally, our conclusions apply and are generic even if a complete scan of the parameter space $(m_{3/2},m_{1/2}, \\omega _Q)$ has not yet been performed.  \\par  The paper is arranged as follows. In the following section, \\ie section~\\ref{Quintessence and Supersymmetry Breaking}, we discuss the coupling of SUGRA quintessence to the hidden sector of supersymmetry breaking and the observable sector. In particular we find that the generic quintessence potential has a minimum with a mass for the quintessence field of the order of the gravitino mass.  On the other hand one can fine--tune the hidden sector dynamics to preserve the form of the SUGRA potential.  In section~\\ref{Implications for Gravity Experiments}, we investigate the implications for gravity experiments of this fine--tuning of the SUGRA potential. In section~\\ref{Implications for Cosmology}, we discuss the cosmological evolution of the generic case where the quintessence potential develops a minimum and show that the field settles at the minimum of the potential before nucleosynthesis. We then examine the cosmological perturbations around the minimum of the potential and show that there are no growing modes for quintessence perturbations. We then conclude and mention possible way-outs in section~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  We have presented a cosmological analysis of quintessence models in supergravity coupled to matter. For the quintessence sector, we have used the SUGRA model whose main feature is to reduce to the Ratra--Peebles potential at small values of the quintessence field. This is effectively the only property which is required here. Hence our results generalize to quintessence models with the same type of potentials. Supersymmetry is broken in a hidden sector which couples gravitationally both to the observable sector, \\ie~the MSSM, and the quintessence sector. Requiring that the gravitino mass is large enough to lead to acceptable masses for the sparticles implies that the quintessence potential is drastically modified by the presence of the hidden sector. In particular, as soon as the hidden sector fields are stabilized and supersymmetry is broken, we find that the quintessence potential develops a minimum acting as an attractor for the quintessence field in the very early universe. Indeed the quintessence field settles down at the minimum of the potential before BBN and both at the background and perturbation level, the model becomes equivalent to a pure cosmological constant scenario. This is the first quintessential difficulty as no observational consequences seems to spring from such a scenario. Moreover, the energy scales of the quintessence sector required to obtain the correct vacuum energy now are minute as a consequence of the tiny value of the quintessence field now. As a result, one must introduce a highly fine-tuned scale in the quintessence sector which is as difficult to explain as the smallness of the vacuum energy. On the other hand, one may hope to preserve the SUGRA potential runaway shape by tuning the hidden sector dynamics. In this case, we find that the local test of gravity (fifth force, weak equivalence principle and proton to electron mass ratio variation) are incompatible with the vev of the quintessence field implied by the presence of an attractor which, again, is required if we want insensitivity to the initial conditions. Hence it seems difficult to build models of quintessence in supergravity where the cosmological, gravitational and particle physics aspects are compatible.  One of the plausible possibilities consists in getting rid of the regularity of the K\\\"ahler potential for small values of the quintessence field. This is what happens in no--scale models for instance. The analysis of no-scale models leading to both supersymmetry breaking and quintessence is left for future work~\\cite{BMnoscale}.    \\clearpage \\nocite{*}"
    },
    "0606/astro-ph0606130_arXiv.txt": {
        "abstract": "s{ The Atmospheric Cherenkov Imaging Technique has opened up the gamma-ray spectrum from 100 GeV to 50 TeV to astrophysical exploration. The development of the technique (with emphasis on the early days) is described as are the basic principles underlying its application to gamma-ray astronomy. The current generation of arrays of telescopes, in particular, VERITAS is briefly described. } ",
        "introduction": "One of the last frontiers of the gamma-ray sky is that characterized by the distribution of TeV photons. These photons can be detected relatively easily with ground-based detectors (constituting a TeV ``window'' in the atmosphere); thus the detection of TeV gamma-ray sources did not have to await the availability of space platforms. In practice although the technology was available at an early date, it required the impetus of gamma-ray space astronomy to justify a major effort in a new discipline. Since it concerns the highest energy photons with which it is yet feasible to map the sky, it is of particular interest to high energy astrophysicists. Any source of TeV photons must be associated with a cosmic particle accelerator and of inherent interest to high energy particle physicists as well as students of the cosmic radiation.  To date almost all the observational results in the energy interval 100 GeV - 100 TeV have come from observations using the so-called ``Atmospheric Cherenkov Imaging Technique (ACIT).''  Although considerable effort has been applied to the development of alternative techniques, they are more specialized and will not be considered here.  In this historical review of the ACIT, emphasis will be on the early days in which the technique was established; a brief outline of the general principles underlying atmospheric Cherenkov telescopes (ACT) will be given and a description, albeit incomplete, of the ACIT as currently used and the present generation of instruments will be described. More complete accounts can be found elsewhere\\cite{porter1981},\\cite{weekes1988}, \\cite{Aharonian:97},\\cite{fegan1997},\\cite{ong},\\cite{weekes2003}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606240_arXiv.txt": {
        "abstract": "We present a measurement of the systemic proper motion of the Small Magellanic Cloud (SMC) made using the Advanced Camera for Surveys (ACS) on the \\textit{Hubble Space Telescope} (\\textit{HST}). We tracked the SMC's motion relative to 4 background QSOs over a baseline of approximately 2 years. The measured proper motion is : $\\mu_W = -1.16 \\pm 0.18 \\masyr, \\ \\mu_N = -1.17 \\pm 0.18 \\masyr$. This is the best measurement yet of the SMC's proper motion. We combine this new result with our prior estimate of the proper motion of the Large Magellanic Cloud (LMC) from the same observing program to investigate the orbital evolution of both Clouds over the past 9 Gyr.  The current relative velocity between the Clouds is $105 \\pm 42 \\kms$. Our investigations of the past orbital motions of the Clouds in a simple model for the dark halo of the Milky Way imply that the Clouds could be unbound from each other. However, our data are also consistent with orbits in which the Clouds have been bound to each other for approximately a Hubble time. Smaller proper motion errors and better understanding of the LMC and SMC masses would be required to constrain their past orbital history and their bound vs.~unbound nature unambiguously. The new proper motion measurements should be sufficient to allow the construction of improved models for the origin and properties of the Magellanic Stream. In turn, this will provide new constraints on the properties of the Milky Way dark halo. ",
        "introduction": "The Large and Small Magellanic Clouds (LMC \\& SMC), both satellites of the Milky Way (MW), provide a unique opportunity to study interacting galaxy pairs and three-body systems. It is expected from current models of hierarchical structure formation (e.g. Zentner \\& Bullock 2003) that the interaction between the Clouds and the MW will have played an important role in the dynamical evolution of the MW's outer parts (e.g. Font \\etal 2006). A resonant interaction between the LMC \\& the MW is thought to be responsible for the MW warp (Weinberg \\& Blitz 2006). In return, the MW has had a major influence on the Clouds' development, including their star formation history (Holtzman \\etal 1997; Harris \\& Zaritsky 2001; Smecker-Hane \\etal 2002), structural and chemical evolution (Mathewson \\etal 1986; Bekki \\& Chiba 2005), and kinematics (Hatzidimitriou \\etal 1993; Cole \\etal 2005).  This three-body interaction is also linked to the origin of the Magellanic Stream, an approximately circum-polar HI feature that trails the Clouds in their orbit around the MW (Wannier \\& Wrixon, 1972; Mathewson \\etal 1974; Putman \\etal 1998), the inter-Cloud bridge (Putman \\etal 1998), and the complicated geometry (Caldwell \\& Coulson 1986; Crowl \\etal 2001; but see also Welch \\etal 1987) and gas distribution of both Clouds (Stanimirovi{\\'c} \\etal 2004; Gardiner \\etal 1994).  Some current dwarf galaxies and globular clusters may have originated from tidal stripping as the Clouds orbited the Galaxy (Lin \\etal 1995).  The Magellanic Stream and any stars that have been tidally stripped from the Clouds as they orbit the MW provide a ``fossil record'' of the history of the build-up of MW mass (e.g. Belokurov \\etal 2006; Pe{\\~n}arrubia \\etal 2005; Johnston \\etal 1999). Decoding this record requires detailed modeling, describing how the internal evolution of the satellites is affected by tides, and sensitive observations that make it possible to falsify theoretical predictions. A major uncertainty in this effort is the orbital motion of the Clouds. While the radial velocities of the Clouds have been measured to high precision (van der Marel \\etal 2002; Harris \\& Zaritsky 2006), the velocity transverse to the line of sight (the proper motion) has been harder to constrain. This uncertainty complicates the efforts to infer the fossil record because the past history of the Clouds is ambiguous.  There have been several groups involved in modeling the orbits of the Clouds around the Galaxy with the intent of reproducing the Magellanic Stream (Murai \\& Fujimoto 1980; Lin \\& Lynden-Bell 1982; Heller \\& Rohlfs 1994; Moore \\& Davis 1994; Lin \\etal 1995; Gardiner \\& Noguchi 1996; Yoshizawa \\& Noguchi 2003; Mastropietro \\etal 2005; Connors \\etal 2005).  The models fall into two main categories: tidal models and ram-pressure stripping models. The former deal with the tidal force exerted by the Galaxy on the MCs.  The Stream is modeled as a product of the tidal disruption of the SMC, and the inter-Cloud region is the result of a close encounter between the Clouds, which are assumed to have been a bound system for the past Hubble time. In general, the conclusions of these studies are that the MCs are near perigalacticon, they are gravitationally bound to the Galaxy with apogalacticon beyond 100 kpc, and it is expected that they will become separated in the next $1$ - $2$ Gyr. The gas in the Stream is thought to have originated from the SMC between $1$ - $2$ Gyr ago, and the mass in the Stream is comparable to the gas content in the SMC itself (Lin \\etal 1995, Putman \\etal 2003).  The ram-pressure models also invoke an encounter between the Clouds to produce the inter-Cloud region, and subsequently the Stream is produced from collisions between the inter-Cloud region gas and either high velocity clouds in the Galactic halo (`discrete ram-pressure stripping'; Wayte 1991) or an extended halo of diffuse ionized gas (`diffuse ram-pressure stripping'; Moore \\& Davis 1994).  The ram-pressure stripping models have some difficulty in producing the leading arm of the Stream, while the tidal models do so naturally (Connors \\etal 2005). Also, the number density of high velocity clouds in the outer halo is uncertain, as is the existence of a sufficiently dense extended gaseous halo. The ram-pressure stripping models have become quite sophisticated (Mastropietro \\etal 2005), but so far only include the LMC in the analysis and not the SMC. The tidal models too, while increasingly detailed, have a few open questions which might only be adequately addressed with the inclusion of more detailed gas-dynamical properties. For instance, the lack of symmetry between the leading and trailing arms of the Stream are indicative of drag on the HI gas, and the apparent lack of stars in the Stream is still poorly understood. While there is some evidence for tidally stripped Magellanic stars far ($22\\degr$) from the LMC center and in the direction of the Carina dwarf (Mu{\\~n}oz \\etal 2006), no stars have yet been associated with the Magellanic Stream itself.    If the space velocities of both Clouds were known accurately enough they would be a valuable tool in understanding evolutionary features of both the Clouds and the MW.  As stated above, the radial velocities of the Clouds have been more readily determined than the transverse velocities, which can only be estimated via proper motions. Since the LMC is closer and larger, there has been much more detailed work on its proper motion. There is now good general agreement between the results from various teams (see Kallivayalil \\etal 2006; hereafter Paper~I). But it is worth noting that while the implied LMC space velocities agree with some of the models of the Stream (Heller \\& Rohlfs 1994), they differ by $\\sim100\\kms$ from others (Gardiner \\& Noguchi 1996). The proper motion of the SMC, has been much harder to constrain. A good inertial reference frame for such a measurement has been hard to come by.  CCD astrometry with a good sample of background QSOs is necessary to get a handle on its proper motion.   Previous work on the SMC proper motion includes: Kroupa \\& Bastian (1997) who used Hipparcos measurements of 11 stars to get $\\mu_W= -1.23 \\pm 0.84, \\mu_N=-1.21 \\pm 0.75 \\masyr$ (here we define the proper motions $\\mu_W, \\mu_N$ in the directions west and north as $\\mu_W \\equiv -\\mu_\\alpha \\cos(\\delta)$ and $\\mu_N \\equiv \\mu_\\delta$). There is also a study by Irwin \\etal (1996) using AAT (Anglo-Australian Telescope) and CTIO (Cerro Tololo Inter-American Observatory) 4m photographic plates covering a baseline of 15 - 20 years. A measurement of $\\mu_W =- 0.92 \\pm 0.2, \\ \\mu_N=-0.69 \\pm 0.2\\masyr$ is quoted in Irwin (1999), but the analysis of these data is unpublished. Anderson \\& King (2004b) measured a very accurate relative proper motion between the SMC and the Globular Cluster 47 \\textit{Tucanae} of $\\Delta \\mu_W = -4.716 \\pm 0.035 \\masyr$ and $\\Delta \\mu_N = -1.357 \\pm 0.021 \\masyr$. When combined with an estimate of the absolute proper motion of 47 \\textit{Tucanae} by Freire \\etal (2003), who report $\\mu_W = -5.3 \\pm 0.6 \\masyr$ and $\\mu_N = -3.3 \\pm 0.6 \\masyr$, this implies that the SMC's proper motion is $\\mu_W = -0.6 \\pm 0.6\\masyr$, $\\mu_N = -1.9 \\pm 0.6 \\masyr$. The unweighted average of these three independent SMC proper motion measurements\\footnote{Momany \\& Zaggia (2005) recently obtained a very different proper motion for the SMC using the USNO CCD Astrograph all-sky Catalog (UCAC2) : $\\mu_W =-4.44 \\masyr; \\ \\mu_N = -2.94 \\masyr$. This is inconsistent with our current knowledge of the MC-MW system and probably indicates the presence of systematic errors in the catalog, as the authors themselves point out. We therefore ignore this measurement in the following.} is $\\langle \\mu_W \\rangle =-0.91 \\pm 0.19 \\masyr, \\langle \\mu_N \\rangle = -1.28 \\pm 0.36\\masyr$.  This is broadly consistent with the current understanding of the Magellanic Stream and the MC-MW system, but the errors are are not accurate enough to significantly constrain its dynamics.  In this paper, we present the results of a project that has allowed us to measure the systemic proper motion of the SMC to 15\\% accuracy using $HST$ observations of a sample of background QSOs, with just a two year baseline. Our results for the LMC from the same $HST$ program were presented in Paper~I. \\S~2 describes the QSO sample; \\S~3 summarizes the analysis of the data; \\S~4 presents the SMC proper motion results. We then move on to an investigation of the global dynamics of both the LMC and the SMC given our new velocity measurements. \\S~5 describes the Clouds' movements within a fiducial model of the Galactic halo and \\S~6 gives a discussion and summary of the results. ",
        "conclusions": "We obtained two epochs of ACS HRC data to determine the proper motions of the LMC \\& SMC with respect to a sample of background QSOs distributed homogeneously behind the central few degrees of both galaxies. Our result for the SMC is presented here. With 4 QSOs and an approximately 2 year-long baseline, we have determined the proper motion of the SMC to be $\\mu_W = -1.16 \\pm 0.18 \\masyr, \\ \\mu_N = -1.17 \\pm 0.18 \\masyr$ . This is the best available measurement of its proper motion and is accurate to 15\\%. We have carried out a suite of tests to robustly quantify both random and systematic errors. We use our LMC (from Paper~I) and SMC proper motion estimates to investigate the past orbital evolution of both Clouds around the Milky Way. We find that while our data are consistent with orbits in which the Clouds have been bound to each other for approximately a Hubble time, there are also many unbound orbits within the error circles. So even though our errors on the SMC's motion are the most accurate thus far, they are not sufficient to uniquely say whether the Clouds are indeed bound. Also, it should be noted that our treatment of the orbits as a restricted few-body interaction within a fixed Milky Way potential is oversimplified. Structure formation in the Universe proceeds hierarchically and it is unrealistic to assume that the mass and properties of the galaxies involved were the same many Gyrs ago as they are now. For example, over the past 9 Gyr the mass of the Galaxy has probably increased by a factor of two (e.g. Bullock \\& Johnston 2005). Full cosmological simulations will be required to take this into account properly.  Now that the space velocities of both Clouds are well-characterized, it will be worthwhile to combine this information with the morphology and radial velocity of the Magellanic Stream to place constraints on the potential of the Galactic halo. The natural next step in this project is therefore to vary the prescription of the Galaxy model and quantify what type of halo would best match the Stream, given the observed LMC \\& SMC three-dimensional velocities. This will be the subject of a forthcoming paper. The parameters that this should constrain are the axial ratio of the halo, the slope of the circular velocity curve beyond $\\sim50$ kpc, and the implied Galactic mass at that distance. There are few other kinematic tracers at this location that can be used for such a measurement. Thus the combination of the orbital information of the Clouds and the morphology of the Magellanic Stream will provide a valuable new constraint. Combination of the orbital information of the MCs and the Magellanic Stream with tidal streams from other Local Group dwarfs such as the Sagittarius Stream (Ibata \\etal 1994; Majewski \\etal 2003) might further constrain the potential of the halo (e.g. Fellhauer \\etal 2006; Belokurov \\etal 2006; Johnston \\etal 2005; Law \\etal 2005; Helmi 2004).   The present-day velocities of the clouds assumed in models of the Stream have spanned a considerable range. However, most authors who have investigated tidal models for the origin of the Stream have used the assumed present-day velocities suggested by GN96. These differ from our observed values (given in Table~3) by $109 \\pm 16\\kms$ for the LMC and $142 \\pm 50\\kms$ for SMC (where each listed value is the length of the three-dimensional residual vector). So our measurements are not consistent with these models. Therefore, even without detailed calculations it is clear that revisions may be necessary to bring existing models into accordance with our observations. Previous models for the Stream have generally assumed a spherical logarithmic potential for the dark halo. It might be necessary to deviate from this simple prescription to obtain a good fit to the Stream given the observed present-day velocities. As mentioned, this would provide new insight into the properties of the Milky Way dark halo. But there are many other uncertainties in existing models as well. Probably the biggest uncertainties reside in the total mass and mass profiles of the Clouds themselves. Future models may also need to incorporate more details of the processes of hydrodynamics, mass loss, and star formation. Conversely, tighter constraints on the dynamical evolution of the Clouds may provide a better picture of the star formation histories of the Clouds.   An important outcome of this work is the evidence that $HST$ is stable enough to provide good proper motions using relatively short baselines. However, future improvements to our estimate of the SMC's proper motion should be possible by using a longer baseline and a larger sample of QSOs. In addition to the Geha \\etal (2003) MACHO sample, QSOs behind the SMC have also been found from the OGLE database (Dobrzycki \\etal 2003). Even though the mass distribution of the MCs remain the largest source of uncertainty, smaller errors on the SMC's motion would greatly reduce the observational phase space. This may allow us to definitively say whether the Clouds are bound to each other.  Hence, an additional epoch of data would be valuable.  With astrometric missions such as $SIM$ and $GAIA$ a rather long time down the road, $HST$ can continue to make important contributions to these subjects for several years to come."
    },
    "0606/astro-ph0606289_arXiv.txt": {
        "abstract": "We study temperature and density sensitivities of ratios of Si XI soft X-ray emission lines, in the wavelength range of 43---54\\AA\\,. The typical temperature of the formation of the analyzed lines is around 1.6~MK, which makes this analysis complementary to the analysis of He-like triplets being sensitive to hotter plasma. We present theoretical calculations and compare them with ratios obtained from high-resolution X-ray spectra of five solar-like stars: Procyon, $\\alpha$~Cen~A$\\&$B, $\\epsilon$~Eri, and Capella. We find that our results are in good agreement with results obtained by other authors through different diagnostics, namely the analysis of density- and temperature-sensitive He-like triplet lines. We further estimate the coronal pressure and filling factors from Si~XI lines in this study. ",
        "introduction": "Determinations of elemental abundances and structures in a various of stars belong to the most important tasks of observational astronomy and the understanding of the evolution of chemical elements. The temperature structure of emitting layers such as the stellar coronae must be known before elemental abundances can be determined. The density is another important physical parameter to describe the magnetically confined plasma in outer atmospheres of late-type stars, which can be used to estimate the spatial information for the stellar coronae. For those pre-main sequence stars, the constrained density can give us clues for the origin of the X-ray production such as magnetics alike the stellar coronae or a shock due to accretion processes.  Before the launch of missions with high resolution spectrograph such as the $Chandra$ and XMM-$Newton$, a determination of the electron temperature of X-ray emitting layers usually adopts a technique of global-fitting to spectra. The determined temperature structure from this technique strongly depends on the determination of continuum level of the observed spectra. The measurement of electron density is not possible for non-solar stars with previous X-ray satellites such as ASCA and ROSAT configured with low- and medium-resolution spectrographs. The observation data provided by high-resolution spectrograph on board new generation satellites $Chandra$ and XMM-$Newton$, allow us to estimate this parameter. Individual emission lines can be resolved from high-quality spectra, further the electron temperature can be derived from line intensity ratios such as $G$ value (between intercombination $i$ and forbidden $f$ lines {\\it vs} resonance line $r$) of He-like ions which was described by \\citet{PMD01} in detail. Since then, the estimation of the electron density becomes possible from X-ray emission line ratios such as $R$ value between $i$ and $f$ lines He-like ions. Using observation data with lower spectral resolution, $n_{\\rm e}$ and emission measures $EM$ can not be obtained independently, such that no emitting volumes $V$ could be estimated from the relation of $EM=n_{\\rm e}^2~V$. The structural information of the outmost atmosphere of stars was accessible from the analysis of lightcurves for a few special systems such as rotational modulation, eclipse mapping, etc. However, the solution of the reconstructed differential EM is always nonunique, and involves the difficulty of disentangling time-dependent versus geometric variability \\citep{SDH96,SF99,Gue04}. Since the operation of the two satellites with high-resolution and high-collecting areas, the two parameters have been derived from spectroscopic methods for every class of X-ray sources by many observers \\citep[etc.]{BGK00,CHD00,KHS02,NSB02,NBD03,RNM03,SS04,SRN05}. Using He-like triplet lines, a systematic investigation of the electron temperature and density for late-type stars with various levels of activity, has been made by \\citet{NSB02,Ness04} and \\citet{Testa04}. However, the UV field significantly affects the He-like triplet lines. Recently, \\citet{NS05} presented a density diagnostic ratio being independent from the UV field, that therefore allows to disentangle $n_{\\rm e}$ and UV field that both affect the O~VII triplet lines. Additionally, the derived density has also been used to estimate the production mechanism (alike stellar corona or accretion) of the X-ray emission in pre-main sequence stars \\citep{SS04,SRN05}.  Soft X-ray emissions of L-shell ions of silicon have been detected extensively in spectra of stellar coronae. These emissions carried out information of the two physical parameters for the solar-like coronae. In our previous work, we investigated the emission lines of Si X, and diagnosed the coronal electron density using them for late-type star---Procyon \\citep{LZS06a,LZS06b}. The soft X-ray emissions arising from 3d---2p, 3s---2p and 3p---2s transitions of Si XI have also been detected obviously for some cool stars, such as Procyon and $\\alpha$ Cen A\\&B \\citep{RMA02,RNM03}. And the temperature (1.6~MK) of maximum formation of Si~XI ion in the ionization equilibrium \\citep{MMC98}, is the typical temperature of corona of stars with low activities, which makes the analysis of Si~XI soft X-ray emission lines interesting.  EUV features arising from 2s$^2$--2s2p and 2s$^2$--2p$^2$ transitions of Si XI have been used to derive the electron temperature and density for solar flares. \\citet{KGF95} presented seven line ratios whereas these ratios are usually sensitive to both the electron temperature and density of the emitting plasma. Moreover measured ratios obtained with the Naval Research Laboratory's S082A spectrograph on board {\\it Skylab}, are larger than theoretical high-temperature and high-density limits \\citep{KGF95}. Later, \\citet{LBM01} calculated these line ratios again and compared with solar observations obtained with normal incidence spectrometer on board SOHO CDS, as well as laboratory measurements performed by \\citet{KKK96}. In contrast to the EUV lines, there has been little work on the soft X-ray transitions of Si XI in astrophysical literatures, probably as a result of the limited availability of high quality spectra of stellar in this wavelength region.  This paper is organized as following, we firstly investigate the Si~XI soft X-ray emission lines and present theoretical calculations of temperature- and/or density-sensitive line intensity ratios in Sect. 2. A brief description of observations and extractions of line fluxes is presented in Sect. 3. Results and discussions are outlined in Sect. 4. Our conclusions are given in Sect. 5. ",
        "conclusions": "In summary, we investigate L-shell soft X-ray emission lines of Si XI by using a more complete kinetic model. In this model, cascade effect from higher excited levels with $n=7$ on the upper levels of the present interested lines has been included. Our analysis reveals that some line intensity ratios such as $R_1$ and $R_2$, provide useful temperature diagnostics in a temperature range of 5.5---6.5 (in log$T_{\\rm e}$(K)), while $R_4$ is useful to diagnose the density in a range of $10^{8-10}$~cm$^{-3}$.  By comparing the observed ratios with theoretical prediction, the electron temperature and density for our sample including normal stars (Procyon and $\\alpha$~Cen~A\\&B) and active binary system---Capella, have been determined. The constrained temperatures by $R_1$ with 1.6---2.4~MK, are consistent with the temperature of maximum Si XI fractional abundance in the ionization equilibrium \\citep{MMC98} and the peak temperature of the DEM constructed from a global fitting of spectra for inactive stars. Since the blending effect in $R_2$ from Ar~IX line at 49.181~\\AA\\,, the correct line flux measurement of 49.207~\\AA\\, may induce a certain of systemical errors. Using the experimental ratio (1.11) of Ar~IX lines between 48.730~\\AA\\, and 49.181~\\AA\\,, we re-derived $R_2$ and the temperatures, yet they show a good agreement with those constrained by $R_1$. Taking into account the contamination of Ar~IX line, densities constrained by $R_3$ also show agreement with those by $R_4$ within $1\\sigma$ statistical error.  Temperatures determined by Si~XI emissions slightly higher than ones constrained by emissions of H- and He-like N, whereas they are in agreement within 1$\\sigma$ error. So we presume the emissions of N~VI-VII and Si~XI are emitted by the same plasma. The yielded temperatures from $G$ ratio are systematically lower than present ones as reported by \\citet{Testa04}. This may be due to the slightly lower peak temperature (1.4~MK) of the contribution function of the N~VI triplets. An agreement of the densities from the two different line ratios of highly charged silicon and nitrogen, is found. This further suggests that the Si~XI and N~VI-VII emitting plasma may be confined in the same loop.  A filling factor of $10^{-3}$ to $10^{-2}$ is derived for our sample from Si~XI, and a upper limit of 4\\% is obtained. Such results are consistent with the findings of \\citet{Testa04} and \\citet{Ness04} that filling factors are much smaller than unity; they also support the findings of \\citet{Testa04} that the filling factor of cooler plasma is directly proportional to the X-ray surface flux of stars."
    },
    "0606/astro-ph0606595_arXiv.txt": {
        "abstract": "We present catalogues of faint 1.4-GHz radio sources from extremely-deep Very Large Array pointings in the Lockman Hole, the Hubble Deep Field North (HDF-N) and ELAIS~N2. Our analysis of the HDF-N data has produced maps that are significantly deeper than those previously published and we have used these to search for counterparts to submm sources. For each of the fields we have derived normalised differential source counts and in the case of the HDF-N find no evidence for the previously-reported under-density of sources; our counts are entirely consistent with those found for the majority of other fields. The catalogues are available as an online supplement to this paper and the maps are also available for download. ",
        "introduction": "Our knowledge of the evolution of galaxies and active galactic nuclei (AGN) arguably owes as much to radio surveys as to those in any other band. Our knowledge of the cosmological evolution of radio-loud AGN is an obvious example \\citep{dunlop90}, and one which foretold with remarkable accuracy the history of accretion luminosity, and even that of star formation \\citep{dunlop98}.  Recent work has shown that deep ($\\mu$Jy) radio surveys have the power to probe star-forming galaxies, unbiased by dust obscuration, out to very high redshifts. Moreover, deep radio data have proved to be extraordinarily complementary to datasets obtained in the submillimetre (submm), infrared (IR) and X-ray wavebands \\citep{ivison02,donley05}.  Unfortunately, the radio waveband has not benefited from the technological improvements and investment seen at shorter wavelengths. This is in stark contrast with our mapping capability in the submm and IR wavebands, which has improved by many orders of magnitude in barely a decade.  At 20\\,cm, the National Radio Astronomy Observatory's (NRAO's) Very Large Array (VLA) has been the dominant survey facility for several decades, with the Australia Telescope (ATCA), MERLIN and Westerbork (WSRT) synthesis arrays all having found significant niches. The most recent major radio surveys were the NRAO VLA Sky Survey (NVSS), covering 82~per~cent of the sky to $\\sigma = 450\\,\\mu$Jy\\,beam$^{-1}$ \\citep{condon98} and FIRST (Faint Images of the Radio Sky at Twenty\\,cm, the radio equivalent of the Palomar Observatory Sky Survey) which mapped 10$^4$\\,deg$^2$ to around $\\sigma = 140\\,\\mu$Jy\\,beam$^{-1}$ \\citep{white97}.  Going deeper, radio surveys in a plethora of fields have reached noise levels of $\\sigma \\sim$10--20\\,$\\mu$Jy\\,beam$^{-1}$, the confusion limit at 20\\,cm for arrays smaller than $\\sim$5\\,km. These surveys, often unpublished but totalling around 10\\,deg$^2$, have been carried out with a variety of instruments (predominantly the VLA, but also the WSRT and ATCA) and include the Hubble Deep Field North \\citep{richards00}, the Hubble Deep Field South \\citep{huynh05}, the {\\em Spitzer} First Look Survey \\citep{condon03,morganti04}, the VLA--VIRMOS Deep Field \\citep{bondi03} and the Phoenix Deep Survey \\citep{hopkins03}.  Given the rarity with which surveys have probed significantly below the $\\sigma \\sim 10\\,\\mu$Jy\\,beam$^{-1}$ level \\citep{ivison02,muxlow05} it is natural to assume that to do so must require enormous allocations of observing time. This is not the case. Here, we describe three radio pointings which, together, required less than one week of VLA time. Each covers an area of $\\sim$320\\,arcmin$^2$ with 1.5-arcsec resolution, two of them to a noise level of $\\sim$5\\,$\\mu$Jy\\,beam$^{-1}$, the other to $\\sim$10\\,$\\mu$Jy\\,beam$^{-1}$.  Two of the VLA pointings lie within two regions of the sky which have very low Galactic backgrounds ($<$0.5\\,MJy\\,sr$^{-1}$ at 100\\,$\\mu$m and $E(B-V)< 0.01$) and which were selected for the {\\em Spitzer} Wide-area InfraRed Extragalactic (SWIRE) survey. These are the `Lockman Hole' and ELAIS N2. The Lockman Hole is a $4\\times3$-deg$^2$ region centred at $\\mathrm{10^h 48^m, +57^{\\circ} 04'}$ (J2000) with $\\sim5 \\times 10^{19}$\\,cm$^{-2}$ of H\\,{\\sc i} \\citep*{lockman86}. This makes it uniquely well suited to panchromatic, deep studies of the Universe. The final VLA pointing is located in the Hubble Deep Field North (HDF-N), part of the Great Observatories Origins Deep Surveys \\citep[GOODS,][]{dickinson03}. The ELAIS N2 and HDF-N data have been previously discussed by \\citet{ivison02} and \\citet{richards00}, respectively. A new reduction of the HDF-N 1.4-GHz VLA data (Morrison et al., in preparation) has also recently featured in \\citet{pope06}.  We describe our observations in \\S2, the data reduction in \\S3, the extraction of sources in \\S4, present differential source counts for all three fields in \\S5 and discuss our findings in \\S6. Throughout we adopt a cosmology, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$\\,km\\,${\\mathrm{s}^{-1}}$\\,Mpc$^{-1}$. ",
        "conclusions": "In this paper we have presented radio maps and catalogues from three legacy fields, the HDF-N, the 8-mJy Survey Region of the Lockman Hole and ELAIS~N2. These are areas of the sky which continue to be observed at all wavelengths and with a variety of telescopes, resolutions, etc; we hope that by making our radio maps available, the continued exploitation of these fields will be made much easier given that the protracted tasks of assembly, calibration and mapping of the data will be not be necessary.  The data presented here have been discussed before by other authors, but the new reduction of the HDF-N data is a marked improvement on that already published \\citep{richards00} and has yielded a new source catalogue with more and fainter sources. This has been used to find radio counterparts to two submm sources which previously had none; both of these identifications have recently also been made by \\citet{pope06} from a separate re-analysis of the HDF-N data (Morrison et al., in preparation). By going deeper still it may be possible to identify the radio counterparts to those submm sources in the HDF-N which remain undetected, sources which otherwise might be considered spurious \\citep[e.g.][]{greve04}. Deeper, more reliable radio maps of the HDF-N should be available in the near future due to further VLA observations at 1.4~GHz in the `A, `B', `C' and `D' configurations (P.I. Glenn Morrison).  Source counts have been presented for each of our fields, down to a flux density limit of 5$\\sigma$. These are broadly consistent with one another, but the observed differences (such as the low counts seen in the Lockman Hole) might be a consequence of cosmic variance. Our most important conclusion is that we find no evidence for the under-density of sources in the HDF-N reported by \\citet{richards00}, instead obtaining results which are entirely consistent with those found for other fields. It has long been noted that the HDF-N source counts were anomalously low and this finding appears to have been spurious. The recent resurgence in interest in the HDF-N radio data, both in the re-analyses of the archived data by ourselves and Morrison et al. (in preparation) and the on-going effort to add new data, should confirm our results and allow the source counts to be extended to still fainter flux densities."
    },
    "0606/astro-ph0606076_arXiv.txt": {
        "abstract": "Based on a spectral analysis of the X-ray continuum that employs a fully relativistic accretion-disk model, we conclude that the compact primary of the binary X-ray source GRS~1915+105 is a rapidly-rotating Kerr black hole.  We find a lower limit on the dimensionless spin parameter of $a_* > 0.98$.  Our result is robust in the sense that it is independent of the details of the data analysis and insensitive to the uncertainties in the mass and distance of the black hole.  Furthermore, our accretion-disk model includes an advanced treatment of spectral hardening.  Our data selection relies on a rigorous and quantitative definition of the thermal state of black hole binaries, which we used to screen all of the available {\\it RXTE} and {\\it ASCA} data for the thermal state of GRS~1915+105.  In addition, we focus on those data for which the accretion disk luminosity is less than 30\\% of the Eddington luminosity.  We argue that these low-luminosity data are most appropriate for the thin $\\alpha$-disk model that we employ.  We assume that there is zero torque at the inner edge of the disk, as is likely when the disk is thin, although we show that the presence of a significant torque does not affect our results.  Our model and the model of the relativistic jets observed for this source constrain the distance and black hole mass and could thus be tested by determining a VLBA parallax distance and improving the measurement of the mass function. Finally, we comment on the significance of our results for relativistic-jet and core-collapse models, and for the detection of gravitational waves. ",
        "introduction": "GRS 1915+105 has unique and striking properties that sharply distinguish it from the 40 known binaries that are believed to contain a stellar-mass black hole (Remillard \\& McClintock 2006, hereafter RM06). It is the most reliable source of highly relativistic radio jets in the Galaxy (Mirabel \\& Rodr\\'iguez 1994; Fender et al.\\ 1999; Miller-Jones et al.\\ 2006), and it is the prototype of the microquasars (Mirabel \\& Rodr\\'iguez 1999).  GRS~1915+105 (hereafter GRS1915) frequently displays extraordinary X-ray variability that is not mimicked by any other black hole system (e.g., Belloni et al.\\ 2000; Klein-Wolt et al.\\ 2002).  Its black hole (BH) primary is unique in displaying a constellation of high-frequency QPOs (HFQPOs), namely, 41~Hz, 67~Hz, 113~Hz and 166~Hz. The 67~Hz QPO is atypically coherent ($Q \\equiv \\nu/\\Delta\\nu \\sim 20$) and relatively strong (rms $ > 1$\\%) compared to the HFQPOs observed for six other accreting BHs (Morgan et al.\\ 1997; McClintock \\& Remillard 2006, hereafter MR06).  Among the 17 transient and ephemeral systems that contain a dynamically confirmed BH (RM06), GRS1915 is unique in having remained active for more than a decade since its discovery during outburst in 1992 (MR06).  GRS1915 has an orbital period of 33.5 days and is the widest of the BH binaries (BHBs), and it likely contains the most massive stellar BH (Greiner et al.\\ 2001; Harlaftis \\& Greiner 2004; RM06).  Zhang et al.\\ (1997) first argued that the relativistic jets and extraordinary X-ray behavior of GRS1915 are due to the high spin of its BH primary.  In their approximate analysis, they found that both GRS1915 and GRO~J1655--40 had high spins, $a_* > 0.9$ ($a_* = cJ/GM^2$, where $M$ and $J$ are the mass and angular momentum of the BH; $a_* = 0$ for a Schwarzschild hole and $a_* = 1$ for an extreme Kerr hole). Subsequently, Gierlin\\'ski et al.\\ (2001) estimated the spin of GRO~J1655--40 and LMC X--3.  Recently, we have firmly established the methodology pioneered by Zhang et al. and Gierlin\\'ski et al. by constructing relativistic accretion disk models (Li et al.\\ 2005; Davis et al.\\ 2005) and by modeling in detail the effects of spectral hardening (Davis et al.\\ 2005, 2006).  We have made these analysis tools publicly available via XSPEC ({\\it kerrbb} and {\\it bhspec}; Arnaud 1996).  Using this modern methodology, spins have now been estimated for several stellar-mass BHs, most notably: GRO~J1655--40 and 4U~1543--47 (Shafee et al.\\ 2006, hereafter S06), GRS1915 (Middleton et al.\\ 2006), and LMC X-3 (Davis et al.\\ 2006).  All of the plausibly reliable estimates of BH spin to date, including the present work, depend on fits to the X-ray continuum and measurements of the X-ray luminosity, coupled with optical measurements of BH mass, orbital inclination, and distance (e.g., S06).  In this paper, we show that GRS1915 does indeed harbor a rapidly-spinning Kerr BH as suggested by Zhang et al.\\ (1997).  However, in the case of GRO~J1655--40 the results obtained by ourselves and others show that the spin of this BH is modest ($a_* \\sim 0.75$; S06; Gierlin\\'ski et al.\\ 2001) and much lower than the value ($a_* \\sim 0.93$) suggested by Zhang et al.  The high spin reported herein for GRS1915 contradicts the modest spin value ($a_* \\sim 0.7$) reported by Middleton et al.\\ (2006), and we discuss this inconsistency in detail in \\S5.3.  Our spin estimates are based on an analysis of the ``thermal state'' of BHBs (MR06) whose remarkably simple properties have been recognized for decades.  Basic principles of physics predict that accreting BHs should radiate thermal emission from the inner accretion disk, and a multi-temperature model of a thin accretion disk was published shortly after the launch of {\\it Uhuru} (Pringle \\& Rees 1972; Shakura \\& Sunyaev 1973; Novikov \\& Thorne 1973; Lynden-Bell \\& Pringle 1974).  A nonrelativistic approximation to this model, now referred to as {\\it diskbb} in XSPEC (Arnaud 1996) was first implemented and used extensively by Mitsuda et al.\\ (1984) and Makishima et al.\\ (1986).  The two parameters of the model are the temperature $T_{\\rm in}$ and radius $R_{\\rm in}$ of the inner edge of the accretion disk.  In their review on BHBs, Tanaka and Lewin (1995) show for a few BHBs (see their Fig.\\ 3.14) that as the thermal disk flux varies by 1--2 orders of magnitude the value of $R_{\\rm in}$ remains constant to within $\\lesssim 20$\\%. This striking result prompted Tanaka \\& Lewin to comment that $R_{\\rm in}$, which was typically found to be $\\sim$ tens of kilometers, must be related to the radius of the innermost stable circular orbit ($R_{\\rm ISCO})$.  The stability of $R_{\\rm in}$ has by now been observed in great detail for many BHBs (e.g., Ebisawa et al.\\ 1994; Sobczak et al.\\ 1999; Sobczak et al.\\ 2000; Park et al.\\ 2004).  Further strong evidence for a thermal disk interpretation is provided by plots of the observed disk flux versus apparent temperature, which track the expected $L \\propto T^4$ relation for a constant inner disk radius (Gierlin\\'ski \\& Done 2004; Kubota \\& Done 2004).  Spin can be determined because it has a profound impact on the behavior and properties of a BH.  Quantitatively and specifically, consider two BHs with the same mass $M$, one a Schwarzschild hole and the other an extreme Kerr hole.  For the Kerr hole, the radius of the ISCO is six times smaller and the binding energy at the ISCO seven times greater than for the Schwarzschild hole.  Relative to the spinless BH, the much deeper gravity well of the extreme Kerr hole hardens the X-ray spectrum and greatly increases its efficiency for converting accreted rest mass into radiant energy.  The continuum fitting approach that we use is based on measuring spectral shape (hardness) and luminosity (efficiency).  This paper is organized as follows.  In \\S\\S2--4 we discuss respectively the selection, reduction, and analysis of the data.  In \\S5 we present our results for GRS1915 and compare them with those of Middleton et al.\\ (2006), and we present a table summary of the spins of GRS1915 and three other BHs.  The discussion topics in \\S6 include a description of our methodology, our rationale for favoring low-luminosity data, the natal origin of BH spin, the significance of measuring BH spin, and a proposed test of our model.  In \\S7 we offer our conclusions. ",
        "conclusions": "GRS~1915+105 is a rapidly-rotating BH with a lower limit on its spin parameter of $a_* > 0.98$.  Finally, we propose an observational test of our spin model.  \\noindent"
    },
    "0606/hep-ph0606204_arXiv.txt": {
        "abstract": "We show that as a result of non-linear self-interactions, scalar field theories that couple to matter much more strongly than gravity are not only viable but could well be detected by a number of future experiments provided they are properly designed to do so. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606507_arXiv.txt": {
        "abstract": "The KM3NeT project is a common European effort for the design of a $\\text{km}^3$-scale deep-sea neutrino telescope in the Mediterranean. For the upcoming Design Study simulations have been done using modified ANTARES software. Several concepts and ideas have been tested for their merits and feasibility. ",
        "introduction": "\\label{intro} The European neutrino telescope experiments \\cite{uli} have joined their efforts in the KM3NeT project \\cite{km3net} to design a $\\text{km}^3$-scale deep-sea neutrino telescope in the Mediterranean. A neutrino telescope of these dimensions on the Northern hemisphere, complementary to IceCube \\cite{icecube}, is necessary for high-energy neutrino astronomy.\\par As a first step for the KM3NeT project, a detailed simulation study of different detector models and photo-detectors is necessary. The software used by the ANTARES collaboration, with some modifications, is flexible enough for this task and was used for this study.\\par The modifications include the adaptation to a larger detector and a new causality filter for the event reconstruction \\cite{zab}. In addition some changes were necessary to implement the different photo-detection systems presented in this work.\\par Nonetheless the reconstruction and the selection cuts are optimized for ANTARES, and are therefore not necessarily optimal for different detector designs. However as the results show similar efficiencies for the detector models studied in this work and for the ANTARES detector, any reconstruction induced inefficiencies must be small.\\par The event sample used contained $10^9$ muon neutrinos with energies distributed between $10$ and $10^7\\,\\text{GeV}$ and incident isotropically from the whole solid angle. Only charged-current $\\nu$N interactions were simulated, and the hadronic component of the final state was neglected. A \\nuc{40}{K} background rate of 91 Hz per $\\text{cm}^2$ of photocathode area was used, corresponding to $40\\,\\text{kHz}$ for a 10\" Photomultiplier.\\par For comparison of different detector models the neutrino effective area was calculated as a function of the neutrino energy. The angular resolution, defined as the median angular deviation between reconstructed and true neutrino direction, was derived from the results, again as a function of energy.\\par ",
        "conclusions": "\\label{Conclusions} Several promising concepts of PM configurations and geometries for the design of the Mediterranean $\\text{km}^3$ neutrino detector were considered.\\par Through simulations it was shown that the use of many small PMs in pressure cylinders can provide better performance as conventional large hemispherical PMs. Inhomogeneous geometries have been shown to increase efficiency at low energies with only small losses at high energies, while drastically reducing the number of necessary strings.\\par The simulations also demonstrated that ring geometries are a possible way to reduce the number of strings without any loss of performance.\\par The decision for a definite configuration has to depend on the physics priorities of the project, as the performance of the different detector concepts depends on the neutrino energy. Further studies with dedicated software for a $\\text{km}^3$-scale detector are absolutely necessary to further clarify the results presented here, since the software used for this work is optimized for the ANTARES detector.\\par"
    },
    "0606/astro-ph0606731_arXiv.txt": {
        "abstract": "We discuss certain issues related to the limitations of density parameters for a ``radiation''-like contribution to the Friedmann equation using kinematical or geometrical measurements. We analyse the observational constraint of a negative $(1+z)^4$-type contribution in cosmological models. We argue that it is not possible to determine the energy densities of individual components of matter scaling like radiation from astronomical observations. We find three different interpretations of the presence of the radiation term: 1) the FRW universe filled with a massless scalar field in a quantum regime (the Casimir effect), 2) the Friedmann-Robertson-Walker (FRW) model in the Randall Sundrum scenario with dark radiation, 3) the cosmological model with global rotation. From supernovae type Ia (SNIa) data, Fanaroff-Riley type IIb (FRIIb) radio galaxy (RG) data, baryon oscillation peak and cosmic microwave background radiation (CMBR) observations we obtain bounds for the negative radiation-like term. A small negative contribution of dark radiation can reconcile the tension in nucleosynthesis and remove also the disagreement between $H_0$ values obtained from both SNIa and the Wilkinson Microwave Anisotropy Probe (WMAP) satellite data. ",
        "introduction": "In constraining the negative energy contribution to the Friedmann equation, a variety of astronomical observations were used, such as SNIa data \\cite{Riess:2004nr,Astier:2005}, FRIIb RG data \\cite{Daly04}, baryon oscillation peak and CMBR observations. Although we can obtain the bounds on the sum of density parameters for all fluids scaling like radiation, it is impossible to determine the contribution coming from each component. In particular it is impossible to determine the negative energy density arising from the Casimir effect or dark radiation and positive energy density arising from radiation because they scale with respect to redshift in the same way in the Friedmann equation. We note that the CMBR and big-bang nucleosynthesis (BBN) offer stringent conditions on this term, which can be regarded as an established upper limit on any individual components of negative energy density, and therefore on the Casimir effect, global rotation, discreteness of space following loop quantum gravity or brane dark radiation. We argue that even if the precise value of the density parameter for fictitious fluid is known from observations, it is still not possible to determine the nature (or, say, the origin) of the global rotation of the Universe using measurements basing on geometrical optics only, i.e. null geodesics $ds^2=0$. For example, if we consider the origin of this negative radiation-like term from global rotation effects, then it is not possible to come up with different forms of contribution leading to the same expression of $-\\Omega_{\\text{r},0}(1+z)^4$ type. There are, in principle, three different interpretations of the presence of the negative radiation term in the $H^2(z)$ relation, as follows.  The Casimir effect is a simple observational consequence of the existence of quantum fluctuation \\cite{Bordag01}. The Casimir force between conducting plates leads to a repulsive force, like the positive cosmological constant. Moreover, different laboratory experiments were designed to measure the Casimir affect with increased precision and thus strengthen the constraints on corrections to Newtonian gravitational law \\cite{Fischback98}. For a survey of recently obtained results in Casimir-energy studies see \\cite{Nesterenko05}. For our purpose, it is important that the effect of Casimir energy, which scales like radiation, can contribute into the $H^2(z)$ relation --- crucial for any kinematical test. It is also interesting that the same type of contribution to the effective energy density can be produced by loop quantum theory effects in semi-classical quantum cosmology \\cite{Vandersloot05,Singh05}. These effects give rise to an evolutionary scenario in which the initial singularity is replaced by a bounce. It is worth mentioning that the Casimir-type contributions arising from tachyon condensation are possible \\cite{McInnes06}.  In the brane world scenario, our universe is a submanifold which is embedded in a higher-dimensional spacetime, called `bulk space'. In the Randall and Sundrum scenario \\cite{Randall99}, the Einstein equations restricted to the brane reduce to a generalization of the FRW equation with additional terms. One of these  terms, called dark radiation, is of considerable interest because it scales like radiation \\cite{Vishwakarma03}. This term arises from the non-vanishing electric part of the five dimensional Weyl tensor. Dark radiation strongly effects both BBN and CMBR. Ichiki et al. \\cite{Ichiki03} gave limits on the possible contribution as $-1.23<\\rho_{dr}/\\rho_{\\gamma} \\le 0.11$ from BBN and $-0.41<\\rho_{dr}/\\rho_{\\gamma} \\le 0.105$ at the $95\\%$ confidence level from CMBR measurements. Let us note that a small negative contribution of dark radiation can also reconcile the tension between the observed ${}^4\\mathrm{He}$ and $D$ abundances \\cite{Ichiki02}.  Let us consider Newtonian cosmology following Senovilla et al. \\cite{Senovilla98}. Then we can define, following the authors, a homogeneous Newtonian cosmology with $\\rho$ and $p$ having no spatial dependence, i.e. $\\rho=\\rho(t)$ and $p=p(t)$, and we require that the velocity vector fields depends linearly on the spatial coordinates. Then the equation, which represents shear-free Newtonian cosmologies with expansion and rotation, well known as the Heckmann-Sch{\\\"u}cking model \\cite{Heckmann59} takes the following form \\begin{equation} \\label{eq:5} \\dot{a}^2=\\frac{\\rho(t_0)}{3a}-\\frac{2\\omega^2}{3a^2}+C \\end{equation} where $\\omega$ is the rotation scalar and $C$ is an integration constant. We interprete it in terms of curvature constant although in the Newtonian spacetime the curvature is zero. For our purpose, it is important that the effect of rotation produces a negative term in the Newtonian analogue of the Friedmann equation.  In the Newtonian cosmology in contrast to general relativity rotation can appear in homogenous and isotropic space. In general relativity the effect of rotation are strictly related to non-vanishing shear. The homogeneous universe with non-vanishing shear may expand and rotate relative to local gyroscopes. The relation between the rotation of the universe and the origin of the rotation of galaxies was also investigated \\cite{Li98,Godlowski03a,Godlowski05,Aryal06}. Additionally, the role of rotation of objects in the Universe, their significance and astronomical measurements was recently addressed by \\cite{Vishwakarma04,Godlowski03b}). ",
        "conclusions": "In our paper we analysed the observational constraints on the negative $(1+z)^4$-type contribution in the Friedmann equation. We discussed different proposals for the presence of such a dark radiation term. Although it is not possible, with present kinematic observations, to determine the energy densities of individual components which scales like radiation, we show that some stringent bounds on the value of this total contribution can be given. Our detailed conclusions are the following.  \\begin{enumerate}  \\item The combined analysis of SNIa data and FRIIb radio galaxies using baryon oscillation peaks and CMBR ``shift parameter'' give rise to the almost flat universe with $\\Omega_{\\text{m},0} \\simeq 0.3$.  \\item From the above-mentioned combined analysis, we obtain an upper bound $-\\Omega_{\\text{dr},0}<0.00035$ at the $95\\%$ confidence level. This is a stronger limit than obtained previously by us from SNIa data only \\cite{Godlowski03b}.  \\item We find new stringent limits on a negative component scaling like radiation from the location of the peak in the CMBR power spectrum, $-1.05 \\times 10^{-5}<\\Omega_{\\text{dr},0}<-0.5 \\times 10^{-5}$ at the $95\\%$ confidence level. This bound is stronger than that obtained from the BBN and CMBR by Ichiki et~al. \\cite{Ichiki03}.  \\item From the limit $-\\Omega_{\\text{dr},0}<1.05 \\times 10^{-5}$ we obtain that $\\Omega_{\\text{dr},0}+\\Omega_{\\text{r},0}>0$. This implies that $H^2(z)$ is always greater than zero ($H^2(z)>0$) and the bounce does not appear which means that the big bang scenario is strongly favoured over the bounce scenario.  \\item The discussed model with a small contribution of dark radiation type can also resolve the disagreement between values of $H_0$ obtained from SNIa and WMAP.  \\end{enumerate}"
    },
    "0606/astro-ph0606488_arXiv.txt": {
        "abstract": "In this paper we present INTEGRAL observations of 7 AGNs: two newly discovered type 1 Seyferts, IGR J18027-1455 and IGR J21247+5058, and five well known Seyferts, NGC 6814 (type 1.5), Cyg A (Type 2), MCG-05-23-16 (type 2), ESO 103-G035 (type2) and GRS1734-292. For IGR J18027-1455 and IGR J21247+5058 only INTEGRAL/IBIS data were available, while broadband spectra are presented and discussed for the remaining 5 sources for which either BeppoSAX or ASCA data were used in conjunction with INTEGRAL measurements. In the cases of NGC 6814 and GRS 1734-292, data taken in different periods indicate variability in the flux: in the case of NGC 6814 by a factor of 16 over a period of about 10 years. Although limited in size, our sample can be used to investigate the parameter space of both the photon index and cut-off energy. The mean photon index is 1.8, while the cut-off energy ranges from 30-50 keV to greater than 200 keV; in the particular case of MCG-05-23-16, ESO 103-G035 and GRS 1734-292 the cut-off energy is well constrained at or below 100 keV. We have also tested an  enlarged sample, which includes INTEGRAL data of 3 more AGNs, against the correlation found by a number of authors between the photon index and the cut-off energy but have found no evidence for a relation between these two parameters. Our analysis indicates that there is a diversity in cut-off energies in the primary continuum of Seyfert galaxies. ",
        "introduction": "The high energy emission of AGNs (Active Galactic Nuclei) is often to the first order well described by a power law of photon index 1.8-2.0, extending from a few keV to over 100 keV; at higher energies there is evidence of an exponential cut-off, the exact value of which is still uncertain \\citep{b25, b29, b41}. Secondary features, which are also commonly present, are considered to be the effects of reprocessing of this primary continuum and are relatively well understood  \\citep{b42, b41}. Modelling of high energy AGNs spectra has so far generally focussed on how to reproduce and explain the observed primary continuum shape. A good fraction of the proposed models ascribe the power law to the inverse Compton scattering of soft photons in a bath of {\\textquotedblleft{hot electrons}\\textquotedblright} (see for example \\citealt{b43}). Variations to this baseline model depend on the energy distribution of these electrons and their location in relation to the accretion disc. Measuring both the primary continuum and its cut-off energy is therefore crucial for understanding models and discriminating between them. While the photon index distribution has been well investigated \\citep{b22}, observational results on the cut-off energy have so far been limited by the scarcity of measurements above 10-20 keV, with most information coming from BeppoSAX broadband spectra. Analysis of type 1 and 2 AGNs \\citep{b25} provides evidence for a wide range of values in the cut-off energy, spanning from 30 to 300 keV and further suggests a possible trend of increasing cut-off energy as the power law index increases; it is not clear however if this effect is due to limitations in the spectral analysis or if it is intrinsic to the sampled source populations.  Direct INTEGRAL measurements enable us to obtain further observational results on the primary continuum and high energy cut-off thus providing more refined parameters for AGNs modelling. The sample of 12 AGNs first detected by INTEGRAL \\citep{b2} can be taken as a case study: they are representative of the larger sample of AGNs recently reported in the 20-100 keV band \\citep{b4} and are detected with adequate statistics to allow more in depth analysis. In particular, they are ideal objects for the study of the primary power law component, the presence of any high energy cut-off and the existence of a relation (if any) between these two parameters. Spectral analysis has so far been reported for 5 Seyfert galaxies listed in this initial sample, 4 of type 2 (NGC 4945, Centaurus A, Circinus Galaxy and NGC 6300, \\citealt{b32}) and one of type 1 (GRS 1734-292, \\citealt{b30}) as well as for the Blazar PKS 1830-211 \\citep{b11}. In the present paper, we concentrate on the remaining 6 objects: the two newly discovered AGNs, IGR J18027-1455 and IGR J21247+5058 both of type 1, and the four known Seyferts NGC 6814 (type 1.5), Cygnus A (type 2), MCG-05-23-16 (type 2) and ESO 103-G035 (type 2). The significance of the INTEGRAL detection for the last two objects was quite low at the time of the paper by \\citet{b32}, but they were analised anyway. Their analysis has therefore been updated here as more exposure is now available. INTEGRAL/IBIS data alone are reported for IGR J18027-1455 and IGR J21247+5058 while broadband spectra are presented and discussed in the other 4 cases since existing BeppoSAX and ASCA data are available for use in conjunction with the INTEGRAL spectra. We have also updated and reanalysed the data of the remaining objects in the \\citet{b2} sample but only in one case, that of GRS 1734-292, are the results of our spectral analysis different from already published values; analysis of the broadband ASCA/INTEGRAL spectrum of this object is therefore re-discussed in the present study.  \\begin{small} \\begin{figure} \\centering \\includegraphics[width=0.8\\linewidth]{fig1.eps} \\caption{ISGRI significance map of the region surrounding NGC 6814 and RX J1940.2-1025} \\label{fig1} \\end{figure} \\end{small} ",
        "conclusions": "Analysis of INTEGRAL data alone, or in conjunction with ASCA or BeppoSAX spectral data of a sample of 7 Seyfert galaxies of both type 1 and 2 provides information on the high energy emission of this class of objects. Only in the case of NGC 6814 and GRS 1734-292 do data taken in different periods indicate variability in the flux: in the case of NGC 6814 by a factor of 16 over a period of about 10 years. Although limited in size, our sample can be used to investigate the parameter space of both $\\Gamma$ and cut-off energy. The mean photon index is 1.83$\\pm$0.07, while the cut-off energy ranges from 30-50 keV to greater than 200 keV; in the case of MCG-05-23-16, ESO 103-G035 and GRS 1734-292 the cut-off energy is well constrained at or below 100 keV. We can also test our sample against the correlation found by a number of authors between the photon index and the cut-off energy. First claimed by \\citet{b27}, it was further analysed and discussed by \\citet{b22} and \\citet{b26}, taking advantage of the broad band BeppoSAX observations of AGN. At present it is still debated, since previous works found that the two parameters are not independent in the fitting procedure, with $\\Gamma$ increasing as the cut-off energy decreases \\citep{b25}. In figure \\ref{cutoff}, we plot the photon index versus the high energy cut-off for the sample of objects discussed here, to which we have also added the results obtained by \\citealt{b32}. Although these authors only used high energy data above 10 keV to constrain the power law parameters, their results on $\\Gamma$ and cut-off energy for NGC 4945, Circinus Galaxy and Centaurus A are quite solid being independently obtained by different instruments (see \\citealt{b32} and references therein). Discussing this correlation in terms of the mechanisms that give rise to the main spectral component goes beyond the scope of this paper. We simply point out here the range of values that these two parameters occupy in our sample: this is an important observational information for modelling of the primary continuum but also for estimating the AGNs contribution to the X-ray Cosmic Diffuse Background \\citep{b10}. From figure \\ref{cutoff} it is clear that our data do not show the correlation found by \\citet{b25}: the values for the photon index, in fact, cluster around 1.8, while the values of the cut-off energy range over a wide interval. We can conclude that the correlation between the high energy cut-off and the photon index remains to be proved and that a diversity in cut-off energy is most likely a property of Seyfert galaxies. INTEGRAL will keep observing AGNs for the duration of the mission and so we expect that more data on this class of objects will become available; this will provide further information on the high energy cut-off in AGNs and a deeper insight into the primary continuum emission.  \\begin{small} \\begin{figure*} \\centering \\includegraphics[width=0.8\\linewidth]{fig10.eps} \\caption{Primary power law photon index against cut-off energy for the sample of AGNs analysed in the present paper plus 3 AGNs (NGC4945, Circinus galaxy and CenA) discussed by \\citet{b32}.} \\label{cutoff} \\end{figure*} \\end{small} \\newpage"
    },
    "0606/astro-ph0606394_arXiv.txt": {
        "abstract": "We measure the 2-1 cumulant correlator power spectrum \\Cl, a degenerate bispectrum, from the second data release of the Wilkinson Microwave Anisotropy Probe (WMAP). Our high resolution measurements with SpICE span a large configuration space ($\\simeq 168\\times999$) corresponding to the possible cross-correlations of the maps recorded by the different differencing assemblies. We present a novel method to recover the eigenmodes of the correspondingly large Monte Carlo covariance matrix. We examine its eigenvalue spectrum and use random matrix theory to show that the off diagonal  terms are dominated by noise. We minimize the  $\\chi^2$ to obtain constraints for the non-linear coupling parameter $f_{NL} = 22 \\pm 52\\, (1\\sigma)$. ",
        "introduction": "Quantifying the non-Gaussianity in the cosmic microwave background (CMB) puts constraints on inflationary models and possibly identifies non-linear effects from large scale structure. In addition, systematics and foregrounds might also produce non-Gaussian signatures; this possibly weakens constraints on the primordial and secondary non-Gaussianities.  A natural phenomenological parametrization of non-Gaussian models is in terms of the perturbative non-linear coupling in the primordial curvature perturbation \\citep{KomatsuSpergel2001}: \\begin{equation} \\label{eq:modelreal} \\Phi({\\mathbf x}) =\\Phi_L({\\mathbf x}) +f_{NL}\\left( \\Phi^2_{L}({\\mathbf x})- \\left<\\Phi^2_{L}({\\mathbf x})\\right> \\right), \\end{equation} where $\\Phi_L({\\mathbf x})$ denotes the linear Gaussian part of the Bardeen curvature and $f_{NL}$ is the non-linear coupling parameter. The resulting leading-order non-Gaussianity is at the three-point level. Thus, the three point correlation function \\cite[e.g.,][and references therein]{ChenSzapudi2005} or its spherical harmonic transform, the bispectrum, directly estimate the leading order effect.  The bispectrum has been used extensively for studying non-Gaussianity \\citep{KomatsuEtal2005, CreminelliEtal2006, MedeirosContaldi2006, LiguoriEtal2006, CabellaEtal2006}. In previous measurements, the pseudo-bispectrum was used, which, like the pseudo-$C_l$'s, ignores in detail the effects of the complicated geometry induced by Galactic cut and cut-out holes. Pixel space statistics, such as the three-point correlation function, trivially deconvolve the geometric effects, as the convolution kernel is diagonal in pixel space. Indeed, SpICE \\citep[][Spatially Inhomogenous Correlation Estimator]{SzapudiEtal2001a} uses this simple fact to estimate the angular power spectrum without explicitly inverting the $M_{ll^\\prime}$ kernel \\citep{HivonEtal2002}. For  the bispectrum, the convolution kernel is even more complex than for the $C_l$'s, therefore pixel space methods are advantagous. We use the fact that the SpICE algorithm can be used to calculate (deconvolved) power spectrum of cumulant correlators $\\avg{\\delta T^N \\delta T^M}$ \\citep{SzapudiEtal1992}. These are degenerate $N+M$-point correlation functions, and their power spectra correspond to integrated $N+M-1$ poly-spectra. In this paper we focus on the 2-1 cumulant correlator power spectrum which is directly related to bispectrum configurations. These contain less configration information than the full bispectrum, although retain more than the skewness \\citep{KomatsuEtal2005}. Note that in terms of $f_{NL}$, with near optimal weighting, the statistical power of the skewness is nearly as optimal as the full bispectrum \\citep{KomatsuEtal2003, SpergelEtal2006}. For models with non-trivial configuration dependence, this is not the case.  We introduce the power spectrum of the 2-1 cumulant correlator together with the corresponding theoretical theoretical predictions. Statistical estimation in the data and simulations, and the theoretical calculations are are described in \\S3. In \\S4 we investigate the covariance matrix of the measurements, the $\\chi^2$ analysis of $f_{NL}$. We summarize and discuss the work in \\S5. ",
        "conclusions": ""
    },
    "0606/astro-ph0606327_arXiv.txt": {
        "abstract": "We report on the results from a five-night campaign of high-speed spectroscopy of the 17-minute binary AM Canum Venaticorum, obtained with the 4.2-m William Herschel Telescope on La Palma.  We detect a kinematic feature that appears to be entirely analogous to the `central spike' known from the long-period, emission-line AM CVn stars GP Com, V396 Hya and SDSS J124058.03$-$015919.2, which has been attributed to the accreting white dwarf. Assuming that the feature indeed represents the projected velocity amplitude and phase of the accreting white dwarf we derive a mass ratio \\q\\ for AM CVn. This is significantly higher than the value found in previous, less direct measurements. We discuss the implications for AM CVn's evolutionary history and show that a helium star progenitor scenario is strongly favoured. We further discuss the implications for the interpretation of AM CVn's superhump behaviour, and for the detectability of its gravitational-wave signal with \\emph{LISA}.  In addition we demonstrate a method for measuring the circularity or eccentricity of AM CVn's accretion disc, using stroboscopic Doppler tomography. We test the predictions of an eccentric, precessing disc that are based on AM CVn's observed superhump behaviour. We limit the effective eccentricity in the outermost part of the disc, where the resonances that drive the eccentricity are thought to occur, to \\e, which is smaller than previous models indicated. ",
        "introduction": "AM Canum Venaticorum was found in a survey of faint, blue objects by \\citet{hz}. It was observed to have peculiarly broad and shallow helium absorption lines, but no hydrogen \\citep{greenstein}. The star was shown to be a possible ultra-compact white dwarf binary by \\citet{smak67}, who discovered photometric variations on an 18-minute period; quickly thereafter, \\citet{paczynski} noted that it would be a prime example of a binary whose evolution is expected to be governed by gravitational-wave radiation, and which could serve as an excellent test for the existence of gravitational radiation. The interpretation as an interacting binary analogous to the Cataclysmic Variables was first proposed by \\citet{warner} upon their discovery of rapid flickering in the light-curve. The discussion as to the true orbital period of the system was finally put to rest by \\citet{nsg} who discovered a kinematic `S-wave' feature in time-resolved spectra of AM CVn, thereby proving that the orbital period is 1028 seconds, while the main photometric signal at 1051 seconds is to be interpreted as a `superhump'.  Here we present phase-resolved spectra of AM CVn with significantly higher spectral resolution and signal-to-noise ratio than those previously used by \\cite{nsg}. Moreover, our current data-set fully samples the proposed 13.37-h accretion disc precession period \\citep{pattersonprec,skillman} for the first time, allowing for a characterisation of the spectroscopic appearance of the system as a function of orbital period, precession period, and superhump (beat) period. ",
        "conclusions": "\\subsection{A central spike in AM CVn?} The mysterious central spike has so far been observed in long-period, low mass transfer rate AM CVn stars, where the accreting white dwarf dominates the optical flux. It is thought to originate on (or very close to) the surface of the accreting white dwarf, since the spike perfectly tracks the expected movement of the accreting white dwarf relative to the bright spot in the accretion disc. In addition, there is no accretion disc component that is expected to move at such low velocities, and the phase of the spike relative to the bright spot does not match the phase of the secondary. See \\cite{lmr} for a detailed analysis of the central spike in GP Com.  The kinematic feature observed in the \\mbox{He\\,{\\sc i}} 4471 line in \\obj, as presented here, has all the characteristics for being the perfect analogy to the central spike in GP Com \\citep{lmr}, V396 Hya (Steeghs et al.\\ in preparation) and SDSS J124058.03$-$015919.2 \\citep{roelofs}: (1) it is intrinsically redshifted; (2) it agrees in amplitude with the expected velocity amplitude of the accreting white dwarf (and not with any other binary component's velocity amplitude); and (3) it agrees in phase with the expected phase of the accreting white dwarf, relative to the bright spot. We therefore conclude that it is \\emph{very} likely that it is the central spike.  The mass ratio for AM CVn implied by the central spike is significantly higher than previously thought, although there has been a study by \\citet{pearson} suggesting, in fact, a mass ratio $q\\sim0.19-0.25$ based on a model in which the secondary star is moderately magnetic. The motivation for that study was to try to reconcile the mass ratio with the $K$-velocity of the \\mbox{He\\,{\\sc ii}} 4686 emission line measured (but discarded) by \\citet{nsg}, which suggested a mass ratio in this regime. It is clear from this work however, in particular figure \\ref{heliumemission}, that the measured $K$-velocities of the \\mbox{He\\,{\\sc ii}} 4686 emission indeed \\emph{cannot} be used reliably as an indicator for the motion of the accretor, since the apparent kinematic signal could easily be dominated by the eccentricity of the disc rather than the actual motion of the primary.  \\subsection{Implications for AM CVn's formation channel}  One of the long-standing questions regarding the AM CVn stars is how they are formed. Theoretically, they can be formed in both double-degenerate and single-degenerate configurations: Roche-lobe overflow from a white dwarf onto another (more massive) white dwarf (e.g.\\ \\citealt{nelemans}), or from a helium-burning star onto a white dwarf \\citep{ibentutukov}. A third suggested formation channel is Roche-lobe overflow from an evolved main-sequence star onto a white dwarf \\citep{podsi}, where mass transfer starts by the time hydrogen core burning ends. This latter scenario, however, has the problem of leaving significant amounts of hydrogen, while our high-S/N average and phase-binned spectra show exclusively helium and metals. For the other two formation channels, the main discriminator is the mass of the secondary star: in the helium-star channel, one expects a relatively hot and massive, only partially degenerate secondary, while in the double-degenerate channel, one expects a `cold' degenerate donor star.  The significantly higher mass ratio found for AM CVn in this paper has important implications for the nature of the donor star, and hence for AM CVn's formation channel, as it suggests a relatively massive donor. If we combine the orbital period with the mass--radius relation for a relatively `cold' (core temperature $T_c \\leq 10^6$\\,K), degenerate helium white dwarf as used in \\citet{nelemans} and refined recently in \\citet{deloye}, we get a secondary mass $M_2\\approx 0.035M_{\\sun}$ and hence, through our minimum mass ratio of \\q, a maximum primary mass of only $M_1=0.19M_{\\sun}$. This is \\emph{very} low; the maximum combined mass of $M_1+M_2=0.22M_{\\sun}$ would, for instance, be significantly lower than the average value $M_1+M_2=0.8M_{\\sun}$ found in detached WD--WD binaries (e.g.\\ \\citealt{nelemansspy}). This casts doubt on the possible white dwarf (WD) nature of the donor star, and thereby on the double-degenerate formation channel for AM CVn, unless the white dwarf donor was still very hot upon Roche-lobe overflow -- that is, unless mass transfer started shortly after the second common-envelope phase.  Concurrent with the measurement of AM CVn's mass ratio presented here, it was discovered that AM CVn's distance is larger than expected, through an \\emph{HST/FGS} parallax measurement ($\\pi=1.65\\pm0.30$ mas, Roelofs et al.\\ in preparation). This means a smaller absolute magnitude $M_V$ and thus, since we expect the accretion disc to dominate the optical flux, a higher accretion luminosity.  In order to link the observed $M_V$ to the mass of the secondary $M_2$, we proceed in the same way as in \\cite{deloye}. We assume conservative mass transfer that is driven entirely by loss of angular momentum due to the emission of gravitational waves and, since accretion proceeds via an accretion disc, we assume that all angular momentum carried by the transferred matter is fed back to the orbit. We then have \\begin{equation} \\frac{\\dot{M_2}}{M_2} = \\frac{\\dot J}{J} \\frac{2}{\\zeta_2+5/3-2q} \\label{mdotmeq} \\end{equation} where $J$ is the orbital angular momentum, $\\dot J$ is the orbital angular momentum loss rate \\begin{equation} \\frac{\\dot J}{J} = -\\frac{32}{5}\\frac{G^3}{c^5} \\frac{M_1 M_2 (M_1+M_2)}{a^4} \\end{equation} due to the emission of gravitational waves \\citep{landau}, and \\begin{equation} \\zeta_2 \\equiv \\frac{\\mathrm{d} \\log R_2}{\\mathrm{d} \\log M_2}. \\end{equation} We take $\\zeta_2=-0.06$ for $M_2 < 0.2M_\\odot$ as in \\citet{nelemans}, based on evolutionary calculations for a mass-losing helium-burning star by \\citet{tf}. Clearly $\\zeta_2 = 0$ around the point where helium burning stops, while it may be somewhat closer to $\\zeta_2=-0.19$ along the semi-degenerate part of the track for a secondary that is more helium depleted, as found by \\citet{savonije}. This represents a $\\lesssim$10\\% uncertainty in eq.\\ (\\ref{mdotmeq}).  The last step is to calculate the bolometric luminosity $L$, for which we use \\begin{equation} L = \\frac{1}{2} \\dot{M_2} \\big( \\Phi(L_1) - \\Phi(R_1) \\big) \\label{Lacc} \\end{equation} where $\\Phi$ is the common Roche potential, at the inner Lagrange point $L_1$ and the surface of the accretor $R_1$. This states that the matter in the disc stays virialised during the accretion process, which should hold for approximately Keplerian particle orbits.  Figure \\ref{mdotmplot} shows $M_V$ as a function of $M_2$ for $q=0.18$, $\\zeta_2=-0.06$ and a bolometric correction $\\mathrm{BC}=-3.0$ (see below). The lower limit to $M_2$ is given by the requirement that the system does not eclipse. The data point $M_V = 5.0\\pm0.4$ as determined from our aforementioned \\emph{HST} parallax is overplotted, which yields a secondary mass $M_2=0.125\\pm0.012\\,M_\\odot$ and, via \\q, a primary mass $M_1 = 0.68\\pm0.06\\,M_\\odot$. The corresponding mass transfer rate is $\\dot{M_2} = 6.7^{+1.9}_{-1.3}\\cdot 10^{-9}M_\\odot$/yr, and the inclination of the binary comes out at $i=43\\pm2^\\circ$. Table \\ref{parameters} lists the system parameters for AM CVn in convenient tabular form.  \\begin{table} \\begin{center} \\begin{tabular}{l l} \\hline Parameter               &Value\\\\ \\hline \\hline $P_\\mathrm{orb}$ (s)    &$1028.7322\\pm0.0003$ \\citep{skillman}\\\\ $q$                     &$0.18\\pm0.01$\\\\ $M_1 (M_\\odot)$         &$0.68\\pm0.06$\\\\ $M_2 (M_\\odot)$         &$0.125\\pm0.012$\\\\[.5ex] $\\dot{M_2} (M_\\odot/\\mathrm{yr})$         &$6.7^{+1.9}_{-1.3}\\cdot10^{-9}$\\\\[.5ex] $i$ (degrees)           &$43\\pm2$\\\\[.5ex] $d$ (pc)                &$606^{+135}_{-93}$ (Roelofs et al.\\ in prep.)\\\\ \\hline \\end{tabular} \\caption{The collection of system parameters for AM CVn.} \\label{parameters} \\end{center} \\end{table}  Our results have been corrected for the geometric effect that the `apparent absolute magnitude' \\citep{warnerbook} of an accretion disc depends on the inclination. A factor \\begin{equation} \\Delta M_V(i) = -2.5 \\log \\left(\\frac{\\cos i}{1/2}\\right) \\end{equation} has been taken into account as the correction from apparent absolute magnitude to absolute magnitude, where the factor $1/2$ represents the direction-averaged fraction of the disc area seen by the observer, $\\langle\\cos i\\rangle$. We neglect the (unknown) effect of limb darkening.  In converting the bolometric luminosity $M_\\mathrm{bol}$ obtained from equation (\\ref{Lacc}) to an absolute visual magnitude $M_V$, the bolometric correction is of some importance. AM CVn's spectral energy distribution from the far-UV to the optical (Roelofs et al.\\ in preparation; see also \\citealt{nasser}) indicates that it is dominated by a blackbody of $\\simeq$$30,000$\\,K. This corresponds to a BC of $-3.0$, with an estimated error of $0.3$. If we calculate the minimum temperature the accretion disc must have in order to be able to radiate away the required luminosity (for any given $M_{1,2}$ we know the maximum radiating surface of the disc), we find $T_\\mathrm{disc}\\gtrsim 30,000$\\,K for accretion rates $\\dot{M_2} \\ge 4\\cdot 10^{-9} M_\\odot$/yr. The BC we use is thus self-consistent with the accretion rate we derive; in particular, it seems unlikely that we are overestimating the mass transfer rate due to an overestimate of the BC.  \\begin{figure} \\centering \\includegraphics[angle=270,width=84mm]{mdotm.ps} \\caption{Relation between $M_2$ and $M_V$ for AM CVn. Our measurement of the absolute magnitude $M_V$ and its implications for the secondary mass $M_2$ are indicated by the dashed line. The open circle represents the minimum secondary mass for which the system does not eclipse (assuming a disc that extends to the $3:1$ resonance radius); the dotted line, towards larger $M_2$ and smaller $M_V$, shows the helium-burning regime where we expect our assumptions regarding the mass-radius relation to break down ($>$10\\% error). The labelled dots indicate $\\log\\dot{M_2}$ in $M_\\odot$/yr.} \\label{mdotmplot} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=270,width=84mm]{mr.ps} \\caption{Relation between $M_2$ and $R_2$ for AM CVn. The upper and lower solid line represent the evolutionary tracks for semi-degenerate helium star and for cold white dwarf secondaries, respectively \\citep{nelemans,tf}. The dotted line is the requirement that the secondary star fills its Roche lobe. The dashed line represents our measurement of $M_2$. The data favour a semi-degenerate helium star secondary, rather than a white dwarf.} \\label{mrplot} \\end{figure}  Summarising, we have two independent pieces of evidence suggesting a relatively massive donor star -- the large mass ratio and the large luminosity. If we compare our results to evolutionary models for helium star and white dwarf secondaries as given by \\citet{nelemans}, we see that our derived secondary mass is compatible with the secondary being the semi-degenerate core of a formerly helium burning star, while a `cold' degenerate donor star is ruled out. See figure \\ref{mrplot}. The secondary mass $M_2=0.125M_\\odot$ is close to the turn-off mass for helium burning, which lies at $M_2\\sim 0.16 M_\\odot$ \\citep{savonije,tf}. We conclude that it is very likely that AM CVn formed through an evolutionary channel in which the donor star was, upon the start of Roche-lobe overflow, still burning helium. These observations thus provide the first observational evidence that the helium star evolutionary channel contributes to the AM CVn population.  \\subsection{Implications for AM CVn's gravitational wave signal}  We expect the orbit of AM CVn to be circular, so that the gravitational wave polarisation amplitudes $A_+$ and $A_\\times$ at twice the orbital frequency, $f$, become (e.g.\\ \\citealt{timpano}) \\begin{align} A_+ &= 2\\frac{(G\\mathcal{M})^{5/3}}{c^4d} \\left(\\pi f\\right)^{2/3} \\left(1+\\cos^2i\\right)\\\\ A_\\times &= -4\\frac{(G\\mathcal{M})^{5/3}}{c^4d} \\left(\\pi f\\right)^{2/3} \\cos i \\end{align} where $\\mathcal{M}=(M_1M_2)^{3/5}/(M_1+M_2)^{1/5}$ is the so-called chirp mass, $i$ is the inclination of the binary, and $d$ its distance. Defining the strain amplitude \\begin{equation} h=\\left(\\frac{1}{2}\\left(A^2_{+} + A^2_\\times\\right)\\right)^{1/2} \\end{equation} and converting to more convenient solar units, one finds \\begin{align} h &= 2.84 \\cdot 10^{-22}\\sqrt{\\cos^4i+ 6\\cos^2i + 1}\\nonumber\\\\ &\\phantom{=}\\times\\left(\\frac{\\mathcal{M}}{M_\\odot}\\right)^{5/3} \\left(\\frac{P_\\mathrm{orb}}{\\mathrm{1\\,hr}}\\right)^{-2/3} \\left(\\frac{d}{\\mathrm{1\\,kpc}}\\right)^{-1} \\end{align}  Filling in the numbers derived above gives $h=2.1^{+0.4}_{-0.3}\\cdot 10^{-22}$. Figure \\ref{gwr} shows the estimated strain amplitudes and frequencies of the `known' AM CVn stars. AM CVn itself is the first potential \\emph{LISA} source which is sufficiently well-constrained that we can plot error bars with some confidence, and we see that it stands out significantly above both the instrumental design sensitivity and the estimated confusion-limited Galactic background due to, mainly, detached WD--WD binaries.  \\begin{figure} \\centering \\includegraphics[angle=270,width=84mm]{gwr.ps} \\caption{Gravitational-wave strain amplitude of AM CVn (solid circle). The upper and lower dashed lines show the design sensitivities of \\emph{LISA} for a S/N of 5 and 1, respectively, in one year of data-collecting \\citep{larson}. The solid line is a population synthesis prediction for the confusion-limited Galactic background \\citep{npy}. The open circles represent `known' AM CVn stars that are as yet not well-constrained in terms of either distance or component masses, or both.} \\label{gwr} \\end{figure}  \\subsection{The superhump excess--mass ratio relation}  It is understood from numerical simulations, and it is observed in outbursting dwarf novae and nova-likes, that there exists a relation between the superhump period excess $\\epsilon$ and the mass ratio $q$ of the binary. See \\citet{patterson}, and references therein, for an overview including the very latest empirical work on the $\\epsilon(q)$ relation. In short, they conclude that for hydrogen-rich systems, \\begin{equation} \\epsilon(q) = 0.18 q + 0.29 q^2 \\label{epsilon} \\end{equation} For AM CVn, this results in $q=0.10$ based on the photometry of \\citet{skillman}. Although it is unknown what the intrinsic spread in $q(\\epsilon)$ is, it appears from figure 9 in \\citet{patterson} to be quite low, no more than about $0.01$. This means that the $q(\\epsilon)$ for AM CVn deviates significantly from our minimum kinematic mass ratio \\q.  Of course, the superhump excess $\\epsilon$ may depend on more parameters than just the mass ratio $q$; it is quite conceivable that the helium nature of AM CVn may be of influence. A recent study by \\citet{goodchild} indicates, for instance, quite a significant dependency of $\\epsilon$ on the thickness of the accretion disc, where thicker discs lead to smaller superhump period excesses for a given mass ratio. If this effect should be at play, our results would suggest a physically thicker disc in AM CVn compared to discs in hydrogen-rich systems of the same mass ratio.  Although AM CVn represents only a single data point, it seems advisable to apply the empirical superhump excess--mass ratio relation with caution for such ultra-compact binaries. An obvious test would be to obtain the mass ratio in HP Librae, the other known helium nova-like, and check whether it, too, shows a discrepancy with relation (\\ref{epsilon}).  \\subsection{Size and eccentricity of the accretion disc}  In our derivation of the minimum mass ratio \\q, we have assumed a maximum accretion disc radius equal to the minimum radius of the primary Roche lobe. It is customary to assume a slightly more stringent limit on the maximum accretion disc radius, caused by tidal truncation of the disc. See e.g.\\ \\citet{warnerbook} and references therein. For reference, we show in figure \\ref{truncation} the often-used tidal truncation radius \\begin{equation} \\frac{R_T}{a} = \\frac{0.6}{1+q}\\qquad (0.03 < q < 1.0) \\label{truncfunc} \\end{equation} for an inviscid flow, which is indeed smaller than the full minimum Roche lobe radius we assume. However, since this tidal truncation radius depends, in particular, on the unknown viscosity of the matter in the disc (and increasing with viscosity), we have not used it as a hard upper limit to the disc's radius. Interestingly, one sees that the best-fitting accretion disc radius we derive for ballistic stream velicities in the bright spot (reproduced in figure \\ref{truncation}) lies very close to, but slightly above, the tidal truncation radius (\\ref{truncfunc}).  \\begin{figure} \\centering \\includegraphics[angle=270,width=84mm]{truncation.ps} \\caption{Same as the left panel in figure \\ref{massratio}, but now including the tidal truncation radius (eq.\\ (\\ref{truncfunc})) for an inviscid disc, indicated by the dot-dashed line.} \\label{truncation} \\end{figure}  The effective eccentricity value \\e\\ we derive for the outermost part of the disc, from the stream--disc impact spot, is smaller than typical values $e\\sim0.1-0.2$ found in numerical simulations. Although the eccentricity is not one of the observables that is usually quoted in numerical studies (the resulting superhump period excess is often the only given quantity), it can be inferred that the effective eccentricity we find is significantly smaller than that found by \\citet{simpsonwood} in their simulations of discs in AM CVn stars.  A possible explanation would be that tidal truncation, discussed in the previous paragraph, acts to circularise the disc at the outer edge to some extent by effectively removing particles from the most eccentric orbits there. This tidal truncation process is essentially a competition between viscosity transporting angular momentum outwards and tidal forces `dissipating' this angular momentum increasingly strongly with radius, and predictions about this tidal truncation process would thus be affected by our lack of understanding of the viscosity in accretion discs.  \\subsection{Spin of the accretor and tidal synchronisation}  If the central spike originates from the surface of the accreting white dwarf, in particular if it rotates with the accretor, its width puts an interesting constraint upon the accretor's spin. The central spike is unresolved in our data, and we therefore place an upper limit of $\\sim$45 km/s, equal to our formal spectral resolution, on the FWHM of the feature. For the system parameters derived for AM CVn, corotation with the orbit would imply a projected equatorial velocity of $\\sim$36 km/s. The central spike we observe is thus consistent with a corotating primary, while significantly faster rotation is ruled out.  We can estimate the tidal synchronisation timescale of the accretor from this constraint. For simplicity, we assume that AM CVn has been steadily accreting matter at a rate $\\dot{M_2}\\sim 7\\cdot 10^{-9}\\,M_\\odot$/yr for a long time. As long as the accretor is near corotation with the orbit, the spin-up timescale $\\tau$ of the primary due to the accretion of disc matter can be written as \\begin{equation} \\tau = \\frac{\\omega}{\\dot\\omega} \\approx \\frac{2\\pi}{P_\\mathrm{orb}}k\\sqrt{\\frac{M_1}{G}}\\frac{R_1^{3/2}}{\\dot{M_2}} \\end{equation} where $k\\approx 0.18$ (see \\citealt{masstransfer}) is the moment of inertia factor ($k=2/5$ for a solid sphere), and $\\omega$ is the angular velocity of the accretor. Requiring that tidal synchronisation proceeds on timescales similar to or shorter than the accretion-induced spin-up of the primary, gives a maximum tidal synchronisation timescale $\\tau_s \\lesssim 2\\cdot10^5$ yr.  A key question in binary evolution theory is whether two white dwarfs, upon Roche-lobe overflow at orbital periods of a few minutes, can manage to stabilise the accretion process and avoid a merger by feeding back angular momentum from the spun-up accretor to the orbit. This has profound implications for the number of systems that may survive the initial phase of mass transfer, when accretion proceeds at a high rate of $\\dot{M_2}\\sim10^{-6}M_\\odot$/yr \\citep{masstransfer}.  Although its absolute magnitude is highly uncertain from a theoretical point of view, tidal synchronisation is understood to scale with the orbital separation and the radius of the accretor as \\begin{equation} \\tau_s \\propto \\left(\\frac{M_1}{M_2}\\right)^2 \\left(\\frac{a}{R_1}\\right)^6 \\end{equation} (see, e.g., \\citealt{masstransfer}). This means that, for a three-minute white dwarf binary in which Roche-lobe overflow is about to commence, the tidal synchronisation timescale will be at least a factor $10^3$ shorter than for AM CVn, or $\\tau_s \\lesssim 200$ yr. This is an interesting result since it is exactly the regime that is needed for an appreciable fraction of binary white dwarf systems to survive the initial phase of mass transfer. See \\citet{masstransfer}, in particular their figure~11.  If the central spike is thus associated with the accreting white dwarf as we think, it implies a tidal synchronisation timescale that is short enough to make the double-degenerate formation channel for AM CVn stars viable \\citep{nelemans}. It would thereby also increase the viability of the direct-impact accretor scenario for the ultra-compact binaries RX J0806.3+1527 and V407 Vul \\citep{directimpact}."
    },
    "0606/astro-ph0606057_arXiv.txt": {
        "abstract": "{Vibrationally excited emission from the SiO and \\hzo\\ molecules probes the innermost circumstellar envelopes of oxygen-rich red giant and supergiant stars. VY CMa is the most prolific known emission source in these molecules.} {Observations were made to search for % rotational lines in the lowest vibrationally excited state of \\hzo.} % {The APEX telescope was used for observations of \\hzo\\ lines at frequencies around 300 GHz.} {Two vibrationally excited  \\hzo\\ lines were detected, % a third one could not be found. In one of the lines we find evidence for weak maser action, similar to known (sub)millimeter \\nuteo\\ lines. % We find that the other line's intensity % is consistent with thermal excitation by the circumstellar infrared radiation field. Several SiO lines were detected together with the \\hzo\\ lines.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606261_arXiv.txt": {
        "abstract": "We present the results of deep Westerbork Synthesis Radio Telescope observations of neutral hydrogen in 12 nearby elliptical and lenticular galaxies.  The selected objects come from a representative sample of nearby galaxies earlier studied at optical wavelengths with the integral-field spectrograph \\sauron.  They are field galaxies, or (in two cases) located in poor group environments.  We detect \\HI\\ - both in regular discs as well as in clouds and tails off-set from the host galaxy - in 70\\% of the galaxies.  This detection rate is much higher than in previous, shallower single-dish surveys, and is similar to that for the ionised gas.  The results suggest that at faint detection levels the presence of \\HI\\ is a relatively common characteristic of field early-type galaxies,  confirming what was suggested twenty years ago by Jura based on IRAS observations.  The observed total \\HI\\ masses range between a few times $10^6 M_\\odot$ to just over $10^9 M_\\odot$.  The presence of regular disc-like structures is a situation as common as \\HI\\ in offset clouds and tails around early-type galaxies.  All galaxies where \\HI\\ isdetected also contain ionised gas, whereas no \\HI\\ is found around galaxies without ionised gas. Galaxies with regular \\HI\\ discs tend to have strong emission from ionised gas. In these cases, the similar kinematics of the neutral hydrogen and ionised gas suggest that they form one structure. The kinematical axis of the stellar component is nearly always misaligned with respect to that of the gas.  We do not find a clear trend between the presence of \\HI\\ and the global age of the stellar population or the global dynamical characteristics of the galaxies.  More specifically, \\HI\\ detections are uniformly spread through the $(V/\\sigma, \\epsilon)$ diagram. If fast and slow rotators - galaxies with high and low specific angular momentum - represent the relics of different formation paths, this does not appear in the presence and characteristics of the \\HI.  Our observations support the idea that gas accretion is common and does not happen exclusively in peculiar early-type galaxies.  The links observed between the large-scale gas and the characteristics on the nuclear scale (e.g., the presence of kinematically decoupled cores, radio continuum emission etc.), suggest that for the majority of the cases the gas is acquired through merging, but the lack of correlation with the stellar population age suggests that smooth, cold accretion could be an alternative scenario, at least in some galaxies. In either cases, the data suggest that early-type galaxies continue to build their mass up to the present.  \\date{} ",
        "introduction": "\\label{sec:introduction}  The currently favoured paradigm for early-type galaxy formation is the so-called hierarchical formation scenario. It is supported by detailed N-body and hydrodynamical simulations, which are able to reproduce the main stellar morphological characteristics and global scaling relations of early-type galaxies (e.g., De Lucia et al.\\ 2006\\nocite{2006MNRAS.366..499D}). The role and fate of the gas component, however, is still not well-understood. This is partly due to the complexity of understanding gaseous processes (star-formation, energetic feedback and reprocessing), and partly because the gas content of early-type galaxies is often thought to be insignificant. Moreover, the parameter space of gas accretion is large, ranging from merging of two large, equal mass gas-rich objects to infall of a tiny gas-rich companion, and perhaps so-called cold accretion, the slow but long-lasting infall of primordial gas (e.g., Keres et al.\\ 2005\\nocite{2005MNRAS.363....2K}).  Recent observations of early-type galaxies show that gas is clearly present in these objects, and that gas processes may play a more important role in shaping the stellar properties than previously thought. In a galaxy merger event, the amount of gas involved can have a profound effect on the merger remnant, with gas-rich events leading to more disc-like objects (e.g., Bekki \\& Shioya 1997\\nocite{1997ApJ...478L..17B}; but see Burkert \\& Naab 2005\\nocite{2005MNRAS.363..597B}). This may explain the `discy' isophote distortions present in many early-type galaxies, and it has been suggested as an explanation of the apparent dichotomy of fast- and slow-rotating galaxies (i.e.  galaxies with high and low specific angular momentum, Bender, Burstein \\& Faber 1992). Dynamically distinct stellar sub-components are often found in early-type galaxies, and taken as evidence for formation via merging. Connecting a significant star-formation event with such a merger has given mixed results, but there are clearly cases where sub-components of the galaxy are both chemically and kinematically distinct (McDermid et al.\\ 2006), strongly suggesting that external gas has entered the system.  In order to make progress on these issues, high-quality observations of the gas content of early-type galaxies are crucial to allow detailed comparisons with the stellar properties. In the optical, various studies in the recent past have explored the characteristics of the ionised gas in early-type galaxies (Phillips et al.\\ 1986\\nocite{1986AJ.....91.1062P}; Buson et al.\\ 1993\\nocite{A&A...280..409B}; Goudfrooij et al.\\ 1994\\nocite{1994A&AS..105..341G}) and found the presence of gas with complex kinematics (e.g., counter-rotating with respect to the stellar component, Bertola et al.\\ 1992\\nocite{1992ApJ...401L..79B}). However, so far the kinematics and ionisation of the gas in early-type galaxies have been studied mostly through long-slit observations, usually along one or two position angles, which limit the correct determination of the morphology and dynamical structure of the ionised gas. The recent systematic survey based on observations with the panoramic integral-field spectrograph \\sauron\\ shows that these objects display a variety of line-strength distributions and kinematic structures which is richer than often assumed (Bacon et al.\\ 2001\\nocite{2001MNRAS.326...23B}; de Zeeuw et al.\\ 2002\\nocite{2002MNRAS.329..513D}). The survey includes 48 representative nearby early-type galaxies classified as E or S0 in the RC3 (de Vaucouleurs et al.\\ 1991\\nocite{1991trcb.book.....D}). Many examples of minor axis rotation, decoupled cores, embedded metal-rich stellar discs, as well as non-axisymmetric and counter-rotating gaseous discs, have been found (Emsellem et al.\\ 2004\\nocite{2004MNRAS.352..721E}; Sarzi et al.\\ 2006\\nocite{2006MNRAS.366.1151S}; Kuntschner et al.\\ 2006\\nocite{keta06}).  At radio wavelengths, our knowledge about the neutral hydrogen content of early-type galaxies is also changing. This is partly due to the growing number of cases where \\HI\\ has been imaged --- instead of using only single-dish data --- and information about the morphology and the detailed kinematics of the gas is now available.  Many \\HI-rich early-type galaxies are now known (e.g. Schiminovich et al.\\ 1995\\nocite{1995ApJ...444L..77S}; van Gorkom \\& Schiminovich 1997\\nocite{1997ASPC..116..310V}; Morganti et al.\\ 1997\\nocite{1997AJ....113..937M}; Sadler et al.\\ 2000\\nocite{2000AJ....119.1180S}; Balcells et al.\\ 2001\\nocite{2001AJ....122.1758B}; Oosterloo et al.\\ 2002\\nocite{2002AJ....123..729O} and refs therein; Oosterloo et al.\\ 2004\\nocite{2004IAUS..217..486O}, 2005). The large amount of neutral hydrogen detected around some of these galaxies (up to more than $10^{10} M_\\odot$) is often distributed in huge (up to 200 kpc in size) regularly rotating discs or rings. The flat rotation curves of the discs indicate the existence of large halos of dark matter.  Given their regular appearance, these discs must be relatively old (several $\\times 10^9$ yr).  Although an external origin of the \\HI\\ has been suggested already in several earlier studies (e.g. Knapp et al.\\ 1985\\nocite{1985AJ.....90..454K}), the imaging of the kinematics of such structures opens the possibility of studying in detail how these galaxies have formed. In particular, although major mergers and small-companion accretions are clearly at the origin of some of the \\HI\\ structures observed (e.g. Serra et al.\\ 2006\\nocite{2006astro.ph..2621S}), recent work has shown that smooth cold gas accretion (e.g., Macci\\'o, Moore \\& Stadel 2006\\nocite{2006ApJ...636..25M}) can also play an important role and should, therefore, be taken into account.  The shallow \\HI\\ surveys available so far are, however, able to study only the most extreme \\HI-rich early-type galaxies. Much deeper observations are needed to explore the complete \\HI\\ mass distribution of early-type galaxies.  Furthermore, the study of \\HI\\ in these systems has lacked the important combination of having both the \\HI\\ data and multi-slit or integral-field optical spectroscopy available for a significant number of objects. For these reasons, we have performed deep \\HI\\ observations, using the recently upgraded Westerbork Synthesis Radio Telescope (WSRT), of a sub-sample of E and S0 galaxies in the \\sauron\\ representative survey.  We present the results here. We describe the sample selection and the WSRT observations in Section~\\ref{sec:observations}. In Section~\\ref{sec:results} we discuss the \\HI\\ maps, and we compare with earlier \\HI\\ surveys in Section~\\ref{sec:comparison_hi}. A discussion of the relation between the presence of \\HI\\ and of a radio loud AGN is given in Section~\\ref{sec:comparison_agn}. We investigate the relation with the characteristics of the stellar component and the ionised gas in Section~\\ref{sec:comparison_sauron}, and comment on the origin of the neutral gas in Section~\\ref{sec:origin}.  We summarise our conclusions in Section~\\ref{sec:conclusions}. In Appendix~\\ref{appendix:a} we report the serendipitous discovery of a megamaser in the field of NGC 4150. ",
        "conclusions": "\\label{sec:conclusions}  The main result of these observations is that, in terms of the neutral gas, {\\sl the class of early-type galaxies is extremely rich and varied and that galaxies that {\\it prima facie} appear very similar show subtle but important differences}. Our detection rate of \\HI\\ ($\\sim 70$\\%) is comparable to that ionised gas (75\\%, Sarzi et al.\\ 2006\\nocite{2005astro.ph.11307S}). This is surprisingly high compared with earlier studies, likely the consequence of the fact that our \\ HI\\/ data are a factor 50-100 deeper. As with the ionised gas, we find a variety of structures, including irregular distributions (likely the remnant of a recent accretion), small polar rings, strongly-warped (up to 90$^\\circ$) structures, and extended re gular discs (in some cases containing as much \\HI\\/ as the Milky Way!). The peak column density in these discs is low - below (and sometimes well below) a few times $10^{20}$ cm$^{-2}$ - so no significant star formation is occurring, implying that these d isks may survive for a very long time.  A very clear correlation exists between the presence and the properties of ionised and neutral gas. They share the same kinematics, showing that they are two phases of the same structure. On the other hand, the correlation of the gas kinematics (neutral and ionised) with that of the stars is quite poor. Most interestingly, we find that neither the amount of gas, nor its kinematics, correlate clearly with the stellar dynamics: the galaxies with slow stellar rotation (likely to be `true' elliptical galaxies) are all detected in \\HI, with some containing regular gas discs. Conversely, many of the fast-rotating galaxies (likely S0-like galaxies) are not detected. This range of gas properties does not intuitively reflect the structural differences between elliptical and S0 galaxies, but suggests that the relationship between the gaseous and stellar components is more complex than previously thought.  Moreover, the amount of neutral gas {\\sl does not} appear to correlate with the stellar population characteristics. In a few galaxies, large amounts of \\HI\\ are detected while the luminosity-weighted stellar population is purely old, showing no evidence of young stars that may have formed from the gas. Conversely, there are galaxies with a young luminosity-weighted age where no gas is detected. This is contrary to the commonly held idea that accretion and merging trigger star formation in the central regions. This suggests that other modes of accretion also exist, and may hint that `cold accretion' does occur in some systems.  In the galaxies detected in \\HI\\ we also detect radio  continuum radiation which may be from a small AGN or result from star formation. In the \\HI\\ non-detections, no such source is detected. This suggests that whatever process brings the neutral gas to early-type galaxies, it can trigger some activity/star formation in the central regions.   Our survey indicates that the presence of neutral gas is common in field early-type galaxies, if sufficiently deep observations are available. This is further supported by the fact that the observed galaxies were not selected based on any peculiarity and are relatively regular objects.  Our results point to an external origin of the gas, suggesting that gas accretion is common and does not happen only in peculiar early-type galaxies. This may be connected to the high incidence of small-scale kinematical features in the \\sauron\\ data. This accretion can happen in many different ways, and can involve a very large range of masses. Also in spiral galaxies, gas accretion events are found to be very common over their lifetime (Naab \\& Ostriker 2006\\nocite{2006MNRAS.366..899N}, van der Hulst \\& Sancisi 2004). From our study it appears that the distribution, more than the amount, of \\HI\\ is an important element in determining the morphology of the galaxy. Even in the very \\HI-rich early-type galaxies, the neutral hydrogen is always distributed over very large areas therefore it has always a very low column density, too low for star formation to occur.  Our current sample includes only 12 galaxies restricted to low density environments, and the statistical validity of our results is therefore limited.  Given the links that seem to be present between the characteristics of the neutral and ionised gas components, it is important to re-examine also the cluster environment to fully explore the effect/importance that the environment has on our results. In cluster galaxies, the fraction of ionised gas detections --- derived from the \\sauron\\ study --- is about 55\\%, and therefore, if the relation holds, a significant number of galaxies may indeed show neutral hydrogen even in this more hostile environment."
    },
    "0606/astro-ph0606275_arXiv.txt": {
        "abstract": "We study the nature of the [\\ti2] and [\\ni2] emission from the so-called {\\it strontium filament} found in the ejecta of \\etacar. To this purpose we employ multilevel models of the \\ti2 and \\ni2 systems which are used to investigate the physical condition of the filament and the excitation mechanisms of the observed lines. For the \\ti2 ion, for which no atomic data was previously available, we carry out {\\it ab initio} calculations of radiative transition rates and electron impact excitation rate coefficients. It is found that the observed spectrum is consistent with the lines being excited in a mostly neutral region with an electron density of the order of $10^7$ cm$^{-3}$ and a temperature around  6000~K. In analyzing three observations with different slit orientations recorded between March~2000 and November~2001 we find line ratios that change among various observations, in a way consistent with changes of up to an order of magnitude in the strength of the continuum radiation field. These changes result from different samplings of the extended filament, due to the different slit orientations used for each observation, and yield  clues on the spatial extent and optical depth of the filament. The observed emission indicates a large Ti/Ni abundance ratio relative to solar abundances. It is suggested that the observed high Ti/Ni ratio in gas is caused by dust-gas fractionation processes and does not reflect the absolute Ti/Ni ratio in the ejecta of \\etacar. We study the condensation chemistry of Ti, Ni and Fe within the filament and suggest that the observed gas phase overabundance of Ti is likely the result of selective photo-evaporation of Ti-bearing grains. Some mechanisms for such a scenario are proposed. ",
        "introduction": "\\etacar is the most luminous star in our part of the Galaxy and indeed is one of the most mysterious known to us (Davidson \\& Humphreys 1997).  It is surrounded by complex ejecta created in several outbursts of the star and an outgoing massive wind with periodic events every 5.54 years (Damineli 1996). The largest mass ejection occurred during the so-called Great Eruption observed in the 1840s -- which the star survived.  In that event, \\etacar released about 10$^{49.6}$ ergs of kinetic energy contained in more than 10 M$_{\\odot}$ of material (Smith et al.\\ 2003), which has since expanded to form the bipolar Homunculus nebula we see today (Gaviola 1950; Smith \\& Gehrz 1998; Morse et al.\\ 2001).  Despite the hot and very luminous central star, the circumstellar nebula is cool enough to produce a remarkably rich emission spectrum of lines from neutral and low ionization species like Fe$^+$ and Ni$^+$ (Davidson et al.\\ 2001; Zethson 2001; Zethson \\etal 2001), line-of-sight absorption from several neutral and singly ionized atoms (Gull et al.\\ 2005), and even emission from H$_2$ and very cool dust (Smith 2002a; Smith et al.\\ 2003).  In addition to being dense and cool, the ejecta of \\etacar are nitrogen rich, representing ashes of the CNO cycle (Davidson et al.\\ 1982, 1986; Smith \\& Morse 2004; Verner et al.\\ 2005).  It is the high density, low-excitation, and CNO abundances of the ejecta that make the Homunculus a unique laboratory among circumstellar nebulae, having a profound effect on the results discussed in this paper.  In February 1999 a series of STIS observations of \\etacar was made of a region about 1.5\" northwest of the star (Zethson \\etal 2001). The spectrum from this region was significantly different from the nebular emission spectra elsewhere in the Homunculus, with two of the lines being identified as forbidden emission of singly-ionized strontium ([\\sr2]). The measured velocity of the filament ($\\sim -100$ km/s) and its distance from the central star are consistent with the gas having been ejected from the same star during the Great Eruption.  Further spectroscopic observations and a complete list of lines identified in the filament were reported by Hartman et al. (2004). They show that lines of Fe~{\\sc i}, V~II, and \\ti2  are strong in the filament, whereas only a few \\fe2 lines  are detected. Other lines in the spectrum at about the same velocity come from C~{\\sc i}, Mg~{\\sc i}, Al~{\\sc ii}, Ca~{\\sc i}, Ca~{\\sc ii}, Sc~{\\sc ii}, Cr~{\\sc ii}, Mn~{\\sc ii}, Co~{\\sc ii}, and \\ni2. Curiously no H, He, S, O, or N lines are present in this region.  The cosmic abundance of strontium is lower than that of iron and nickel by factors of 30,000 and 3,000 respectively. Thus, strontium emission is normally undetectable in nebular spectra. The fact that [\\sr2] emission was detected with strengths comparable to those of [\\fe2] and [\\ni2] lines indicates that either strontium is locally overabundant in \\etacar by orders of magnitude, which is inconsistent with the evidence of CNO-cycle-processes ejecta,  or  very unusual conditions exist in this region of the nebula.  In a recent study of the \\sr2  and [\\sr2] spectra  of the filament (Bautista \\etal 2002) we found that [\\sr2] emission may be strongly enhanced by continuum fluorescence, which partially explains the strength of these lines. Although it was not possible at that time to establish the precise  abundance of Sr, this result did imply that Sr might not be as overabundant as previously thought. This point opened up new questions. If [\\sr2]  lines can be efficiently excited by fluorescence, why are the lines seen only in this structure, but not elsewhere in the Homunculus, and not in any other emission line nebulae?  The observed spatial  and velocity distribution of the Sr-filament suggests that it is located in the equatorial region between the hourglass-shaped bipolar Homunculus (Davidson et al 2001; Smith 2002a; Ishibashi et al, 2003). The Sr-filament is located at a projected distance of 10 lightdays (at 2300 pc distance), where it should receive intense UV radiation from \\etacar. The spectra of the Weigelt Blobs B and D, located a few light days from \\etacar, require the spectral flux distribution of an O star with Te=35,000K during the 5-year broad maximum (Verner \\etal 2005). The Sr-filament should receive a diluted amount of this same UV radiation, but its emission line spectrum indicates a much lower ionization condition. The Sr-filament appears to be shielded not only by neutral hydrogen (13.6 eV), but also singly-ionized iron (7.9 eV) which requires a substantial column density of \\fe2 between \\etacar and the Sr-filament. One possibility is that the Sr-filament is located at the outer side of an apparent optically thick, dusty, torus seen in the mid-IR (Smith et al. 2003). In order to demonstrate this scenario, we must determine the nature of a hot star continuum filtered out by an \\fe2  blocking column and predict the physical conditions within the filament that would sustain such a low ionization region. Then we can determine the actual abundances of all species observed in this structure.  In this paper we study the emission spectra of two prominent ions in the filament: \\ti2  and \\ni2. \\ti2  is studied for the very first time as the relevant atomic data has never before been obtained. \\ti2  emission lines are rarely seen in emission nebulae as its ionization potential, 13.58 eV, leads to Ti~{\\sc iii} in most H~II regions. \\ni2, often seen in nebular emission spectroscopy, is much better understood and herein serves as a reference ion.  The remainder of the paper is organized as follows: In Section 2 we present the atomic models built for the \\ni2 and \\ti2 systems, with particular detail in the calculations of atomic parameters that had to be performed for \\ti2. In Section~3 we show the results of the nebular diagnostics. In Section~4 we derive the relative abundances of \\ti2 and \\ni2 in gas phase, while in Section~5 we discuss possible scenarios of Ni and Ti depletion in the nebula. Finally, our discussion and conclusions are given in Section~6. ",
        "conclusions": "We have studied the \\ni2 and \\ti2 spectra of the Sr-filament in \\etacar, as taken in three different epochs with the HST. In doing this we build spectral models for the two ions using the best published data for \\ni2, while we carry out extensive calculations of radiative and electron impact rates for \\ti2.  The spectral models were employed for diagnosing the physical conditions of the emitting region. The results are generally consistent with previous determinations using [\\sr2] and \\sr2 lines. Moreover, the [\\ti2] lines allow us to constrain the conditions for observations with the aperture across the \\etacar's central source to $N_e\\approx 10^7$ \\cm3, $T_e=6000\\pm1000$~K, and dilution factor $log_{10}w=-9.0\\pm0.5$ for an assumed blackbody continuum radiation with temperature of 35~000~K. For observations with the slit perpendicular to the radial direction the dilution factor seems to be somewhat lower.  Study of the observed [\\ti2] emission reveals a very large overabundance of Ti in gas phase with respect to Ni and Fe relative to the cosmic values.  We study the chemistry of Ti, Ni, and Fe in CNO-cycled gas and find that Ti condenses as TiC, or alternatively as TiN, while Ni and Fe are locked into (Ni,Fe)Si. Moreover, the TiC condenses at higher temperatures than Ni- and Fe-bearings. Later, we suggest that the TiC would be selectively destroyed by the stellar UV radiation leading to a relative overabundance of Ti in gas phase.  This explanation for the unusual Ti/Ni ratios in the Sr-filament requires that significant dust and molecules had formed in this material shortly after the Great Eruption, but have since been largely photoevporated here.  In this context, it is very interesting that near-IR emission from molecular hydrogen is seen throughout the polar lobes of the Homunculus Nebula (Smith 2002a), but this same H$_2$ emission is not seen at comparable levels from the equatorial skirt in the vicinity of the Sr-filament (Smith 2002a, 2004).  Today, the gas densities we infer in the Sr filament ($\\sim$10$^7$ cm$^{-3}$) are very high compared to many other circumstellar nebulae around massive stars, and are similar to the extremely high densities seen in the walls of the Homunculus (e.g., Gull et al. 2005; Nielsen et al. 2005).  A weaker radiation field at larger radii is not the only explanation for this difference: H$_2$ emission is seen both in the polar caps of the Homunculus at larger radii from the star than the Sr-filament, and in the side walls of the polar lobes near the equator, even {\\it closer} to the star than the Sr filament (Smith 2004).  In the past when the Sr-filament ejecta were even denser, it is conceivable that H$_2$ or other molecules may have existed, but have since been destroyed.  If true, this would seem to support the scenario of unusual chemistry resulting from selective photoevaporation and CNO processed ejecta, as we have suggested here."
    },
    "0606/astro-ph0606569_arXiv.txt": {
        "abstract": "We investigate the viscosity driven instability in rotating relativistic stars by means of an iterative approach.  We focus on polytropic rotating equilibrium stars and impose an $m=2$ perturbation in the lapse.  We vary both the stiffness of the equation of state and the compactness of the star to study those effects on the value of the threshold.  For a uniformly rotating star, the criterion $T/W$, where $T$ is the rotational kinetic energy and $W$ is the gravitational binding energy, mainly depends on the compactness of the star and takes values around $0.13 \\sim 0.16$, which differ slightly from that of Newtonian incompressible stars ($\\sim 0.14$).  For differentially rotating stars, the critical value of $T/W$ is found to span the range $0.17 - 0.25$.  This is significantly larger than the uniformly rotating case with the same compactness of the star. Finally we discuss a possibility of detecting gravitational waves from viscosity driven instability with ground-based interferometers. ",
        "introduction": "Stars in nature are usually rotating and subject to nonaxisymmetric rotational instabilities (see \\citep{Sterg03}, \\citep{Ander03}, \\citep{AnderC06} or \\citep{Villa06} for recent reviews).  An exact treatment of these instabilities exists only for incompressible equilibrium fluids in Newtonian gravity, (e.g. \\citep{Chandra69,Tassoul78,ST83}).  For these configurations, global rotational instabilities arise from non-radial toroidal modes $e^{im\\varphi}$ ($m=\\pm 1,\\pm 2, \\dots$) when $\\beta \\equiv T/W$ exceeds a certain critical value. Here $\\varphi$ is the azimuthal coordinate and $T$ and $W$ are the rotational kinetic and gravitational binding energies.  In the following we will focus on the $m=\\pm 2$ bar-mode, since it is the fastest growing mode when the rotation is sufficiently rapid.  There exist two different mechanisms and corresponding timescales for bar-mode instabilities.  Uniformly rotating, incompressible stars in Newtonian theory are {\\em secularly} unstable to bar-mode formation when $\\beta \\gtrsim \\beta_{\\rm sec} \\simeq 0.14$.  This instability can grow in the presence of some dissipative mechanism, like viscosity or gravitational radiation, and the growth time is determined by the dissipative timescale, which is usually much longer than the dynamical timescale of the system.  By contrast, a {\\em dynamical} instability to bar-mode formation sets in when $\\beta \\gtrsim \\beta_{\\rm dyn} \\simeq 0.27$.  This instability is independent of any dissipative mechanism, and the growth time is the hydrodynamic timescale of the system.  There are two representative dissipative mechanisms that drive the secular bar mode instability, viscosity and the gravitational radiation, in the absence of thermal dissipation.  The viscosity driven instability sets in when a mode has a zero-frequency in the frame rotating with the star \\citep{RS63}, and the first unstable mode in terms of $m$ is the $m=2$ bar mode.  The quasi-static evolution of the star due to viscosity driven instability, which varies the circulation of the star, deforms the Maclaurin spheroid to the Jacobi ellipsoid in Newtonian incompressible stars.  On the other hand, the gravitational radiation induced instability \\citep{Chandra70, FS78} sets in when the backward going mode is dragged forward in the inertial frame (see \\citep{Sterg03,AnderC06,Villa06,Schutz87} for reviews), and the mode in terms of $m$ are all unstable when it exceeds a certain $m$.  The quasi-static evolution of the star due to gravitational radiation induced instability, which varies the angular momentum of the star, deforms the Maclaurin spheroid to the Dedekind ellipsoid in Newtonian incompressible stars.  The viscosity driven instability, especially to determine the critical value either in an incompressible star or in an ellipsoidal equilibrium, has been studied in Newtonian gravity \\citep{IM81}, in post Newtonian gravity \\citep{SZ98, GV02}, and in full general relativity \\citep{BFG96, BFG98, GG02} by an ellipsoidal approximation (e.g. \\citep{Chandra69}) or by an iterative evolution approximation (e.g. \\citep{BFG96}), and shows that the viscosity drives the instability to higher rotation rates $\\beta_{\\rm sec} \\gtrsim 0.14$ as the configurations become more compact.  There is also a study of Newtonian compressible stars to determine the critical value of viscosity driven instability \\citep{BFG96}: It is found that the star becomes secularly unstable at $\\beta_{\\rm sec} \\approx 0.135 \\pm 0.02$, depending on the stiffness of the polytropic equation of state.  The aim of the paper is twofold.  One is to investigate the critical value of viscosity driven instability in the compressible stars rotating uniformly.  The argument that the viscosity driven instability plays a role to deform a star from a Maclaurin spheroid to a Jacobi ellipsoid is only true in the absence of internal energy, since the total energy has a chance to transfer it to the internal energy without emission from the star (e.g. \\citep{Shapiro04}).  Here we assume that the cooling timescale of the star is shorter than the thermal heating timescale so that the thermal energy generated by viscosity is immediately radiated away. Therefore the picture of the deformation process due to viscosity is quite similar to the case of incompressible stars.  The stars are usually considered as compressible bodies, and therefore it is worth to take such effect into account whether there is a significant change on the threshold.  In this respect the present work extends that of Ref.~\\citep{GG02} to the compressible case.  The other aim of the present work is to investigate the effect of differential rotation on the viscosity driven secular bar mode instabilities.  For a strong viscosity or strong magnetic field circumstances, the star maintains uniform rotation.  However, in nature the star may rotate differentially as the Sun.  Stellar collapses and mergers may also lead to differentially rotating stars (e.g. \\citep{Ott04}).  For the coalescence of binary irrotational neutron stars \\citep{SU00,SU02,STU03}, the presence of differential rotation may temporarily stabilize the ``hypermassive'' remnant, which constructs a differential rotation.  Therefore it is also worth to take differential rotation into account to study viscosity driven instabilities in rotating relativistic stars.  This paper is organized as follows.  In Sec.~\\ref{sec:bequation} we present the basic equations of our treatment of general relativity. We discuss our numerical results in Sec.~\\ref{sec:rigid} and in Sec.~\\ref{sec:diff}, focusing on the viscosity driven instability in uniformly and differentially rotating stars.  In Sec.~\\ref{sec:Discussion} we briefly summarize our findings. Throughout this paper, we use the geometrized units with $G=c=1$ and adopt polar coordinates $(r,\\theta,\\varphi)$ with the coordinate time $t$.  Note that Latin index takes $(r,\\theta,\\varphi)$. ",
        "conclusions": "\\label{sec:Discussion} We have studied the viscosity driven instability in both uniformly and differentially rotating polytropic star by means of iterative evolution approach in general relativity. We have focussed on the determination of the critical value of viscosity driven instability.  We find that relativistic gravitation stabilizes the star from the viscosity driven instability, with respect to Newtonian gravitation.  Also the critical value is not sensitive to the stiffness of polytropic equation of state for a given compactness of the star.  In a previous study devoted to compressible stars, \\citet{BFG98} showed that relativistic gravitation does stabilize the uniformly rotating polytropic stars by investigating the mass-shedding sequence.  Our study shows the concrete value of $(T/W)_{\\rm crt}$ in the case of a uniformly rotating star. Also we have improved the numerical code with respect to the study \\citep{BFG98} by making use of surface-fitted coordinates.  Besides we have found that differential rotation also stabilizes the star from the viscosity driven instability.  If we fix the compactness of the star, we find a significant increase of the critical value of $T/W$, which supports the above statement.  We also confirm this statement by changing the rotation profile due to viscosity and find that the differential rotation still stabilizes the threshold of viscosity driven instability.  However to confirm the statement in differentially rotating stars, the study of other approaches such as implicit hydrodynamical evolution or eigenmode analysis should be necessary and helpful.  Finally let us mention the characteristic amplitude and frequency of gravitational waves emitted throughout the viscosity driven secular bar mode instability, which produces quasi-periodic gravitational waves detectable in ground based interferometers.  The characteristic amplitude $h$ of gravitational waves estimated from the evolution of a Jacobi-ellipsoid to a Maclaurin spheroid is (Eq.~[4.2] of Ref.~\\citep{LS95}) \\begin{widetext} \\begin{equation} h \\approx 9.1 \\times 10^{-21} \\left( \\frac{30 {\\rm Mpc}}{d} \\right) \\left( \\frac{M}{1.4 M_{\\odot}} \\right)^{3/4} \\left( \\frac{R}{10 {\\rm km}} \\right)^{1/4} f^{-1/5}, \\end{equation} \\end{widetext} where $d$ is the distance to the source and the characteristic frequency $f = \\Omega / \\pi$ is $f \\gtrsim 1000 {\\rm [Hz]}$, depending on $T/W$ of the star.  Note that the frequency increases throughout this process.  Although the frequency regime of the source is slightly higher than the best sensitivity regime of the ground based detectors to follow all the deformation process, we may have a chance to detect them when it happens in the Virgo cluster, for example."
    },
    "0606/astro-ph0606043_arXiv.txt": {
        "abstract": "We present $\\sim3''$ resolution imaging of the $z=4.7$ QSO BR1202-0725 at 900~$\\micron$ from the Submillimeter Array. The two submillimeter continuum components are clearly resolved from each other, and the positions are consistent with previous lower frequency images. In addition, we detect [\\ion{C}{2}] line emission from the northern component. The ratio of [\\ion{C}{2}] to far-infrared luminosity is 0.04\\% for the northern component, and an upper limit of $< 0.03\\%$ is obtained for the southern component. These ratios are similar to the low values found in local ultraluminous galaxies, indicating that the excitation conditions are different from those found in local field galaxies. X-ray emission is detected by {\\it Chandra} from the southern component at L$_{0.5-2keV}=3\\times10^{45}$~erg~s$^{-1}$, and detected at 99.6\\% confidence from the northern component at L$_{0.5-2keV}\\sim$3$\\times$10$^{44}$erg~s$^{-1}$, supporting the idea that BR1202-0725 is {\\em a pair} of interacting galaxies at $z=4.7$ that each harbor an active nucleus. ",
        "introduction": "The far infrared (FIR) fine structure lines of abundant species (e.g. [\\ion{C}{1}], [\\ion{C}{2}], [\\ion{O}{1}]) are recognized as one of the most important coolants in the warm ISM. In particular, the emission from singly ionized carbon (i.e. [\\ion{C}{2}], $\\lambda_{rest} = 158 \\micron$) is known to trace warm ($\\geq 200$ K) and dense ($n_{cr} = 3 \\times 10^3$~cm$^{-3}$) photo-dissociation regions \\citep[PDRs; ][]{kaufman99}, directly identifying the gas surrounding active star forming regions. The [\\ion{C}{2}] line is the dominant coolant in the Galactic ISM, accounting for 0.1 -- 1 \\% of the FIR luminosity \\citep{hollenbach97}. The  brightness of the line makes it an attractive tracer of star formation in galaxies, and a relation holds between [\\ion{C}{2}] luminosity and the massive star formation rate derived from H$\\alpha$ emission for galaxies that have L$_{\\rm FIR}$~$< 3 \\times 10^{10}$~L$_{\\odot}$ \\citep{boselli02}. In the FIR luminous galaxies (ULIRGs), the [\\ion{C}{2}]-to-FIR flux ratio is an order of magnitude smaller than in normal star forming galaxies \\citep{luhman98,luhman03}. The reason for the deficiency in  the [\\ion{C}{2}] line luminosity is not well understood, but several possible scenarios have been offered \\citep{luhman98, malhotra01,pierini03, luhman03}.   Despite the large line luminosity expected, detecting [\\ion{C}{2}] line emission from luminous high redshift sources using a single-dish submillimeter telescope has proved difficult \\citep[e.g.][]{isaak94,vanderwerf98,bollato04,marsden05}. The low [\\ion{C}{2}]-to-FIR ratio suggests that the physical properties in high-$z$ sources resemble those of the local ULIRG population \\citep[e.g.][]{blain02}. Detection of the \\CII\\ line in the $z=6.4$ QSO J1148+5251 has been reported recently by \\citet[][]{maiolino05}. Here we report a first detection of the [\\ion{C}{2}] line emission from the $z=4.7$ hyper-luminous ($\\sim 10^{13}$~L$_{\\odot}$) QSO BR1202-0725 observed using the Submillimeter Array \\citep[SMA; ][]{ho04}\\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.}.  BR1202-0725 is an optically bright QSO ($M_B = -28.5$) with a $2\\farcs3$ Lyman~$\\alpha$ extension toward the northwest \\citep{hu96, petitjean96, fontana98, ohyama04}. Radio continuum emission was detected at 1.4, 4.9 and 43~GHz \\citep{yun00,carilli02}.  The emission at 1.4~GHz was resolved into two components; one at the location of the QSO and the other $\\sim4''$ northwest.  This apparent companion is also seen in millimeter continuum emission \\citep{omont96,guilloteau99}. Abundant CO emission  (M$_{H_2} \\geq 10^{10}$~M$_{\\odot}$) was detected toward both the optical QSO and the northwestern companion \\citep{omont96,ohta96, guilloteau01,carilli02}.   The CO-derived redshifts are $z=4.6947$ and 4.6916 for the QSO and the northwestern companion, respectively \\citep{omont96}.    We adopt $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, which gives $D_L = 43.27$~Gpc ($1''$ = 6.5~kpc) for the luminosity distance to BR1202-0725. ",
        "conclusions": "\\subsection{Properties of the [\\ion{C}{2}] Line Emission}  The exceedingly low $\\rm L_{[CII]}/L_{FIR}$ ratios for both BR1202N and BR1202S are consistent with the similarly low ratios found in local ULIRGs where the average was found to be $\\rm L_{[CII]}/L_{FIR} = 0.04\\%$ \\citep{luhman98,luhman03}. This is in stark contrast to the values found in the local galaxies by \\citet{malhotra01} in which $\\sim 70\\%$ of the sources show $\\rm L_{[CII]}/L_{FIR} > 0.2\\%$, with the majority in the range $0.1 - 1\\%$. Past studies have reported $\\rm L_{[CII]}/L_{FIR} \\sim 0.02\\%$ for the $z=6.42$ QSO SDSS~J1148+5251 \\citep{maiolino05}, and $\\rm L_{[CII]}/L_{FIR} < 0.4\\%$ in a $z=4.926$ QSO CL~1358+62 \\citep{marsden05}. Our new measurement suggests that the [\\ion{C}{2}]-to-FIR ratio (0.04\\%) in the high-$z$ hyper-luminous source BR1202-0725 is comparable to local ULIRGs. Table~\\ref{ratios} summarizes these results.    The integrated [\\ion{C}{2}] line intensity is known to be a tracer of star formation activity in normal galaxies. We derive the star formation rate (SFR) of 2900~$\\rm M_{\\odot}~yr^{-1}$ using the calibration of \\citet{maiolino05}. This is 1 -- 2 orders of magnitude higher than the SFRs derived from optical recombination lines observed near BR1202N \\citep[i.e. 10 -- 230 M$_\\odot$~yr$^{-1}$;][]{hu96,petitjean96, ohta00,ohyama04}, but consistent with the SFR derived from the SED fit to the millimeter/submillimeter continuum measurements (i.e. SFR = 2000 $\\rm M_{\\odot}~yr^{-1}$ in BR1202N; see \\S 4.2).   Another important diagnostic of the physical properties is the [\\ion{C}{2}]-CO flux ratio. Early studies suggest the close spatial and kinematical association between the [\\ion{C}{2}] and CO~(1--0) emissions in nearby galaxies, suggesting that this ratio is a good tracer of star formation activity that is independent of the beam filling factor \\citep{crawford85,stacey91}. Assuming the luminosity ratio between the CO~(5--4) and CO~(1--0) lines in BR1202N is similar to that of the Cloverleaf \\citep{barvainis94,barvainis97}, the CO~(1--0) luminosity is estimated to be L$_{\\rm CO(1-0)} \\sim 10^{6}$~L$_{\\odot}$, yielding L$_{\\rm [CII]}$/L$_{\\rm CO(1-0)} \\sim 4500$. This value is highly uncertain since the CO~(5--4) -- CO~(1--0) ratio may be different in BR1202-0725.  Nevertheless, this is a factor of 2.5 larger than the average values of ULIRGs \\citep[i.e. L$_{\\rm [CII]}$/L$_{\\rm CO(1-0)} \\sim$ 1700;][]{vanderwerf98}, but comparable to the average values of normal galaxies \\citep[i.e. L$_{\\rm [CII]}$/L$_{\\rm CO(1-0)}$ = 4400 -- 6300;][]{crawford85, stacey91}. The higher observed ratio may suggest a factor of few higher G$_0$/n ratio (where G$_0$ and n are the far UV flux and the cloud density respectively) in BR1202N compared to the local ULIRGs \\citep[see Figure~9 of][]{kaufman99}.   \\subsection{Properties of the Continuum Emission}  The spatially resolved submillimeter images allow us to investigate the properties of the individual components in BR1202-0725. The radio to FIR SED of BR1202-0725 is shown in Figure~\\ref{fig2}, where the individual flux densities are given for BR1202N and BR1202S for the 1.4~GHz \\citep{carilli02}, 1.3~mm \\citep{omont96}  and $900\\micron$ observations (see figure caption for the adopted parameters).  The SEDs of the two sources are analyzed using the model by \\citet{yun02}, yielding L$_{\\rm FIR} = 1.2 \\times 10^{13}$ and $2.6 \\times 10^{13}$~L$_{\\odot}$, and SFR = 2000 and 4500 $\\rm M_{\\odot}~yr^{-1}$ for BR1202N and BR1202S respectively. The discrepancy between the SFR derived by us and \\citet{yun02}(2300 $\\rm M_{\\odot}~yr^{-1}$ for the entire BR1202-0725) is due to the difference in the adopted cosmology. Note that while the millimeter/submillimeter continuum flux densities of BR1202N and BR1202S are comparable (\\S3), the radio continuum flux density of BR1202N is 3 times larger. This suggests a presence of a highly obscured AGN in BR1202N (see \\S4.3), or a strong effect of a highly energetic jet from BR1202S \\citep{klamer04}.  Assuming a dust absorption coefficient of $\\kappa_d = 1.5$~m$^2$~kg$^{-1}$ at $160\\micron$ \\citep{hildebrand83,draine84}, the dust mass of the individual components is estimated to be M$_{\\rm d}~(\\rm BR1202N) = (5.4 \\pm 0.8) \\times 10^8$~M$_\\odot$ and M$_{\\rm d}~(\\rm BR1202S)= (9.2 \\pm 1.1) \\times 10^8$~M$_\\odot$. Using the H$_2$ mass derived from the CO~(5-4) emission \\citep{omont96}, we find M$_{\\rm H_2}$/M$_{\\rm d} \\sim 50 \\pm 10$ for both components. This ratio is an order of magnitude lower than the mean value for LIRGs/ULIRGs \\citep[M$_{\\rm H_2}$/M$_{\\rm d} = 540 \\pm 290$;][]{sanders91}, nearby spiral galaxies\\footnote{The CO to H$_2$ conversion factor used by \\citet{omont96}, \\citet{sanders91} and \\citet{devereux90} are all consistent to within 20\\%. In addition, the rest frame $160\\micron$ emission for BR1202-0725 traces the same warm dust (and hence the same M$_d$) as measured by {\\it IRAS} \\citep{yun06}.} \\citep[M$_{\\rm H_2}$/M$_{\\rm d} \\sim 500$; see Fig. 2b of ][]{devereux90} and a sample of {\\it Spitzer} 160 $\\micron$ selected galaxies \\citep[M$_{\\rm H_2}$/M$_{\\rm d} = 330$;][]{yun06}. These gas-to-dust ratios are highly uncertain because dust emissivity, temperature, and CO-to-H$_2$ conversion factor are all poorly constrained. If the low ratio for BR1202-0725 is correct, however, it may indicate that the H$_2$ mass is significantly underestimated in BR1202-0725, and/or the warm component of dust dominates the far infrared emission, suggesting that the entire ISM is involved in the starburst activity.   \\subsection{BR1202-0725: A Colliding Galaxy System at $z=4.7$} Our analysis of previously unpublished archival {\\em Chandra} data (ObsID 3025) clearly shows that BR1202S is an X-ray source and gives 26.4 net counts in 9.64~ksec. The derived flux from BR1202S is $f_{0.5-2keV} = 6\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$  which yields a luminosity of L$_{0.5-2keV}=3\\times10^{45}$~erg~s$^{-1}$ \\citep[assuming $\\Gamma$ = 2.0 and Galactic foreground N$_H$ = $3.35\\times10^{20}$cm$^{-2}$;][]{stark92}. X-ray emission is also seen within $1''$ of the position of BR1202N. With just 2 counts in a single pixel, this detection seems marginal, but the local background of 0.03 counts/pixel implies that the likelihood of a random occurrence is only 0.04\\%. The derived flux from BR1202N is $f_{0.5-2keV} = 4.8^{+6.5}_{-4.3}\\times$10$^{-16}$erg~cm$^{-2}$s$^{-1}$ \\citep[95\\% confidence;][]{regener51}, implying L$_{0.5-2keV}\\sim$3$\\times$10$^{44}$erg~s$^{-1}$.   If the X-ray emission from BR1202N is real, its high X-ray luminosity suggests that BR1202N also hosts a luminous AGN. In addition, our new [\\ion{C}{2}] line detection, large dust emission, and the presence of abundant molecular gas suggest that BR1202N is forming stars at a high rate.  BR1202S and BR1202N thus form a system of two massive galaxies undergoing an interaction or a merger, separated by a projected distance of $\\sim$25~kpc, observed when the universe was only 1.2 Gyrs old. The presence of a highly obscured AGN (and an optically bright QSO) during a massive merger is consistent with the scenario predicted in recent galaxy collision simulations that investigate the evolution of a central black hole \\citep[e.g.][]{hopkins05}.  A high degree of clustering near massive galaxies is expected in the theoretical models of large scale structure formation \\citep[e.g.][]{white78}. This is supported by observations such as the higher surface density of submillimeter sources in the fields surrounding high-$z$ radio galaxies \\citep{stevens03}. Although an earlier study by \\citet{giallongo98} found little evidence for clustering of galaxies near BR1202-0725 with $R' \\leq 25$, two massive galaxies in proximity of each other makes BR1202-0725 a particularly interesting test ground for the hierarchical scenarios, soon after the reionization."
    },
    "0606/astro-ph0606619_arXiv.txt": {
        "abstract": "%  An upper limit  of 16\\% (at 95\\% c.l.) is derived for the photon fraction in cosmic rays with energies greater than $10^{19}$~eV, based on observations of the depth of shower maximum performed with the hybrid detector of the Pierre Auger Observatory. This is the first such limit on photons obtained by observing the fluorescence light profile of air showers. This upper limit confirms and improves on previous results from the Haverah Park and AGASA surface arrays. Additional data recorded with the Auger surface detectors for a subset of the event sample support the conclusion that a photon origin of the observed events is not favored. ",
        "introduction": "\\label{sec-intro}  The origin of ultra-high energy (UHE) cosmic rays above $10^{19}$~eV is still unknown~\\cite{reviews}. Their energy spectrum, arrival directions and composition can be inferred from air shower observations. However, agreement has not yet been reached on whether there is a break in the energy spectrum around $E_{\\rm GZK} \\sim 6 \\times 10^{19}$~eV (= 60 EeV). Such a steepening in the energy spectrum  is expected if UHE cosmic rays come from cosmologically distant sources~\\cite{gzk}, as is suggested by their overall isotropy. There have been claims, as yet unconfirmed, for clustering on small angular scales, and correlations with possible classes of sources. Moreover, results concerning the nuclear composition are still inconclusive.  While this deficit of robust observational results is partly due to the extremely small fluxes and, correspondingly, small numbers of events at such high energies, discrepancies might arise also from the different experimental techniques used. For instance, the determination of the primary energy from the ground array alone relies on the comparison with air shower simulations and is thus prone to uncertainties in modelling high energy interactions. Therefore it is essential to test results from air shower observations independently. The present work provides just such a cross-check for the upper limit derived previously from ground arrays on the photon fraction in UHE cosmic rays. An upper limit is set on the photon fraction above 10~EeV which is twice as strong as those given previously.  Photons are expected to dominate over nucleon primaries in non-acceleration (``top-down'') models of UHE cosmic-ray  origin~\\cite{bhat-sigl,sarkar03,models} which have been invoked in particular to account for a continuation of the flux above $E_{\\rm GZK}$ without a spectral feature as indicated by AGASA data~\\cite{agasa-gzk}. Thus, the determination of the photon contribution is a crucial probe of cosmic-ray source models. Separating photon-induced showers from events initiated by nuclear primaries is experimentally much easier than distinguishing light and heavy nuclear primaries. As an example, average depths of shower maxima at 10~EeV primary energy are predicted to be about 1000~g~cm$^{-2}$, 800~g~cm$^{-2}$, and 700~g~cm$^{-2}$ for primary photons, protons, and iron nuclei, respectively. Moreover, analyses of nuclear composition are uncertain due to our poor knowledge of hadronic interactions at very high energies. Photon showers, being driven mostly by electromagnetic interactions, are less affected by such uncertainties and can be modelled with greater confidence. To avoid the uncertainty from modelling hadronic interactions, we adopt an analysis method that does not require the simulation of nuclear primaries but compares data to photon simulations only.  So far limits on the UHE photon fraction in cosmic rays have been set by ground arrays alone. By comparing the rates of near-vertical showers to inclined ones recorded by the Haverah Park shower detector, upper limits (95\\% c.l.) of 48\\% above 10~EeV and 50\\% above 40~EeV were deduced~\\cite{ave}. Based on an analysis of muons in air showers observed by the Akeno Giant Air Shower Array (AGASA), the upper limits (95\\% c.l.) to the photon fraction were estimated to be 28\\% above 10~EeV and 67\\% above 32~EeV ~\\cite{shinozaki}. An upper limit of 67\\% (95\\% c.l.) above 125~EeV was derived in a dedicated study of the highest energy AGASA events~\\cite{risse05}.  In this work, we obtain a photon limit from the {\\em direct} observation of the shower profile with fluorescence telescopes, using the depth of shower maximum $X_{\\rm max}$ as the discriminating observable. To achieve a high accuracy in reconstructing the shower geometry, we make use of the ``hybrid'' detection technique, i.e. we select events observed by both the ground array and the fluorescence telescopes~\\cite{mostafa}. For a subset of the event sample, a sufficient number of ground detectors were also triggered, yielding a variety of additional shower observables. Considering as example the signal risetime measured with the ground array, we demonstrate the discrimination power of these independent observables to photon-induced showers.  The plan of the paper is as follows. In Section~\\ref{sec-photons}, predictions for the UHE photon fraction in cosmic-ray source models and features of photon-initiated air showers are summarized. Section~\\ref{sec-data} contains the description of the data and of photon simulations. In particular, the data selection criteria are discussed. A careful choice of the quality and fiducial volume cuts is required to control a possible experimental bias for photon primaries. In Section~\\ref{sec-results}, the method for deriving a photon fraction is described and applied to the data. An example of the discrimination power of observables registered by the surface array is shown in Section~\\ref{sec-sd}. Finally in Section~\\ref{sec-outlook}, we discuss the prospects for improving the bound on UHE photons. ",
        "conclusions": ""
    },
    "0606/astro-ph0606333_arXiv.txt": {
        "abstract": "Recent wide area hard X-ray and soft Gamma ray surveys have shown that the fraction of X-ray obscured AGNs in the local universe significantly decreases with intrinsic luminosity. In this letter we point out that two correction have to be made to the samples: 1) radio loud AGNs have to be excluded since their X-ray emission might be dominated by the jet component; 2) Compton thick sources have to be excluded too since their hard X-ray and soft gamma ray emission are also strongly attenuated by Compton scattering. The soft gamma-ray selected AGN samples obtained by $SWIFT$ and $INTEGRAL$ provide the best opportunity to study the fraction of obscured AGN in the local universe in the least biased way. We choose these samples to check if the corrections could alter the above result on the fraction of obscured AGNs. We find that before the corrections both samples show significant anti-correlation between $L_X$ and $N_H$, indicating obvious decrease in the fraction of obscured AGNs with luminosity. However, after the corrections, we find only marginal evidence of anti-correlation (at 98\\% confidence level) in the SWIFT sample, and no evidence at all in the $INTEGRAL$ sample which consists of comparable number of objects. We conclude that current samples only show a marginal decrease in the fraction of obscured AGNs in the local universe, and much larger samples are required to reach a more robust conclusion. ",
        "introduction": "A large population of obscured powerful quasars called type-2 quasars has long been predicted by the widely accepted unified model for active galactic nuclei (AGNs; Antonucci 1993). They are believed to be intrinsically the same as type-1 quasars but with strong obscuration in both optical and X-ray band due to the presumable torus. Most of such type-2 quasars, which might dominate the black hole growth (e.g., Mart\\'\\i nez-Sansigre et al. 2005), have been missed by optical surveys for quasars. Doubt has been expressed at times on the existence of type-2 quasars (e.g. Halpern, Turner \\& George 1999), mainly due to their rareness and the fact that they are easily mimicked by other types of AGNs. The hard X-ray emission is less affected by the photon-electric absorption, making hard X-ray surveys good approach to search for type-2 quasars. Another advantage of X-day data is that X-ray spectral fitting can yield intrinsic luminosity and absorption simultaneously, which are both required to study the population of obscured quasars. A number of obscured quasars have been successfully revealed by recent deep and wide-area X-ray surveys performed by $Chandra$ and XMM (Norman et al. 2002; Stern et al. 2002; Mainieri et al. 2002; Fiore et al. 2003; Caccianiga et al. 2004).  However, many studies have claimed that the fraction of obscured quasars is much smaller than expected from the simplified AGN unified model. Ueda et al. (2003) computed the X-ray luminosity function for 2 -- 10 keV selected AGN samples, and found that the fraction of X-ray absorbed AGNs (with $N_H$ $>$ 10$^{22}$ cm$^{-2}$) drops from $\\sim$ 0.6 at intrinsic $L_X$ around 10$^{42}$ ergs s$^{-1}$ to around 0.3 at $L_X$ above 10$^{44}$ ergs s$^{-1}$. Several other studies based on hard X-ray surveys also support the scheme that the fraction of type II AGNs (or X-ray absorbed AGNs) decreases with intrinsic luminosity (Steffen et al. 2003; Barger et al. 2005; La Franca et al. 2005; Shinozaki et al. 2006). However, contrary results are also reported. Perola et al. (2004) found that the fraction of obscured AGNs in HELLAS2XMM does not change with X-ray luminosity, although the fraction (40\\%) appears smaller than expected. Eckart et al. (2006) found that half of the AGNs identified by the SEXSI program are X-ray obscured with $N_H$ $>$ 10$^{22}$ cm$^{-2}$, and the fraction of obscured AGNs is independent of the unobscured luminosity. The discrepancy can be attributed to the redshift evolution in the fraction of obscured quasars (La Franca et al. 2005; Wang et al. 2006): while shallow and wide surveys are probing quasars in the local universe where probably only a small fraction of quasars are obscured, deep surveys are detecting quasars at much higher redshift where more than half of them are obscured (Wang et al. 2006).  Soft gamma-ray selected AGN samples (Markwardt et al. 2005, hereafter M06; Bassani et al.  2006, hereafter B06; Beckmann, Gehrels \\& Shrader 2006) obtained by $SWIFT$ and $INTEGRAL$ provide the best opportunity to study the fraction of obscured AGN in the local universe in the least biased way. Note that for Compton-thick cases (N$_H$ $>$ 10$^{24}$ cm$^{-2}$), even soft gamma-ray emission could be significantly attenuated, thus soft gamma-ray selected samples are not complete to Compton-thick sources. Claims have been made based on the soft gamma-ray selected samples, that the fraction of obscured AGNs in the local universe significantly decrease with increasing luminosity. In this letter we point out that two correction have to be made to the samples to study of the fraction of obscured quasars. The corrections are: a) radio loud AGNs have to be excluded since their X-ray emission might be dominated by the jet component thus the measured luminosity and the obscuration does not reflect the intrinsic values in the nuclei; b) Compton thick sources have to be excluded too since their soft gamma ray emission are also strongly attenuated by Compton scattering and their intrinsic luminosities are hard to estimate. After the corrections, we find only marginal decrease in the fraction of obscured AGN with luminosity in the local universe. Larger samples are required to reach a more robust conclusion. ",
        "conclusions": ""
    },
    "0606/astro-ph0606105_arXiv.txt": {
        "abstract": "In this paper, we describe the performance of an $N$-body simulation of star cluster with 64k stars on a Cray XD1 system with 400 dual-core Opteron processors. A number of astrophysical $N$-body simulations were reported in SCxy conferences. All previous entries for Gordon-Bell prizes used at least 700k particles. The reason for this preference of  large numbers of particles is the parallel efficiency. It is very difficult to achieve high performance on large parallel machines, if the number of particles is small. However, for many scientifically important problems the calculation cost scales as $O(N^{3.3})$, and it is very important to use large machines for relatively small number of particles. We achieved 2.03 Tflops, or 57.7\\% of the theoretical peak performance, using a direct $O(N^2)$ calculation with the individual timestep algorithm, on 64k particles. The best efficiency previously reported on similar calculation with 64K or smaller number of particles is 12\\% (9 Gflops) on Cray T3E-600 with 128 processors. Our implementation is based on highly scalable two-dimensional parallelization scheme, and low-latency communication network of Cray XD1 turned out to be essential to achieve this level of performance. ",
        "introduction": "Many astronomical objects, such as the Solar system, star clusters, galaxies, and clusters of galaxies and the large scale structure of the universe, are expressed as the system of many ``particles'' interacting through the Newtonian gravity. The numerical simulation of such systems, usually called ``Gravitational $N$-body simulations'', proved itself to be one of the most useful tools in computational astrophysics. It has been used to study systems of all scales, from the formation of the Moon \\cite{KMI99} to the entire visible universe \\cite{MilleniumRun}.  Many of these simulations are computationally demanding, because of three reasons. The first reason is that the calculation cost per timestep, at least with naive implementation, scales as $O(N^2)$, where $N$ is the number of particles in the system.  The second reason is that in many cases the evolution timescale of the system is many orders of magnitude longer than the orbital timescale of typical particles in the system. An example is the long term simulation of the Solar system. We need to follow the orbits of planets and other objects for $4.5\\times 10^9$ years, and we need around 100 timesteps per year or more to achieve reasonable accuracy. In the case of the simulation of the formation of planets, we need to follow the evolution of planetesimals for several million orbits. This requirement of large number of orbits limits the number of particles we can handle to around 1,000 or less \\cite{KokuboIda02}, even on special-purpose computers such as GRAPE-4 \\cite{Grape4} or GRAPE-6 \\cite{Grape6}. The same is true for the simulation of star clusters. Here, the evolution timescale depends linearly on the number of particles, and the number of timesteps per orbits increases as $O(N^{1/3})$, since the distance which one particle can move in a single timestep must be a small fraction of the distance to its nearest neighbor. In total, the calculation cost increases as $O(N^{3.3})$.  The third reason is that different particles can have very different orbital timescales. Since the gravitational force is an attracting force, two particles can become arbitrary close to each other. Thus, the timestep need to be adaptive. Moreover, it is not enough to just change the timestep. If we use the variable, but single, timestep to integrate all particles, when we increase $N$ the average step-size for star cluster simulation shrinks as $1/N$ or somewhat faster. Thus, the calculation cost becomes $O(N^4)$ instead of $O(N^{3.3})$.  We can solve the third problem by allowing individual particles to have their own time and timestep \\cite{Aarseth63, Aarseth85}. However, with this scheme the number of particles we can integrate in parallel becomes much smaller than $N$, and it becomes more difficult to use large scale parallel computers even when $N$ is not very small.  Dorband et al. \\cite{DHM03} reported the performance of parallel implementation of this individual timestep (or block timestep) algorithm, combined with simple $O(N^2)$ direct force calculation, on a Cray T3E. They have fine-tuned the communication scheme so that there will be maximal overlap between the communication and calculation. Even so, if the initial distribution of particles is highly anisotropic, they found that performance did not increase for more than 64 processors, for $N=64{\\rm k}$. More recent works on PC clusters \\cite{Gualandris2004} generally gave much worse scaling, because the communication performance of PC clusters, when normalized by the calculation speed, is orders of magnitude worse than that of Cray T3E.  The use of special-purpose GRAPE processors does improve the performance of single-node calculation quite significantly and made it possible to perform long simulations with relatively small number of particles such as 64K or less. However, in this case the overall performance was limited by the speed of the single host computer, since the parallelization over multiple GRAPE nodes involves much bigger relative overhead compared to that of PC clusters, and it turned out to be difficult to achieve any gain unless the number of particles is quite large \\cite{Grape6A}.  In this paper, we report the result of the first calculation which used nearly 1,000 processors for the simulation of star clusters with just 64K stars and achieved more than 50\\% of the theoretical peak speed of the machine. Average speed measured on a 400 CPU Cray XD1 (2.2GHz Dual-Core Opteron) was 2.03 Tflops, while the nominal theoretical peak speed is 3.52 Tflops.  Previous works, even when the parallel efficiency was good, achieved less than 15\\% of the theoretical peak speed of the machine. Thus, our implementation is remarkable in two ways: it is highly scalable for small number of particles, and it can achieve quite high absolute performance.  In the rest of this paper, we describe how we achieved the good scalability and high absolute peak performance at the same time. In section 2 we outline the individual (block) timestep algorithm. In section 3 we describe the two-dimensional parallelization scheme we used. It was originally proposed by Makino\\cite{Makino2002}. We describe the modifications we made to achieve high performance. In section 4 we briefly describe the highly optimized force calculation loop. The details are given in Nitadori et al.\\cite{NMH06}. In section 5 we describe the result of a long calculation of a two-component star cluster with 64K particles. Section 6 sums up. ",
        "conclusions": "We developed  a new highly scalable parallel algorithm for direct $N$-body simulation with individual (block) timestep integration method.  We implemented this new algorithm on a Cray XD1 and achieved quite good performance scaling for up to 800 processors (the largest existing configuration of Cray XD1), for small number of particles such as 16k or 64k. Parallel efficiency better than 80\\% is obtained for 512 processors, and sustained speed of 2.03 Tflops is achieved for a long calculation on 800 processors.  For the first time, we demonstrated that a modern MPP with $\\sim 1000$ processors can be used for astrophysical $N$-body simulations with less than $10^5$ particles. This capability to achieve high efficiency for small $N$ is quite important, since the total calculation cost scales as $O(N^{3.3})$ and $N=10^5$ is the practical limit of the number of particles we can handle with the effective speed of several Tflops.  We thank Atsushi Kawai and Toshiyuki Fukushige of K\\&F Computing Research for their stimulating discussions.  We are grateful to Hideki Matsumoto of SONY and staffs of Cray Japan for valuable technical supports on the use of the XD1 system.  K.N. thanks Shigeo Mitsunari at u10 Networks for discussions on speeding-up of the calculations with assembly-language programming."
    },
    "0606/astro-ph0606380_arXiv.txt": {
        "abstract": "{  We present the results of a near-infrared monitoring programme of a selected sample of stars, initially suspected to be Mira variables and OH/IR stars, covering more than a decade of observations. The objects monitored cover the typical range of IRAS colours shown by O-rich stars on the asymptotic giant branch and show a surprisingly large diversity of variability properties. Sixteen objects are confirmed as large-amplitude variables. Periods between 360 and 1\\,800\\Day\\, and typical amplitudes 1\\Mag\\,$\\la$\\,$\\Delta$K\\,$\\la$\\,2\\Mag could be determined for nine of them. In three light curves, we find a systematic decrease in the mean brightness, and two light curves show pronounced asymmetry. One source, IRAS\\,07222--2005, shows infrared colours typical of Mira variables, but it pulsates with a much longer period ($\\approx$1200\\Day) than a normal Mira. Two objects are either close to (IRAS\\,03293+6010) or probably in (IRAS\\,18299--1705) the post-AGB phase. In IRAS\\,16029--3041 we found a systematic increase in the H--K colour of $\\approx$1\\Mag, which we interpret as evidence of a recent episode of enhanced mass loss. IRAS\\,18576+0341, a heavily obscured luminous blue variable was also monitored. The star showed a continued decrease in brightness over a period of 7 years (1995 -- 2002). ",
        "introduction": " ",
        "conclusions": "Conclusions}  The monitoring programme has shown that the variability properties along the `O-rich AGB sequence' are diversified. In general, periods and amplitudes increase with redder colours, but the scatter is large, and AGB stars in evolutionary phases of non- or low-amplitude variability are found at the borders of the sequence. The transition to post-AGB stars in the IRAS colour-colour diagram appears to be smooth. The variability properties certainly hold important clues to the evolutionary stage of the stars on the AGB. To uncover the links between these properties and particular phases of AGB evolution requires, however, larger samples and therefore long-term monitoring programmes to study the variability properties, especially at the red end of the `O-rich AGB sequence', should be encouraged."
    },
    "0606/astro-ph0606349_arXiv.txt": {
        "abstract": "This paper studies the long term evolution of planetary systems containing short-period planets, including the effects of tidal circularization, secular excitation of eccentricity by companion planets, and stellar damping. For planetary systems subject to all of these effects, analytic solutions (or approximations) are presented for the time evolution of the semi-major axes and eccentricities. Secular interactions enhance the inward migration and accretion of hot Jupiters, while general relativity tends to act in opposition by reducing the effectiveness of the secular perturbations. The analytic solutions presented herein allow us to understand these effects over a wide range of parameter space and to isolate the effects of general relativity in these planetary systems. ",
        "introduction": "Starting with the discovery of extrasolar planets (Mayor \\& Queloz 1995; Marcy \\& Butler 1996), a substantial fraction of the planetary orbits have been found close to their stars, with periods $P \\sim 4$ days. These objects are often referred to as ``hot Jupiters''.  With $\\sim200$ planets detected to date\\footnote{http://www.ucolick.org/\\~{$\\,$}laugh/}, the distribution of orbital periods shows a measurable pile-up at periods $P$ = 3--5 days, i.e., roughly 10 percent of the currently detected planets have periods $P < 5$ days. Although the close planets are the most easily detected, this finding is not a selection effect: The current statistics indicate that 1.2 percent of all FGK stars have hot Jupiters within 0.1 AU of their stars (Marcy et al. 2005).  These hot Jupiters are subject to tidal interactions with their central stars (e.g., Goldreich \\& Soter 1963), and this star-planet coupling can influence the long term evolution of these systems. If the system contains additional bodies, then planet-planet interactions can also affect the long term evolution of the close planets.  In our companion paper (Adams \\& Laughlin 2006a; hereafter Paper I), we present a treatment of secular interactions for (non-resonant) systems with multiple planets.  In this paper, we combine this formulation of secular interactions with star-planet interactions to study the long term fate and evolution of close planets.  The basic theory of secular interactions, as applied to extrasolar planetary systems and used herein, is reviewed in Paper I (see also Murray \\& Dermott 1999; hereafter MD99). These interactions allow for planetary orbits to exchange angular momentum so that orbital eccentricities change on secular time scales that are long compared to both the orbit time and observational monitoring times (tens of years), but short compared to the ages of the systems (a few Gyr). The characteristic secular interaction time scales provide a simple metric of the importance of planet-planet perturbations within a given system.  For a collection of observed multi-planet systems, these secular time scales fall in the range 100 -- 50,000 yr, as listed in Table 1 of Paper I.  On longer time scales, tidal interactions with the star act to circularize close orbits and can lead to continued inward migration.  This process might have a bearing on the observed pile-up of planets with periods of 3--4 days and can lead to considerable energy input into the planetary atmospheres.  Although this process has been considered previously (e.g., Trilling 2000; Bodenheimer et al. 2001, 2003; Yu \\& Goldreich 2002; Mardling \\& Lin 2002, 2004; Gu et al. 2003; Faber et al. 2005), this paper takes the additional step of providing analytic expressions for the evolution time for various classes of systems, including both tidal damping and secular interactions.  This work thus provides additional analytic insight into the problem as well as elucidating the role played by secular interactions.  The version of secular interaction theory formulated in Paper I includes the leading order corrections for general relativity (GR), which causes the periasta of planets to precess forward in their orbits.  As a result, this semi-analytical treatment also allows us to explicitly delineate the role played by general relativity in the long term evolution of close planets.  For a single planet in a close orbit, the tidal circularization effect is usually written in terms of the time scale for eccentricity decay, $\\taucirc = -e / {\\dot e}$. For close orbits in multiple planet systems, however, the eccentricity is excited via interactions that we model according to secular theory. The first approximation to the long term behavior is to assume that the orbit decays with constant angular momentum and that the semi-major axis decreases on the same time scale, i.e., \\be { {\\dot a} \\over a} = {2 e {\\dot e} \\over 1 - e^2} = - {2 e^2 \\over (1 - e^2)}\\, {1 \\over \\taucirc} \\, , \\ee where the time scale $\\taucirc$ for circularization can be written in the form \\be \\taucirc \\approx {4 Q_P \\over 63} \\Bigl( {a^3 \\over GM_\\ast} \\Bigr)^{1/2} {m_P \\over M_\\ast} \\Bigl( {a \\over R_P} \\Bigr)^5 \\, (1 - e^2)^{13/2} [ F(e^2) ]^{-1} \\, , \\label{eq:circone} \\ee where $Q_P \\approx 10^5 - 10^6$ is the tidal quality factor and $R_P$ is the planet radius. For extrasolar planets, the quality factor $Q_P$ and the radius $R_P$ depend sensitively on the planetary mass, temperature, and composition (Bodenheimer et al. 2003).  This general form for the time circularization scale is well known (e.g., Goldreich \\& Soter 1966), but includes additional factors to account for the effects of nonzero eccentricity (see Hut 1981). The choice $F \\approx 1 + 6 e^2 + {\\cal O}(e^4)$ provides a reasonable approximation to the results for close binaries (Hut 1981) for moderate eccentricities. Inserting representative values, and taking the limit $e \\to 0$, we write the circularization time scale in the form \\be \\tau_0 = \\taucirc (e=0) \\approx  \\, \\, 1.6 \\, {\\rm Gyr} \\, \\, \\Bigl( {Q_P \\over 10^6} \\Bigr) \\, \\, \\Bigl( {m_P \\over m_J} \\Bigr) \\, \\, \\Bigl( {M_\\ast \\over M_\\odot} \\Bigr)^{-3/2} \\, \\, \\Bigl( {R_P \\over R_J} \\Bigr)^{-5} \\, \\, \\Bigl( {a \\over 0.05 \\, {\\rm AU}} \\Bigr)^{13/2} \\, \\, . \\label{eq:circtwo} \\ee Throughout this paper we work in terms of the dimensionless time scale \\be \\tz \\equiv { t / \\tau_0} \\, . \\ee  These systems are also affected by a stellar damping effect in which energy is dissipated in the star due to tides raised by the planet. The effectiveness of this process is determined by an analogous parameter $Q_\\ast$, the tidal quality factor of the star. The net result of this process is to cause the semi-major axis of the inner planet to decay on a time scale $\\taups$ (again, see Goldreich \\& Soter 1966, Hut 1981) that can be written in the form \\be \\taups = \\tau_0 \\Gamma^{-1} \\Bigl( {a \\over a_0} \\Bigr)^{13/2} \\, , \\qquad {\\rm where} \\qquad \\Gamma \\equiv {2\\over 7} \\, \\Bigl( {R_\\ast \\over R_P} \\Bigr)^5 \\, \\Bigl( {Q_P \\over Q_\\ast} \\Bigr) \\, \\Bigl( {m_P \\over M_\\ast} \\Bigr)^2 \\, . \\label{eq:taupdef} \\ee As defined here, $\\Gamma$ is generally a small parameter. For Jupiter-like planets, $R_P \\sim 0.1 R_\\ast$ and $m_P \\sim 10^{-3} M_\\ast$ so that $\\Gamma_J \\sim 0.03 (Q_P/Q_\\ast)$. For Neptune-like planets, we find $\\Gamma_{Nep} \\sim 0.015 (Q_P/Q_\\ast)$, but we expect the $Q_P$ value to be much smaller so that $\\Gamma_{Nep} \\ll \\Gamma_J$.  This paper considers hot Jupiters subject to tidal circularization effects (with strength determined by $Q_P$), both in single planet systems (\\S 2.1) and in two-planet systems where the eccentricity of the inner planet is excited through secular interactions (\\S 2.2). Next we consider hot Jupiter systems with additional stellar damping effects (with strength determined by $Q_\\ast$), both for single planet systems (\\S 3.1) and two-planet systems (\\S 3.2). Our results are summarized in \\S 4 along with a discussion of their ramifications. ",
        "conclusions": "This paper presents solutions for the the long term evolution of four types of planetary systems --- one and two planet systems with tidal circularization, both with and without the inclusion of stellar damping effects. For the two planet systems, planet-planet interactions are modelled using secular interaction theory (MD99) including leading order relativistic corrections (Paper I).  The solutions are presented in implicit form; for example, $\\tz (f)$ gives the time as a function of the inner planet's semi-major axis, as specified by $f = a_1(t)/a_1(0)$. In the absence of stellar damping, we find exact analytic solutions to the long term evolution, as given by equations (\\ref{eq:jsol}) and (\\ref{eq:twosol}). For systems in which the stellar damping term is important, we find approximate analytic solutions given by equations (\\ref{eq:timesol}) and (\\ref{eq:tdsol}).  In multiple planet systems, secular interactions enhance the accretion of inner planets by the central star, provided that the outer planet has a sufficiently large mass, small semi-major axis, and/or large eccentricity. We have presented a quantitative assessment of the orbital elements of a second planet required to drive a hot Jupiter into the central star within 5 Gyr, for two-planet systems both with (Fig. \\ref{fig:longterm2}) and without (Fig. \\ref{fig:longterm1}) stellar damping. General relativity acts to delay the accretion of planets, i.e., the region of parameter space for which the inner planet would be accreted within a given time would be much larger in the absence of relativistic effects (see the dashed curves in Figs. \\ref{fig:longterm1} and \\ref{fig:longterm2}).  For multiple planet systems with hot Jupiters as inner planets, secular interactions tend to give the inner planet nonzero eccentricity. This continual addition of eccentricity, in conjunction with the circularization processes experienced by such planets, leads to large amounts of energy dissipation within the hot Jupiters (eqs. [\\ref{eq:energydiss}] and [\\ref{eq:energyratio}]).  Such extreme dissipation, in turn, may lead to mass loss and thereby explain the relatively small masses observed for close planetary companions. This process should be studied in greater detail in the future.  The framework of analysis presented here can be readily applied to individual systems whenever the orbital elements of the constituent planets are known to reasonable accuracy. We maintain an up-to-date catalog of the known extrasolar planets\\footnote{http://www.ucolick.org/\\~{$\\,$}laugh/} where the time scales derived here are tabulated using the best available fits to the radial velocity data sets. This data base, in conjunction with the analytic solutions presented here, should provide a useful resource for further research on close planetary systems.  \\bigskip  We would like to thank an anonymous referee for suggesting that this paper be published separately from Paper I. This work was supported at U. Michigan (FCA) by the Michigan Center for Theoretical Physics and by NASA through the Terrestrial Planet Finder Mission (NNG04G190G) and the Astrophysics Theory Program (NNG04GK56G0).  This material is based in part upon work supported by the National Science Foundation CAREER program under Grant No. 0449986 (GL), and was also supported at U. C. Santa Cruz (GL) by NASA through the Terrestrial Planet Finder Precursor Science Program (NNG04G191G) and through the Origins of Solar Systems Program (NAG5-13285).  \\newpage"
    },
    "0606/astro-ph0606725_arXiv.txt": {
        "abstract": "We investigate the difference between the host galaxy properties of core-collapse supernovae and long-duration gamma-ray bursts (LGRBs), and quantify a possible metallicity dependence of the efficiency of producing LGRBs. We use a sample of 16 CC SNe and 16 LGRBs from \\citet{F06} which have similar redshift distributions to eliminate galaxy evolution biases. We make a forward prediction of their host galaxy luminosity distributions from the overall cosmic metallicity distribution of star formation. The latter is based on luminosity functions, star-formation rates (SFR) and luminosity--metallicity ($L$--$Z$) relations of galaxies. This approach is supported by the finding that LGRB hosts follow the $L$--$Z$-relations of star-forming galaxies. We then compare predictions for metallicity-dependent event efficiencies with the observed host data. We find that UV-based SFR estimates predict the hosts distribution of CC SNe perfectly well in metallicity-independent form. In contrast, LGRB hosts are on average fainter by one magnitude, almost as faint as the Large Magellanic Cloud. Assuming this is a metallicity effect, the present data are insufficient to discriminate between a sharp cutoff and a soft decrease in efficiency towards higher metallicity. For a sharp cut-off, however, we find a best value for the cutoff metallicity, as reflected in the oxygen abundance, $12+\\log {\\rm (O/H)}_{\\rm lim}\\simeq 8.7\\pm 0.3$ at 95\\% confidence including systematic uncertainties on the calibration of \\citet{KK04}. This value is somewhat lower than the traditionally quoted value for the Sun, but is comparable to the revised solar oxygen abundance \\citep{A05}. LGRB models that require sharp metallicity cutoffs well below $\\sim 1/2$ the revised solar metallicity appear to be effectively ruled out, as they would require fainter LGRB hosts than are observed. We also discuss the likely implications of the still ongoing metallicity 'calibration debate'. ",
        "introduction": "Since the first optical afterglow to a gamma-ray burst was identified \\citep[ GRB 970228]{vP97}, the restframe properties of about 50 host galaxies of long-duration gamma-ray bursts (LGRBs) have been identified. Already with a small sample it became clear that the hosts are preferentially faint blue irregular galaxies \\citep{F99,LeF03,CHG04}. This trend was recently confirmed by a comparison between hosts of LGRBs and hosts of core-collapse supernovae (CC SNe). CC SNe were found in late-type galaxies of all morphologies and luminosities, in contrast to LGRB hosts, which were still found to be almost exclusively irregular and on average fainter \\citep[ hereafter F06]{F06}.  Since both CC SNe and LGRBs are considered explosions of massive young stars with ages below $\\sim 30$~Myr and probably below the typical timescale of star formation in most galaxies, there is {\\em apriori\\/} no reason to expect a considerable difference in their host galaxies. CC SNe should be direct and unbiased tracers of very recent star formation, although complications may arise from dust extinction intrinsic to the host. Furthermore, we expect a dependence of the mass threshold for core collapse/black-hole formation on metallicity \\citep{ H95,H03,ET04} and binary evolution effects \\citep{B99,B01,Pf02,P04b}.  The suggested explanation for the host preference of LGRBs is that they arise from low-metallicity stellar populations which are predominantly found in low-luminosity galaxies according to the luminosity--metallicity relation \\citep[e.g][ hereafter KK04]{KK04}. Also, in several progenitor models of GRBs \\citep[in particular `collapsar' models;][]{MW99}, wind mass loss plays an important role in slowing down the rotation of massive stellar cores. Since lower metallicity leads to weaker stellar winds and hence less angular-momentum loss, the stars are more likely to retain rapidly rotating cores at the time of the explosion, as required in the collapsar model.   LGRBs discovered by the {\\it Swift} mission have a median redshift of 2.75 and apparently have a distribution expected for an unbiased tracer of star formation \\citep{J06}. If both the cosmic star formation history and the cosmic evolution of metallicity in cold star-forming gas were known with high accuracy, this redshift distribution could provide clues about the metallicity dependence of the LGRB rate. However, the LGRBs discovered by BeppoSAX have a median redshift near $z\\simeq 1$, and predictions for the completeness function are very difficult to make, owing to the complexity of GRB trigger algorithms and detector intricacies. Even with a much larger sample, the redshift distribution of GRBs might be affected by systematic uncertainties such that no strong constraints on metallicity can be found.  Instead, first spectroscopic measurements of the metallicity of LGRB hosts at redshifts $0.4<z<1$ from the GHost Studies \\citep{Sav06} indicate that LGRB hosts follow the luminosity--metallicity relation of star-forming galaxies. A preference for low metallicity would then require LGRBs to favour less luminous galaxies more than general star formation does. When working with gas-phase metallicities of galaxies, we need to be aware of the {\\it calibration debate} between two basic methods which result in numbers that differ by a factor of $\\sim 2$: spectra of low signal-to-noise only facilitate an indirect strong-line method (Pagel's $R_{23}$ indicator, 1979) to measure metallicity, which is calibrated against photoionization models. In contrast, high-S/N spectra allow for a direct measurement of metallicity taking into account the electron temperature measured from weak auroral lines \\citep[see e.g. ][]{K03}. The $T_e$-based methods find typically 0.2--0.5 dex lower metallicities, but have been critized to be biased by not taking inhomogeneities in the temperature into account.  In this paper, we wish to derive constraints on the metallicity dependence of LGRB progenitors by comparing forward predictions of host luminosity distributions with the observed host sample. We apply a minor modification to the F06 samples by restricting them to identical redshift ranges $z= [0.2,1]$ (see Sect.~2). Our quantitative predictions of the host luminosity distribution is based on luminosity functions and star-formation rates of galaxies and their luminosity--metallicity relations (see Sect.~3). Here, we need to make sure that all ingredients are on the exact same metallicity calibrator and opt for the strong-line method on the KK04 calibrator for consistency. We then fold in different prescriptions for the metallicity dependence of the LGRB event rate and constrain these by comparing our predictions to the observed sample in Sect.~4. In Sect.~5 we discuss the results and the implications of the calibration debate. Throughout the paper, we use Vega magnitudes and the cosmological parameters $(\\Omega_{ \\rm m},\\Omega_\\Lambda) =(0.3,0.7)$ and $H_0=h_{70}\\times 70$~km/(s~Mpc). We use the revised solar oxygen abundance of $12+\\log {\\rm (O/H)}=8.66$ \\citep[see also \\citet{ALA01}]{A05}. ",
        "conclusions": "\\subsection{Caveats}  One difficulty with the present analysis could be host confusion. If the GRBs took preferentially place in low-luminosity companions to the formally identified hosts, then the observed host luminosity distribution would be biased to higher luminosities, suggesting higher metallicity cutoffs. It is beyond the scope of this paper to investigate whether HST observations could resolve such dwarfs from their larger, mostly LMC-type companions at $z<1$.  However, if most LGRBs occurred in such small galaxies, one would also expect a large population of sufficiently isolated small dwarfs, where no confusion would take place. But the small number of observed isolated hosts with $M_V>-17$ leaves little room for any adjustment of our results.  On the contrary, it is possible that more massive LGRB hosts are missing from the sample: luminous dust-obscured star-bursting galaxies harbour a significant fraction of the overall cosmic star formation. If LGRBs occur in these, their optical afterglows may be invisible, which would depopulate the host sample specifically at the bright end. The CC SNe host sample would be less affected by this effect, as CC SNe have a lower average progenitor mass and hence longer average time delay from formation to explosion. This allows the progenitors to leave their immediate formation environment and travel into less obscured regions. The effect of progenitor mass on the distance to the likely formation site has been shown to apply in the comparison of SNe II and SNe Ib/c \\citep{JA06}. The sample analysed here lacks by definition all so-called 'dark bursts', which may be related to massive dusty galaxies: GRB 020127 only produced an X-ray afterglow, and the host was identified as a $z\\sim 1.9$ dusty ERO galaxy with 5~$L^*$ luminosity \\citep{Be06}. \\citet{Le06} observed an extremely red ($R-K\\sim 6$) afterglow to GRB 030115 in an ERO host galaxy with $R-K\\sim 5$, suggesting that this burst was highly obscured by dust. The afterglow was only identified with NIR observations and would have been missed by optical follow-up alone. Again, for the more recent GRB 050223 only an X-ray afterglow was seen by {\\it Swift}, but no optical afterglow. The host was shown to be a dusty galaxy with $A_V>2$ \\citep{Pel06}. {\\it If there is a significant observational bias against LGRBs in massive dust-enshrouded galaxies, then the metallicity-dependence of LGRB rates may be rather weaker than our findings presently suggest.}  Uncertainties also exist with respect to the calibration of oxygen gas-phase metallicities and on how the oxygen abundance is representative of other elements. E.g., $\\alpha$-elements are often found to be enhanced in systems where the star-formation timescale is much shorter than the timescale for Type Ia supernovae. Depending on the progenitor model for LGRBs, different elements might be most responsible for affecting opacities in stellar atmospheres/winds and hence stellar evolution.  However, $\\alpha$-enhancements are typically observed in old ellipticals which are not at all common LGRB hosts or typical star-forming galaxies, and thus should not affect our analysis appreciably.  In summary, solar metallicity appears to mark a pivotal point of roughly constant efficiency apparently independent of the softness of the break.  At $z\\sim 0.7$, this metallicity seems typical for galaxies of $M_B-5 \\log h_{70}=-20$ in either $L$--$Z_{\\rm O}$-relation. The number ratio of CC SNe to either side of this host luminosity cut is 1:1, whereas it is 1:3 for LGRBs. If there was any doubt about the calibration of the oxygen abundance at $z\\sim 0.7$, or the relevance of oxygen in comparison to other elements, then we could rephrase our main conclusion such that a most likely cutoff for LGRBs is around the mean metallicity (by any measurement) of $M_B=-20$ galaxies at $z\\sim 0.7$.   \\subsection{Are GRB hosts on the luminosity-metallicity relation?}  Several recent studies have aimed at constraining LGRB progenitors from direct measurements of host metallicities \\citep{So05,Sav06,Sta06} using very small samples:  \\citet{So05} present measured metallicities for three LGRB hosts with values of $\\sim 8.2\\ldots 8.7$. \\citet{Sav06} consider seven hosts at $z=[0.4,1.0]$ and find mean galaxy metallicities of 8.3 to 8.55 for five of them, while two hosts are clearly super-solar. They furthermore claim, that the hosts fall right onto the regular mass--metallicity relation of normal star-forming galaxies.  This host sample with measured metallicities lacks the lowest-luminosity hosts and so contains preferentially higher metallicity galaxies. Hence, their median metallicity is not representative, but the statement on consistency with the mass--metallicity relation is highly relevant.  \\citet{Sav06b} give a number of host metallicity determinations, which are collected from the literature and converted onto the same (KK04) calibrator, which is also the calibrator used in this paper for the $L$--$Z_{\\rm O}$-relation. Five of the above hosts are actually part of the LGRB host sample used here. Their spectroscopic metallicities are compared to the estimated values we have assigned to them using $L$--$Z_{\\rm O}$-fit B. We find them all to be within $\\pm 1/3$~dex of the relation with a mean of $-0.034$~dex, which is consistent with no bias within the statistical power of five objects.  This results is not surprising. A physical model for GRB progenitors may incorporate an explicit metallicity dependence, but we can not expect the progenitor to know explicitely about galaxy parameters such as luminosity or morphology. At a fixed metallicity level, we would then expect bursts predominantly from those galaxy mass bins containing the highest star-formation rate. Here, both the shape of the mass function and the SFR trends with mass play a role. Presumably, the declining mass function and the increasing SFR with mass cancel to some extent. KK04 searched for trends of galaxy properties within the scatter of their $L$--$Z_{\\rm O}$-relation. While they found no such trends, they particularly did NOT find strongly star-forming galaxies on the bright or metal-poor side of the relation. However, among very low-mass dwarf galaxies it is conceivable, that strong variations in SFR with time lead to such strong variations in mass-to-light ratio, such that CC SNe and LGRBs are found predominantly on the bright side of the relation. Presumably, this would have no effect for our analysis, as the mean metallicity bias would be small and remain within the flat portion of the parametric efficiency functions considered in this paper.   \\subsection{Global vs. local metallicity measurements}  We would like to point out, that even spectroscopic, but spatially unresolved, measurements of a mean host metallicity do not directly reflect the metallicity of the young stellar population near the LGRB progenitor. In fact, the error in the progenitor metallicity is dominated by the internal $Z_{\\rm O}$ dispersion of the galaxy, which is estimated to be as large as the dispersion of galaxy-averaged $Z_{\\rm O}$ values in the $L$--$Z_{\\rm O}$-relation ($\\la 0.2$~dex).  If we make no prior assumptions on LGRB hosts and accept the result that LGRB hosts follow the usual $L$--$Z_{\\rm O}$-relation, then we can in turn estimate even inidividual host metallicities from this relation. These individual estimates then have errors on the order of the dispersion in the $L$--$Z_{\\rm O}$-relation, i.e. $\\la 0.2$~dex.  In other words, owing to the internal metallicity dispersion, our indirect estimate of the progenitor metallicity via the $L$--$Z_{\\rm O}$-relation of the host galaxy should be almost as useful as an unresolved spectroscopic observation, with an error that is roughly larger by a factor of $\\sqrt{2}$. If the internal metallicity dispersions exceeded 0.2~dex, then the determination of integrated metallicities of individual LGRB hosts would have very little value, because it could as well be estimated from an $L$--$Z_{\\rm O}$-relation, while the error was mostly internal in origin.  Therefore, major progress can mostly be expected from spatially resolved metallicity measurements that focus on the immediate burst environment. However, such observations are currently not possible for GRBs at cosmological distances.   \\subsection{A metallicity--LGRB energy relation?}  \\citet{Sta06} suggest a tentative relation between host metallicity $Z_{\\rm O}$ and the isotropic energy associated with the LGRB $E_{\\rm iso}$, which they interpret as a metallicity cutoff around 0.15 solar for regular cosmological bursts. This conjecture is based on one object out of five, where a high LGRB energy coincides with low host metallicity.While this object was assigned a much higher host metallicity in a previous analysis \\citep{So05}, the mass of the galaxy makes that interpretation very unlikely. While the claimed relation relies on the significance of this single object, it is unclear whether the other (underluminous) LGRBs are representative of the more energetic and distant cosmological bursts. GRB 060218, e.g., could not even have been detected by {\\it Swift} at redshifts $z>0.05$.  Even in the absence of spectroscopic metallicity measurements for the individual hosts, we can tentatively investigate such a relation for the full sample with known $E_{\\rm iso}$ and host luminosities using the $L$--$Z_{\\rm O}$-relation. Fig.~\\ref{EisoZ} shows the result for 13 GRBs from our sample (black data points) which have $E_{\\rm iso}$ values in the literature \\citep{ Am06}. We also add the five local objects discussed by \\citet{Sta06} based on their actual host $Z_{\\rm O}$-data as grey data points.  At least for the cosmological bursts (black) no relation is apparent. However, due to the expected significant variations in jet geometry and viewing angle among the bursts, we would not expect the fiducial isotropic energy estimate to show very clear trends with other parameters. \\citet{GGL04} present LGRB energies $E_{\\rm corr}$ that are corrected for jet beaming effects. Correspondingly, their Ghirlanda relation $E_{\\rm corr}$ vs.  $E_{\\rm peak}$ is much more clearly defined than the original Amati relation $E_{\\rm iso}$ vs.\\ $E_{\\rm peak}$ \\citep{Am02}, where $E_{\\rm peak}$ is the restframe photon energy at the peak in the $\\nu F_\\nu$ spectrum.  In the bottom panel of Fig.~\\ref{EisoZ}, we show $E_{\\rm corr}$ vs.\\ the host metallicity for the bursts with available data. We find no trend in these data.   \\subsection{Theoretical implications}  From a theoretical point of view, it is not surprising that metallicity plays an important role for LGRB progenitors. In one of the most promising progenitor scenarios, the collapsar model (Woosley 1993; MacFadyen \\& Woosley 1999), the progenitor is the rapidly rotating core of a massive star. On the other hand, the core of a massive star is believed to efficiently lose angular momentum during its evolution: both within the star by hydrodynamical (e.g. Heger, Langer \\& Woosley 2000) and magnetohydrodynamical processes (e.g. Spruit 2002) and subsequently from the surface of the star in the form of a stellar wind. As a consequence, it seems unlikely that most massive stars will still have cores at the time of core collapse that rotate as rapidly as is required in the collapsar model (Heger, Woosley \\& Spruit 2005).  Indeed, this is also not necessary, since it is clear that LGRBs are rare events, associated with less than 1 in 1000 core-collapse supernovae (Podsiadlowski et al.\\ 2004b).  One of the key unresolved questions is what are the special circumstances that produce LGRB progenitors: are these special conditions in single stars (e.g. Yoon \\& Langer 2005; Woosley \\& Heger 2006) or do they require particular binary channels (e.g. Fryer \\& Woosley 1998; Izzard, Ramirez-Ruiz \\& Tout 2003; Fryer \\& Heger 2005; Petrovic et al.\\ 2005; Detmers et al.\\ 2006; Fryer et al.\\ 2006; Podsiadlowski et al.\\ 2006)? One of the major problems in many of these models is that mass loss, both in the red-supergiant phase and, in particular, in the Wolf-Rayet phase, leads to very efficient angular-momentum loss. Since these wind mass-loss rates are dependent on the metallicity (typically $\\dot{M} \\propto Z^{0.5-0.7}$; e.g. Vink \\& de Koter 2005), massive stars with lower metallicity are more likely to have rapidly rotating cores at the time of core collapse. This is a generic advantage for many of the proposed models, both single and binary.  Specifically, Yoon \\& Langer (2005) and Woosley \\& Heger (2006) recently proposed that a low-metallicity, very rapidly rotating star may evolve essentially homogeneously during its early evolutionary phases and may avoid a red-supergiant phase altogether, in which most of the angular-momentum loss from the core tends to take place. As a consequence, these stars would still preserve rapidly rotating cores at the time of collapse, fulfilling one of the key conditions in the collapsar model. Similarly, in some of the most promising binary scenarios, in which two massive stars (almost) merge or interact tidally, it is advantageous or even necessary that this occurs late in the evolution of one of the stars, i.e. after helium core burning (so-called case C mass transfer). The range of masses for which case C mass transfer occurs and allows the formation of a black hole is a strong, non-linear function of metallicity (Justham et al.\\ 2006), again favouring low-metallicity progenitors.  As this discussion illustrates, the dependence of the LGRB rate on host metallicity can provide an important constraint on the various proposed progenitor models. Indeed, \\citet{HMM05} and \\citet{YL06} have estimated that, when the effects of magnetic fields are included, the single, low-metallicity, high-rotation model \\citep{YL05,WH06} requires a metallicity less than 1/5th solar. This already appears inconsistent with the constraints derived in this paper, which seem to rule out any models that {\\em require} a metallicity significantly less than 1/2 solar.  The ongoing metallicity {\\it calibration debate} might render our limit to lie $\\sim 0.3$~dex lower, if $T_e$-based methods were correct. There are now explorations of a third method for measuring metallicities, the O~{\\sc II}$_{RL}$ method \\citep{Pe06}, which is aimed at overcoming the shortcomings of both previous methods. Early results indicate that this method gives values in between the two debated versions, but closer to the $R_{23}$ method. This would leave our results largely unchanged."
    },
    "0606/astro-ph0606513_arXiv.txt": {
        "abstract": "{} {We present the Fourier parameter fit method, a new method for spectroscopically identifying stellar radial and non-radial pulsation modes based on the high-resolution time-series spectroscopy of absorption-line profiles. In contrast to previous methods this one permits a quantification of the statistical significance of the computed solutions. The application of genetic algorithms in seeking solutions makes it possible to search through a large parameter space.} {The mode identification is carried out by minimizing $\\chi^2$, using the observed amplitude and phase across the line profile and their modeled counterparts. Computations of the theoretical line profiles are based on a stellar displacement field, which is described as superposition of spherical harmonics and that includes the first order effects of the Coriolis force.} {We made numerical tests of the method on a grid of different mono- and multi-mode models for $0 \\le \\ell \\le 4$ in order to explore its capabilities and limitations. Our results show that whereas the azimuthal order $m$ can be unambiguously identified for low-order modes, the error of $\\ell$ is in the range of $\\pm 1$. The value of $m$ can be determined  with higher precision than with other spectroscopic mode identification methods. Improved values for the inclination can be obtained from the analysis of non-axisymmetric pulsation modes. The new method is ideally suited to intermediatley rotating $\\delta$ Scuti and $\\beta$ Cephei stars.} {} ",
        "introduction": "Asteroseismology relies on a detailed comparison between the set of frequencies observed in a pulsating star and those predicted by theory. To do this successfully, we need to know the quantum numbers associated with each observed pulsation mode, and thus we need {\\it mode identifications} for as many frequencies as possible. The combined approach of spectroscopic identification methods, multicolor photometry, and frequency pattern recognition can identify both the harmonic degree $\\ell$ and the azimuthal order $m$ of a pulsation mode. The parameters of a stellar seismic model, such as the mass, the luminosity, the rotation rate, the effective temperature, the metallicity, and the age, can be tuned to obtain a fit between the theoretical and observed frequencies. In most cases the range of possible parameters is very large. By only considering the models that are consistent with mode identification, one is able to put a strong constraint on the internal structure of the star. For an extensive overview paper about photometric mode identification methods we refer the reader to Garrido (2000), and to Balona \\& Evers (1999) for a paper about the state-of-the-art approach of photometric mode identification.  Profile variations of absorption lines formed in the photosphere can be used to study the pulsational behavior of a star. The line-profile variations (LPV) are caused by the Doppler shift of absorption lines originating from different parts of the surface. High-resolution time-series spectroscopy is a powerful tool for identifying the $\\ell$ and $m$ values of a non-radially oscillating star by analyzing its LPV. In addition, these techniques allow determination of the inclination of the pulsation/rotation axis, the pulsational temperature variations, and the intrinsic oscillation amplitude. If photometric amplitudes in different passbands are also known, it is possible to calculate the ratio of flux perturbation to radial displacement, which is very sensitive to convection, thereby permitting selection between different convection models (Daszynska-Daszkiewicz et al. 2003).  In the past decade, considerable effort has been expended to develop and test {\\it reliable} spectroscopic mode identification methods. Basically, two kinds of approaches have been developed: one based on quantities integrated across the line profile and one based on the intensity variations within each wavelength bin. The first approach is utilized in the moment method (Balona 1986, Aerts 1996, Briquet \\& Aerts 2003) and is applicable to the identification of low-degree modes ($\\ell \\le 3$). Different variants exist of the second approach, which is based on the Doppler image principle and requires sufficient rotational broadening of the profile. These include the intensity-period search (IPS) method (Schrijvers et al. 1997), where $\\ell$ and $m$ are determined from the amplitude and phase distribution across the line profile by applying empirically derived relations, and the pixel-by-pixel method (PPM) described by Mantegazza (2000).  These spectroscopic mode identification methods have one thing in common: no statistical criterion that can quantify the significance of the obtained solutions can be calculated yet. For instance, the moment method (MM) quantifies the goodness-of-fit in elaborately weighted discriminants, which are not related to the observational uncertainties. Therefore it is often not possible to select between different solutions if their discriminants have similar values. De Ridder et al. (2003) report pessimistic prospects for defining such error estimates in the future because of the complex nature of the moments.  We have developed a spectroscopic mode identification method, the Fourier parameter fit (FPF) method, which relies on a fit of the observational and modeled Fourier parameters across the line profile. For every detected pulsation frequency, the zero point, amplitude, and phase are computed for every wavelength bin across the profile by a multi-periodic least-square fit. The observational errors of these parameters can be calculated analytically from the least-square fit. This enables us to carry out a classical chi-square ($\\chi^2$)-test and to derive the statistical significance of our fits.  The paper is constructed in the following way: we first provide the formalism of the displacement field used for the computation of the LPV. The FPF method is presented in the next section, followed by the description of numerical experiments applied for testing the method. ",
        "conclusions": "We have presented a new spectroscopic mode identification method, the FPF method, based on fitting the Fourier parameters of the intensity variations across the line profile. Our description of the displacement field is valid for a slowly rotating non-radially pulsating star and includes toroidal motion, i.e., first-order effects due to the Coriolis force. The implementation allows a modeling of pulsational geometries where the pulsational symmetry axis is not aligned with the stellar rotation axis. Temperature variations of the stellar atmosphere and their impact on the equivalent width variations of a line can be taken into account.  The strength of the FPF method is its ability to estimate the significance of the derived mode parameters by means of an $\\chi^2$ test. For a given error threshold it is possible to exclude a number of tested $\\ell$ and $m$ values for a pulsation mode. Furthermore, the stability of the derived fits can be determined by analyzing the trend of the $\\chi^2$ values across the tested parameter ranges. Compared to previous methods, this permits a statistical quantification of the obtained solutions. We have also shown that the FPF method can determine $m$ with better precision than other methods. This is especially important for improving the treatment of rotational effects in theoretical seismic models (Pamyatnykh 2003).   The presented method is optimized for slowly to intermediately rotating pulsating stars that have a large projected rotational velocity compared to the pulsational radial velocity amplitude. Such objects can be found especially among $\\delta$ Scuti and $\\beta$ Cephei stars (Aerts \\& De Cat 2003). In a subsequent paper we demonstrate the successful application of the FPF method to high-dispersion time-series spectra of the non-radial pulsating $\\delta$ Scuti star FG Vir (HD 106384).  \\appendix"
    },
    "0606/astro-ph0606039_arXiv.txt": {
        "abstract": "We performed a 24 $\\mu$m $2\\degr \\times 1\\degr\\,$ survey of the Pleiades cluster, using the MIPS instrument on \\textit{Spitzer}. Fifty four members ranging in spectral type from B8 to K6 show 24 $\\mu$m fluxes consistent with bare photospheres. All Be stars show excesses attributed to free-free emission in their gaseous envelopes. Five early-type stars and four solar-type stars show excesses indicative of debris disks. We find a debris disk fraction of 25\\% for B-A members and 10\\% for F-K3 ones. These fractions appear intermediate between those for younger clusters and for the older field stars. They indicate a decay with age of the frequency of the dust-production events inside the planetary zone, with similar time scales for solar-mass stars as have been found previously for A-stars. ",
        "introduction": "After a few hundred million years, a planetary system is expected to have assumed its final configuration and has either set the stage for life, or will probably remain barren forever. It is difficult to probe this era. Most of its traces have been obliterated in the Solar System. Only a minority of the nearby stars are so young. Even for them, planets -- and particularly those in the terrestrial planet/asteroidal region -- are faint and are lost in the glare of their central stars. However, when bodies in this zone collide, they initiate cascades of further collisions among the debris and between it and other members of the system, eventually grinding a significant amount of material into dust grains distributed in a so-called debris disk. Because the grains have larger surface area per unit mass compared to larger bodies, they (re)radiate more energy and therefore are more easily detected in the IR compared to their parent bodies. By studying this signal, we can probe the evolution of other planetary systems through this early, critical stage.  Debris disks are found around stars generally older than $\\sim$10 Myr, with no signs of gas accretion (as judged from the absence of emission lines or UV excess) \\citep{Lagrange00,Hillenbrand05}. In the absence of gas drag, a 10 micron-size dust grain from the primordial, protoplanetary nebula cannot survive longer than $\\sim$ 1 Myr within  10 $AU$ of a star due to a number of clearing processes, like sublimation, radiation pressure, Poynting-Robertson and stellar wind drag \\citep{Backman93,Chen05a}. Therefore, any main sequence star older than 10 Myr with an infrared excess is a candidate to have circumstellar material supplied through debris disk processes.  In our Solar System there are two main sources of dust: colliding asteroids and evaporating comets. The dust (grains $<$ 1 mm in size) is concentrated in two belts -- the asteroid belt at 2-3.5 $AU$, $L_{dust}=10^{-7}L_{\\sun}, T_{dust} = 150-200$ $K, M_{dust} =10^{-8}M_{\\earth} $ \\citep{Low84,Dermott02,Hahn02}; and the Kuiper belt at 30-50 $AU$, $L_{dust}=10^{-7}L_{\\sun}, T_{dust} =40-100$ $K, M_{dust} = 10^{-5} M_{\\earth}$ \\citep{Backman95,Morbidelli03}. Asteroid collisions that produce warm dust are driven by perturbations by planets, a process that has cleared most of the interplanetary material from the inner 30 $AU$ of the system. In contrast, there is substantial mass in the Kuiper Belt, concentrated in reasonably stable structures controlled by giant planets \\citep{Liou99, MoroMartin02}. Similarly, rings, blobs, warps and other non-uniformities are frequently observed in the spatially resolved images of nearby extrasolar debris disks and inferred from the spectral energy distributions in the distant unresolved disks, implying the presence of \\textit{planetary mass bodies} in these systems.  The first debris disks were discovered by \\textit{IRAS} in the 1980s. Only a handful of nearby ones have been resolved in optical and near-IR scattered light or detected in submm thermal emission; \\textit{ISO} has extended the \\textit{IRAS} results, but only to a modest extent \\citep[e.g.,][]{Laureijs02}. The limited sensitivity of \\textit{IRAS} and \\textit{ISO} and the even more severe limitations of other methods have prevented a complete census, particularly to photospheric emission levels. In addition, some alleged disks have been confused with background objects or cirrus in the large beam of \\textit{IRAS}. The attempt to assess the percentage of stars with debris-disk emission therefore appeared premature from pre-\\textit{Spitzer} data \\citep{Zuckerman01}.  With superior sensitivity and a smaller beam size, \\textit{Spitzer} is providing many new insights to debris disk behavior. For example, \\citet{Rieke05} demonstrated both the overall decline of debris disks with age (first noted in \\citet{Holland98,Spangler01}) and the large scatter of disk properties at any given age (as previously noted in \\citet{Decin03}). By probing excesses to within 25\\% of the photospheric emission, \\citet{Rieke05} found a surprising number of non-excess stars even at ages as young as 10 - 20 Myr, implying very rapid clearing of the inner $10-60 AU$ region in these systems.  The study of Rieke et al. was purposely confined to A stars, with masses $\\sim 2.5 M_{\\odot}$. It poses the question of whether the same patterns of disk evolution hold for less massive stars, especially in the \\textit{solar} mass range -- spectral types FGK. The study of excesses in young solar analogs could corroborate lunar geological evidence of the bombardment history in the Solar System. Two major surveys of solar type stars in the field are under way with \\textit{Spitzer} -- the solar star survey by the MIPS GTO team, searching for disks around $\\sim$150 field stars within 25 pc \\citep{Bryden06}, and the Legacy survey by the FEPS group, searching for disks around $\\sim$330 stars, both field and members of nearby open clusters and associations \\citep{Meyer04}. Interpretation of these surveys is still in progress; nevertheless, preliminary results indicate that among solar analogs typically older than 1 Gyr $\\sim$10-15\\% have excess emission at 70 $\\mu$m, indicative of Kuiper Belt analogs $\\sim$10-100 times more massive than in the Solar System \\citep{Kim05,Bryden06}. On the other hand, 24 $\\mu$m excesses that would be indicative of massive zodiacal belts are much more rare around these stars -- $\\sim$1\\% \\citep{Beichman05b,Beichman06}.  To explore younger systems with good number statistics and small uncertainties in age determination, an open cluster survey has been initiated by the MIPS and IRAC GTO groups. The first clusters studied in this program were the 25-35 Myr old NGC 2547 \\citep{Young04} and the 60-100 Myr old M47 (NGC 2422) \\citep{Gorlova04}. At 24 $\\mu$m photospheres were detected in NGC 2547 down to mid-F and in M 47 to early G spectral types, beginning to probe disks in the solar-mass regime. A decrease with age in both the amount and the fraction of excess (from $\\sim$40\\% to 25\\%) is indicated for B-A stars between the two clusters (as included in Rieke et al. 2005). For F-K stars a similar trend is only suggested: quantitative estimates are impossible because photospheres start dropping below the detection limit. Thus, it is not clear whether the high incidence of 24 $\\mu$m excesses around A stars is a product of their high luminosities and efficiency in heating dust, or simply results because they are by definition relatively young ($<$ 1 Gyr).  To probe the evolution of debris disks around solar-mass stars, we have investigated the Pleiades. The Pleiades is one of the best studied clusters with membership extending into the substellar regime \\citep[e.g.,][]{Moraux03,Pinfield03}. At 24 $\\mu$m {\\it Spitzer} is potentially sensitive enough to detect photospheres in this cluster down to early M spectral type. To exploit these advantages, one must contend with an interstellar cloud of gas and dust, whose thermal emission beyond 10 $\\mu$m dominates over starlight in the South-West part of the cluster. There is no consistent picture of the structure and kinematics of this nebulosity, except that it might be a part of the $\\sim$10 Myr Taurus-Auriga starforming complex through or behind which the Pleiades happens to be passing at the present epoch \\citep[e.g.,][]{Breger87}. Recent studies however agree that at least some of the material is located within the cluster and possibly interacts with the brightest members \\citep{Gibson03,White03}. Dust in the molecular cloud is sheared by the radiation pressure and assembles into filaments that wrap around stars, forming a characteristic pattern of reflection nebulosity known as the \"Pleiades phenomenon\" \\citep{Arny77,Herbig01}. The unresolved thermal emission from the interstellar dust shell can mimic emission from a debris disk, as was for example demonstrated with coronagraphic imaging of the \\textit{IRAS-}selected Vega-like stars by \\citet{Kalas02}. Special measures were undertaken in our survey to avoid confusion with interstellar dust (\\S\\ref{red},\\ref{sample}).  Pre-\\textit{Spitzer} IR and submm \\citep{Zuckerman93} searches for Pleiades disks were largely unsuccessful. The few reported detections by \\textit{IRAS} \\citep{Castelaz87} and \\textit{MSX} \\citep{Kraemer03} concern the brightest B type members, which we find to be either surrounded by heavy cirrus or emission line stars with gaseous disks. \\citet{Spangler01} observed 14 low mass members with \\textit{ISO}, and claim detection (and excesses) at 60 and/or 90 $\\mu$m in the two G0 and K2 members HII 1132 and HII 3163. The fractional dust luminosity in these stars, if accurate, is unusually high -- about $10^{-2}$, comparable to the transitional disks found around $\\lesssim$10 Myr-old stars. Unfortunately both stars are outside of our \\textit{Spitzer} fields. Recently \\citet{Stauffer05} within the framework of the FEPS project investigated 19 solar-type members outside of the cluster core. They found excess emission at 24 $\\mu$m at levels up to 40\\% above the photosphere in a few members. Compared to the FEPS study, ours is centered on the cluster core, covers a larger area, and considers all members independent of spectral type or binarity. It is a magnitude- and background-limited survey based on the membership derived from the optical studies. For the first time, it provides adequate number statistics to constrain the disk incidence around Pleiades members similar in mass to the Sun. ",
        "conclusions": "We have conducted a photometric survey for warm dusty disks in the Pleiades, the nearest medium-age ($\\sim$100 Myr) open cluster. At this age, terrestrial planet formation should be complete, and the inner few $AU$ region is expected to be largely cleared of dust. The dust however can be temporarily replenished, either from occasional asteroid collisions or from the outer regions where giant planet migration may still be occurring. Indeed, we find 24 $\\mu$m excess emission around nine B9-K0 non-emission stars. The fraction of B-A stars with excesses is 25\\%, comparing well to a similar age cluster M 47 \\citep{Gorlova04}.  Because of the proximity of the Pleiades, for the first time we can explore excesses in solar-mass stars down to the photospheric limit. Combined with the \\citet{Stauffer05} sample, we find that the incidence of 24 $\\mu$m excesses for these stars, at 10\\%, is significantly higher than that for old field solar-mass stars, 1\\%. Thus, the clearing of debris in the planetary zone ($\\sim$1 -- 20 $AU$) for these stars takes a similar time as the similar process in A-stars, of order 100 million years. Comparing with the incidence of 70 $\\mu$m excess in $>$1 Gyr old solar mass stars, 13$\\pm$3\\%, it appears that the outer, Kuiper Belt - like zones of planetary debris clear much more slowly that the planetary zones seen at 24 $\\mu$m."
    },
    "0606/hep-ph0606246_arXiv.txt": {
        "abstract": "We examine the dependence of event rates at neutrino telescopes on the neutrino-nucleon cross section for neutrinos with energy above 1~PeV, and contrast the results with those for cosmic ray experiments. Simple scaling of the Standard Model cross sections leaves the rate of upward events essentially unchanged. Details, such as detector depth and cross section inelasticity, can influence rates. Numerical estimates of upward shower, muon and tau event rates in the IceCube detector confirm these results. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606635_arXiv.txt": {
        "abstract": "We present the results of rest-frame ultraviolet spectroscopic observations of a sample of 14 $z\\sim 3$ star-forming galaxies in the SSA22a field. These spectra are characterized by unprecedented depth in the Lyman-continuum region. For the first time, we have detected escaping ionizing radiation from individual galaxies at high redshift, with two of the 14 objects showing significant emission below the Lyman limit. We also measured the ratio of emergent flux density at 1500~\\AA\\ to that in the Lyman-continuum region, for the individual detections (C49 and D3) and the sample average. If a correction for the average IGM opacity is applied to the spectra of the objects C49 and D3, we find $f_{1500}/f_{900}$$,_{corr,C49}=4.5$ and $f_{1500}/f_{900}$$,_{corr,D3}=2.9$. These numbers imply Lyman-continuum escape fractions at least as large as that presented in \\citet{steidel2001}, measured from a composite spectrum of 29 Lyman Break Galaxies (LBGs). The average emergent flux-density ratio in our sample is $\\langle f_{1500}/f_{900},_{corr} \\rangle $$=22$, implying an escape fraction $\\sim 4.5$ times lower than inferred from the \\citeauthor{steidel2001} composite spectrum. If this new estimate is representative of LBGs, their contribution to the metagalactic ionizing radiation field is $J_{\\nu}(900) \\sim  2.6 \\times  10^{-22} \\mbox{ erg s}^{-1}\\mbox{cm}^{-2}\\mbox{Hz}^{-1}\\mbox{sr}^{-1}$, comparable to the contribution of optically-selected quasars at the same redshift. The sum of the contributions from galaxies and quasars is consistent with recent estimates of the level of the ionizing background at $z\\sim 3$, inferred from the H~I Ly$\\alpha$ forest optical depth. There is significant variance among the emergent far-UV spectra in our sample, yet the factors controlling the detection or non-detection of Lyman-continuum emission from galaxies are not well-determined. Specifically, the two objects with detections are not differentiated in a consistent manner from the remainder of the sample in terms of their spectroscopic or photometric properties at wavelengths longer than the Lyman limit. There are also differences in the average emergent far-UV spectra of current and previous samples used to estimate the escape fraction from star-forming galaxies. Because we do not yet understand the source of this variance, significantly larger samples will be required to obtain robust constraints on the galaxy contribution to the ionizing background at $z\\sim 3$ and beyond. ",
        "introduction": "\\label{sec:intro} The metagalactic ionizing background determines the evolution of the physical state of the intergalactic medium (IGM). The ionization rates of H~I and He~II are controlled by both the amplitude of the radiation field and its spectral shape between 1 and 4 Rydbergs. These properties of the ionizing background reflect the total and relative contributions of quasars and star-forming galaxies. It is possible to estimate the nature of the background from the quasar proximity effect \\citep{bajtlik1988,scott2000}, and also from the evolution of optical depths of H~I and He~II in the Ly$\\alpha$ forests of both those species, which are probed by the spectra of bright quasars \\citep{shull2004,bolton2005,bolton2006, mcdonald2001,fan2005}. Such studies indicate that, while the comoving ionizing emissivity decreases between $z=3$ and $z=5$, the decrease in the comoving space density of quasars is much steeper, and that, by $z\\sim 5$, quasars cannot dominate the ionizing background \\citep{fan2001,fan2005,mcdonald2001}. Even at $z\\sim 3$, the mean H~I optical depth, and the mean and variation in the ratio of He~II and H~I column densities, $N(\\mbox{He~II})/N(\\mbox{H~I})$, may indicate a substantial contribution from star-forming galaxies \\citep{bolton2005,bolton2006,shull2004}. Since Lyman-continuum radiation from galaxies appears to constitute an increasingly important ingredient of the ionizing background at redshifts beyond $z\\sim 3$, and may dominate the radiation field at the epoch of H~I reionization \\citep{fan2001,spergel2003}, estimating the galaxy contribution directly represents a fundamental goal for observational cosmology.  One inherent complexity in such an estimate stems from the unknown value of $f_{esc}$, the escape fraction of ionizing photons from the interstellar medium of a star-forming galaxy. Knowledge of this escape fraction is required to convert the star-formation rate density of a population of galaxies into a comoving ionizing emissivity, a crucial component of models of reionization \\citep{furlanetto2004,loeb2001}. Theoretical models of the propagation of ionizing radiation through the ISM of disk galaxies similar to the Milky Way predict an escape fraction of $f_{esc} \\leq 0.1$ \\citep{dove1994,dove2000}, though this fraction may increase for high-redshift galaxies with active star-formation rates, in which supernova explosions increase the ISM porosity \\citep{clarke2002} and large-scale outflows may clear out chimneys through which ionizing radiation can easily escape \\citep{fujita2003}. The model predictions for $f_{esc}$ depend on the detailed structure and dynamics of the ISM, driven by the effects of supernova explosions and stellar winds. Currently, we have only fairly crude data on the properties of the ISM in distant galaxies \\citep{pettini2001,shapley2003}, with minimal spatially-resolved information.  Another approach is to measure directly the emergent Lyman-continuum radiation from galaxies. {\\it Hopkins Ultraviolet Telescope}  and {\\it Far-Ultraviolet Spectroscopic Explorer} ($FUSE$) observations of a small sample of nearby starburst galaxies mostly yielded upper limits for $f_{esc}$ of less than 0.10 \\citep{leitherer1995,hurwitz1997,deharveng2001}. Only recently, \\citet{bergvall2006} reported the first {\\it FUSE} detection of Lyman-continuum emission from a local dwarf starburst galaxy, Haro~11, which implies $f_{esc} \\sim 0.04-0.10$, and will provide a useful low-redshift baseline for comparing with the observations at high redshift.  There has been some degree of controversy over the value of $f_{esc}$ in galaxies at $z\\sim 3$. Escaping Lyman-continuum flux was apparently detected by \\citet{steidel2001} in a composite spectrum of 29 Lyman Break Galaxies (LBGs) drawn from the high-redshift tail of the $z\\sim 3$ LBG selection function. If the ratio of Lyman-continuum to $1500$~\\AA\\ flux density observed in this composite spectrum is typical for the entire LBG sample, it implies $f_{esc}\\sim 0.1$, after correction for the average IGM opacity and the estimate that only $15-20$\\% of 1500 \\AA\\ photons escape from typical LBGs \\citep{adelberger2000}. However, the spectrum may not be representative of typical LBGs, as it reflects the mean emergent spectrum for LBGs with bluer than average rest-frame UV colors and, therefore, objects less affected by dust extinction. Other attempts to measure escaping ionizing radiation from LBGs have yielded only upper limits. \\citet{giallongo2002} targeted two objects drawn from the bright end of the LBG luminosity function, obtaining an upper limit on the escape fraction that is 4 times lower than the detection in \\citet{steidel2001}. This limit may not be as stringent if the two individual sightlines pass through higher than average H~I optical depth in the Lyman-continuum region -- i.e. application of the average IGM correction factor for $z\\sim 3$ may not necessarily be correct in individual cases. \\citet{fernandez2003} fit the broad-band spectral energy distributions (SEDs) of 27 galaxies at $1.9 \\leq z \\leq 3.5$, including the effects of intrinsic Lyman-limit absorption, and intervening Lyman series and Lyman limit absorption, and found on average $f_{esc} \\leq 0.04$. It should be noted that there are  significant uncertainties associated with this method, because of the use of the broad {\\it HST}/WFPC2 $F300W$ filter to probe the Lyman limit at $z\\sim 2-3$. For galaxies towards the low redshift end of the sample, galaxy flux from above the Lyman limit falls in this filter, while for galaxies towards the high redshift end of the sample, the level of flux within the filter will be modulated primarily by the effects of IGM absorption, providing only crude information on the flux level just below the Lyman limit. Using a narrow-band filter more finely tuned to probe just below the Lyman limit for two galaxies at $z\\sim 3$, \\citet{inoue2005} also found only upper limits to the escape fraction, but imaging with greater depth will be required to test if the limits are in contradiction to the \\citet{steidel2001} detection.  Because of the increasing IGM optical depth at higher redshift, it is not possible to measure escaping ionizing radiation directly from galaxies much beyond $z\\sim 3$. Understanding how $f_{esc}$ relates to other galaxy properties that {\\it can} be measured at $z\\sim 6$ is required for assessing the contribution of galaxies to the ionizing background at the epoch of reionization. Therefore it is crucial to resolve the controversy over the escape fraction from galaxies at $z\\sim 3$. In an effort to address this question, we have undertaken a program of deep spectroscopy of a sample of 14 LBGs, drawn from the heart of the LBG redshift selection function ($\\langle z \\rangle = 3.06$). The mean and range of UV colors in this sample are representative of the total sample of LBGs, in contrast to the bluer than average sample of \\citet{steidel2001}. Spectra were obtained with the upgraded LRIS spectrograph on the Keck~I telescope, which provides unparalleled sensitivity in the $3500-5000$~\\AA\\ wavelength range \\citep{steidel2004}. The depth in the Lyman-continuum region for individual spectra in our sample is comparable to that in the composite spectrum of \\citet{steidel2001}. In order to correlate the physical properties of the ISM in these galaxies with the degree of Lyman-continuum escape, we have also obtained sensitive observations redwards of Ly$\\alpha$ for our sample, covering strong low- and high-ionization interstellar absorption features at twice the resolution of typical LBG discovery spectra.  For the first time, we present the detection of escaping ionizing radiation for {\\it individual} galaxies at $z\\sim 3$. Considering our two detections and 12 non-detections, we find significant variation in the escape fraction from galaxy to galaxy, and an average escape fraction which is significantly lower than the one presented in \\citet{steidel2001}. Samples an order of magnitude larger are required to understand the cause of variation among the emergent far-UV spectra of LBGs in terms of stellar populations, morphology, orientation, and ISM physical conditions, and to generalize to the global contribution of galaxies to the ionizing background. In \\S~2, we present our observations and reductions, while \\S~3 contains the main empirical results. In \\S~4, we discuss the emergent far-UV spectra of galaxies in our sample, and conclude in \\S~5 with the implications for the contribution of galaxies to the ionizing background. A cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $h=0.7$ is assumed throughout. ",
        "conclusions": ""
    },
    "0606/astro-ph0606403_arXiv.txt": {
        "abstract": "{ We have studied the optical spectra of a sample of 31 O- and early B-type stars in the Small Magellanic Cloud, 21 of which are associated with the young massive cluster \\ngc346. Stellar parameters are determined using an automated fitting method (Mokiem et al.\\ 2005), which combines the stellar atmosphere code \\fastwind\\ (Puls et al.\\ 2005) with the genetic algorithm based optimisation routine \\pikaia\\ (Charbonneau 1995). Comparison with predictions of stellar evolution that account for stellar rotation does not result in a unique age, though most stars are best represented by an age of 1--3~Myr. The automated method allows for a detailed determination of the projected rotational velocities. The present day \\vsini\\ distribution of the 21 dwarf stars in our sample is consistent with an underlying rotational velocity (\\vr) distribution that can be characterised by a mean velocity of about $160 - 190~\\kmsec$ and an effective half width of $100 - 150~\\kmsec$. The \\vr\\ distribution must include a small percentage of slowly rotating stars. If predictions of the time evolution of the equatorial velocity for massive stars within the environment of the SMC are correct (Maeder \\& Meynet 2001), the young age of the cluster implies that this underlying distribution is representative for the initial rotational velocity distribution. The location in the Hertzsprung-Russell diagram of the stars showing helium enrichment is in qualitative agreement with evolutionary tracks accounting for rotation, but not for those ignoring \\vr. The mass loss rates of the SMC objects having luminosities of $\\log \\lstar/\\lsun \\gtrsim 5.4$ are in excellent agreement with predictions by Vink, de Koter \\& Lamers (2001). However, for lower luminosity stars the winds are too weak to determine \\mdot\\ accurately from the optical spectrum. Two of three spectroscopically classified Vz stars from our sample are located close to the theoretical zero age main sequence, as expected.  Three additional objects of lower luminosity, which are not given this classification, are also found to lie near the ZAMS. We argue that this is related to a temperature effect inhibiting relatively cool stars from displaying the spectral features characteristic for the Vz luminosity class.} ",
        "introduction": "Mass loss and rotation play a crucial role in the evolution of the most massive stars. These two processes are linked. To illustrate this, fast rotation may lead to an enhanced mass loss, as is, for instance, observed in the case of Be stars \\citep{lamers91} and $\\eta$\\,Carinae \\citep{smith03, vanboekel03}. Conversely, the stellar winds of these stars will cause loss of angular momentum, leading to spin down. In a Galactic environment the spin down of massive stars is expected to be relatively rapid, essentially ``erasing'' their initial rotational velocity properties within only a few million years \\citep{meynet00}. For an environment characteristic of the Small Magellanic Cloud (SMC) it is predicted that the initial conditions of rotational velocity remain preserved during the main sequence life of O-type stars \\citep{maeder01}, as the strength of stellar winds diminishes with decreasing metal content \\citep[e.g.][]{abbott82, kudritzki87, vink01}.  The study of a young SMC cluster containing a substantial population of O-type main sequence stars would allow the role of metallicity on mass loss and angular momentum loss during the early evolution of massive stars to be better constrained. This provides a better insight into the current and perhaps even initial rotational conditions from which this evolution unfolds. Among others, this would provide important constraints on the physics of the formation of massive stars, for instance on the role of magnetic fields and the lifetime of accretion disks \\citep[see e.g.][]{porter03}.  The distribution of projected rotational velocities \\vsini\\ of Galactic early-type stars has been studied by \\cite{conti77b}, \\cite{penny96} and \\cite{howarth97}. The first extragalactic studies of rotational velocities have been presented by \\cite{keller04} for a sample of Large Magellanic Cloud (LMC) B-type stars in the vicinity of the main sequence turnoff and by \\cite{penny04} for predominately giant and supergiant SMC and LMC O-type stars. So far, no systematic study for early main sequence O-type stars in a low metallicity environment has been attempted.  The young cluster \\ngc346 in the Small Magellanic Cloud (SMC) contains a substantial population of O-type main sequence stars. In the context of the {\\em VLT-FLAMES Survey of Massive Stars} \\citep{evans05}, we have used the {\\em Fibre Large Array Multi-Element Spectrograph} at the {\\em ESO Very Large Telescope} to obtain spectra of several dozen dwarf O-type stars in this cluster. Our genetic algorithm based fitting method \\citep{mokiem05} is used to perform a homogeneous spectroscopic analysis of this set of stars, including the projected rotational velocity as one of the fitting parameters. The \\vsini\\ information obtained from this analysis may be compared to models of the present and possibly initial (see above) \\vsini\\ distribution. This is a main aim of this paper.  The strength of radiatively driven winds is predicted to be reduced for decreasing metal content. In the last two decades this prediction has been quantified by different groups who found a metallicity dependence of the mass loss rate of $\\mdot(Z)\\propto Z^{0.5 - 0.7}$ \\citep[e.g.][]{kudritzki87, puls00, vink01}. Qualitatively this metallicity dependence has been confirmed by several authors \\citep[e.g.][]{puls96, bouret03, evans04b, massey05}. However, until now a quantitative comparison of the theoretically predicted and the empirically determined $Z$ dependence is still lacking. The analysis of our SMC targets will provide insight in the wind characteristics of objects with a metal content approximately five times lower than in Galactic objects. Consequently, given the large number of objects that we have analysed in a homogeneous way, the current study will for the first time be able to provide such a quantitative comparison.  While comparing mass loss rate predictions with observed values it is important to realise that recent studies of the wind strengths of not too luminous ($\\log \\lstar/\\lsun \\lesssim 5.3$) Galactic and SMC dwarf O-type stars \\citep{bouret03, martins04, martins05b} seem to indicate that the mass loss rates are significantly (up to two orders of magnitude) lower than expected from theory \\citep*{vink01}. If this is indeed the case it would imply that our chances of actually observing a \\vr\\ pattern that is representative of the initial distribution will increase, as less angular momentum will be carried away in the stellar wind. However, from the viewpoint of our understanding of the physics of relatively weak stellar winds ($\\sim$$10^{-8}$ -- $10^{-6} \\msunyr$) the situation is obviously less desirable.  The mass-loss rates reported by \\citeauthor{bouret03} and \\citeauthor{martins05b} are based on an analysis of unsaturated ultraviolet resonance lines. The \\mdot\\ sensitivity that can be obtained using such lines is $\\sim$$10^{-9}$ -- $10^{-8} \\msunyr$. This is, in principle, much better than is possible using \\ha\\ as the mass loss diagnostic. Using our automated fitting method and spectra having a signal-to-noise of 50--200 (typical for those secured within the context of our VLT-FLAMES programme) we can push the \\ha\\ method as low as $\\sim$$2\\times 10^{-7}\\,\\msunyr$. This provides some overlap with the weak wind regime so far reserved for the UV line method. It thus allows for a first investigation of the question whether the so-called ``weak wind problem'' points to missing physics in the theoretical predictions of mass loss or that there may be problems with the mass loss diagnostic, i.e.\\ that we do not fully understand the formation of UV resonance lines.  The structure of this paper is as follows: in Sect.~\\ref{sec:data} we describe the data set that has been analysed using our genetic algorithm based fitting method, which is discussed in Sect.~\\ref{sec:method}. The stellar properties of our sample determined in this analysis are presented in Sect.~\\ref{sec:fun_par}. In sections~\\ref{sec:par_vr_dist}, \\ref{sec:weak_winds} and \\ref{sec:age} we investigate and discuss, respectively, the underlying rotational velocity distribution of our sample, the mass loss rates we have determined in the context of weak winds and the evolutionary status of the cluster \\ngc346. Sect.~\\ref{sec:sum} summarises and lists the conclusions of our study. Finally, in the appendix the fits and comments on the individual objects are provided. ",
        "conclusions": "  \\begin{enumerate}  \\item[\\it i)] The average effective temperature for a given spectral type for SMC dwarf stars is hotter by typically 3 to 4~kK than for their Galactic counterparts \\citep{martins05a}. We attribute this hotter temperature scale to a reduced line blanketing in the SMC. In contrast to \\cite{massey05} we find that this blanketing effect on \\teff\\ extends to spectral types as late as B0.  \\item[\\it ii)] The helium enrichment found in our stars does not appear to be consistent with predictions of the evolution of massive stars ignoring rotation. Comparison with evolutionary tracks that do account for rotation shows a qualitative agreement. In the HR-diagram the region in which a helium enrichment is predicted roughly coincides with the location of evolved objects showing evidence of \\yhe\\ enrichment. The actual degree of enrichment is in many cases still strongly under predicted by the most recent evolutionary models.  \\item[\\it iii)] From our three programme stars with a Vz luminosity class designation, we find that only two are true ZAMS objects. Moreover, three additional objects, which are not classified as Vz stars, are found to lie on or left of the ZAMS. We argue that this is related to a temperature effect inhibiting relatively cool stars located on the ZAMS to exhibit the spectral features characteristic for the Vz luminosity class.  \\item[\\it iv)] We find good agreement between our spectroscopically derived masses and those derived from evolutionary calculations. A mild mass discrepancy is still found for stars with $\\Ms \\lesssim 20\\,\\msun$. This discrepancy correlates with the surface helium abundance and is qualitatively consistent with the predictions of chemically homogeneous evolution.  \\end{enumerate}  \\noindent The stars in the young cluster \\ngc346 allow to study the role of mass loss and angular momentum loss during the early evolution of massive stars in a low metallicity environment. Regarding these issues we can conclude that:  \\begin{enumerate} \\item[\\it v)] The observed distribution of \\vsini\\ for evolved stars (luminosity classes I-II) contains relatively less fast rotators and slow rotators compared to the distribution for unevolved stars (luminosity classes IV-V). These findings are in agreement with analyses of Galactic samples and can be interpreted as a spin down due to an increased radius and angular momentum loss through the stellar wind, and excess turbulent broadening among the evolved stars \\citep[cf.][]{howarth97}. We also find that compared to the velocity distribution of unevolved Galactic objects the distribution for unevolved SMC objects shows a relative excess of stars rotating with projected velocities $120 \\gtrsim \\vsini \\lesssim 190\\,\\kmsec$. This excess can be interpreted as evidence of a reduction of dissipation of angular momentum in weaker stellar winds resulting from the low $Z$ environment of the SMC.  \\item[\\it vi)] Using a simple modelling approach we have constrained the underlying \\vr\\ distribution of the unevolved SMC stars. The exact shape of this underlying distribution -- for which we have tried a block, Gaussian and Maxwellian shape -- is found to be degenerate, however, it can be characterised by a mean velocity of about $150-180$~\\kmsec\\ and an effective half width of roughly $100-150$~\\kmsec. The Maxwellian distribution appears a somewhat poorer representation of the \\vr\\ distribution, indicating that 5--15 percent of the stars must be slow rotators ($\\vr \\lesssim 50$~\\kmsec).  \\item[\\it vii)] It is not possible to derive a unique age for \\ngc346 using moderately rotating or non-rotating evolutionary models. Though the bulk of the stars indicate a lifetime of 1--3~Myr, part of the investigated stars seem somewhat older (3--5 Myr). Clearly the ages depend on the adopted evolutionary tracks. However, we point out that predictions for non-rotating stars and those accounting for the rotational evolution of stars with an initial $\\vr = 300$~\\kmsec\\ essentially give the same result. Three main sequence stars, all with an age determination of less than 2 Myr, show significant helium enhancement. Assuming that these stars have the same initial \\yhe\\ this implies a lower limit to their age of 2 to 3.5~Myr.  \\item[\\it viii)] Based on the predictions of the time evolution of the equatorial velocity for low metallicity \\citep{maeder01} the above age considerations imply that the underlying \\vr\\ distribution derived for our sample (conclusion {\\it vi}) is representative for the initial rotational velocity distribution.  \\item[\\it ix)] Mass loss rates determined for the SMC objects are found to be systematically lower than for Galactic objects. {\\em For $\\log\\lstar/\\lsun \\gtrsim 5.4$ we find that the wind strengths are in excellent agreement with the theoretical predictions of \\cite*{vink01}}. The empirical modified wind momentum luminosity relation constructed for this luminosity regime agrees to within the error bars with the predicted WLR.  \\item[\\it x)] Down to $\\log \\lstar/\\lsun \\approx 5.3$ the mass loss rates determined from the optical spectra are in agreement with UV rates. For luminosities lower than this value the winds become too weak for an accurate determination of \\mdot\\ from the optical spectrum only. Consequently, we cannot (yet) conclude whether or not the UV and \\ha\\ diagnostics agree or disagree in weak wind situations.  \\end{enumerate}"
    },
    "0606/astro-ph0606129_arXiv.txt": {
        "abstract": "Vega may have a massive companion in a wide orbit, as evidenced by structure in its cold dust debris.  We have tested this hypothesis by direct imaging with adaptive optics in the M band.  The observations were made with a newly commissioned thermal infrared camera, Clio, on the 6.5 MMT AO system with low-background deformable secondary.  The observations constrain a planet to be less than 7 M$_J$ at the approximate position angle expected from the dust structure and at a radius $>$ 20AU (2.5 arcsec) .  This result is more stringent than similar previous near-infrared observations of Vega, that achieve limits of 20 and 10 M$_J$ at separations of 7 arcsec.  The higher sensitivity is due both to the more favorable contrast of gas giant planets at M band and to the higher Strehl and more stable point spread function at longer wavelengths.  Future L' or M band observations could provide a powerful approach for wide separation planet detection, especially for cooler, and thus older or less massive planets. The natural best targets are nearby stars where planets in the range of 5-15 M$_J$ and as old as several Gyr are expected to be detectable with this technique. ",
        "introduction": "Direct detection of extrasolar planets is a highly desirable goal for a range of reasons.  Although planet detection through its gravitational influence on its star has yielded important information about planets around other stars, many parameters of a planet are most easily derived from its spectral energy distribution (SED) including its temperature, size, and composition.  For planets in orbit beyond approximately 5-10 AU, the length of time needed to detect a planet through indirect means also becomes prohibitive.  Yet these are precisely the radii where we see massive planets in our own solar system.  Initial success in the area of direct detection is just now occurring, although the detectable objects are still significantly different from what we think of as a typical solar system.  Measurements of secondary eclipses of transiting hot Jupiters can begin to constrain the temperature and albedo of these objects (Deming et al. 2005,  Charbonneau et al. 2005, Rowe et al. 2005).  For younger stars and wider separations faint companions have been detected (Chauvin et al. 2004, Neuhauser et al. 2005) which are consistent with being planetary mass objects. However, these objects are unlikely to have been formed by either core accretion or gravitational instability (Boss 2006) suggesting that we are either seeing scattered planets or low mass objects that have formed by fragmentation.  Direct detection searches are typically focusing on very young stars where giant planets are early in the process of cooling and contracting and thus are expected to be relatively bright in the near infrared ( see Metchev et al. 2004, Oppenheimer et al. 2004, Mugrauer et al. 2005, Lowrance et al. 2005, and Biller et al. 2005 as examples of ongoing surveys).  Of these bands the H band is particularly attractive due to the methane absorption feature which creates an identifiable spectral feature, as well as its relative brightness compared to the K band.  The majority of the direct detection plans with large telescopes focus in this spectral region, utilizing techniques such as simultaneous spectral difference imaging (Marois et al. 2000, Biller et al. 2005) and higher order adaptive optics (Oppenheimer et al. 2004, Macintosh et al. 2003a) to optimize the achievable contrast.  Several observatories (see Gratton et al. 2004 as an example) are currently designing ambitious near-infrared instruments to pursue the detection of young planets.  Although near-infrared sensitivity from the ground is better than at longer wavelengths, a significant contrast advantage can be achieved by looking for cool giant planets closer in wavelength to the peak of their expected SED.  Giant planets have very non-blackbody SEDs, requiring a knowledge of their spectral characteristics to best choose an optimum wavelength.  Additionally, the sensitivity of a ground-based system is dominated by the spectral dependence of the atmospheric transparency and brightness, limiting possible detection to discrete spectral windows.  It has long been known  that Jupiter's SED displays an anomolous peak at approximately 4-5 microns (Gillett, Low, and Stein  1969). This is due to the lack of absorption features in this spectral region which allows the observation of thermal emission from much deeper in the planet's atmosphere than at other wavelengths.  A similar, broader peak can be seen in T type brown dwarfs (Oppenheimer et al. 1998) and is an expected generic feature of objects for objects with effective temperatures below approximately 1200 K, as modeled by Burrows et al. (1997, 2003) and Baraffe et al. (2003).  The high flux appears in the L' and M bands for hotter object and narrows to a feature which is well matched to the M band window for the coolest objects.  The relative flux in L' versus M bands for hotter objects, as measured by Golimowski et al. (2004) appear to be dependent on the amount of carbon monoxide in the atmosphere due to non-equlibrium effects, lessening the expected brightness at M compared to equilibrium models. Such effects may determine whether the M or L' atmospheric window is preferable, but the existence of a broad hump in the SED in the 4-5 micron region appears to be secure both from model predictions and observed cool objects.  To make use of this expected contrast enhancement it is necessary to be able to detect the relatively faint planet flux in the presence of background from the sky as well as any warm optics in the system.  The relatively small expected separations require diffraction-limited performance of the optics, which in turn requires an adaptive optics system to correct for atmospheric turbulence.  Typical AO systems add 5-10 warm surfaces. For wavelengths longward of approximately 2 microns the infrared glow from the warm optics can dominate the background and, thus, the noise of an optical system.  For planet detection at longer wavelengths an optimum system is one which is integrated into the telescope itself, such as the deformable secondary mirrror of the MMT AO system (Riccardi et al. 2002, Brusa et al. 2003).  Optimal targets for direct detection of extrasolar planets are stars which are both young and nearby.  Unfortunately most radial velocity targets do not satisfy both of these criteria.  Target stars are typically chosen to avoid potential noise in a Doppler signal from stellar activity (Wright et al. 2004).  For young stars, a possible alternative indicator of the existence of a planetary companion is substructure in a debris disk (Roques et al. 1994, Liou \\& Zook 1999). Giant planets will tend to set up resonances in debris material around a star with the chracterisitics of the resonances dependent on parameters of the planet's mass and obrbit  (Kuchner \\& Holman 2003).  If this substructure can be modeled properly it may be possible to use such information to guide where and which stars to search for planetary companions.  In this letter we present initial observations with Clio (Freed et al. 2004), a camera designed specifically to detect giant planets through their 3-5 micron infrared radiation.  We have used the imager to constrain the existence of a planetary companion to Vega, at the orientation expected from resonances in the cold debris disk.    Section 2 describes the instrumental setup.   In section 3 we present the observations.  Section 4 details the data reduction and achievable limits versus separation.  In Section 5 we discuss the implications of these initial observations for future planet detection.  \\subsection{The Vega System}  Vega provides an interesting target for planet searches as one of the nearest young stars with a debris disk indicative of planetesimals. For the purpose of planet models used in this paper we adopt an age of 300 Myr and a distance of 7.8 pc consistent with evolutionary tracks (Song et al. 2001) and Hipparcos astrometry. The debris disk around Vega has been intensively studied since its discovery with IRAS (Aumann et al. 1984).   It appears that we are viewing a cold (Backman \\& Paresce 1994), probably transiently bright (Su et al. 2005) dust disk equivalent to the Kuiper belt at nearly pole-on orientation around Vega.   Sub-mm and millimeter observations of Vega have revealed a double-lobed enhancement to the ring (Holland et al., 1998, Koerner et et al., 2001, Wilner et al., 2002, hereafter W02).    The offset nature of the lobes seen at long wavelengths have been modeled by W02. They reproduce the observed structure for a planet which is approximately 90 $\\deg$ from the line connecting the lobes, on the same side of the star as the offset of the lobes.  They suggest that the planet is 7\" from the star in the northwest direction.  Interestingly, the observed structure seen around Vega is similar to one of the four broad classes of dust structure described by Kuchner and Holman (2003) which can arise from giant planet perturbations to debris disks.  The class which fits the double-lobed pattern of Vega is indicative of a high-mass planet on an eccentric orbit. The mass of the planet is not well-constrained by the observable resonant structure, although W02 achieve good correspondence between their model and the observations using a 3 M$_J$ planet.  Another hypothesis for the lobed structure has been proposed by Wyatt (2003). where the resonance arises from outward planetary migration of a less massive (Neptune-mass) planet from 40 to 65 AU. Direct imaging has the potential to detect massive planets around Vega, potentially providing an important constraint to modelling of the resonant structure.  Detection of a companion to Vega has previously been attempted with the Palomar system in H band. Metchev et al. (2003) set a 5 sigma limit of approximately 20 M$_J$ based on their observations at the separation of the expected perturber.  Macintosh et al. (2003) also used Keck at K band to carry out a similar observation. They estimate their 5 sigma limit corresponded to 10 M$_J$ at the expected separation of the planet. More recent deep integrations have been achieved by Marois et al. (2006) reaching a limit of approximately 4 M$_J$ at a separation of 7 $\\arcsec$. ",
        "conclusions": "We present initial M band observations of Vega with a newly commissioned thermal infrared camera, Clio, on the 6.5 MMT using the integrated deformable secondary.  The observation constrains a planet to be less than 7 Jupiter masses in the orientation expected from the dust structure into approximately 2.5 arcsec (20 AU) from the star.    Extrapolation to longer integration times suggests that observations similar to these will be sensitive to wide companion planets at the top end of the mass distribution seen by radial velocity searches, for nearby stars."
    },
    "0606/astro-ph0606259_arXiv.txt": {
        "abstract": "The recent detection of the anomalous X-ray pulsar (AXP) 4U 0142+61 in the mid-infrared with the {\\it Spitzer} observatory by Z.Wang and coworkers constitutes the first instance of a disk around an AXP. We show, by analyzing earlier optical and near-IR data together with the recent data, that the overall broadband data set  can be reproduced by a single model of an irradiated and viscously heated disk. ",
        "introduction": "The zoo of young neutron stars contains a number of categories recognized by the distinct properties of such stars discovered or identified in the last decade. The existence of these different types of neutron stars strongly suggests that not all neutron stars are born with the same initial conditions, nor do they follow the same evolutionary path as the familiar radio pulsars, of which the typical young example is usually taken to be the Crab pulsar. Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) \\citep[see][for a review of AXPs and SGRs]{WT05} are widely accepted to be magnetars \\citep{DT92}. The \"radio-quiet\" neutron stars (also known as central compact objects [CCOs]; (Pavlov et al. 2004) and dim isolated (thermal) neutron stars \\citep[DTNs or DINs;][]{haberl05} are probably related classes, although the latter are relatively older. Measured rotation periods of AXPs, SGRs and DINs cluster in the narrow range of $3-12$ s. \\citet{Alpar01} proposed that the presence or absence, and properties, of bound matter with angular momentum that is, a fallback disk,  may be among the initial parameters of newborn neutron stars, in addition to magnetic dipole moment and initial rotation rate. According to this scenario, the differences between isolated pulsars, AXPs, SGRs, and DINs, as well as the CCOs in certain supernova remnants, are due to different initial conditions, including the presence or absence and properties of fallback disks.  The suggestion that the X-ray luminosity of AXPs is due to mass accretion from a fallback disk (Chatterjee et al. 2000; Alpar 2001) has triggered searches for evidence of such disks. Searches for fallback disks around AXPs have been conducted primarily in the optical and near IR bands. A realistic model for a putative fallback disk should take into account the effects of irradiation by the neutron star's X-ray luminosity. Irradiation is the dominant source of disk luminosity for the outer disk at IR and longer wavelengths. The radial temperature profile of a fallback disk has been computed by Perna et al. (2000)  and Hulleman et al. (2000) using the prescription given by \\citet{Vrt90}. To estimate the temperatures, both Hulleman et al. and Perna et al. assumed the same particular irradiation efficiency and found that their estimated optical flux values lie well above the values indicated by observations of the AXPs 4U 0142+61 and 1E 2259+586. While Perna et al. suggested that the difference might be due to a probable advection-dominated flow at the inner disk, Hulleman et al. drew the conclusion that this model can only fit the data for an unrealistically small outer disk radius and that fallback-disk models were ruled out (see \\S\\ 4 for further discussion). Recent studies show that irradiated-disk model fits to existing detections and upper limits from AXPs allow the presence of accreting fallback disks with reasonable irradiation parameters for all AXPs \\citep{ErCa06} when the fits are not restricted to a particular irradiation efficiency.  Recent results by Wang et al. (2006, hereafter WCK06) from {\\it Spitzer} data on the AXP 4U 0142+61 complement earlier data in the near-infrared and optical (Hulleman et al. 2000, 2004, hereafter HvKK00 and HvKK04) and the detection of pulsed emission in the optical \\citep{KM02}, thus providing the first instance of a multiband data set that can be tested against disk models. WCK06 interpret their data with a dust disk radiating in the infrared, while the optical data are ascribed a magnetospheric origin. They take the origin of the X-ray luminosity to be dissipation of magnetic energy in the neutron star's crust, in accordance with the standard magnetar model.  Here we show that the entire unpulsed optical, near-IR, and far-IR spectrum can be explained as due to a gaseous disk whose luminosity is provided by viscous energy dissipation due to mass transfer through the disk, as well as irradiation by the X-ray luminosity $L_{\\mathrm{X}}$ from the neutron star. Taking $L_{\\mathrm{X}}$ to be an accretion luminosity, as we do, links the irradiation contribution to the disk luminosity with the mass accretion rate $\\dot{M}_{\\mathrm{acc}}$ and, thereby, the mass inflow rate $\\dot{M}$ through the disk and the viscous dissipation contribution of viscous dissipation to the disk luminosity. Luminosity in the optical band comes from the inner regions of the disk and is supplied, to a significant extent, although not fully, by the local viscous energy dissipation in the inner disk. The innermost disk, which is dominated by dissipation, radiates mostly in the UV band but also contributes to the total optical emission. Extension of the same irradiation model to outer regions of the disk explains the IR radiation. The disk luminosity is dominated by the contribution of irradiation at outer radii, where the thermal radiation peaks in the longer wavelength bands (see Table 1).  Pulsed emission in the optical \\citep{KM02} is due to magnetospheric emission powered by a disk-shell dynamo in our model. It has been shown by \\citet{CR91} that magnetospheric pulsar emission in optical and higher energy bands can be sustained by an outer gap in the presence of a disk protruding into the magnetosphere. This disk model is different from the first disk-magnetosphere models \\citep{michel81}. More recently, \\citet{EC04} showed that the pulsed optical emission from 4U 0142+61 can be explained in terms of a magnetospheric model with a disk and a neutron star surface dipole magnetic field of $B_{\\mathrm{dipole}}\\sim 10^{12}$-$10^{13}$ G. The magnetospheric emission is nearly 100\\% pulsed in this magnetosphere-disk model. The discovery of strongly pulsed optical emission \\citep{KM02} has generally been perceived as excluding disk models, on the basis of the notion that magnetospheric emission cannot survive with a disk protruding into the light cylinder. The disk-magnetosphere model of  \\citet{CR91} and the application by \\citet{EC04} are important for fallback-disk models, since they demonstrate that a disk can actually be part of a magnetospheric mechanism that generates strongly pulsed emission.  The clustering of AXP and SGR rotational periods between 6 and 12 s was investigated by \\citet{PM02}, who found that statistical clustering with a spin-down law of braking index $n\\sim$ 2-4 implies that the final period of these systems must be about 12 s. This final period does not correspond to reaching the ``death line'' in the $P$-$\\dot{P}$ plane which reflects the cut-off voltage for magnetospheric emission. Within isolated dipole-magnetar models, the alternative possibility for reaching a final clustering period, asymptotically, is by magnetic field decay on timescales commensurate with the AXP and SGR ages \\citep{PM02}. An investigation of various magnetar field decay models (Colpi et al. 2000) has shown that having the X-ray luminosity data be associated with magnetic field decay is consistent with only one model, involving crustal Hall cascade, and for decay times less than about $10^{4}$ yr. The presence of a fallback disk provides a natural explanation for period clustering of the AXPs \\citep{chat}, as well as dim isolated (thermal) neutron stars (see, e.g., Haberl 2005) and SGRs \\citep{Alpar01, EkA03}. Fallback-disk models do not address the origin of bursts, while magnetar models explain the bursts by employing the energy stored in very strong magnetic fields. The basic workings of the magnetar models require only local magnetic stresses and energies in the neutron star crust, and they can thus be sustained with higher multipole components of the field, while the period clustering and explanation of torques with a fallback disk require the dipole component of the magnetic field to be on the order of $10^{12}$-$10^{13}$ G. Hybrid models \\citep{EkA03,ErA03} with $B_{\\mathrm{multipole}}\\sim10^{14}$-$10^{15}$ G and $B_{\\mathrm{dipole}}\\sim 10^{12}$-$10^{13}$ G are a possibility that incorporates a fallback disk.  In this paper, we evaluate the recent observations by WCK06 in conjunction with older results (\\S\\ 2). We show that the data agree quite well with a gaseous disk. In \\S\\ 3, we estimate the characteristic lifetime of such a disk that has evolved through a propeller phase and find that the disk's estimated age, depending on the torque model, is comparable to the $P_{\\ast }/2\\dot{P}_{\\ast }$ characteristic age estimate taken as if dipole spin-down were effective. The discussion and conclusions are given in \\S\\ 4. ",
        "conclusions": "By combining earlier optical data (HvKK00, HvKK04; Morii et al. 2005) with  recent {\\it Spitzer} infrared data (WCK06), we have shown that the full data set is consistent with the luminosity and spectrum expected from a gas disk. The optical radiation is powered to a significant extent by viscous dissipation in the inner regions of the disk while the IR radiation comes from outer regions where irradiation by the central source is the primary source of the disk's local luminosity.  What is the maximum dipole magnetic moment $\\mu _{\\ast }$ that is compatible with these disk model fits to the data? From equations (1) and (2), we obtain \\begin{equation} \\mu _{\\ast }=f_{1}^{-7/4}f_{2}^{-1/2}r_{\\mathrm{in}}^{7/4}\\dot{M}_{\\mathrm{acc% }}^{1/2}(GM_{\\ast })^{1/4}. \\end{equation} Using $f_{1}=0.5$ and $f_{2}=1.0$ and employing the maximum $r_{\\mathrm{in}}\\cong 8 \\times 10^{9}$ cm that is compatible with our fits, for a minimum $A_{V}=2.6$ we find the maximum possible $\\mu _{\\ast }$ to be $5.8\\times 10^{31}$ G cm$^3$ corresponding to a maximum surface dipole field of $1.2 \\times 10^{14}$ G on the poles of a neutron star of mass $1.4\\,M_{\\odot }$ and radius $R_{\\ast}=10^{6}$ cm. A further increase in the upper limit for the magnetic field can be obtained with smaller values of $f_{2}$. The minimum $f_{2}$ that allows a fit within 60\\%  V-band deviation is around 0.3,  together with $r_{\\mathrm{in}}\\simeq 1.0 \\times 10^{10}$ cm. For these extremes of parameter values, we obtain $B_{\\ast}\\simeq 3 \\times 10^{14}$ G on the pole of the neutron star. These estimates should be noted together with the uncertainties due to disk-magnetosphere interaction, possible shielding effects, and time-dependent variations as discussed in \\S\\ 2.  We note that if the disk inner radius is large, $r_{\\mathrm{in}} \\gg r_{\\mathrm{co}}$, the neutron star is a fast rotator, far from rotational equilibrium with the disk, and likely to be in the strong propeller regime. If the simple description of propeller spin-down given by equation (9) is roughly correct, as is likely the case, the large moment arm provided by a large $r_{\\mathrm{in}}$, together with mass outflows at least comparable to (and probably larger than) the mass accretion rate onto the star will yield spindown rates much larger than that observed: \\begin{equation} |\\dot{\\Omega}_{\\ast}| \\sim \\dot{M}_{\\mathrm{acc}}\\Omega_{\\ast}r_{\\mathrm{in}}^{2}. \\end{equation} A disk with $r_{\\mathrm{in}}\\simeq 1.0 \\times 10^{10}$ cm, which allows for a magnetic dipole as large as $B_{\\ast}\\simeq$ 3 $\\times 10^{14}$ G, requires a spindown rate of $|\\dot{\\Omega}_{\\ast}| \\sim $ 10$^{-10}$ rad s$^{-2}$. This is much larger than the observed $|\\dot{\\Omega}_{\\ast}| \\simeq $ 1.7 $\\times$ 10$^{-13}$ rad s$^{-2}$. Here we have a strong indication against values of $r_{\\mathrm{in}}$ large enough accommodate dipole magnetic fields in the magnetar range.  As our disk models depend only on the mass inflow rate $\\dot{M}$ inside the disk and the irradiation parameter $C$, it could be argued that the entire mass inflow actually turns into an outflow and none of the mass flowing through the disk is accreted onto the neutron star, the X-ray luminosity being due to some energy release mechanism other than accretion. The observation that the inner disk radius in our best fit is so close to the co-rotation radius argues against this. Indeed our fits yield the inner radius required by the optical observations. The actual inner radius may be even smaller and closer to $r_{\\mathrm{co}}$, but we do not have the possibility of observing such a highly absorbed source in the UV band to check. The rather small value of the spin-down rate compared with what one would expect with a substantial mass outflow from a larger inner radius also indicates that this is a weak propeller, close to equilibrium. It is therefore likely that there is mass accretion.  Could the disk of our model fits be passive, with no viscous dissipation, no mass inflow, and no accretion at all onto the neutron star? Under this hypothesis, the disk luminosity would be generated entirely by irradiation by the neutron star's X-ray luminosity, which in turn is not due to accretion. It is unlikely that there is no accretion at all, if there is a gas disk with the properties indicated by our model fits, which are sensitive to the inner disk radius and require temperatures above $ 10^{5}$ K in the inner disk. The disk cannot be neutral and passive at these temperatures; it has to be ionized, whether it is pure hydrogen or contains heavy elements. An ionized disk would have viscosity generated by the magnetorotational instability \\citep{BH91}, which will result in mass inflow through the disk and interaction between the disk and the neutron star magnetosphere, as well as some mass accretion onto the neutron star, depending on the closeness to equilibrium. Ionization temperatures for hydrogen disks are of the order of $6000$ K. For fallback disks with heavy metals (Werner et al. 2007), the ionization temperatures could be as low as $\\sim 2500$ K. In hydrogen disk simulations of soft X-ray transients and dwarf novae, two different viscosity parameters are employed above and below the ionization temperatures to account for the observations of these systems. In these models, the viscosity parameter decreases by a factor of 0.01-0.5 as the temperatures decrease to below the ionization temperature, $T_{\\m{i}}$.  This indicates that even the disk regions with $T < T_{\\m{i}}$ remain active and that the critical temperature below which the entire disk enters a passive state is likely to be much lower than the ionization temperatures. Early studies by Gammie (1996) on protoplanetary disks suggested the possibility that a weakly ionized disk might become passive in its dense regions in the mid-plane while accretion proceeds over its surface. Recent work by Inutsuka $\\&$ Sano (2005), however, indicates the existence of various physical mechanisms that can provide, sustain, and homogenize a sufficient degree of ionization for the magnetorotational instability (MRI) to work through most regions of a protoplanetary disk even at temperatures as low as 300 K. According to Inutsuka $\\&$ Sano, once MHD turbulence starts to prevail in a disk, it seems quite difficult to switch it off. The state of turbulent viscosity keeps the temperature and other parameters such that the ionization level required for MRI is self-sustained.  The presence of the disk imposes constraints only on the dipole component of the magnetic field. The very strong magnetic fields that are employed to explain SGR and AXP bursts in the magnetar model may well reside in the higher multipole structure of the neutron star's surface magnetic field. In this hybrid situation, dissipation of the magnetar fields' energy in the neutron star crust may  contribute to the surface luminosity along with accretion. In this case,  $\\dot{M}_{\\mathrm{acc}}$ will have a value less than that inferred by taking the X-ray luminosity to be fully due to accretion, as in equation (2).  The object 4U 0142+61 is the anomalous X-ray pulsar with the most detailed optical and IR observations. It would  be of great interest to check and extend the current database with further observations, in particular with {\\it Spitzer}. It is important to realize that magnetar fields in higher multipoles and a dipole field with  $B_* \\la 10^{13}$ G, to accomodate a disk with the equilibrium period  range of observed period clustering, can coexist in the neutron star. The search for disks in other AXPs has also provided some data in the optical and near-infrared that are consistent with gas disk models, as discussed in \\citet{ErCa06}. Time-dependent variations in the X-ray luminosity could be smeared out during the X-ray reprocessing at the outer disk. On the other hand, variations in the mass inflow rate modify the disk emission in the optical bands before the changes have been observed in X-rays, because of varying accretion rates. Any observed signatures in the optical bands followed by similar variations in X-rays with viscous time delays (minutes to days) would be a clear indication of ongoing disk accretion onto the neutron star.  At present, 4U 0142+61 seems to be the best source with which to test this idea.  Finally, we comment on the passive dust-disk model proposed by Wang et al. (2006). Their detection of 4U 0142+61 in the mid-infrared has firmly indicated, for the first time, the presence of a disk around an isolated neutron star. They fitted the data with a two-component model, using a power law to fit the optical and near IR data, which they ascribe to magnetospheric emission. This leaves only their recent discovery, the {\\it Spitzer} detections in the mid-IR band, to be explained by a disk. Hence, while the evidence in that band clearly points to a disk, the association of only this narrow-band emission with the disk constrains it to be a dust disk beyond the magnetosphere. Questions arising from this model concern why this disk did not generate some viscosity to extend into an active gaseous disk, how it cooled and remained cold and confined, and whether it is stable against the radiation pressure of the pulsar. With an inner disk radius $r_{\\m{in}}\\sim 2 r_{\\m{lc}}$, where $r_{\\m{lc}}$ is the light cylinder's radius, such a passive dust disk could be stable against the radiation pressure of the magnetospheric pulsar emission (Ek\\c{s}i \\& Alpar 2005). This stability is valid for $r_{\\m{lc}} ~\\la ~r_{\\m{in}}~\\la ~2 r_{\\m{lc}}$ for orthogonal rotators and for even wider range for non-orthogonal rotators according to the investigation of Ek\\c{s}i \\& Alpar (2005), the stability issue should be kept in mind for dust disks beyond the light cylinder."
    },
    "0606/astro-ph0606545_arXiv.txt": {
        "abstract": "We present a new determination of the cluster mass function and velocity dispersion function in a volume $\\sim$10$^7 h^3$Mpc$^{-3}$ using data from the Fourth Data Release of the Sloan Digital Sky Survey (SDSS) to determine virial masses.  We use the caustic technique to remove foreground and background galaxies.  The cluster virial mass function agrees well with recent estimates from both X-ray observations and cluster richnesses.  Our determination of the mass function lies between those predicted by the First-Year and Three-Year WMAP data.  We constrain the cosmological parameters $\\Omega_m$ and $\\sigma_8$ and find good agreement with WMAP and constraints from other techniques.  With the CIRS mass function alone, we estimate $\\Omega_m=0.24^{+0.14}_{-0.09}$ and $\\sigma_8=0.92^{+0.24}_{-0.19}$, or $\\sigma_8=0.84\\pm$0.03 when holding $\\Omega_m=0.3$ fixed.  We also use the WMAP parameters as priors and constrain velocity segregation in clusters.  Using the First and Third-Year results, we infer velocity segregation of $\\sigma_{gxy}/\\sigma_{DM}\\approx$0.94$\\pm$0.05 or 1.28$\\pm$0.06 respectively.  The good agreement of various estimates of the cluster mass function shows that it is a useful independent constraint on estimates of cosmological parameters.  We compare the velocity dispersion function of clusters to that of early-type galaxies and conclude that clusters comprise the high-velocity end of the velocity dispersion function of dark matter haloes.  Future studies of galaxy groups are needed to study the transition between dark matter haloes containing individiual galaxies and those containing systems of galaxies.  The evolution of cluster abundances provides constraints on dark energy models; the mass function presented here offers an important low redshift calibration benchmark. ",
        "introduction": "Clusters of galaxies are the most massive gravitationally relaxed systems in the universe, so the observed cluster mass function is a sensitive probe of cosmological parameters.  Galaxy cluster abundances are most sensitive to the matter density of the universe $\\Omega_m$ and $\\sigma_8$, the rms fluctuations in spheres of radius 8$\\Mpc$ and the usual normalization of the linear power spectrum \\citep{henry91,bahcall93}.  The evolution of the mass function is an excellent probe of the growth of structure, and thus can constrain the properties of dark energy \\citep[e.g.,][]{haiman01,hu03b,vikhlinin03,henry04,majumdar04,vv04}.  Robust constraints from cluster evolution are only possible with an accurate determination of the cluster mass function in the nearby universe.  Most recent estimates of $\\sigma_8$ from the local cluster mass function find values of $\\sigma_8$$\\approx$0.6-0.8 for $\\Omega_m$=0.3 \\citep{pierpaoli01,hiflugcs,seljak02,viana02,pierpaoli03,schuecker03,viana03,bahcall03a,vv04,eke05,dahle06} (there is a strong degeneracy between $\\sigma_8$ and $\\Omega_m$).  Estimates based on cosmic shear usually find larger values of $\\sigma_8$$\\approx$0.8-1.0 \\citep{massey05,hoekstra06,semboloni06}, although some investigators find smaller values \\citep{hamana03,heymans05}.  Recent attempts to model the relation between dark matter haloes and galaxies \\citep{2003MNRAS.345..923V,tinker05} suggest that smaller values of $\\sigma_8$$\\approx$0.6-0.8 are required to match cluster mass-to-light ratios \\citep{cnoc96,bahcall2000,cairnsii}.  Various estimates based on bulk flow models seem to be converging on intermediate values \\citep[see Table 3 of][]{pike05}.  The WMAP satellite has produced tight constraints on many cosmological parameters \\citep{spergel03,spergel06}.  Among the largest changes between the First-Year and Three-Year WMAP results (hereafter WMAP1 and WMAP3) was a decrease in both $\\sigma_8$ and $\\Omega_m$ \\citep{spergel06}.  The WMAP3 results provide a much better fit to the mass function of X-ray clusters from the ROSAT HIghest X-ray FLUx Galaxy Cluster Survey \\citep[HIFLUGCS,][]{hiflugcs} than the WMAP1 results \\citep{reiprich06}.  The greatest difficulty in measuring the cluster mass function is obtaining sufficiently accurate mass estimates.  Three well-established methods exist for measuring cluster masses.  The dynamics of cluster galaxies provided the first evidence for dark matter when \\citet{zwicky1933,zwicky1937} applied the virial theorem to the Coma cluster.  The second method is to use the properties of the hot intracluster medium (ICM), whose distribution and temperature probe the gravitational potential of the cluster.  Finally, gravitational lensing can provide very accurate mass estimates in the centers of clusters and somewhat less accurate estimates at large radii.  All three techniques have been used extensively, but they are each subject to potentially large systematic uncertainties.  There may be velocity bias between cluster galaxies and dark matter particles, although the magnitude and direction of this bias is under debate \\citep[e.g.,][]{kauffmann1999b,diemand04,faltenbacher05,benatov06}.  {\\em Chandra} and {\\em XMM-Newton} observations show that the physical properties of the ICM are complex and involve significant interaction with AGN \\citep[e.g.,][]{2000ApJ...534L.135M,2000ApJ...541..542M,fabian00,2001ApJ...551..160V,fabian06}. Finally, mass estimates from gravitational lensing (especially at large radii) are subject to significant uncertainties due to lensing by other structures along the line-of-sight \\citep[e.g.,][]{2001ApJ...547..560M}, although these might be overcome by combining information from strong and weak lensing \\citep{natarajan02,bradac04}.  These potential systematics have led some investigators to invert the problem and determine cluster properties by matching the mass function predicted by a cosmological model to the observed richness function or luminosity function \\citep{berlind06,stanek06}.  Despite these challenges, the cluster mass function is the subject of much study.  Most recent measurements of the cluster mass function have been made with either X-ray data \\citep{hiflugcs,allen03,vv04} or optical richness data \\citep{bahcall03a}.  The most significant uncertainty in determining $\\sigma_8$ from X-ray observations is the normalization of scaling relations, either the $M-T_X$ or the $M-L_X$ relation.  The normalizations are often determined from hydrodynamical simulations \\citep[e.g.,][]{borgani01}, but different simulations produce a range of normalizations \\citep{pierpaoli03,viana03}. Another potential concern in determining scaling relations is Malmquist bias \\citep{stanek06}.  \\citet{vv04} avoid these scaling relations by estimating the mass of the intracluster medium (ICM) and measuring the baryonic mass function \\citep[a similar analysis was presented by][]{shimasaku97}.  This method requires assumptions about: the relative contribution of stars and gas to the total baryon mass, the ratio of the baryon fraction in clusters to the global value, and the mass dependence of this ratio.  A potentially large systematic uncertainty in all of these X-ray studies is a possible offset between the temperature $T_{spec}$ measured by X-ray satellites and the emission-weighted temperature $T_{ew}$ calculated in hydrodynamical simulations \\citep{rasia05}.  This effect can be understood as the excess contribution of line emission from cooler gas in an ICM with a variety of temperatures, thus causing $T_{spec}$ to be an underestimate of $T_{ew}$.  Correcting for this effect would increase the estimated cluster masses and thus increase the best-fit values of $\\sigma_8$ \\citep{rasia05,shimizu06}.  A similar systematic effect is suggested by the higher normalization of the mass-temperature relation found by \\citet{vikhlinin06} from a study of relaxed clusters with temperature profiles from {\\em Chandra} data.  An increase in this normalization would increase the inferred cluster masses and hence the overall mass function, shifting the constraints towards higher values of $\\Omega_m$ and $\\sigma_8$.  \\citet{viana02} use a hybrid approach of combining the X-ray luminosity function of clusters with a mass-luminosity relation determined from weak lensing \\citep{sheldon01}.  They find results consistent with other X-ray studies, but note that large systematic uncertainties are possible.  Recently, \\citet{dahle06} used weak lensing mass estimates of an X-ray selected cluster sample to measure the mass function.  He found good agreement with many of the constraints from X-ray estimates but lower values of $\\Omega_m$ and $\\sigma_8$ than recent cosmic shear studies.  To our knowledge, the most recent cluster virial mass function is that of \\citet{girardi98b}.  The advent of large-scale redshift surveys, especially the Sloan Digital Sky Survey \\citep[SDSS,][]{sdss}, enables more accurate virial mass estimates due to their size and uniformity. Compared to X-ray studies, virial mass estimates are sensitive to larger scales ($r_{200}$ rather than $r_{500}$, where $r_{\\Delta}$ is the radius within which the enclosed density is $\\Delta$ times the critical density), which allow for comparison with theoretical mass functions with significantly less extrapolation \\citep{white02}. Also, virial masses can be estimated for poor clusters and rich groups, whereas X-ray mass estimates of these systems are complicated by possible energy input from supernovae and AGN \\citep[e.g.,][]{loewenstein00}.  Probing these smaller masses enables a direct constraint on fluctuations on the scale 8$\\Mpc$, rather than the $\\sim$14$\\Mpc$ scales probed by $\\sim$10$^{15}\\msun$ clusters \\citep{pierpaoli01}. \\citet{eke05} used a simplified version of the virial theorem to estimate the group mass function from an optically-selected group catalog in the Two Degree Field Galaxy Redshift Survey \\citep{2df}.  Differences in our analysis include: X-ray selection rather than optical selection of clusters, use of the full virial theorem with corrections for the surface pressure term, much better sampling of individual systems (the average number of cluster members is 50), and more conservative rejection of interlopers (thus significantly reducing scatter in the mass estimates). Estimates of the group mass function using virial masses are usually limited by systematics in group selection, mass estimation, and cosmic variance \\citep{girardi00b,martinez02b,heinamaki03}.    Here we use the virial theorem to estimate cluster masses of an X-ray selected sample of clusters with data from the Sloan Digital Sky Survey.  The use of X-ray selection reduces the impact of projection effects on cluster selection.  Perhaps the greatest advantage of X-ray selection of our sample is that the selection function is well understood and can be computed directly.  The cluster sample we use is the Cluster Infall Regions in SDSS (CIRS) sample of \\citet{cirsi} with some minor modifications described below.    We describe the data and the cluster sample in $\\S$ 2.  In $\\S$ 3, we estimate the mass function using both virial masses and caustic masses and discuss potential systematic effects.  We compute the velocity dispersion function of clusters in $\\S$4 and compare it to the velocity dispersion function of early-type galaxies in SDSS.  We discuss our results and conclude in $\\S 5$.  We assume $H_0 = 100 h~\\kms$, and a flat $\\Lambda$CDM cosmology ($\\Omega _\\Lambda = 1- \\Omega _m$) throughout.  Where not stated explicitly, we assume $\\Omega_m=0.3$ for initial calculations. ",
        "conclusions": "We use the Fourth Data Release of the Sloan Digital Sky Survey to measure the mass function of galaxy clusters in the local universe. We select clusters based on their X-ray flux and a redshift limit ($z<0.1$).  X-ray selection has the crucial advantage over optical selection of yielding a well-defined selection function.  We measure cluster masses from the dynamics of their member galaxies using both the virial theorem and the caustic technique.  The resulting mass functions agree well with previous determinations using X-ray properties and the mass-richness relation.  We conclude that the cluster mass function is measured robustly in the local universe.  We compare this mass function to predictions of various cosmological models to obtain constraints in the ($\\Omega_m$,$\\sigma_8$) plane. With the CIRS virial mass function (and assuming $h$=0.7), we estimate $\\Omega_m$=$0.24^{+0.14}_{-0.09}$ and $\\sigma_8$=$0.92^{+0.24}_{-0.19}$, or $\\sigma_8$=$0.84\\pm$0.03 when holding $\\Omega_m$=0.3 fixed.  The CIRS virial mass function agrees well with the WMAP1 and lies slightly above the WMAP3 results, although still within the 95\\% confidence levels.  This result contrasts slightly with \\citet{reiprich06}, who finds that the HIFLUGCS mass function lies much closer to WMAP3 than WMAP1.  However, other investigators have found mass functions using X-ray data that lie closer to the CIRS mass function, and there remain open questions about the systematic uncertainties in the X-ray mass estimates \\citep{rasia05,stanek06}. Virial masses have the following advantages: they measure the masses at larger radii, where less extrapolation is required to compare with theoretical results, they are less sensitive to the complex physics in the centers of clusters, and they can be applied to lower mass systems for which nongravitational physics greatly complicates X-ray mass estimates, thus eliminating any uncertainty due to possible scale dependence of the estimate of $\\sigma_8$ \\citep[the clusters of, e.g.,][sample fluctuations on larger scales, $\\sigma_{14}$, which must then be converted to $\\sigma_8$]{pierpaoli01}.  We investigate several potential sources of systematic error and conclude that the most significant effect would be a large underestimate of the uncertainties in the virial mass estimates.  Recent investigations with numerical simulations \\citet{biviano06} suggest that the scatter in virial mass estimates of well-sampled clusters is smaller than the amount required to bring our results into agreement with the WMAP3 parameters, although further confirmation of this result would be useful.  Although obtaining virial mass estimates from galaxy dynamics for large numbers of clusters requires much observational effort, the reward is an independent measurement of the mass function and a powerful confirmation of the current cosmological model. This work further demonstrates the power of the mass function as a fundamental constraint on cosmological parameters.  The SDSS-II Legacy Survey should substantially enlarge the sample of local clusters and reduce both statistical and systematic uncertainties in the mass function.  The CIRS mass function provides a critical normalization of the mass function in the nearby universe.  An accurate normalization of the local mass function is a prerequisite for future studies of the evolution of the mass function.  Future determinations of the mass function at higher redshifts will directly probe the growth rate of structure and provide further constraints on cosmological parameters \\citep[e.g.,][]{haiman01,hu03b,vikhlinin03,henry04,majumdar04,vv04}. We conclude that a large spectroscopic program to measure virial masses of an X-ray selected sample at moderate redshift would provide an independent measurement of the evolution of the mass function."
    },
    "0606/astro-ph0606688.txt": {
        "abstract": "We present a satellite mission concept to measure the dark energy equation of state parameter $w$ with percent-level precision. The Very Ambitious Dark Energy Research satellite (VADER) is a multi-wavelength survey mission joining X-ray, optical, and IR instruments for a simultaneous spectral coverage from 4\\,$\\mu$m (0.3\\,eV) to 10\\,keV over a field of view (FoV) of 1 square degree. VADER combines several clean methods for dark energy studies, the baryonic acoustic oscillations in the galaxy and galaxy cluster power spectrum and weak lensing, for a joint analysis over an unrivalled survey volume.  The payload consists of two XMM-like X-ray telescopes with an effective area of 2,800\\,cm$^2$ at 1.5\\,keV and  state-of-the-art wide field DEPFET pixel detectors (0.1-10\\,keV) in a curved focal plane configuration to extend the FoV. The X-ray telescopes are complemented by a 1.5\\,m optical/IR telescope with 8 instruments for simultaneous coverage of the same FoV from 0.3\\,$\\mu$m to 4\\,$\\mu$m. The 8 dichroic-separated  bands (u,g,r,z,J,H,K,L) provide accurate photometric galaxy redshifts, whereas the diffraction-limited resolution of the central z-band allows precise shape measurements for cosmic shear analysis.  The 5 year VADER survey will cover a contiguous sky area of 3,500 square degrees to a depth of $z$$\\sim$2 and will yield accurate photometric redshifts and multi-wavelength object parameters for about 175,000 galaxy clusters, one billion galaxies, and 5 million AGN. VADER will not only provide unprecedented constraints on the nature of dark energy, but will additionally extend and trigger a multitude of cosmic evolution studies to very large ($>$10 Gyrs) look-back times. ",
        "introduction": "\\label{sect:intro}  % \\label{} allows reference to this section  %0.5 pages: %Rene, Jutta %(DE, other planned surveys \\& missions)  %The DE riddle The nature of dark energy (DE), as the driving force of the observed acceleration of the Universe, is currently one of the deepest mysteries in astrophysics. Whereas a physically motivated theory of DE is still out of reach, upcoming observational data sets will be able to put improved constraints on the DE equation of state parameter $w$, which relates the fundamental properties pressure and energy density of DE. A precise knowledge of $w$ and in particular its redshift dependance, usually parameterized as $w(z)=w_0+w_1 \\cdot z$, is required to allow at least a phenomenological understanding of the dominant energy contribution in the Universe. The answer to the fundamental question whether DE is equivalent to Einstein's cosmological constant ($w$=-1), phantom energy\\cite{Caldwell2003a} ($w$$<$-1), or time varying as in quintessence models\\cite{Wetterich1988a} (-1$<$$w$$<$-1/3) will not only shed light on the fate of the Universe but will also lead the way to a unified physical theory.  %1st \\& 2nd generation experiments Current observational data can constrain $w_0$ to about 5-10\\%\\cite{Schuecker2005a}, and are in agreement with the value of the cosmological constant. A first generation of designated dark energy surveys (e.g. KIDS, APEX-SZ) is currently upcoming or in the final planning stages. Second generation ground-based experiments (HETDEX, DES, Pan-STARRS, LSST) and planned satellite missions (eROSITA, SNAP, DUNE) are foreseen to yield valuable results after the year 2010. However, constraining $w$ with high precision will be a very challenging observational task. A final distinction between DE models, in particular with $w_0$ being close to -1, is likely to require sub-percent precisions\\cite{Caldwell2005a} that are beyond the capabilities of the currently planned experiments mentioned above. Our proposed VADER survey will bridge this gap. The VADER project originated from a case study of  future DE space missions and is mainly targeted at the feasibility of the experiment without being constrained by a specific mission framework.   %why not good enough %Organization The paper is organized as follows. In Sect.\\,2 we describe the scientific goals and the requirements for the mission. The survey design, sensitivities, and expected results are discussed in Sect.\\,3. The scientific payload is presented in Sect.\\,4 and the spacecraft in Sect.\\,5. The conclusions are given in Sect.\\,6. We assume concordance model cosmological parameters with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ and $h=0.7$, unless otherwise stated.  %Outlook: %use concordance model ($O_m=0.3, O_lam=0.7, h=0.7$)     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%0.5 pages: %Rene, Jutta %(comparison to other missions) %============ %Total: %10 pages  %Competitiveness / costs  We have presented  a satellite mission concept with which the dark energy equation of state parameter $w$ and its time variation can be constrained with unprecedented precision. VADER will allow detailed studies using different cosmological probes and independent DE tests with low systematics out to redshifts of 2 and beyond.  The special mission strengths are (i) the cross-check capabilities of several independent DE methods, (ii) the high redshift leverage, (iii) the high sensitivity multi-wavelength coverage, and (iv) the versatility of the survey data set for numerous additional astrophysical studies. Whereas the X-ray part of the mission bears only low technological risk, the innovative multiplexing capability of the wide-field optical/IR telescope system will require a challenging optical design. The main R\\&D activities for VADER will thus have to focus on the different optical components, the wide-field detector technology, and an optimized thermal structure. However, compared to other planned missions (e.g. JWST, LISA) the overall technological risks of VADER are quite moderate.  Unveiling the very nature of the dominant dark energy component in the Universe will be VADER's main challenging task. Considering the fundamental importance and the observational difficulty, the scope of the proposed mission concept poses an adequate answer to this pressing question. Even though the notion of dark energy might be replaced in the theories at some point, VADER's contribution to a profound understanding of the Universe will remain and could guide the way to new insights in fundamental physics.   %Strength \\& Risks %Strength: self contained data (no follow-up or other missions / telescopes necessary) %Risk: optical system (so far simultaneous in small, single cameras in big), thermal control of different components %R \\& D necessary %VADER --> most precise and cross checks  %DES Whitepaper: the most exciting prospect for combining muliple probes is the possiblity to distinguish DE from alternative %explanations of cosmic acceleration, such as modifications of GR or local inhomogeneities on horizon scales that cause apparent %accelerated expansion  %Weinberg 2005: a demonstration that the DE density has changed over time (w ne -1) would be an achievement comparable to the %discovery of DE. A precise measurement of w would rule out many models fo DE and provide a clear target of predictive models that %might emerge from fundamental physics.  a demonstration that the effective value of w has changed with time would be more %remarkable still, eliminating most models and providing a strong clue to the physics of DE. %The level of precision and control of systematics required to detect such small departures from w=-1 can only be achieved from %space. %--> need space to reach high z and control systematics for measurement of time varying w %DE statement at end    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%\\appendix    %>>>> this command starts appendixes %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0606/astro-ph0606111.txt": {
        "abstract": "We discuss Monte-Carlo techniques for addressing the 3-dimensional time-dependent radiative transfer problem in rapidly expanding supernova atmospheres.  The transfer code \\Code\\ has been developed to calculate the lightcurves, spectra, and polarization of aspherical supernova models.  From the onset of free-expansion in the supernova ejecta, \\Code\\ solves the radiative transfer problem self-consistently, including a detailed treatment of gamma-ray transfer from radioactive decay and with a radiative equilibrium solution of the temperature structure.  Line fluorescence processes can also be treated directly. No free parameters need be adjusted in the radiative transfer calculation, providing a direct link between multi-dimensional hydrodynamical explosion models and observations. We describe the computational techniques applied in \\Code, and verify the code by comparison to existing calculations.  We find that convergence of the Monte Carlo method is rapid and stable even for complicated multi-dimensional configurations.  We also investigate the accuracy of a few commonly applied approximations in supernova transfer, namely the stationarity approximation and the two-level atom expansion opacity formalism. ",
        "introduction": "\\subsection{Motivations}  Most of what we know about supernovae (SNe) has been learned from observations of the lightcurve, spectrum, and polarization of the supernova light during the months and years following the explosion. Except for most nearby events, the explosion process itself is never directly observed, and the progenitor star system only rarely.  What we do see is emission from the hot, radioactive material ejected in the explosion.  Theoretical radiative transfer modeling of the emission is needed to discern the physical conditions in the SN ejecta, offering insight into the physics of the explosion itself and the progenitor star system which gave rise to it.  One dimensional (1D) explosion models of SNe have been used to synthesize emergent spectra and light curves in reasonable agreement with observed ones.  Because nearly all observed SNe are too distant to be resolved in the early phases, deviations from spherical symmetry can not be directly imaged.  Nevertheless, three lines of observational evidence imply that the ejecta possess an interesting multi-dimensional structure: (1) The detection of a non-zero intrinsic polarization in SNe of all types, indicating a preferred direction in the scattering medium \\citep[e.g.,][]{Cropper_87A,Kawabata_02ap,Wang_01el,Leonard_SNIa}; (2) The appearance of unusual flux features in some SNe, naturally explained by a clumpy ejecta structure, e.g., the ``Bochum event'' in SN~1987A, \\citep{Hanuschik_87A, Utrobin_87A}; (3) The complex morphologies of SN remnants, which show clumpy, filamentary, or jet-like structures \\citep{Fesen_1996, Hwang_2000,Decourchelle_2001}.  Most current theoretical SNe explosion scenarios involve multi-dimensional effects in an essential way.  In hydrodynamical explosion models, ejecta asymmetries arise, for example, from random instabilities in the explosion physics, e.g., Rayleigh Taylor instabilities, convective mixing \\citep{Chevalier-Klein, Burrows_1995, Kifonidis_2000, Gamezo_3D, Roepke_fullstar}; from anisotropic energy ejection mechanisms, e.g., jets, off-center ignition \\citep{khokhlov_jet, Macfadyen_1,Plewa_GCD}, or from asphericities in the progenitor star or its surrounding medium, e.g., rapid rotation of the progenitor, the presence of a binary companion star \\citep{Marietta,Yoon_2005}.  The multi-D explosion models make predictions as to the velocity, composition and geometry of the material ejected in the SN explosion. To confront such predictions with observations, the multi-D radiative transfer problem in expanding SN atmospheres must be addressed. Detailed radiative transfer codes synthesize model spectra, light curves and polarization which can be compared directly to observations.  Such transfer codes can also be usefully applied in an empirical ``inverse'' approach, in which hand-tailored, parameterized ejecta configurations are used to extract model-independent information directly from the observations.  Here we describe a Monte Carlo approach to the multi-dimensional time-dependent radiative transfer problem in expanding SN atmospheres, embodied in the transfer code \\Code.  Given an arbitrary 3-dimensional ejecta structure (i.e., the density, composition and velocity structure of freely expanding SN material) \\Code\\ self-consistently calculates the emergent broadband light curves, spectral time-series (in both optical and gamma-rays) and polarization spectra from various viewing angles.  No free parameters need be adjusted in the transfer calculations, providing a direct link between multi-dimensional hydrodynamical explosion models and observations. In this paper, we describe the radiative transfer techniques used in \\Code, and give some examples of its verification and application.  \\subsection{Monte Carlo Radiative Transfer}  In the Monte-Carlo (MC) approach to radiative transfer, packets of radiant energy (``photons'') are emitted from within the SN envelope and tracked through random scatterings and absorptions until they escape the atmosphere.  Each photon packet possesses a specific wavelength and polarization state which are updated at each interaction event.  Tallies of photon packets can be used to construct estimators of the local radiation field properties and the emergent spectrum at infinity.  The calculated quantities possess statistical noise, which is reduced as the number of propagated packets is increased.  The MC approach has several advantages over direct numerical solution of the radiative transfer equation.  MC codes are intuitive, relatively easy to develop, and generally less likely to fall victim to subtle numerical errors \\citep{Auer_MC}.  The method generalizes readily to multi-dimensional time-dependent problems and the inclusion of polarization.  As discussed in detail below (\\S\\ref{rtm:Converge}), convergence of MC calculations is found to be stable and rapid even for complicated configurations.  Finally, although MC techniques can be computational expensive, they parallelize well and can be profitably run on multi-processor supercomputers.  Monte Carlo approaches have been applied to a wide range of astrophysical radiative transfer problems, including multi-dimensional polarization problems \\citep[e.g.,][]{Daniel_MC,Code_blobs,Wood_EScatterI}. The MC code described in \\cite{Hoeflich_93J} has been used to calculate the continuum polarization and polarization spectrum of 2-D SN \\citep{Hoeflich_1991pol,Wang_96X,Howell_99by}.  In addition, the 1-D MC code of \\cite{Mazzali_MC} has been used in numerous studies of SN flux spectra \\citep[e.g.,][]{Mazzali_91T, Mazzali_98bw}.  The papers of Leon Lucy have been particularly important in developing the MC technique for astrophysical applications \\citep{Lucy_Radeq,Lucy_1999,Lucy1,Lucy2,Lucy3,Lucy_Timedep,Lucy_2005}. The new techniques, many of which are applied here, make it feasible for MC codes to match the physical accuracy of formal solutions of the radiative transfer equation. ",
        "conclusions": "In this paper we have described the computational techniques, and demonstrated applications and verifications of the multi-dimensional MC transfer code \\Code.  We have also explored the validity of several common approximations made in SN transfer models.  These findings will be drawn upon in future applications of the code.  Although computational efficiency is often considered a drawback of MC codes, the effective parallelization and favorable convergence properties of the present techniques allow \\Code\\ to immediately address complicated transfer problems using currently available computing resources.  Moreover, as the physics included in the radiative transfer simulation becomes increasingly complicated, including all of multi-dimensionality, NLTE effects, polarization and time-dependence, it is not clear that difference equation techniques will outperform MC approaches in their execution times.  With further work, and with the advance of computing power, some of the presently applied approximations can be relaxed.  Most outstanding is the inclusion of an NLTE treatment of the occupation numbers, including excitation/ionization by non-thermal electrons scattered from radioactive gamma-rays.  Inclusion of nebular physics and cooling by forbidden line transitions would also improve model accuracy at late times. Eventually, coupling of the transfer code to a multi-dimensional hydrodynamical solver would allow the calculations to be generalized to non-homologous flows.  Fortunately, the foreseeable improvements are readily incorporated into the MC framework in a straightforward and intuitive manner."
    },
    "0606/astro-ph0606709_arXiv.txt": {
        "abstract": "We present a fast, accurate, robust and flexible method of accelerating parameter estimation. This algorithm, called Pico, can compute the CMB power spectrum and matter transfer function as well as any computationally expensive likelihoods in a few milliseconds. By removing these bottlenecks from parameter estimation codes, Pico decreases their computational time by $1$ or $2$ orders of magnitude. Pico has several important properties. First, it is extremely fast and accurate over a large volume of parameter space. Furthermore, its accuracy can continue to be improved by using a larger training set. This method is generalizable to an arbitrary number of cosmological parameters and to any range of $\\ell$-values in multipole space. Pico is approximately $3000$ times faster than CAMB for flat models, and approximately $2000$ times faster then the WMAP $3$ year likelihood code. In this paper, we demonstrate that using Pico to compute power spectra and likelihoods produces parameter posteriors that are very similar to those using CAMB and the official WMAP3 code, but in only a fraction of the time. Pico and an interface to CosmoMC are made publicly available at \\verb+www.astro.uiuc.edu/~bwandelt/pico/+. ",
        "introduction": "With WMAP's second data release \\citep{Hinshaw:2006ia,Page:2006hz,Spergel:2006hy}, there is a wealth of new CMB data available to further constrain the cosmological parameters. The major computational burden in parameter estimation remains the calculation of the theoretical power spectrum for a large number of cosmological models as well as the likelihood based on these spectra. Generally the power spectrum is computed with codes such as CMBfast \\citep{Seljak:1996is} or CAMB \\citep{Lewis:1999bs}, which evolve the Boltzmann equation using a line of sight integration approach. While this provides a $1$ or $2$ order of magnitude decrease in the computation time over the full Boltzmann codes, power spectrum calculations remain a bottleneck of parameter estimation. Other software such as CMBwarp \\citep{Jimenez:2004ct} and DASh \\citep{Kaplinghat:2002mh} have found ways to improve the efficiency of power spectrum calculations at the cost of a loss of accuracy against the full Boltzmann codes and/or placing restrictions on the parameters which are available as input. In particular, CMBwarp builds upon the method introduced in \\citep{Kosowsky:2002zt}, where a new set of nearly uncorrelated ``physical'' parameters were defined which have nearly independent effects on the power spectrum. CMBwarp uses a modified polynomial fit whose coefficients are based on a fiducial model. It allows rapid calculation of the temperature (TT), E-mode polarization (EE) and temperature-polarization (TE) cross power spectra. CMBwarp, however, requires the use of specific cosmological parameters and the accuracy of the computed power spectra quickly diminishes as one moves away from the fiducial model in parameter space. Another code, CMBFit \\citep{Sandvik:2003ii}, attempts to avoid the need to compute the power spectrum by fitting the likelihood function. This idea is particularly important for the WMAP $3$ year data \\citep{Hinshaw:2006ia,Page:2006hz,Spergel:2006hy} whose likelihood is time consuming to compute.  In this paper we introduce Pico, a computational technique to accelerate both power spectrum and likelihood computations. This approach removes the $2$ major bottlenecks in parameter estimation. While in a similar spirit as CMBwarp, and providing a similar speedup over CMBfast and CAMB, Pico has several important advantages over CMBwarp and DASh. First, it allows the calculation of power spectra from an arbitrary number of cosmological parameters and in any range of $\\ell$-values in multipole space. Because of this flexibility, it is  easily incorporated into parameter estimation codes. Secondly, Pico allows the simultaneous computation of all scalar, tensor and lensed power spectra as well as the transfer functions. Pico provides more than an order of magnitude increase in accuracy over CMBwarp, and about $2$ orders of magnitude increase in speed over DASh. Lastly, Pico is generic enough to allow the direct fitting of any likelihood functions. Due to the computational expense in computing the likelihood of certain experiments, e.g. WMAP3, any power spectrum acceleration scheme will at most provide a speed up of order $1$ to $10$ in parameter estimation. However, using Pico to also compute the likelihood results in speed ups of order $10$ to $100$. As an additional bonus, using Pico to compute the likelihood directly provides more accurate results then using it to fit the power spectra and computing the likelihood from these approximate spectra. This is important for current and next generation all-sky CMB data. Meanwhile, the power spectra computed by Pico are more than accurate enough for suborbital experiments with smaller sky coverage and coarser $\\ell$-resolution.  This paper is organized as follows. We examine the CPU and memory requirements of Pico in section \\ref{cpu}. Section \\ref{results} presents several tests of the performance of Pico. This includes comparisons of power spectra computed using Pico and CAMB as well as results of parameter estimation runs using Pico to compute the power spectra and the WMAP3 likelihood. In section \\ref{concl} we summarize and discuss the future of Pico. The details of the algorithm used by Pico are presented in the appendix. ",
        "conclusions": "This paper provides a fast, accurate and robust method of calculating CMB power spectra and likelihood functions using local polynomial interpolation. A K-means clustering algorithm is used to partition the cosmological parameter space into local regions. Over each region we approximate the CMB power spectra as a polynomial in the cosmological parameters. This method, which we have named Pico, provides several orders of magnitude increase in speed over CAMB and the WMAP $3$ year likelihood code, while proving accurate enough for the analysis of data from the current and next generation of cosmic microwave background experiments. The flexibility of our algorithm enables it to handle any reasonable number of cosmological parameters. It has been generalized to allow the fast computation of any observables relevant to a particular data set, e.g. the transfer functions and the power spectrum of B-mode polarization anisotropies. Even higher order correlation functions, such as the reduced bispectrum, could be added. Pico is able to compute accurate power spectra over a large volume of parameter space consistent with the WMAP data. Furthermore, Pico's performance will only improve as the volume of space it must fit and the uncertainties in the parameters shrink.  Pico is easily inserted into parameter estimation codes such as CosmoMC. It can be used to compute the power spectra, transfer function and the WMAP3 likelihood, resulting in a significant decrease in computational time. In fact, when Pico is used, CosmoMC spends $80\\%$ of its time on tasks other then computing the power spectra and likelihood. This time is spent generating random numbers, evaluating internal and derived parameters, etc. We envision that CAMB will only be needed for nonstandard cosmological models outside the scopes of our training sets. While it is likely possible to further improve the accuracy of our code by using less generic techniques, we have chosen to keep Pico as generic as possible to allow it to grow and adapt to the parameter estimation tasks of the next generation of experiments.  We have made a Fortran 90 implementation of this algorithm publicly available.\\footnote{ http://www.astro.uiuc.edu/\\~{}bwandelt/pico/ } Here the user will find regression coefficients to use the algorithm for various parameter sets, as well as short and straight forward instructions for incorporating Pico into CosmoMC or using it as a front end for CAMB. The authors also welcome requests for regression coefficients for specific combinations and ranges of parameters. Enabling Pico on a new parameter set simply involves running CAMB to generate a new training set."
    },
    "0606/astro-ph0606223_arXiv.txt": {
        "abstract": "In the model of low-energy quantum gravity by the author, cosmological redshifts are caused by interactions of photons with gravitons. Non-forehead collisions with gravitons will lead to an additional relaxation of any photonic flux. Using only the luminosity distance and a geometrical one as functions of a redshift in this model, theoretical predictions for galaxy number counts are considered here. The Schechter luminosity function with $\\alpha =-2.43$ is used. The considered model provides a good fit to galaxy observations by Yasuda et al. (AJ, 122 (2001) 1104) if the same K-corrections are added. It is shown that observations of $N(z)$ for different magnitudes $m$ are a lot more informative than the ones of $N(m).$ ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606663.txt": {
        "abstract": "We bring you, as usual, the sun and moon and stars, plus some galaxies and a new section on astrobiology. Some highlights are short (the newly identified class of gamma-ray bursts, and the Deep Impact on Comet 9P/Tempel 1), some long (the age of the universe, which will be found to have the Earth at its center), and a few metonymic, for instance the term ``down-sizing'' to describe the evolution of star formation rates with redshift. ",
        "introduction": "The ApXX series celebrates its quincea\\~{n}era by adding a new, astrobiology, section and co-opting an additional author to write it and to cope with many other problems.\\footnote{Astrophysics in 1991 to 2004 appeared in volumes 104--117 of PASP. They are cited here as Ap91, etc.} Used in compiling sections 3--6 and 8--13 were the issues that arrived as paper between 1 October 2004 and 30 September 2005 of \\textsl{Nature}, \\textsl{Physical Review Letters}, the \\textsl{Astrophysical Journal} (plus \\textsl{Letters} and \\textsl{Supplement Series}), \\textsl{Monthly Notices of the Royal Astronomical Society} (including \\textsl{Letters} up to December 2004 only), \\textsl{Astronomy and Astrophysics} (plus \\textsl{Reviews}), \\textsl{Astronomical Journal}, \\textsl{Acta Astronomica}, \\textsl{Revista Mexicana Astronomia y Astrofisica}, \\textsl{Astrophysics and Space Science}, \\textsl{Astronomy Reports}, \\textsl{Astronomy Letters}, \\textsl{Astrofizica}, \\textsl{Astronomische Nachrichten}, \\textsl{Publications of the Astronomical  Society of Japan}, \\textsl{Journal of Astrophysics and Astronomy}, \\textsl{Bulletin of the Astronomical Society of India}, \\textsl{Contributions of the Astronomical Observatory Skalnate Pleso}, \\textsl{New Astronomy} (plus \\textsl{Reviews}), \\textsl{IAU Circulars}, and, of course, \\textsl{Publications of the Astronomical Society of the Pacific}. Journals read less systematically and cited irregularly include \\textsl{Observatory}, \\textsl{Journal of the American Association of Variable Star Observers}, \\textsl{ESO Messenger}, \\textsl{Astronomy and Geophysics}, \\textsl{Mercury}, \\textsl{New Scientist}, \\textsl{Science News}, \\textsl{American Scientist}, \\textsl{Scientometrics}, \\textsl{Sky \\& Telescope}, \\textsl{Monthly Notes of the Astronomical Society of South Africa}, and \\textsl{Journal of the Royal Astronomical Society of Canada}. Additional journals provided material for Sections 2 and 7 and are mentioned there.  A few papers are mentioned as deserving of gold stars, green dots, and other colorful recognition. This is as nice as we get. Among the people who appear in the following pages are Jack Benny, the Keen Amateur Dentist, the Faustian Acquaintance, and the Medical Musician. All are pseudonyms, for Benjamin  Kubelsky, and three colleagues, left as an exercise for readers, not intended to include the also-pseudonymous Mr.~H., who is supposed to be completing a thesis in X-ray astronomy.  \\subsection{Terminations} We record here a number of things that came to an end, surely or probably, in the index year. Beginnings appear in \\S 1.2 and more complicated relationships in 1.3. (1) The last launch of a Skylark sounding rocket happened on 30 April 2005; the first was from Woomera during International Geophysical Year 1957. (2) Both the \\textsl{Letters} section of \\textsl{MNRAS} and the newsletter of the European Space Agency went e-only during the year (January and July 2005 respectively) and so are no longer accessible to the least electronic author.  (3) SLAC, like Brookhaven and Los Alamos before it, was shut down by an accident on 1 October (Science 306, 809). (4) The percentage of women in computer sciences has actually declined over the past 20 years (Science 306, 809), while that in physical sciences and engineering has crept slowly up. (5) NASA's KC--135 jet, the ``vomit comet,'' used to produce brief experiences of weightlessness, free-fell for the last time on 29 October. (6) Kodak produced its last carousel projector in November 2004 but will support existing ones  (light bulbs and things) until 2007. (7) The Yerkes Observatory (where the modern sequence of spectral types was established by Morgan and Keenan) is to be turned into, but turned into what is not entirely clear (Sky \\& Telescope 108, No.~3, p.~19). (8) The Mohorovi\\v{c}i\\'{c} discontinuity was not reached again at 1.4 km where the mid-atlantic ridge was expected to be only 0.7 km thick (Science 307, 1707); of course every child who attempted to dig to China with a tablespoon (from California, or to California from China with chopsticks) can claim a similar result (Scientific American 293, 94). (9) Ball Aerospace provided (at least) its second out-of-focus telescope, this one on \\textsl{Deep Impact} (Nature 434, 685). (10) Another try at \\textsl{Cosmos-1}, a solar sail vehicle, was lost on launch 21 June, probably due to a pump failure. (11) \\textsl{Gravity Probe B} gathered its last data on 1 October, with, up to mid January, no results announced and no triumphant renaming of the craft as, for instance, Lens-Thirring to indicate success.  \\subsection{Inceptions} Most of these beginnings are, we think, good news or at least progress toward good news. (1) \\textsl{Smart-1} on its way to the moon was probably the most fuel-efficient vehicle in history, achieving the equivalent of $2\\times10^6$ km per liter of gasoline (ESA Space Science News No.~7, p.~2); the only competition would seem to be something that simply rides along with tectonic plate spreading. It arrived on 15 November after an October 2003 launch, and was so also only marginally faster than tectonic plates. (2) Some of the antennas for LOFAR are in place (10 near the beginning of the reference year, Sky \\& Telescope December 2005, p.~24). (3) The \\textsl{Spitzer Space Telescope} released its first large package of papers on the cusp of the year (September 2004 issue of Astrophysical Journal supplement Series). (4) \\textsl{SWIFT} caught its first gamma ray burst on 17 December, after a 20 November launch. (5) The first hole for Ice Cube was dug in January 2005, an advanced muon and neutrino detector in Antarctica (and a confusing notebook entry, because we generally use upward pointing arrows for good news items, while the hole probably went down). (6) The Ariane heavy lifter had its first successful launch from French Guiana on 12--13 February. It will be needed for \\textsl{James Webb Space Telescope}, unless very considerable additional descoping occurs. (7) A million dollar Kavli astrophysics prize will be given starting in 2008 (Nature 435, 37) along with ones in neuroscience and nano-technology. (8) The street lighting in the Canary Islands has been considerably dimmed (Nature 435, 41), partly in anticipation of first light at the GTC (out of period). (9) The Sloan Digital Sky Survey is being extended, primarily to examine stars for their own sake and for galactic structure (Nature 436, 316). (10) The \\textsl{Mars Reconnaisance Orbiter} got off the ground in August, and \\textsl{Mars Express} will continue to operate for another year or two (a funding extension), though one spectrometer has gone funny\\footnote{Expert readers may wish to translate this into  some more technical term, but the published description sounded to us as if the spectrometer had gone funny.} (Nature 437, 465). (11) The US is once again planning to develop a major underground laboratory facility for high energy physics and other purposes, perhaps back at Homestake (Ray Davis's old site) or at Henderson, an active molybdenum mine (Science 309, 682). (12) France has removed an assortment of popular topical antibiotics from  over-the-counter sale (Science 309, 872). Not astronomy perhaps, but anything that slows down the development of resistent organisms has to be good for us all! (13) SALT, the South African Large Telescope (near relative of the Hobby-Eberly Telescope) has collected some photons and a dedication (Nature 437, 182). (14) The Pierre Auger facility, which looks for very high energy things, released its first data on 6 July. (15) The Atacama Pathfinder Experiment, APEX, at Chajnator (a 12-meter telescope for millimeter radio astronomy and submillimeter submillimeter astronomy in the in the Atacama desert) carried out its first scientific observations in July 2005. (16) The Las Cumbres Observatory has acquired its first 2-meter telescopes (in Hawaii and Australia) and its first GRB afterglow (in a galaxy long ago and far away, 051111). Six telescopes and observations of many more afterglows and other transient phenomena are planned.  \\subsection{More Complicated Stories} \\textsl{Astro-E}, a Japanese X-ray satellite, was successfully launched on 10 July 2005 and named \\textsl{Suzaku} (red bird of the south). But its coolant was lost some time before 8 august and, although its CCD X-ray detectors will still work, the spectrometer that was one of its main goals will not. This was, sadly, the second try at \\textsl{Astro-E}. The first, also unsuccessful, attempt was roughly contemporaneous with the launches of the \\textsl{Chandra} and \\textsl{XMM} X-ray satellites, and the projects were meant to be complementary.  \\textsl{Muses-C}, which became \\textsl{Hayabusa} after launch on 9 May 2003, was supposed to examine asteroid Itokawa, break off a bit, and bring it back in the summer of 2007. The bad news/good news items have been coming out more or less monthly on beyond the end of the reference year. It got to the asteroid. Two gyros failed. It managed to touch down, or not? It picked up some stuff, or not? And it may or may not be able to turn around and come home.  The American Physical Society, worried about folks saying, ``Gee, you don't look very strong,'' like a Canadian border guard many years ago, attempted to change its name to American Physics Society (Science 309, 378). Loud and long blew the arguments through the summer and fall, both rational and  irrational. But ``they fought the law, and the law won.'' Changing the name would require redoing the incorporation papers, a major hassle and expense. The intention, however, is to use the APS most of the time and the full name only on journals and other items for internal consumption. (Oh dear; the gastric capacity required to internally consume a year of Physical Reviews terrifies.)  All the other items on this list belong more or less to NASA. These consist of existing operational devices that may be turned off early, planned ones that may be greatly delayed or shrunk, items pulled  out of various over-committed\\footnote{In first draft, this item said ``over-committeed,'' which may also be true.} queues and occasional/partial reversals of such decisions, recisions of the reversals, and so forth that have appeared in this context during the year. (1) \\textsl{TRMM} and other sources of climate data (Science 307, 186 \\& 189). (2) \\textsl{JIMO}, \\textsl{Kepler}, \\textsl{SIM}, \\textsl{Beyond Einstein}, and even some of the projects to examine long-term effects of space on people (Science 307, 833). (3) \\textsl{Voyager 1} and \\textsl{2}, \\textsl{Ulysses}, \\textsl{Polar}, \\textsl{Wind}, \\textsl{Geotail}, \\textsl{TRACE}, \\textsl{FAST} (Nature 434, 108). (4) The five-year career grants and archival data program (Science 308, 486). (5) Shuttle flights, the \\textsl{International Space Station}. and \\textsl{Hubble Space Telescope} in multiple stories (Science 309, 540; Nature 436, 603 \\& 163), and, in even more multiple stories, not all on paper, \\textsl{JWST} (Science 309, 1472; Science 308, 935).  %%%%%%%%% Section 2 %%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0606/astro-ph0606015_arXiv.txt": {
        "abstract": "We investigate the evolution of a population of 100 dark matter satellites orbiting in the gravitational potential of a realistic model of M31. We find that after 10 Gyr, seven subhalos are completely disrupted by the tidal field of the host galaxy. The remaining satellites suffer heavy mass loss and overall, 75\\% of the mass initially in the subhalo system is tidally stripped. Not surprisingly, satellites with pericentric radius less than 30 kpc suffer the greatest stripping and leave a complex structure of tails and streams of debris around the host galaxy. Assuming that the most bound particles in each subhalo are kinematic tracers of stars, we find that the halo stellar population resulting from the tidal debris follows an $r^{-3.5}$ density profile at large radii. We construct B-band photometric maps of stars coming from disrupted satellites and find conspicuous features similar both in morphology and brightness to the observed Giant Stream around Andromeda. An assumed star formation efficiency of 5-10\\% in the simulated satellite galaxies results in good agreement with the number of M31 satellites, the V-band surface brightness distribution, and the brightness of the Giant Stream. During the first 5 Gyr, the bombardment of the satellites heats and thickens the disk by a small amount. At about 5 Gyr, satellite interactions induce the formation of a strong bar which, in turn, leads to a significant increase in the velocity dispersion of the disk. ",
        "introduction": "In the current cosmological paradigm, large-scale structure forms via hierarchical clustering wherein small systems composed of dark matter and gas merge to form larger objects. These systems originate from the primordial density fluctuation spectrum of cold dark matter (e.g., \\citealt{dav85}). The hierarchical clustering hypothesis leads to a picture in which subhalos are incorporated into larger systems. A number of processes such as tidal stripping and dynamical friction can lead to the destruction of substructure \\citep{whi78}. However, many subsystems survive and remain in orbit about the parent halo. This phenomenon gives birth to a large population of small satellites orbiting around galaxies and clusters of galaxies \\citep{moo99,kly99,ghi00,die04,gao04,ben05}. Typically this population consists of several hundred subhalos and contributes about 10\\% of the total dark halo mass in a galaxy \\citep{fon01,ghi00}.  From an observational point of view, there is no clear evidence that such a large population of dark matter satellites exists. There are only $\\approx$ 40 observed satellites in the Local Group out of which 13 belong to the Milky Way \\citep{mat98}. This number is an order of magnitude less than predicted from cosmological simulations because the stellar content and hence surface brightness is below current detection limits.  How then is it possible to reconcile the low number of observed satellites with models predictions?  There are several possibilities that we examine below.  The satellites may still be there but difficult to observe. Many of the satellites may be associated with high-velocity clouds \\citep{bli99} or may not be detected because the stellar component is not important enough and the surface brightness of these objects is below current detection limits. There are also theoretical reasons why star formation may be suppressed or inefficient in low mass satellites. Several studies explore  mechanisms that should operate in the early stages of galaxy formation to suppress star formation in low-mass satellites. Gas ejection by supernovae-driven winds from the shallow potential wells of satellites may quench star formation (e.g. \\citealt{dek86}). Also a strong intergalactic ionizing background can prevent the collapse of gas into low-mass systems \\citep{tho96} with circular velocities $v_c \\lesssim 30$ km $s^{-1}$ as revealed in  numerical simulations. Recently, \\citet{min05} claim to have detected an HI source in the Virgo Cluster that could be a dark halo without an accompanying bright stellar galaxy. If the observed system is bound, this implies a dark halo mass up to $10^{11}$ $M_{\\odot}$. In addition, the census of the Milky Way satellites may be incomplete at low latitude due to obscuration. \\citet{wil04} estimated a 33\\% incompleteness in the total number of dwarfs due to obscuration which represents only a small increase of the total number compared to cosmological predictions. All these possible explanations and considerations are plausible solutions to the satellite abundance problem.  A large satellite population may produce a strong dynamical effect and modify the structure and kinematics of the galactic disk. Many studies look at the dynamical response of the host galaxy from an infalling satellite \\citep{tot92,qui93,wal96,hua97,vel99,ben04}. \\citet{mih95} show that the accretion of a low-mass satellite by a disk galaxy can generate a strong bar. This bar can buckle vertically and eject disk material out of the disk giving rise to an X-structure similar to structures observed in some bulges (e.g. \\citealt{whi88}). Angular momentum transfer from the satellite to the disk may also occur during the infall phase. \\citet{hua97} demonstrate that satellites with a mass 0.3$M_{\\rm disk }$ change the orientation of the disk up to $10^{\\circ}$ and generate warps as the satellites, under dynamical friction, sink toward the center of the galaxy. In addition, \\citet{qui86} and \\citet{qui93} note that the effects on the vertical structure of the disk are not uniform across the disk. They observe some flaring of the disk at large radii, i.e. particles at large radius are orbiting at larger $z$ than those at small radius. As suggested by these authors, a thick disk produced by minor mergers should have a scale height that increases with radius. Repeated interactions of a large population of satellites with a stellar disk also lead to disk heating. \\citet{qui86} shows that the disk is not heated isotropically by the infall of a satellite. They note that the heating is largest near the center of the disk and that most of the kinetic energy is distributed in the disk plane, i.e. $\\sigma_z$ receives a small fraction of the available energy. In fact, most of the radial and azimuthal heating comes from the disk spiral response to an infalling satellite \\citep{qui93}. In addition, the disk spreads radially due to an angular momentum exchange between the satellite and the disk. \\citet{vel99} demonstrates that the heating and thickening rate of the disk differs for satellites on prograde and retrograde orbits. The former heat the disk and the second tilt it. \\citet{ben04} show that there are significant differences between heating rates for prograde and retrograde orbits and those differences are amplified by increasing the mass and concentration of the satellites. As noted by some authors \\citep{ben04,fon01}, only satellites with orbits that bring them close to the center of the galaxy affect the disk significantly.  Most of the results described above were based on controlled numerical experiments where a single satellite interacts with a \"live\" disk. The exception was the pioneering study by \\citet{fon01} who examine the evolution of a stellar disk embedded in a system of several hundred subhalos. The work presented here improves upon that study in a number of ways. First, the mass resolution of our simulations is a factor of 250 higher than in the \\citet{fon01} study. This is high enough to adequately follow the evolution of the satellites and disk and to suppress two-body effects. Second, the base model for the galaxy is the self-consistent equilibrium model of M31 designed to fit the observed photometric and kinematic data for the Galaxy. Finally, the assumed mass distribution of the subhalo population and internal structure of individual subhalos are motivated by results from cosmological simulations.  Our aim is to study the effect of tidal stripping on the satellite population. In particular, we attempt to quantify the number of satellites that survive with a detectable stellar surface brightness. (Our subhalos have a single component, but we can infer the surface brightness by modelling the stellar content, as described below.) In addition, we study the stellar halo formed from tidal debris paying particular attention to extended structures. As we will see, these structures are similar to observed features of M31 such as the Giant Stream \\citep{iba01}.  In \\S 2, we describe our N-body model of M31 and describes a control experiment without satellites designed to test the model's inherent stability. In \\S 3, we describe the initial conditions of our satellite system. In \\S 4, we discuss the evolution of the satellite population and follow with a discussion of the hypothesis that the metal-poor stellar halos arise in the tidal stripping of satellite galaxies.  Our main conclusions are summarized in \\S 5. ",
        "conclusions": "This study presents the evolution of a self-consistent population of satellites in the presence of a parent galaxy. The internal structure, mass function and spatial and orbital distribution of the satellite distribution is motivated by cosmological simulations. The work improves upon previous self-consistent studies (e.g. \\citealt{fon01}) by using a more realistic galaxy model and much improved numerical resolution. It follows the lead of current work on satellite tidal disruption to make quantitative predictions of the tidal debris field of galaxies (e.g. \\citealt{joh01}).  This work shows that members of a typical population of subhalos orbiting in a galactic will come close to the disk with a ten billion years timescale. In addition, it is possible to model these satellites and make direct predictions on the number one would expect to detect around a typical galaxy. We also agree that it is possible to maintain a disk in spite of substructure having a mass fraction of about 0.1$M_{\\rm{halo}}$.  Here we summarize the main results of this paper. \\begin{enumerate}  \\item Dynamical friction plays only a minor role in the evolution of the satellite system and the number density profile of the system is relatively unchanged over 10 Gyr. The orbits of the most massive satellites do show some decay but since they suffer severe tidal stripping, dynamical friction quickly becomes unimportant.  \\item The vast majority of the satellites survive in the galactic environment for more than a Hubble time; only 7 are completely disrupted. However, most of satellites lose a significant fraction of their mass due to tidal interactions.  \\item The satellite mass loss due to tidal stripping is described by two distinct phases. The first one is characterized by a sharp mass decline with $t_{1/2}$ = 3.5 Gyr. During that period, a small amount of mass loss, $5.3 \\times 10^8$ $M_{\\odot}$, is due to a transient state where satellites initially located inside 50 kpc are overfilling their tidal radius and start losing mass instantly. During the second phase, the satellites lose about 25\\% of their initial mass with $t_{1/2} = 9.3$ Gyr. Over 10 Gyr, the satellite population mass diminishes by about 75\\% and turns into tidal debris. These debris form long tidal tails around the galaxy and can be detected in photometric maps of M31. The final density profile of the light satellites can be approximated by the fitting function given in \\citet{hay03}.   \\item The spatial distribution of stars associated with infalling satellites follow a power law $\\rho \\propto r^{-3.5}$ at large radii. This result is in agreement with recent numerical studies and observations of the stellar halo in edge-on disk galaxies.  \\item No obvious Holmberg Effect is observed.  \\item The mock B-band photometric maps, computed under the assumption that the most tightly bound particles are kinematic tracers of the stars, show conspicuous features of bridges and tails comparable to what is actually observed. A star formation efficiency of 10\\% is necessary to match the morphological and photometric properties of our Giant Stream proxy with the real one.  \\item The number of satellites detected around M31 below $\\mu_V$ $\\approx$ 26.8 mag $\\rm{arcsec}^{-2}$ can match our results with a star formation efficiency of 10\\%. For a value of 1\\%, the number of satellites detected goes down to 0. A value of 5\\% can not be discarded since dissipation is not taken into account. Many of the satellites between 26.8 and 29.4 mag $\\rm{arcsec}^{-2}$ in a 5\\% star formation efficiency scenario would probably be detected under the limit of 26.8 mag $\\rm{arcsec}^{-2}$ if dissipation was included in the simulations.  \\item A star formation efficiency of about 5\\% is necessary to account for the actual M31 value of the surface brightness of the stellar halo measured at 55 kpc. Comparisons for radii larger than 55 kpc are observationally challenging because of the very low surface brightness in the outer parts of the galaxy. The simulated maps show a flattening in the inner part of the surface brightness profile (for R $<$ 55 kpc, $\\mu_B \\approx$ 28 mag $\\rm{arcsec}^{-2}$) showing that dynamical friction is unable to bring the satellites close to the center. Most of the mass is lost at larger radii.  \\item The formation of a bar around 5 Gyr is the result of interactions of satellites with the disk. This intriguing result will be discussed in a forthcoming paper.  \\end{enumerate}"
    },
    "0606/astro-ph0606686_arXiv.txt": {
        "abstract": "We report on our search for the optical/infrared counterparts to the central compact objects in four young supernova remnants: Pup A, PKS 1209$-$52, RCW 103, and Cas A.  The X-ray point sources in these supernova remnants are excellent targets for probing the existence of supernova fallback disks, since irradiation of a disk by a central X-ray source should lead to an infrared excess.  We used ground-based optical and near-infrared imaging and {\\em Spitzer Space Telescope} mid-infrared imaging to search for optical/infrared counterparts at the X-ray point source positions measured by the {\\em Chandra X-Ray Observatory}.  We did not detect any counterparts, and hence find no evidence for fallback disks around any of these sources.  In PKS 1209$-$52, we are able to exclude a nearby optical/infrared candidate counterpart. In RCW 103, a blend of 3 faint stars at the X-ray source position prevents us from deriving useful limits.  For the other targets, the upper limits on the infrared/X-ray flux ratio are as deep as (1.0--1.7)$\\times 10^{-4}$.  Comparing these limits to the ratio of $\\approx 6\\times10^{-5}$ measured for 4U~0142+61 (a young pulsar recently found with an  X-ray irradiated dust disk), we conclude that the non-detection of any disks around young neutron stars studied here are consistent with their relatively low X-ray luminosities, although we note that a similar dust disk around the neutron star in Pup~A should be detectable by deeper infrared observations. ",
        "introduction": "Neutron stars (NSs) are formed when massive stars end their lives during core-collapse supernovae.  During such an explosion, ``fallback'' may occur when the reverse shock, caused by the impact of the shock wave with the outer stellar envelope, reaches the newly formed NS \\citep{chevalier89, ww95}.  Depending upon the rotation rate of the massive progenitor, some of the fallback may have sufficient angular momentum to form a disk (\\citealt{md81}; \\citealt*{lwb91}).  Although widely predicted by contemporary numerical studies of supernovae \\citep{ww95}, there are few observational constraints on the existence of fallback disks or even of supernova fallback generally.  Fallback could have profound implications for the endgame of massive star evolution, particularly since it could lead a newborn NS to collapse into a black hole \\citep{chevalier89}.  Such an outcome might yet explain the lack of a detectable pulsar in SN 1987A (\\citealt{zam+98}; \\citealt*{fcp99}). Supernova fallback is also a promising mechanism for producing the debris disks necessary to form planets \\citep{ph93, pod93}, which are known to exist around at least one pulsar \\citep{wf92, wol94}.  The compact stellar remnants of supernovae obviously probe supernova physics.  While it was once thought that young radio pulsars like those in the Crab and Vela supernova remnants (SNRs) were prototypical of newborn NSs, it has recently been realized that the young NS population is unexpectedly diverse.  For example, there are a few young SNRs currently known to contain radio-quiet, non-plerionic X-ray point sources near their geometric center (see \\citealt*{pst04}, hereafter \\citetalias{pst04}, for a review).  These X-ray point sources, often called central compact objects (CCOs), have X-ray spectra roughly consistent with the thermal-like surface (but smaller than the expected neutron star surface) emission from young cooling NSs.  However, they differ from ``typical'' young NSs (e.g., the Crab pulsar) in SNRs in that they lack strong non-thermal emission components and most of them do not have detectable pulsations.  Along with other young NS classes including the anomalous X-ray pulsars (AXPs), and soft gamma-ray repeaters \\citep{wt04}, these objects challenge our standard framework for understanding young NSs.  Because CCOs and AXPs do not have the strong, broad-band, non-thermal emission typically observed from young Crab-like pulsars and their associated synchrotron nebulae, they offer an attractive opportunity for testing the existence of fallback disks by detecting their expected thermal emission.  \\citet*{phn00} showed that the combination of an untruncated outer edge of a fallback disk (in contrast to a binary accretion disk) and irradiation of the disk by the central pulsar should lead to a strong thermal excess in the infrared and submillimeter bands.  In a previous paper, we reported the detection of a mid-infrared counterpart to the AXP 4U 0142+61, which we interpreted as dust emission from an X-ray heated debris disk around the pulsar \\citep*{wck06}.  Here, we report our ground-based optical/near-IR and {\\em Spitzer Space Telescope} 4.5/8.0 $\\mu$m observations of four CCOs. \\begin{figure*} \\includegraphics[scale=0.9]{f1.eps} \\figcaption{Optical and IR images of \\pup.  The {\\em Chandra} error circle (solid circle) of the CCO is indicated in the optical $R$ (right panel) and near-IR $K_s$ (middle panel) bands images.  The left panel shows the {\\em Spitzer}/IRAC 8.0 $\\mu$m band image, with the box and circle % indicating the size of the $R$ and $K_s$ images and the CCO position respectively.  The circle is not to scale.  No object was found within the error circle. \\label{fig:pup}} \\end{figure*} These targets were selected from a total of 8 known CCOs because they are relatively young and their distances are well determined (\\citetalias{pst04}). We give brief summaries of our targets in \\S~\\ref{sec:targ}.  We then present the optical/IR observations in \\S~\\ref{sec:obs}.  Finally, we present our results in \\S~\\ref{sec:res} and our conclusions in \\S~\\ref{sec:disc}. Part of our ground-based observations have previously been reported \\citep{wc02}.    \\subsection{Target Summary} \\label{sec:targ} Of the four CCOs that we observed, the basic properties of these sources are similar: young, radio quiet, and X-ray bright, with roughly similar X-ray spectra (again, see \\citetalias{pst04}). However, some have distinguishing characteristics, as we note below. We also describe observations of their host SNRs used to determine ages and distances, fundamental quantities that we use in the interpretation of the observations.  A summary of the general properties of the CCOs that we observed is given in Table~\\ref{tab:xps}.  \\begin{description} \\item[\\pup\\ (Pup A)] The X-ray source \\pup\\ (\\citealt{pkwc82};  \\citealt*{pbw96}) is located in the SNR Puppis~A (Pup A; G260.4$-$3.4).  The SNR is the remnant of a Type II explosion that occurred $<5000$~yrs ago (\\citealt{wtki88};  \\citealt*{adp91}).  The distance to the SNR, determined from associated \\ion{H}{1} absorption, is 2.2$\\pm0.3$~kpc \\citep{rgj+03}. The CCO is unremarkable, except that \\citet{hb05} recently reported a possible detection of 0.22-s X-ray pulsations.  \\item[{\\pks} (PKS 1209--52)] The source \\pks\\ \\citep{hb84} is in the SNR PKS~1209$-$52 (G296.5+10.0).  This SNR has a distance of 2.1$^{+1.8}_{-0.8}$~kpc based on  associated \\ion{H}{1} gas \\citep{gbg+00} and an age of $\\approx 7$~kyr \\citep{rmk+88}. The CCO is the only confirmed pulsator among the sources considered here ($P_{\\rm spin}$=424 ms; \\citealt{zpst00,pzst02}) although its spin does not change monotonically \\citep*{zps04}, and its X-ray spectrum is quite unique with a number of related X-ray absorption features \\citep{spzt02,bcdlm03}.  \\item[{\\rcw} (RCW~103)] The CCO \\rcw\\ \\citep{tg80} is in the SNR RCW~103 (G332.4$-$0.4).  The distance to the SNR of 3.1~kpc has been estimated from an \\ion{H}{1} association \\citep{rgj+04}, while the age of $\\approx 2$~kyr was determined from measurement of the expansion of optical filaments \\citep*{cdb97}. Unlike the other objects, the CCO has shown considerable X-ray variability on a timescale of years \\citep*{gpv99,ba02}, and at least three faint near-infrared sources have been found within the 0\\farcs7 {\\em Chandra} X-ray position uncertainty (\\citetalias{pst04}). Together, these facts suggest that \\rcw\\ is a low-mass X-ray binary with a 6.4 hr orbital period (\\citealt{ba02}; \\citetalias{pst04}).  \\item[{\\cas} (Cas~A)] \\cas\\ \\citep{t99} is in the SNR Cassiopeia A (Cas A; G111.7-2.1), the youngest known Galactic core-collapse SNR ($\\approx$300 yr; \\citealt*{tfvdb01}), but the CCO was only revealed recently (compared to the others, which had been identified by the {\\em Einstein Observatory}) by the {\\em Chandra} X-ray Observatory \\citep{t99,pza+00,cph+01}.  Comparing proper motions and radial velocities of ejecta, \\citet{rhfw95} estimate a distance of $3.4_{-0.1}^{+0.3}$~kpc.  The X-ray properties are, like those of \\pup, unremarkable, and this CCO serves as a benchmark for the class. Deep optical/near-IR observations of the source field were conducted \\citep*{kkm01,fps06}, but no counterpart was found. \\end{description}    \\begin{figure*} \\includegraphics[scale=0.9]{f2.eps} \\figcaption{Optical/IR images of the 1E 1207.4$-$5209 region.  In the left panel, the {\\em Spitzer}/IRAC 8.0 $\\mu$m band image is shown, with the box and circle % indicating the size of the middle and right panel images and the CCO position respectively.  The circle is not to scale. In the middle panel, a $K_s$ band image is shown. The two solid circles indicate the {\\em Chandra} HRC and ACIS positions, which are about 1\\farcs0 away from the object labeled {\\em X}.  The right panel shows the optical $i'$ band image of the target field.  The near-IR object was not detected by our optical observations, but was found with $V\\simeq 26.8$ by Pavlov et al.\\ (private communication).  Previous observations have identified the two blended stars ($A$ and $B$) as main sequence stars (\\citealt*{mlt88}; \\citealt*{bcm92}; see also \\citealt{dlmc+04}). \\label{fig:pks}} \\end{figure*} ",
        "conclusions": "\\label{sec:disc} Based on the current observed properties, it is likely that the CCOs represent elements of a new class of young NSs, or even of several new classes (given their heterogeneous properties). Their overall characteristics --- radio-quietness, distinct thermal X-ray emission, relatively long spin periods (when detected) for their ages, and lack of pulsar wind nebulae --- indicate that they distinctively differ from other classes of young NSs, especially the young rotation-powered pulsars.  The possible origin of the CCOs has been well discussed by comparing them to different types of NS objects and models, which include the rotation-powered pulsars, AXPs, magnetars, accreting binaries, accreting NSs from a residual disk (e.g., \\citealt{cph+01}; \\citetalias{pst04}; \\citealt{ghs05}).  Among them, one intriguing model is that the CCOs are NSs accreting from a residual disk, probably formed from fallback of supernova material.  This model was primarily proposed as an alternative to the magnetar model for AXPs in order to explain their ``anomalous'' natures: that their X-ray luminosities are much higher than their rotational energy loss rates (e.g., \\citealt*{chn00}; \\citealt{mlrh01}).  However, the X-ray fluxes of most of the CCOs, especially our targets (e.g., \\citealt{zps04,fps06}), have been steady based on the observations over the past few years, contrary to the fact that accreting X-ray sources are highly variable (this is in contrast to \\rcw, which \\citetalias{pst04} therefore exclude from the CCO category).  In addition, the optical observations of CCOs have failed to detect any counterparts that could be an accretion disk (e.g., \\citealt{fps06,ghs05}). Particularly for \\pks, which is the least reddened among the CCOs, the optical observations (including ours) were deep enough to exclude the existence of an accretion disk around it (except an edge-on disk, which has very low probability; \\citealt{zps04}).  Therefore, even though the CCOs are probably not a uniform class of objects, it is very unlikely that their X-ray emission is powered by accretion. \\begin{figure*}[t] \\begin{center} \\includegraphics[scale=0.80]{f5.eps} \\figcaption{Upper limits on the optical/IR to X-ray flux ratio for the CCOs: ($a$) \\pup\\ ; ($b$) \\pks\\ ; ($c$) \\cas\\ (the optical/near-IR $RJHK$ upper limits are from \\citealt{fps06}). Fluxes are unabsorbed, and $A_V=$1.6, 0.5, and 7 are respectively used for dereddening the optical/IR data (according to the reddening laws of \\citealt*{sfd98} for the $BVRI$ data, \\citealt{fan99} for the $u^{\\prime}g^{\\prime}r^{\\prime}i^{\\prime}$ data, \\citealt{rl85} for the $JHK$ data, and \\citealt{imb+05} for the IRAC data). ($d$) Optical/IR to X-ray flux ratios (for comparision) of the AXP 4U 0142+61 (plus sign; \\citealt{wck06}) and the Crab pulsar (long-dashed curve; \\citealt{efr+97}; Temim et al.\\ 2006, in preparation).  The IR component of 4U 0142+61 is likely to be due to an X-ray irradiated dust disk (dash-dotted curve). The optical/IR observations should also have detected a main-sequence star down to M5 (dotted curve; from \\citealt{allen} and \\citealt{kurucz93}, and normalized to a typical 0.4--8.0 keV CCO X-ray flux 2$\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$) even at a distance of 3.4 kpc. \\label{fig:all} } \\end{center} \\end{figure*}  However, it is still possible that isolated, young NSs have cool, passive disks. Such a disk would be farther away from a central NS and cooler than an accretion disk. If they exist, the disks would not contribute to generating X-ray emission from young NSs such as the CCOs and AXPs, but their existence might help explain various phenomena of radio pulsars \\citep{md81,bp04,cs06}, planets formation around NSs \\citep{lwb91}, or the diversity of young NSs \\citep{alpar01}.  The recent discovery of the mid-IR emission from the AXP 4U~0142+61 probably indicates the existence of such a passive disk around that young pulsar \\citep{wck06}.  The disk, irradiated by the X-rays from the central X-ray pulsar, emits mainly in the IR (Figure~\\ref{fig:all}).  Using {\\em Spitzer}/IRAC observations similar to those discussed here, its unabsorbed IR--to--X-ray flux ratio was found to have a maximum value of $(\\nu_{4.5\\mu}F_{4.5\\mu})/F_{\\rm X} \\simeq$6.3$\\times 10^{-5}$ at 4.5 $\\mu$m.  This ratio, as it depends largely on the geometry of the disk and not on the peculiarities of the central object, might be typical for any dust disks around young NSs. In Figure~\\ref{fig:all}, we compare the upper limits on the IR/X-ray flux ratios of the CCOs to the same ratios of 4U~0142+61. The upper limits on the IR/X-ray flux ratios of the CCOs are not sufficiently deep to detect a disk similar to that found in 4U 0142+61. As the X-ray luminosities of the CCOs are typically lower ($\\simeq$$2\\times 10^{33}$ erg s$^{-1}$; see Table 1), deeper IR exposures are needed in order to reach a comparable IR/X-ray flux ratio. Accordingly, in order to further explore any existence of dust disks around the CCOs, \\pup\\ would be a plausible target. For example, a 4.5 $\\mu$m flux of 0.6 $\\mu$Jy would be expected from a dust disk, which is detectable by deep {\\em Spitzer}/IRAC observations (several hours exposure time).  In comparison, for the very young CCO \\cas, although mid-IR observations are less affected by reddening ($A_V\\simeq 7$; \\citealt{fps06}), diffuse emission of the SNR highly contaminates the source field, making it extremely difficult to have a sensitive search.   It is interesting to compare the flux ratio limits to those of the Crab-like, rotation-powered pulsars, even though it is clear that their optical/IR emission is nonthermal, originating from the magnetospheres of the pulsars.  For several well-studied rotation-powered pulsars, optical/IR emission can be connected to nonthermal hard X-rays with a same power-law (e.g., PSR B0656+14; \\citealt*{pzs02}).  The connection may be further indicated by a similar optical to nonthermal--X-ray flux ratio, $F_{\\rm opt}/F^{\\rm nonth}_{\\rm X}\\simeq 0.001$--0.01, where $F_{\\rm opt}$ is the optical flux between 4000--9000 \\AA\\ and $F^{\\rm nonth}_{\\rm X}$ is the non-thermal X-ray flux between 1--10 keV \\citep{zp04}.  In Figure~\\ref{fig:all}, we also show the unabsorbed SED of the Crab pulsar, normalized by its X-ray flux (1.5$\\times$10$^{36}$ ergs s$^{-1}$; \\citealt{zp04}). The SED is above most of the upper limits on the CCOs, indicating that our observations would have detected Crab-like young pulsars.  This comparison has additionally shown that the CCOs are a different class of young NSs from the typical Crab-like pulsars.  Our optical/IR observations have additionally provided strong constraints on the binary model proposed for the CCOs.  As the upper limits of our observations of \\pup\\ and \\pks\\ are 24--26 mag in the optical and 20--23 mag in the near-IR (or $\\nu F_{\\nu}\\leq 5\\times 10^{-15}\\mbox{ erg s}^{-1}\\mbox{ cm}^{-2}$; see Figure~\\ref{fig:all} and note that the X-ray fluxes are around 10$^{-12}\\mbox{ erg s}^{-1}\\mbox{ cm}^{-2}$), they would have detected main-sequence stars down to spectral type M5 regardless of extinction or white dwarfs down to effective temperature $T_{\\rm eff}\\simeq$ 10$^4$ K \\citep*{bwb95}, assuming relatively low extinction (e.g., $A_V\\leq$1--2; \\citealt{ps95}; see Table~1), out to distances of 2--3~kpc.  In summary, we have obtained  deep multi-wavelength observations of four CCOs in young SNRs. The observations have allowed us to find a near-IR counterpart to the CCO in RCW 103, which probably corresponds to three faint late type stars resolved by  high-resolution observations.  In addition, the observations have detected a low mass star close to, but inconsistent with, the current {\\em Chandra} positions of \\pks. The non-detection of any IR emission from debris disks around three of the CCOs could be due to their relatively low X-ray luminosities, compared to 4U~0142+61.  While high background prohibits additional observations of two of the sources, deeper IR observations of \\pup\\ could probe the existence of a fallback disk around this CCO."
    },
    "0606/astro-ph0606365_arXiv.txt": {
        "abstract": "We present a 3-dimensional study of the local  ($\\leq 100 \\; h^{-1}$ kpc) and the large scale ($\\leq$ 1 $h^{-1}$ Mpc) environment of Bright IRAS Galaxies (BIRGs). For this purpose we use 87 BIRGs located at high galactic latitudes (with 0.008$\\leq z \\leq$0.018) as well as a control sample of non-active galaxies having the same morphological, redshift and diameter size distributions as the corresponding BIRG sample. Using the Center for Astrophysics (CfA2) and Southern Sky Redshift Survey (SSRS) galaxy catalogues ($m_b\\lesssim 15.5$) as well as our own spectroscopic observations ($m_b\\lesssim19.0$) for a subsample of the original BIRG sample, we find that the fraction of BIRGs with a close neighbor is significantly higher than that of their control sample. Comparing with a related analysis of Sy1 and Sy2 galaxies of Koulouridis et al. (2006) we find that BIRGs have a similar environment as Sy2s, although the fraction of BIRGs with a bright close neighbor is even higher than that of Sy2 galaxies. An additional analysis of the relation between FIR colors and the type of activity of each BIRG shows a significant difference between the colors of strongly-interacting and non-interacting starbursts and a resemblance between the colors of non-interacting starbursts and Sy2s. Our results support the view where close interactions can drive molecular clouds towards the galactic center, triggering starburst activity and obscuring the nuclear activity. When the close neighbor moves away, starburst activity is reduced with the simultaneous appearance of an obscured (type 2) AGN. Finally, the complete disentanglement of the pair gives birth to an unobscured (type 1) AGN. ",
        "introduction": "The IRAS Revised Bright galaxy sample by Sanders et al. (2003) includes all galaxies with total 60 $\\mu m$ flux density greater than 5.24 Jy. The sample is the result of a highly complete flux-limited survey conducted by IRAS covering the entire sky at galactic latitudes $|b|\\geq 5^o$ and was  compiled after the final calibration of the IRAS Level 1 Archive. It offers far more accurate and consistent measurements of the flux of objects with extended emission. In addition, the infrared fluxes of over 100 sources from the sample were recalculated by the IRAS High Resolution (HIRES) processing, which allowed the deconvolution of close galaxy pairs (Surace, Sanders, Mazzarela 2004). The latter provides a more-than-ever reliable database of the IRAS galaxies which can be proved crucial for statistical studies like this one.  While the relation between Ultra Luminous IRAS Galaxies (ULIRGs) and strong interactions has been thoroughly studied (e.g. Sanders, Surace \\& Ishida 1999, Wang et al. 2006), this is not the case for the environment of moderately and low luminous infrared galaxies. A 2-dimensional analysis of Krongold et al. (2002) showed a trend for a Bright IRAS Galaxy (BIRG) sample on having neighbors in excess of normal galaxies and Sy1 galaxies, but in relative agreement with Sy2 galaxies. However, the BIRG population consists of various types of active galaxies, including starbursts (the majority), Seyferts, Liners and normal galaxies and thus it would be of great interest to clarify the connection between infrared emission, interactions and different types of active galaxies.  During the last decade, many studies have investigated the relation among interacting galaxies, starbursting and nuclear activity (eg. Hernandez-Toledo et al. 2001;  Ho 2005). Despite the plethora of available information, searches for correlations between the above physical processes are inconclusive, the only exception being the coupling between interactions and starbursting. However, there is evidence that AGN galaxies host a post-starburst stellar population (eg. Boisson et al. 2000, Gonz\\'alez Delgado et al. 2001) while Kauffmann et al. (2003) showed that the fraction of post-starburst stars increases with AGN emission. Proving a relation of this type between starburst and AGN galaxies would simultaneously solve also the problem of the AGN triggering mechanism. Interactions would be the main cause of such activities, being starbursting and/or the feeding of a central black hole. However, this is not a trivial task. The main difficulty arises from the fact that the Star Formation Rate (SFR) estimation in AGN host galaxies is still problematic. All SFR estimation methods present complications and even those based on the FIR continuum are doubtful, since the contribution of the active nuclei is unknown (eg. Ho 2005).  Despite the difficulties, some studies, based on different diagnostics seem to conclude that there is indeed an evolutionary sequence from starburst to type 2 and then to type 1 AGN galaxies (e.g. Oliva et al. 1999, Krongold et al. 2002). In addition, Kim, Ho and Im (2006), using the [OII] emission line as a SFR indicator, reach the conclusion that type 2 are the precursors of type 1 quasars supporting the previous claims. These studies are based on the observed differences between different types of AGNs and resemblance of type 2 objects to starbursts. This raises doubts about the simplest version of the unification scheme of AGNs. It is true that the recent discovery of 10$\\mu m$ silicate emission in two luminous quasars implies the presence of dust, but it is not clear yet what is the spatial distribution of this material (Siebenmorgen et al. 2005). In addition, silicate emission is not yet detected in other type 1 objects and thus more observations are needed to establish the existence of the dusty torus.  We can summarize all the previous in two statements: (1) the starburst-AGN connection is still not well established and (2) the AGN unification model, although successful in interpreting many observational facts, remains fragile. From our point of view, our BIRG sample offers a homogeneous and complete database, which is ideal for a statistical study on these issues.  We will discuss our galaxy samples in section \\S 2. Our data analysis and results will be presented in \\S 3, while in section \\S 4 we will discuss our results and  present 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$ Mpc. ",
        "conclusions": "We have compared the 3-dimensional environment of a sample of local BIRGs with that of a well defined control sample, selected in such a way as to reproduce the redshift, morphological type and diameter size distributions of the BIRG sample. We searched for close neighbors around each BIRG and control sample galaxy using the distribution of CfA2 and SSRS galaxy catalogues as well as our own spectroscopic observations reaching a fainter magnitude limit (but for a restricted BIRG subsample). We also compared our results with those of a similar analysis of Seyfert galaxies (Dultzin-Hacyan et al. 1999; Koulouridis et al. 2006).  We find that the fraction of BIRG galaxies having a close neighbor, within a projected separation of 75 $h^{-1}$ kpc and radial velocity difference of $\\delta v \\le 600$ km/s, is significantly higher than the corresponding fraction of its control sample and that of Sy1 galaxies while it is comparable to that of Sy2 galaxies. This result is in accordance with some previous two-dimensional studies (eg. Krongold et al. 2002). We reach similar results regarding the large scale environment of BIRGs (within a projected radial separation of 1 $h^{-1}$ Mpc and a radial velocity difference $\\delta v \\le 1000$ km/s). Once more their behavior resembles that of Sy2's but not of Sy1's. We also find a statistically significant difference between highly-interacting and non-interacting BIRGs, based on their FIR color properties. Sy2s appear to display a similar behavior to that of non-interacting starburst galaxies, introducing new evidence for the starburst/AGN connection scenarios.  Our results can be accommodated in a simple evolutionary scenario, starting with an interaction, and ending in a Sy1 phase. First, close interactions would drive molecular clouds towards the central area, creating a circumnuclear starburst. Then, material could fall even further into the innermost regions of the galaxy, feeding the black hole, and giving birth to an AGN which at first cannot be observed due to obscuration. At this stage only a starburst would be observed. As starburst activity relaxes and obscuration decreases, a Sy2 nucleus would be revealed (still obscured by the molecular clouds from all viewing angles). As a final stage, a Sy1 phase could appear. In this case, the molecular clouds, initially in a spheroidal distribution, could flatten and form a ``torus'' (as in the unification scheme for Seyferts). As more material is accreted, it is possible that the AGN strengthens driving away most of the obscuring clouds, and leaving a ``naked'' Sy1 nuclei.  If indeed interactions play a role in triggering activity, as suggested by the above picture, then the lack of close companions among Sy1 galaxies implies that 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 ($\\sim10^9$ years, see Krongold et al. 2002).  It should be noted that the evolutionary scenario does not contradict the unification scheme. It implies that Sy1s and Sy2s are the same objects (as the unification model proposes) but not necessarily at the same evolutionary phase. However, there could be a phase where only orientation could define if an object appears as a Sy1 or as a Sy2, which is the stage where molecular clouds form a torus but have not been swept away yet.  Evidently, more detailed observations of large samples of galaxies are needed to resolve this important issue. However, such a picture is consistent with the evolutionary scenario suggested by Tran (2003). He studied a sample of Sy2 in polarized light, and found that only 50\\% of them showed the presence of a hidden broad line region (HBLR). He suggested that non-HBLR Sy2 galaxies could evolve into HBLR Sy2 galaxies. In this case, the appearance of the BLR could be related to the accretion rate (e.g. Nicastro 2000), and thus with the evolutionary stage of the object. The trend is also consistent with the finding that 50\\% of Sy2 galaxies also show the presence of a strong starburst (Gu et al. 2001; Cid-Fernandez et al. 2001).  On the brightest end, there is also growing evidence showing (a) that ULIRGs can be the precursors of quasars and (b) ULIRGS are found in very strong interacting systems or in mergers (e.g. Sanders, Surace and Ishida 1999; Wang et al. 2006). Therefore, the evolutionary sequence proposed above could be general to all nuclear activity and independent on luminosity (we note that Krongold et al. 2003 suggested a similar scheme for LINERs that can be considered as the very low luminosity extension of the evolutionary model suggested here). Further evidence comes from the fact that type 2 quasars also tend to be in interaction more often than type 1 quasars (Serber et al. 2006).  In order to better understand the role of interactions in driving starburst and nuclear activity (and the validity of the evolutionary trend), we are in  the process of studying AGN and starburst manifestations in  the nearest neighbors of the active galaxies in our samples, since the same physical processes should act on both members of the pair (host and nearest neighbor)."
    },
    "0606/astro-ph0606479_arXiv.txt": {
        "abstract": "We present LVG and non-local radiative transfer calculations involving the rotational and hyperfine structure of the spectrum of N$_2$H$^+$ with collisional rate coefficients recently derived by us. The goal of this study is to check the validity of the assumptions made to treat the hyperfine structure and to study the physical mechanisms leading to the observed hyperfine anomalies.  We find that the usual hypothesis of identical excitation temperatures for all hyperfine components of the $J$=1-0 transition is not correct within the range of densities existing in cold dense cores, i.e., a few 10$^4$ $\\textless$ n(H$_2$) $\\textless$ a few 10$^6$ cm$^{-3}$. This is due to different radiative trapping effects in the hyperfine components. Moreover, within this range of densities and considering the typical abundance of N$_2$H$^+$, the total opacity of rotational lines has to be derived taking into account the hyperfine structure. The error made when only considering the rotational energy structure can be as large as 100\\%.  Using non-local models we find that, due to saturation, hyperfine anomalies appear as soon as the total opacity of the $J$=1-0 transition becomes larger than $\\simeq$ 20. Radiative scattering in less dense regions enhance these anomalies, and particularly, induce a differential increase of the excitation temperatures of the hyperfine components. This process is more effective for the transitions with the highest opacities for which emerging intensities are also reduced by self-absorption effects. These effects are not as critical as in HCO$^+$ or HCN, but should be taken into account when interpreting the spatial extent of the N$_2$H$^+$ emission in dark clouds. ",
        "introduction": "N$_2$H$^+$ was one of the first molecular ions detected in interstellar space \\citep{tha75}. The $J$=1-0 line of this species has been extensively observed toward cold dark clouds and protostellar cores to get some estimates of the physical conditions of the gas (see, e.g., \\citet{ber02,taf04,hot04,bel04,cas02}). These observations indicate that N$_2$H$^+$ is a good tracer of the highest density regions of dark clouds. It seems that N$_2$H$^+$ is less depleted onto dust grain surfaces than CO and other molecular species. This is probably related to the fact that N$_2$, the chemical mother species of N$_2$H$^+$, is more volatile and condensates at lower temperatures than carbon monoxide. In addition, the complex hyperfine structure of N$_2$H$^+$ increases the odds to have at least one optically thin hyperfine line component to probe the innermost regions of these clouds. Therefore, this species is in principle an interesting tool to study cold dark clouds. However, a drawback has been the lack of collisional rate coefficients between N$_2$H$^+$ and molecular hydrogen (or helium). The observational data indicate some hyperfine intensity anomalies that could be due to selective collisional processes or to radiative transfer effects (see \\citet{gon93} for the case the analogous case of HCN hyperfine intensities in dark clouds).  For current research, it is clear that astronomers need to know the state to state collisional rates of N$_2$H$^+$ with H$_2$ and He. This will be even more necessary for ALMA due to the much higher angular resolution and higher sensitivity observations that it could provide of protostellar cores in several rotational transitions of N$_2$H$^+$ (up to $J$=9-8).  N$_2$H$^+$ has also been detected in warm molecular clouds (e.g. \\citet{tur77}) where the lines are broader and very strong. In these objects only the hyperfine structure due to the external N-atom could be noticed as the splitting produced by the internal N-atom is lower than the intrinsic line width. Nevertheless, in order to correctly model the N$_2$H$^+$ intensities emerging from these clouds, astronomers need a complete set of state to state collisional rate coefficients for high temperatures.  A detailed study on molecular ions excitation was carried out by \\citet{gre75} and a set of collisional rate coefficients for N$_2$H$^+$ colliding with He was provided. In that work, the rate coefficients were given for transitions among N$_2$H$^+$ rotational energy levels. Recently, \\citet{dan05} extended the previous study by computing a new set of collisional rate coefficients for transitions among hyperfine energy levels and using a new potential energy surface. The range of kinetic temperatures is between 5-50 K and in a future paper, rate coefficients for temperatures up to 300 K will be provided.  The paper is organized as follows: Section \\ref{spectroscopy} is devoted to the spectroscopy of N$_2$H$^+$. In section \\ref{LVG} we present the results obtained with a Large Velocity Gradient (LVG) model to discuss the treatment of molecular hyperfine transitions. Comparisons with the cases of HCN and HCl are made. In section \\ref{non-local} we present results from non-local radiative transfer models applied to N$_2$H$^+$ for different cloud structures. ",
        "conclusions": "We have used numerical methods in order to investigate the excitation properties of molecules with a hyperfine energy structure, focusing specially on N$_2$H$^+$. Our conclusions are summarized as follows : \\begin{enumerate} \\item The hyperfine structure must be taken into account in the radiative transfer calculations in order to derive the total amount of molecules present in a given rotational level. When the energy structure of the molecule is restricted to its rotational energy structure, high-$J$ levels are more efficiently populated and the opacities of the low-$J$ transitions are underestimated. The error increases with the column density of N$_2$H$^+$. Errors as large as a factor 2 could be induced if the hyperfine structure of N$_2$H$^+$ is neglected. \\item For the typical temperature of dark clouds, i.e. T$_K$=10 K, and n(H$_2$) $\\textless$ 10$^6$ cm$^{-3}$ the ratio of hyperfine brightness temperatures derived in the LVG approximation are always smaller than the ones predicted using the LTE approximation. Thus, this latter method is inadequate to assess densities or column densities from observed N$_2$H$^+$ lines for the typical conditions prevailing in such clouds. \\item The assumption of equal excitation temperatures for all hyperfine components belonging to the same rotational transition is valid for low opacities and fails down for high opacities due to radiative processes. For low opacities, we do not find any difference in the excitation temperatures of the different hyperfine transitions induced by collisional rate coefficients. Our calculations indicate that for temperatures in the range 5-50 K and for all volume densities, the excitation temperature of these lines will be identical if the total opacity of the $J$=1-0 line is small. We stress that such collisional excitation effects, often found in the literature, may not be invoked to explain the intensity anomalies reported for N$_2$H$^+$. \\item Non-local radiative transfer results show that the $J$=1-0 intensity anomalies arise while the opacity of this line increases. Moreover, photon scattering in the low density envelopes affect the $J$=1-0 hyperfine lines differently and tend to reduce the emerging intensity of the thicker lines whereas the intensity of the thinnest lines is not modified by the envelope. Thus, in order to reproduce the intensities of the $J$=1-0 hyperfine lines, the only valid method is a non-local computation of the radiative transfer. \\item For the typical conditions of cold dark clouds, the $J$=2-1 line is a good probe of low velocity fields. Nevertheless, this transition is difficult to observe from ground based observatories due to the high opacity of the atmosphere at 186 GHz, although some work can be carried out in dry weather conditions. \\end{enumerate}"
    },
    "0606/astro-ph0606153_arXiv.txt": {
        "abstract": "A solar photospheric ``thermal profiling'' analysis is presented, exploiting the infrared (2.3--4.6~\\micron) rovibrational bands of carbon monoxide (CO) as observed with the McMath-Pierce Fourier transform spectrometer (FTS) at Kitt Peak, and from above the Earth's atmosphere by the Shuttle-borne ATMOS experiment.  Visible continuum intensities and center-limb behavior constrained the temperature profile of the deep photosphere, while CO center-limb behavior defined the thermal structure at higher altitudes. The oxygen abundance was self consistently determined from weak CO absorptions (for C/O$\\equiv\\,0.5$). Our analysis was meant to complement recent studies based on 3-D convection models which, among other things, have revised the historical solar oxygen (and carbon) abundance downward by a factor of nearly two; although in fact our conclusions do not support such a revision.  Based on various considerations, an $\\epsilon_{\\rm~O}=\\,700{\\pm}100$~ppm (parts per million relative to hydrogen) is recommended; the large uncertainty reflects the model sensitivity of CO.  New solar isotopic ratios also are reported: $^{12}$C/$^{13}$C=$80{\\pm}1$, $^{16}$O/$^{17}$O=$1700{\\pm}220$, and $^{16}$O/$^{18}$O=$440{\\pm}6$; all significantly lower than terrestrial. CO synthesis experiments utilizing a stripped down version of the 3-D model---which has large temperature fluctuations in the middle photosphere, possibly inconsistent with CO ``movies'' from the Infrared Imaging Spectrometer (IRIS), and a steeper mean temperature gradient than matches visible continuum center-limb measurements---point to a lower oxygen abundance ($\\sim$500~ppm) and isotopic ratios closer to terrestrial.  A low oxygen abundance from CO---and other molecules like OH---thus hinges on the reality of the theoretically predicted mid-photospheric convective properties. ",
        "introduction": "For nearly three decades, the ``solar neutrino problem'' haunted Solar Physics. The apparent dearth of the elusive subnuclear byproducts of proton--proton cycle fusion presented severe challenges for theoretical descriptions of interior conditions in the Sun; models that heretofore had been regarded as ironclad and unassailable (Bahcall, Basu, \\& Pinsonneault 2003, and references to previous work therein).  The neutrino problem thankfully was resolved in recent years, ironically from the particle-physics rather than astrophysics side (e.g., Bahcall, Krastev, \\& Smirnov 1999), and solar (and stellar) interior models once again seemed secure.  A new challenge, however, has confronted the even more refined contemporary generation of solar interior models crafted for---and to a large extent by---helioseismology; what one might call the ``solar oxygen crisis.'' Sophisticated analyses, employing time dependent 3-D simulations of convection driven solar surface velocity fields and thermal inhomogeneities, have pointed to a surprisingly low oxygen abundance, based on an impressive collection of tracers including forbidden and weak permitted \\ion{O}{1} lines in the visible spectrum, and hydroxyl (OH) rovibrational and pure rotational bands in the 3--10~\\micron\\ thermal infrared (Asplund et al.\\ 2004, and references to previous work therein).  The new low O abundance ($\\epsilon_{\\rm O}=$\\,460$\\pm 50$ parts per million [ppm] relative to hydrogen versus the most recent previously recommended value 650$\\pm 100$ ppm [Grevesse \\& Sauval 1998, hereafter GS98]: see Figure~1) plays an important role in interior models of the Sun because oxygen ranks third after hydrogen and helium by number in a solar chemical mixture, and contributes very substantially to the interior radiative opacity.  Furthermore, the next most abundant species---C, N, and Ne---are very difficult to measure (as is O itself) in the visible spectrum owing to lack of suitable absorption lines, and changes in their abundances often follow closely any alterations to $\\epsilon_{\\rm O}$.  In other words, abundance ratios such as C/O ($\\sim 0.5$ [Allende Prieto, Lambert, \\& Asplund 2002]) or Ne/O (derived from solar energetic particles [e.g., Meyer 1985] or high resolution spectra of Ne and O ions in the solar transition zone [Warren 2005]) often are regarded as better determined than the absolute abundances themselves. Thus, a revision in $\\epsilon_{\\rm O}$ could have a domino-like effect on the overall solar chemical composition and heavy element mass fraction $Z$, one of the crucial parameters governing the interior structure of a star and its evolution.  The revised low oxygen abundance is said to resolve outstanding issues in interstellar medium and young stellar population studies (Asplund et al.\\ 2004), which previously had pegged the Sun as an ``oxygen-rich'' dwarf; at odds with its middle-aged status in a galaxy undergoing steady metal enrichment owing to successive generations of stellar birth and death (Timmes, Woosley, \\& Weaver 1995). At the same time, the solar oxygen crisis has provoked the helioseismology community by casting doubt on the previous spectacularly good agreement between simulated interior sound speed profiles, and the depth of the convection zone, with values deduced from detailed, high precision measurements of surface $p-$mode oscillations from the ground and space (Bahcall et al.\\ 2005).  Ironically enough, the new oxygen crisis also is having an impact on understanding solar neutrino fluxes (Bahcall \\& Serenelli 2005).  There appears to be no easy way out of the dilemma from the helioseismology side: the upward revisions in interior opacities needed to accommodate the new low $\\epsilon_{\\rm O}$ are larger than permitted by uncertainties in the best contemporary theoretical and laboratory atomic physics data  (Antia \\& Basu 2005). On the low O side, the latest generation of time dependent 3-D solar convection models have become quite sophisticated, and are said to allow the precise matching of intensity profiles of weak velocity-sensitive photospheric lines without the {\\em ad hoc}\\/ specification of extraneous dynamical parameters (such as the micro- or macroturbulent velocity fields familiar from classical 1-D abundance modeling [e.g., Gray 1976]).  This is an extremely important advance in the {\\em ab initio}\\/ modeling---and understanding---of the structure of the solar photosphere, and by extension of all late-type stars dominated by strong convective heat fluxes in their surface layers.  Furthermore, the analyses by Asplund and collaborators of the fundamental atomic physics of the various oxygen abundance diagnostics, of blending issues, and of departures from Local Thermodynamic Equilibrium (LTE) in the permitted \\ion{O}{1} lines are persuasive. In fact, the three major facets of the low O problem---3-D model, atomic physics (including non-LTE effects), and blending---play roughly equal, positively reinforcing roles in explaining the nearly 0.3~dex (factor of $\\sim 2$) decrease of the revised oxygen abundance relative to the value recommended little more than a decade ago (e.g., $850{\\pm}70$ ppm [Grevesse \\&  Anders 1991, hereafter GA91]).  Thus the solar low O problem truly is vexing.  The purpose of the present paper is to add, in our opinion, two key perspectives: (1) the importance of accurately ``calibrating'' the reference photospheric model, be it 3-D or 1-D, against, say, absolute visible continuum intensities and center-limb behavior; and (2) a complementary oxygen abundance (and thermal profile) diagnostic overlooked in the recent work, namely the rovibrational bands of solar carbon monoxide with its fundamental ($\\Delta{v}= 1$) near 4.6~\\micron\\ and first overtone ($\\Delta{v}=2$) near 2.3~\\micron.  The former issue---properly scaling the reference model in temperature---was discussed specifically in a pioneering comprehensive study of the solar carbon, nitrogen, and oxygen abundances by Lambert (1978, hereafter L78), and more generally by Ayres (1977, hereafter A77; 1978a). The solar CO first overtone spectrum, and its significance for the carbon abundance (and thus for $\\epsilon_{\\rm O}$ if C/O is known) was described by Tsuji (1977) and Ayres (1978b), and an extension to the fundamental bands by Ayres \\& Testerman (1981, hereafter AT81).  To preview our conclusions, the solar CO analysis does not support the recently proposed substantial lowering of the oxygen (and carbon) abundances, but instead favors a higher value for oxygen (and carbon), closer to the recent recommendations of GS98, and the earlier work of L78 and GA91. In addition, we derive new accurate abundance ratios for $^{13}$C, $^{17}$O, and $^{18}$O.  The stable isotopes of C and O are signatures of galactic chemical evolution, and should increase over time as the galaxy becomes progressively enriched in the nuclear detritus of successive generations of star formation. The oxygen isotopes, in particular, can trace fractionation processes in the solar nebula associated with the establishment (or not) of dust-gas equilibrium.  It is commonly lamented in the solar system literature, however, that solar astronomy has failed to provide definitive ratios for either $^{16}$O/$^{17}$O or $^{16}$O/$^{18}$O (Wiens et al.\\ 2004). This was one of the major goals, in fact, of NASA's ill-fated {\\em Genesis}\\/ Discovery mission (Burnett et al.\\ 2003). Nevertheless, we demonstrate that it is straightforward to derive accurate oxygen (and carbon) isotopic abundance ratios from CO lines, and further that the per~mil (parts per thousand) differences---e.g., ${\\delta}^{18}$O---are an order of magnitude, or more, larger than anticipated by theoretical models of the gas-dust chemistry in primitive solar system material. ",
        "conclusions": "Through a multifaceted analysis of CO infrared rovibrational bands, we have investigated the thermal structure of the solar photosphere, and as a byproduct derived abundances of C, O, and their main isotopes.  The $\\epsilon_{\\rm O}$, in particular, is similar to values recommended a decade ago, compatible with helioseismological constraints, but is significantly larger than reported in recent studies of OH and \\ion{O}{1} exploiting sophisticated simulations of solar surface convection. We now can understand at least the molecular side of the dichotomy straightforwardly: as was demonstrated in \\S{3.3.1}, the temperature gradient of the mean Asplund model is too steep in the visible continuum forming layers.  Consequently, the model is somewhat too cool in the higher layers where the infrared OH lines arise.  This leads to the inference of a smaller oxygen abundance from the hydroxyl spectrum (and from CO, as well).  On the other hand, the behavior of the \\ion{O}{1} 0.630~\\micron\\ forbidden line---main oxygen diagnostic in the previous work---cannot be explained similarly. The transition arises from the ground state of a majority species, and thus should be relatively immune to thermal effects (see L78). Perhaps the 0.630~\\micron\\ $gf$-value is not as securely established as believed, or the \\ion{Ni}{1} blend is less significant than thought.  One certainly should be wary, in any case, when an important abundance is tied so strongly to essentially a singular spectral feature.  At the same time, one should ask what is the proper role of {\\em ab initio}\\/ models such as that of Asplund and colleagues, versus semiempirical thermal stratifications such as derived by E.\\ H.\\ Avrett and collaborators?  Clearly, an {\\em ab initio}\\/ model is intended to provide direct insight into the fundamental physical processes that govern, say, the structure and dynamics of the solar photosphere.  The Asplund model meets this requirement by accurately reproducing the details of velocity induced lineshape distortions (encoded in the so-called C-shaped line bisectors), and general characteristics of the time dependent brightness pattern of granular convection cells.  Yet, the model is not a perfect representation of the solar atmosphere because it does not predict other key properties, such as the chromospheric temperature rise (or, the temperature {\\em fall}\\/ of the COmosphere), or magnetic-related phenomena such as G-band bright points and \\ion{Ca}{2} network elements.  This, of course, is not the point of the Asplund model, and it should not be faulted for failing tests outside of its intended purview.  Nevertheless, because the model is in some respects not a flawless representation of the solar atmosphere, it should be used with caution in roles outside of its defined domain.  Deriving solar abundances would be a good example of an area where it would be prudent to exercise such caution.  Semiempirical models, on the other hand, provide little in the way of direct physical insight, in and of themselves; but can be used to calculate ancillary quantities (such as chromospheric radiative losses) which are of interpretive value (VAL81).  A semiempirical model, nevertheless, to the extent that it successfully captures mean thermal properties of the atmosphere, could be considered more satisfactory for abundance work.  Deficiencies of 1-D models relative to, say, a fully time dependent spatially resolved convection simulation probably are not as important as accurately mimicking the mean thermal profile, as far as most abundance studies are concerned; and the influence of thermal heterogeneity can be investigated to some extent by imposing 1.5-D thermal perturbations around the basis model, as in the present work. Again, the Asplund et al.\\ study, itself, demonstrated that the thermal inhomogeneities of the deep photosphere have a relatively minor influence on the oxygen abundance, at least as derived from the \\ion{O}{1} spectrum.   Nevertheless, it would be foolish for us to claim a definitive solution to the oxygen crisis. Deriving $\\epsilon_{\\rm O}$ from CO clearly is highly model dependent, and there are enough inconsistencies for the optimum COmosphere model---with regard to the center-limb behavior of weak CO lines like 7--6~R68, and the small but significant discrepancies between abundances derived from $\\Delta{v}=1$ versus $\\Delta{v}=2$ lines---that we view our contrary result mainly as an indication that a low solar oxygen abundance is not fully consistent with all the available diagnostics, and certainly merits further consideration.  One important issue is the C/O ratio.  Increasing it above the assumed 0.5 would allow a decrease of the inferred oxygen abundance, although to achieve the low proposed $\\epsilon_{\\rm O}$ would require C/O\\,$\\gtrsim 1$.  Elevating the Sun to the status of a ``carbon star'' certainly would be more controversial than the low O problem itself.   With regard to the apparent variations of the theoretical rovibrational oscillator strengths with lower state excitation energy, which we have treated as a systematic error and corrected:  the viability of our proposal could be pursued in refined calculations of the dipole matrix elements, and new laboratory measurements of high temperature CO.  However, it also is possible that the consistent behavior of the predicted CO equivalent widths with excitation energy for the Asplund class of models, which---coincidentally or not---yields a low O abundance of $\\sim$500~ppm, is not accidental, but rather is conveying a fundamental message concerning our ability to constrain the thermal properties of the middle photosphere empirically.  The Double Dip model is one answer to that possibility, but its physical reality remains to be tested by independent observations, such as of the \\ion{Ca}{2} wings (AL76), or other potential mid-photosphere thermal tracers.  Finally, while atmospheric thermal profiling and abundance studies might seem less fashionable than searches for exosolar planets or the elusive dark matter (or dark energy, for that matter), the fundamental importance of such basic quantities across diverse areas of astrophysics should not be questioned."
    },
    "0606/astro-ph0606147_arXiv.txt": {
        "abstract": "{} {The classic Weber-Davis model  of the solar wind is reconsidered by incorporating alpha particles and by allowing the solar wind to flow out of the equatorial plane in an axisymmetrical configuration. } { In the ion momentum equations of the solar wind, the ion gyro-frequency is many orders of magnitude higher than any other frequency. This requires that the difference between proton and alpha velocity vectors be aligned with the background magnetic field. With the aid of this alignment condition, the governing equations of the multi-fluid solar wind are derived from the standard transport equations. The governing equations are numerically solved along a prescribed meridional magnetic field line located at colatitude $70^\\circ$ at 1AU and a steady state fast solar wind solution is found. } { A general analysis concludes, in agreement with the Weber-Davis model, that the magnetic field helps the coronal plasma to achieve an effective corotation out to the Alfv\\'enic radius, where the poloidal Alfv\\'enic Mach number $M_T$ equals unity ($M_T$ is defined by equation~(\\ref{eq:mach})). The model computations show that, magnetic stresses predominate the angular momentum loss of the Sun. For the fast wind considered, the proton contribution to the angular momentum loss, which can be larger than the magnetic one, is almost completely canceled by the alpha particles that develop an azimuthal speed in the direction opposite to the solar rotation. The Poynting flux associated with the azimuthal components is negligible in the energy budget. However, the solar rotation can play some role in reducing the relative speed between alpha particles and protons for low latitude fast solar wind streams  in interplanetary space. } {} ",
        "introduction": "The solar angular momentum loss rate ${\\cal L}$ consists of the particle contribution ${\\cal L}_P$ and that contained in magnetic stresses ${\\cal L}_M$. The comparison of measurements of these quantities with models, the Weber-Davis analysis (\\cite{WD:67}) in particular, has yielded divergent results. Missions before Helios measured a total angular momentum flux ${\\cal L}$ consistent with the Weber-Davis model (about $10^{30}$~dyne cm sr$^{-1}$), but the measured azimuthal angle of the bulk flow was generally greater than $1^\\circ$ at 1~AU (or equivalently 7~km~s$^{-1}$ for an average slow wind of 400~km~s$^{-1}$) (see Pizzo et al.~\\cite{Pizzo.etal:83} and references therein). Such a large azimuthal flow speed implies that particles play a far more important role than magnetic stresses in reducing the angular momentum of the Sun. However, in the Weber-Davis model, 3/4 of the angular momentum flux at 1~AU is due to magnetic stresses. The Helios data show that ${\\cal L}$ is $0.2-0.3 \\times 10^{30}$ dyne cm sr$^{-1}$, in which ${\\cal L}_M$ is about $0.15-0.2\\times 10^{30}$ dyne cm sr$^{-1}$ (Pizzo et al.~\\cite{Pizzo.etal:83}). Although the measured magnitude of ${\\cal L}$ is smaller than that computed in the Weber-Davis model, the distribution of angular momentum flux between particles and magnetic stresses is largely compatible with their prediction. An equally important finding concerns further distribution of ${\\cal L}_P$ between two major ion species in the solar wind, namely, protons and alpha particles. Alpha particles are found to carry an angular momentum flux in the direction of counter-rotation with the Sun. This flux is substantial enough to offset the proton contribution which could be comparable to the magnetic one. This finding cannot be addressed by the Weber-Davis model where the solar wind was treated as a bulk flow and only protons were considered.   Apart from being essential in the problem of solar angular momentum loss, the azimuthal ion motions may also provide a possible means to resolve a long  standing  observational puzzle, namely that alpha particles are observed to flow faster than protons in the fast solar wind. The differential streaming in the fast wind could be as pronounced as $150$~km~s$^{-1}$ at 0.3~AU before decreasing to about $40$~km~s$^{-1}$ at 1~AU (Marsch et al.~\\cite{Marsch.etal:82}). Such a behavior has yet to be understood. One possible mechanism is the coupling between the azimuthal and meridional motions, facilitated by the strong magnetic field (McKenzie et al.~\\cite{McKenzie.etal:79}; Hollweg \\& Isenberg~\\cite{HI:81}). Although the Poynting flux may still be negligible (cf. Acuna \\& Whang~(\\cite{AcunaWhang:76}), Alexander \\& {de La Torre}~(\\cite{AT:95}), Hu et al.(\\cite{Hu.etal:03b})), the proposed coupling is expected to limit, at least to a non-trivial extent, the ion differential streaming. As pointed out by Hollweg \\& Isenberg~(\\cite{HI:83}), one shortcoming of the formulation of McKenzie et al.~(\\cite{McKenzie.etal:79}) is that protons are privileged over alphas: The azimuthal magnetic field is assumed to be determined entirely by the protons whose azimuthal flow is neglected. Hence the formulation cannot properly account for the azimuthal dynamics of protons or alphas. In addition, the formulae are applicable only to the equatorial plane where the slow solar wind prevails at solar minimum.   The goal of this paper is to extend the Weber-Davis model by including alpha particles self-consistently. This approach allows us to assess the individual contributions of ion flows and magnetic stresses to the angular momentum loss of the Sun. The effect of the coupling between azimuthal and meridional motions in limiting the proton-alpha differential streaming will also be explored quantitatively. Given that the differential streaming is more prominent in the fast wind, which in general flows out of the equatorial plane, it is necessary to formulate the model such that it treats both protons and alphas on an equal footing, and allows the solar wind to flow  outside the equatorial plane. In this sense, this paper also extends the model of McKenzie et al.~(\\cite{McKenzie.etal:79}).    The paper is organized as follows. The derivation of the governing equations is given in the appendix. Section~\\ref{sec:model} details the physical model and also describes the assumptions on the background poloidal magnetic field and the ion heating mechanism. A general analysis is then given in section~\\ref{sec:analysis}. Section~\\ref{sec:numres} presents the numerical results and the effect of the solar rotation. In section~\\ref{sec:conc}, the main results are summarized. ",
        "conclusions": "\\label{sec:conc}  The main aim of this paper is to extend the Weber-Davis analysis (Weber \\& Davis~\\cite{WD:67}) on the transport of the angular momentum from the Sun by including alpha particles and by allowing the solar wind to flow out of the equatorial plane in an axisymmetrical configuration. Following McKenzie et al.~(\\cite{McKenzie.etal:79}), we exploit the fact that the gyro-frequency of ions is many orders of magnitude higher than any other frequency in ion momentum equations. From this it follows that the difference between proton and alpha velocities must be in the direction of the magnetic field. Using this alignment condition, the governing equations are then derived from the standard five-moment transport equations.  The model equations also enable us to examine quantitatively the effect of azimuthal components in limiting the proton-alpha differential streaming in the fast wind. For simplicity, we choose to solve the governing equations on a prescribed poloidal magnetic field line located at a colatitude of $70^\\circ$ at 1~AU, corresponding to the edge of the fast stream observed by Ulysses at solar minimum conditions (McComas~\\cite{McComas.etal:00}). The effects of the azimuthal components on the meridional dynamics, if any, are optimal in this regard. These effects are directly shown by a comparison of two models with and without azimuthal components.   The main results can be summarized as follows: \\begin{enumerate} \\item The general analysis concludes that, in agreement with the Weber-Davis model, the magnetic field helps the coronal plasma to achieve an effective corotation from the coronal base to the Alfv\\'enic radius, where the poloidal Alfv\\'enic Mach number $M_T=1$ . $M_T$ has to include the contribution from alpha particles (Eq.~(\\ref{eq:mach})).  \\item In the low latitude fast solar wind, the angular momentum loss from the Sun is almost entirely due to magnetic stresses. The proton contribution, which can be as important as the magnetic one in interplanetary space, is offset by alpha particles that develop an azimuthal speed in the direction of counter-rotation with the Sun. \\item The Poynting flux associated with the azimuthal components is negligible. Nevertheless, the solar rotation has an appreciable effect in limiting the proton-alpha differential streaming in fast solar wind streams at low latitudes in interplanetary space. \\end{enumerate}    Although the fast solar wind solution is largely compatible with in situ measurements in terms of the ion mass fluxes and terminal speeds, it fails in a detailed fashion. For instance, the model is not able to predict a proton temperature profile consistent with UVCS measurements in the inner corona, nor does it predict an ion differential speed as large as 150~km~s$^{-1}$ at 0.3~AU to be comparable with  the Helios observation (Marsch et al.~\\cite{Marsch.etal:82}). Hence, including the azimuthal components cannot solely account for the deceleration of alphas relative to protons in interplanetary space. More sophisticated mechanisms, the ion-cyclotron resonance for instance, are expected to alleviate the discrepancies (e.g., Li~\\cite{Li:03}), but can hardly help achieve a satisfactory match (e.g., Hu \\& Habbal~\\cite{HuHabbal:99}). Nevertheless, such a direction is for sure worth pursuing and is left for a future study.  The model also suffers from the inconsistency that the force balance in the direction perpendicular to the poloidal magnetic field is replaced by prescribing a background magnetic field. In a more rigorous treatment, the poloidal magnetic field should be derived self-consistently. In principle, such a task can be accomplished by adopting an iterative approach: the parallel and perpendicular force balance are solved alternately until a convergence is met (Pneuman \\& Kopp~\\cite{PK:71}; Sakurai~\\cite{Sakurai:85}). By doing so, the angular momentum loss from the Sun can be obtained self-consistently for all poloidal flux tubes. An accurate estimate of the duration over which the angular momentum of the Sun is completely removed is then possible (see Hu et al.~\\cite{Hu.etal:03b}).  The present paper is aimed at presenting a rather general analysis of the angular momentum loss from a magnetized rotating object for flows assuming axial symmetry and incorporating two major ion species. Although for the present Sun, the centrifugal and magnetic forces are so weak that they have little impact on the meridional dynamics (especially below the Alfv\\'enic point), a similar study as presented in the text can be carried out for stars that rotate at a faster rate or have a stronger magnetic field than the Sun."
    },
    "0606/astro-ph0606371_arXiv.txt": {
        "abstract": "% We discuss the main observational facts on the eruption of V838 Monocerotis  in terms of possible outburst mechanisms. We conclude that the stellar merger scenario is the only one, which can consistently explain the observations. ",
        "introduction": "The outburst of V838 Monocerotis  was discovered at the beginning of January 2002 (Brown 2002). After about a month it suffered another outburst and reached a luminosity of $\\sim 10^6 L_\\odot$. Initially thought to be a nova, the object soon appeared unusual and enigmatic in its behavior. The eruption, as observed in the optical, lasted about three months. After developing an A-F supergiant spectrum at the optical maximum at the beginning of February 2002, the object showed a general tendency to evolve to lower effective temperatures. In April 2002 it almost disappeared from the optical but remained very bright in infrared, becoming one of the coolest M-type supergiants yet observed. A detailed discussion in Tylenda \\& Soker (2006) shows that novae (or novae-like) outburst models and born-again asymptotic giant branch (AGB) stars cannot account for basic properties of V838~Mon. The merger of two stars, on the main sequence or evolving toward the main sequence, emerged as the most promising model. Our view is that the observations and their analysis (although not completed) that were presented at the first meeting devoted to this object (held in La Palma May 2006), are all in support of the stellar binary merger model.  Stellar mergers have been recognized for a long time as events which can be important for evolution of binary systems. Also in discussions of globular clusters stellar mergers are usually considered as the most probable source of blue stragglers e.g. De Marco et al. (2005, and earlier references therein; Sills, Adams, \\& Davies 2005). However, the main interest in these cases has been directed towards understanding the nature of the final product of a merger in terms of its mass, chemical structure and farther evolution. Little attention has usually been paid to direct observational appearances of these events. As we think that the appearance itself of merger events will attract more attention in the future, we would like to term these type of violent merger events that lead to V838~Mon type events, {\\it mergebursts}.  In this paper written to the proceedings of the meeting, we summarize the main properties of V838~Mon that are relevant to its modelling, mainly as they where presented or cited during the meeting. We then summarize some of our results published elsewhere regarding the merger model. As we summarize the introduction to the final discussion of the meeting as presented by one of us (Soker), we mainly cite talks from the meeting as was done during the original presentation, assuming each paper will present the most updated results, and have a complete list of relevant references to its subject. We also hope that all of these papers will be posted on astro-ph, therefore be freely available. ",
        "conclusions": ""
    },
    "0606/astro-ph0606692_arXiv.txt": {
        "abstract": "Type Ia supernovae (SNe Ia) are the largest thermonuclear explosions in the Universe.  Their light output can be seen across great distances and has led to the discovery that the expansion rate of the Universe is accelerating.  Despite the significance of SNe Ia, there are still a large number of uncertainties in current theoretical models. Computational modeling offers the promise to help answer the outstanding questions.  However, even with today's supercomputers, such calculations are extremely challenging because of the wide range of length and time scales.  In this paper, we discuss several new algorithms for simulations of SNe Ia and demonstrate some of their successes. ",
        "introduction": "The standard theoretical picture of a Type Ia supernova is the thermonuclear explosion of a carbon/oxygen white dwarf that accretes material from a companion star (see \\cite{hillebrandtniemeyer2000} for a good review).  As the mass of the white dwarf approaches the Chandrasekhar limit---the maximum mass that can be supported by electron degeneracy pressure---the temperature and density reach the point where carbon fusion can take place.  When the energy release due to the burning exceeds the local cooling rate (due to expansion and neutrino loss), ignition occurs.  The burning front propagates outward from the center of the star as a thermonuclear flame, wrinkling via the Rayleigh-Taylor instability and interactions with turbulence.  As it continues to burn through the star, the flame accelerates to a large fraction of the sound speed (and possibly transitions into a detonation), producing large amounts of nickel.  The energy release is sufficient to unbind the star.  Large-scale computing has already led to advances in the theoretical understanding of Type Ia supernovae (SNe Ia); however, much more is still unknown. An accurate model of SNe Ia needs to capture physical processes from the scale of the carbon flame, $O(10^{-5}~\\mathrm{cm}) - O(10~\\mathrm{cm})$, to the scale of the star, $O(10^8~\\mathrm{cm})$. The temporal scales are equally impressive, from the century of convection leading up to the ignition of the flame to the seconds-long explosion.  This is truly a multiscale problem, and even at the dawn of petascale computing, a fully resolved simulation of all these processes is beyond the available resources.  Simulations of SNe Ia can be broken down roughly into two types---large eddy simulations operating on the scale of the star (see \\cite{roepkehillebrandt2005,gamezo2005,plewa:2004,garciasenz:2005}), and small-scale simulations that resolve the thermonuclear flame width and model a small, $ < O(100~\\mathrm{cm})$, region of the star (see \\cite{niemeyerhillebrandt1995,roepke2003,SNrt,SNrt3d}). These two types of simulations are complementary in constructing a full picture of the explosion.  The large-scale simulations follow the evolution of the full star down to typical resolutions of 1~km.  On scales smaller than this, they rely on subgrid models to describe the physics of the flame propagation.  The majority of research groups performing full star explosion simulations use the fully compressible piecewise parabolic method \\cite{ppm} for the hydrodynamics. Additionally, they require a way to represent and advance the flame front---typically either thickened flames (see, e.g., \\cite{khokhlov:1995}) or level-sets \\cite{levelset}.  These simulations have successfully showed that a thermonuclear carbon deflagration can release enough energy to unbind the star. However, these pure deflagration models tend to produce weak explosions and leave behind too much unburned carbon.  A transition to detonation at the late stages of the explosion has been proposed \\cite{niemeyerwoosley1997,khokhlov1997}, and may produce more energetic explosions \\cite{gamezo2005}.  Detonations are not without their problems, however \\cite{niemeyer1999}.  In particular, it is unknown if a detonation can develop at all in these environments. This is one of the major outstanding problems in the modeling of SNe Ia.  Another major uncertainty concerns the initial conditions for the explosion. The spatial distribution of the hot spots that seed the flame can have an enormous impact on whether the explosion is successful \\cite{niemeyer:1996,plewa:2004,garciasenz:2005}.  In practice, large-scale calculations begin by randomly initializing one or more hot spots on the grid and propagating the flame outward from there.  In reality, the white dwarf is convecting from the century or more of smoldering burning that precedes the ignition \\cite{Woosley:2004}, and burning fronts can ignite over a finite time interval. With few exceptions \\cite{livne:2005}, the large scale calculations ignore the pre-existing convective velocity field. Long-time three-dimensional simulations of the convective period are required to generate more realistic initial conditions for the ignition and resultant explosion.  Simulations on the scale of the flame can be used to formulate and calibrate the subgrid models employed by the full-star calculations. As the flame propagates outward from the center of the star, it encounters lower density fuel and the laminar flame speed decreases and the flame becomes thicker.  For most of its life, the flame is in the flamelet regime---characterized by a sharp interface between the fuel and the ash.  At densities less than $3\\times 10^7~\\gcc$, the flame becomes broad enough that turbulent eddies on this scale can disrupt the structure of the flame directly, without burning away.  At this point, a mixed region of fuel and ash develops, and the flame is said to be burning in the distributed burning regime. Small-scale simulations of the distributed burning regime can help answer the question of whether a transition to detonation is possible late in the evolution of the explosion.  In this paper, we discuss new algorithms that allow for efficient simulation of both the small scales and the full star.  Our approach exploits the fact that the pre-explosion evolution and flame propagation are subsonic;  only in the very late stages of the explosion does the Mach number approach unity.  Filtering sound waves allows for much longer time evolution than fully compressible codes and forms the basis for a new generation of SNe Ia evolution codes. ",
        "conclusions": "The multiscale nature of Type Ia supernovae makes them challenging to simulate numerically.  As for many multiscale problems, advances in our understanding of the physical phenomena require not just advances in computer hardware, but new algorithms that respect the physical and mathematical nature of the phenomena. This paper outlines several new algorithms for studying SNe Ia based on the low Mach number approach, and several of the successful computations which have resulted.  Future work includes applying these algorithms in the hope of answering questions about the nature of ignition in SNe Ia and whether transition to detonation is possible.  \\ack  This work was supported by DOE grant No.\\ DE-FC02-01ER41176 to the Supernova Science Center/UCSC and Applied Mathematics Program of the DOE Office of Mathematics, Information, and Computational Sciences under contract No.\\ DE-AC03-76SF00098.  The simulations presented were run on Columbia machine at NASA Ames and the jacquard machine at the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC02-05CH11231."
    },
    "0606/astro-ph0606237_arXiv.txt": {
        "abstract": "The tidal interaction between a disk and a planet leads to the planet's migration. It is widely believed that this mechanism explains the variety of orbital radii of extrasolar planets. A long-standing question regarding this mechanism is how to stop the migration before planets plunge into their central stars.  In this paper, we propose a new, simple mechanism to significantly slow down planet migration, and test the possibility by using a hybrid numerical integrator to simulate the disk-planet interaction. The key component of the scenario is the role of low viscosity regions in protostellar disks known as dead zones. A region of low viscosity affects planetary migration in two ways. First of all, it allows a smaller-mass planet to open a gap, and hence switch the faster type I migration (pre-gap-opening migration) to the slower type II migration (post-gap-opening migration). Secondly, a low viscosity slows down type II migration itself, because type II migration is directly proportional to the viscosity.  We present numerical simulations of planetary migration in disks by using a hybrid symplectic integrator-gas dynamics code. Assuming that the disk viscosity parameter inside the dead zone is \\(\\alpha=10^{-4}-10^{-5}\\), we find that, when a low-mass planet (e.g. 1 - 10 Earth masses) migrates from outside the dead zone, its migration is stopped due to the mass accumulation inside the dead zone. When a low-mass planet migrates from inside the dead zone, it opens a gap and slows down its migration. A massive planet like Jupiter, on the other hand, opens a gap and slows down inside the dead zone, independent of its initial orbital radius.  The final orbital radius of a Jupiter mass planet depends on the dead zone's viscosity.  For the range of \\(\\alpha\\)'s noted above, this can vary anywhere from 7 AU, to an orbital radius of 0.1 AU that is characteristic of the hot Jupiters. ",
        "introduction": "Since the first discovery of an extrasolar planet around a solar-type star in 1995 \\citep{Mayor95}, surveys for extrasolar planets have sampled several thousand solar-type stars. Out of these, at least \\(5-25\\%\\) harbor Jupiter mass planets within \\(\\sim 5\\) AU of their central stars \\citep[e.g.][]{Lineweaver03,Butler06}. Most extrasolar planets are of Jupiter mass \\(M_{J}\\), or even larger (up to \\(\\sim 10 M_{J}\\)), and tend to orbit very close to the central star --- sometimes even closer than Mercury is to the Sun \\citep[e.g.][]{Udry03}. These close-in planets, which have orbital radii less than about 0.1 AU, are called ``hot Jupiters\" \\citep[e.g.][]{Mayor95,Marcy96,Santos05}. Arguably the best explanation that we have for these planets is that they were formed further out in the disks, and then migrated to the current locations \\citep[e.g.][]{Goldreich80,Lin93,Ward97,Masset03,Artymowicz04}. This raises several interesting questions about the evolution of planetary systems. As an example, does the current estimate of the percentage of stars harboring planets represent the overall probability of a planetary system to survive, or does it rather suggest that only in \\(5-25\\%\\) of cases planetary systems undergo enough migration to place Jupiter-like planets within \\(5\\) AU? In this latter view, there would be more planets beyond \\(5\\) AU, waiting to be detected in future programs.  The problem with migration as it is currently understood, is that planet-disk interaction is too efficient. Planets rapidly migrate within a gaseous disk and plunge into their central stars in less than a few million years for disk models with standard values of the disk viscosity. This raises the question of whether or not planetary systems typically survive at all. In this work, we will demonstrate that dead zones (low viscosity regions in disks) play a crucial role in saving planetary systems by significantly slowing planetary migration.  Planets migrate as they tidally interact and exchange angular momentum with their disk at Lindblad or corotation resonances. The different kinds of planet migration that arise are generally divided into three types (type I, II, and III). Type I migration occurs for planets that are not massive enough to open a gap in protostellar disks (e.g. Earth mass and 10 Earth mass planets). Such planets most strongly interact with the disk at Lindblad resonances, and hence the effect of corotation torque is negligible \\citep[e.g.][]{Goldreich80}. Usually, planets feel a stronger outer torque than an inner torque, and thus the total torque is negative and the overall migration is inward, relative to the disk. The migration speed is proportional to the planetary mass \\(M_p\\), disk surface mass density \\(\\Sigma\\), and inversely proportional to the square of the disk aspect ratio \\(h/r\\), where \\(h\\) is the pressure scale height and \\(r\\) is the disk radius \\citep[e.g.][]{Goldreich80,Lin93,Ward97}.  Type II migration occurs for planets that are massive enough to open a clear gap in the disks (roughly Jupiter mass and up). Such planets are locked in the gap, and coupled to the disk through the Lindblad resonances which are distant enough from the planets to fall beyond the gap edge. They migrate as the gaseous disk accretes toward the central star due to the disk's viscosity. As long as the disk mass is large compared to the planet mass, the migration speed is proportional to the disk's viscosity parameter \\(\\alpha\\) and the square of the disk's aspect ratio. Type II migration is generally one to two orders of magnitude slower than type I migration \\citep{Ward97}.  Finally, type III migration pertains to planets with an intermediate mass (e.g. Saturn mass planets), which are just heavy enough to start opening a gap. These planets feel the effect of a sharp transition of the surface mass density at the edges of gaps. Since the corotation torque is proportional to the surface mass density gradient \\citep[e.g.][]{Goldreich79,Tanaka02}, its effect dominates the exchange of angular momentum between a planet and a disk. The angular momentum is transferred as the disk material which is trapped in horseshoe orbits passes by the planet. This may lead to a rapid inward or outward migration depending on the initial migration direction of the planet \\citep[e.g.][]{Goldreich80,Ward92,Masset03,Artymowicz04}. The migration speed is comparable to the type I migration \\citep{Masset03}.  For standard disk models, all types of migration are primarily inward, and their timescale is at least one to two orders of magnitude shorter than the disk's lifetime (\\(10^6-10^7\\) years). Therefore, in order to avoid being swallowed by their central stars, the planets' migration has to be significantly slowed down or stopped somehow. Otherwise, giant planets would only have a chance to form once the gas disk is already significantly depleted \\citep{Thommes06}.  A number of scenarios have been proposed to solve this problem. Some scenarios stop planet migration very close to the central stars either by an inner magnetospheric cavity in the disk \\citep[e.g.][]{Shu94,Lin96}, or by the tidal interaction with the star \\citep[e.g.][]{Lin96}. Also, an overflow of the Roche lobe between the central star and a planet could reverse planet migration \\citep{Trilling98}, because a planet loses mass to the central star and therefore moves outward to conserve the angular momentum of the system. However, these mechanisms cannot explain the planets orbiting farther than the immediate vicinity of the star.  Another scenario predicts that the planets stop migration when they move from a non-magnetized region of a disk into a magnetized region \\citep{Terquem03}. This is because the outer torque due to the magnetic resonances disappears in such a case, and because the contribution from the magnetic resonances dominates that from the Lindblad resonances. This mechanism, however, cannot stop type II planet migration.  The torque fluctuations due to the magnetorotational instability (MRI) turbulence which were recently discovered in numerical simulations \\citep[e.g.][]{Laughlin04,Nelson04,Nelson05} may be able to prolong the lifetime of some planets \\citep{Johnson06b}.  This mechanism, however, won't be able to explain the survival of planets heavier than about 10 Earth masses.  We show in this paper that dead zones can stop or slow down planetary migration very efficiently. The underlying assumption here is that the major source of disk's viscosity is the magnetorotational instability (MRI) turbulence \\cite[e.g.][]{Balbus91}. A dead zone is a dense region in a disk which is poorly ionized, and hence MRI inactive and nearly inviscid \\citep[e.g.][]{Gammie96}. Although typical disks are expected to have dead zones \\citep[e.g][hereafter MP03, MP05, and MP06respectively]{Gammie96,Glassgold97,Sano00,Fromang02,Semenov04,Matsumura03,Matsumura05,Matsumura06}, their effect on the evolution of protostellar systems has received little attention. Since the disk's viscosity affects the rate of planetary migration, and since planet formation occurs concurrently with migration, it is important to take account of such variations in disk's viscosity.  In this paper, we study planet migration in disks with dead zones by means of state of the art numerical simulations as well as analytical calculations. We first introduce our disk models and numerical methods, and show the disk evolution with a dead zone in \\S \\ref{chap2}. Then we consider the case of planetary migration beginning from beyond the dead zone, and compare it with the case of disks with no dead zone (\\S \\ref{chap3}). We then present the planet migration from inside the dead zone, focusing on a width of the gap opened by a planet in \\S \\ref{chap4}. The migration in disks with various viscosity parameters is briefly discussed in \\S \\ref{chap5}. Finally, the discussion and conclusion are presented in \\S \\ref{chap7}. ",
        "conclusions": "\\label{chap7} We have studied planet migration in disks with dead zones, and compared the results with the migration in disks with no dead zones. Specifically, we studied the evolution of planets of mass \\(M_{E}\\), \\(10M_E\\), and \\(M_J\\), which are initially placed in orbits both in well coupled, magnetically active, and hence viscous regions as well as inside the dead zone.  Throughout, we have only considered planets on circular orbits. In the case of type I migration, this is a reasonable assumption.  Planets are subject to strong eccentricity damping, on an even much shorter timescale than the orbital decay \\citep[e.g.][]{Artymowicz93b,Papaloizou00}, and therefore core-sized bodies are likely to spend almost all their time with very low eccentricities. Damping becomes weaker once a gap opens (i.e. in type II regime).  In fact, the disk-planet interaction for a gap-opening planet could lead to eccentricity enhancement \\citep[e.g.][]{Artymowicz93b,Goldreich03}, but the issue of gas giant eccentricities is beyond the scope of this work.  For planets migrating toward the dead zone in the active part of the disk, there are some issues that we have not explicitly calculated. First of all, as we have already mentioned, we have not included the effect of the corotation torque.  This is potentially important when planets are migrating toward the density jump generated by a dead zone.  If the corotation torque is stronger than the Lindblad torque, planets would be brought into the dead zone instead of being repelled.  However, we showed that this is unlikely from the results of Appendix B.  For steep density cases relevant to our simulations, Fig. \\ref{fig13} and \\ref{fig14} show that the Lindblad torque changes the sign from negative to positive at \\(0.7\\lesssim r_{\\rm edge}/r_p \\lesssim 0.8\\) (see panels with a density jump \\(F=100\\), and a jump width \\(\\omega\\sim h-2h\\)). Therefore, the net Lindblad torque forces a planet to migrate {\\it outward} at several pressure scale heights away from the dead zone edge. This can be easily seen by noting \\(5\\lesssim (r_p-r_{\\rm edge})/h_{\\rm edge} \\lesssim 8.6\\) with \\(r_{\\rm edge}=r_{DZ}=13 \\ {\\rm AU}\\) and \\(h_{\\rm edge}/r_{\\rm edge}\\sim 0.05\\). Since the corotation torque has a significant effect on a length scale of the density jump width, our results ensure that the Lindblad torque starts acting on a planet {\\it before} the corotation torque becomes effective. Even if the corotation torque becomes comparable to the Lindblad torque, it is unlikely the former exceeds the latter. This is because the magnitude of the net Lindblad torque close to the density jump (positive torque) is much larger than that of the original Lindblad torque far from the jump (negative torque).  Secondly, turbulent fluctuations in active disk may cause corresponding fluctuations in the torque felt by a planet during type I migration \\citep{Laughlin04,Nelson04,Nelson05,Johnson06b}. In principal, if such fluctuations are large enough, this may make it possible for planets/cores to ``jump the barrier\" provided by the edge of a dead zone, instead of being repelled by it. Such behavior can be expected if the density fluctuations, $\\delta \\Sigma/\\Sigma$, become comparable to the density jump between active and dead zone. When this happens, the planets would migrate just like the ones inside the dead zone, i.e. type I migrators open a gap, and slow their migration as seen in \\S 4.  Thirdly, in our simulations a density jump becomes steep enough so that the disk is locally Rayleigh unstable (\\(\\kappa^2<0\\), see Appendix B for the expression).  This will smooth out the density contrast to some degree. Therefore, a density jump in a more realistic disk would not be as steep as our models.  This could affect the migration starting from outside the dead zone --- a planet may not be reflected as strongly as we have shown. We leave an investigation of this issue to future work.  In all three of these instances, we see that although the details of the migration will change, the dead zone will still provide an effective barrier through the onset of type II migration.  We list the conclusions of our paper below: \\\\  (i) Type II migration can be significantly slowed down inside the dead zones (see \\S \\ref{chap3} and \\ref{chap4}). This is because the type II migration speed is directly proportional to the disk's viscosity, and because the dead zones are expected to have a very low viscosity. Jupiter mass planets open a gap inside dead zones and almost stop their migration at around \\(4-8\\) AU when the viscosity parameter is \\(\\alpha=10^{-5}\\). \\\\  (ii) Type I migrators are stopped at the edge of the dead zone (see \\S \\ref{chap3}, as well as Fig. \\ref{fig3} and \\ref{fig4}). This is because type I migration depends on the difference between inner and outer torque strengths.  Due to the mass accumulation at the edge of the dead zone, the magnitude of the inner torque becomes larger as a planet approaches the edge, eventually canceling the outer torque. This mechanism leaves both 1 Earth and 10 Earth mass planets at \\(\\sim 20\\) AU at early times, though they would move inward as the dead zone contracts over time. \\\\  (iii) Type I migration inside the dead zone can be changed to a slower type II migration (see \\S \\ref{chap4}, as well as Fig. \\ref{fig7}). Type I migration itself won't be slowed down by the presence of the dead zones, since the migration speed does not depend on the disk's viscosity. The gap-opening masses, however, become much smaller inside the dead zones, and therefore even small mass planets (e.g. \\(M_E\\) and \\(10M_E\\)) can open a gap there. Although we assume that density waves damp immediately, and hence do not include the wave damping effect proposed by \\cite{Rafikov02}, our simulations show that a 10 Earth mass planet may open a gap at \\(\\sim 4\\) AU, which slows down its migration dramatically. \\\\  (iv) We also showed that the analytically calculated gap-width agrees well with the results of simulations (see \\S \\ref{chap4}, as well as Fig. \\ref{fig9}). \\\\  (v) The dependence of migration on the viscosity parameter is significant. If the viscosity outside the dead zone is measured with \\(\\alpha\\sim 10^{-2}-10^{-3}\\), our simulations indicate that, after \\(10^7\\) years, the orbital radius of planets could vary from our Jupiter-like orbital radius \\(\\sim 3.5-7\\) AU (see Fig. \\ref{fig5}, \\ref{fig10}, and \\ref{fig12}) to hot Jupiter-like radius \\(\\sim 0.1\\) AU (Fig. \\ref{fig11}). \\\\   Our approach so far has ignored the evolution of the dead zone which arises from the accretion of their surface layers onto the central star.  We also ignored the fact that planets accrete a significant fraction of their mass while they migrate. We will study planetary accretion during migration in a disk with an evolving dead zone in our next paper. \\\\  We thank an anonymous referee for useful comments on our manuscript. S. M. is supported by McMaster University, R. E. P. is supported by a grant from the National Sciences and Engineering Research Council of Canada (NSERC), and E. W. T. is supported by NSERC and CITA. \\appendix"
    },
    "0606/hep-ph0606228_arXiv.txt": {
        "abstract": "\\noindent Models of leptogenesis are constrained by the low reheat temperature at the end of reheating associated with the gravitino bound. However a detailed view of reheating, in which the maximum temperature during reheating, $\\Tmax$, can be orders of magnitude higher than the reheat temperature, allows for the production of heavy Majorana neutrinos needed for leptogenesis. But then one must also consider the possibility of enhanced gravitino production in such scenarios. In this article we consider gravitino production during reheating, its dependence on $\\Tmax$, and its relevance for leptogenesis. Earlier analytical studies of the gravitino abundance have only considered gravitino production in the post-reheating radiation dominated era. We find that the gravitino abundance generated during reheating is comparable to that generated after reheating. This lowers the upper bound on the reheat temperature by a factor of 4/3.\\\\ \\vskip 1cm \\noindent {\\bf Key words:} Inflationary cosmology, gravitino abundance, leptogenesis ",
        "introduction": "\\noindent Gravitinos in supersymmetric theories can have important cosmological consequences. Stable gravitinos can overclose the universe while unstable gravitinos can affect the expansion rate of the universe during eras prior to their decay. The decay products of unstable gravitinos can also overclose the universe or affect light element abundances generated during nucleosynthesis. These cosmological consequences are a function of the gravitino energy density, $\\rhoG=\\mG\\nG$, where $\\mG$ and $\\nG$ are the mass and number density of gravitinos.  In a non-inflationary universe, $\\nG\\sim T^3$ and therefore cosmological constraints on the energy density of gravitinos provide bounds on $\\mG$, and equivalently on the scale of supersymmetry breaking~\\cite{ pagelsprimack,weinberg}. In an inflationary universe, $\\nG$ is also a function of the reheat temperature, and so for a fixed $\\mG$, often taken to be $O(100\\gev-1\\tev)$, cosmological constraints on the energy density of gravitinos provide an upper bound on the reheat temperature~\\cite{nos,krauss,khlopovlinde,EKN,falomkinetal,jss,ENS,ks, kletal,moroi.95,cefo}.  In Refs.~\\cite{nos,krauss,khlopovlinde,EKN,falomkinetal,jss,ENS,ks, kletal,moroi.95,cefo} the number density of gravitinos is obtained by considering gravitino production in the radiation dominated era following reheating. It is presumed that $n_{\\tilde G}=0$ at the beginning of the radiation dominated era. Gravitinos are then produced through thermal scattering and $n_{\\tilde G}$ is found to be proportional to the reheat temperature, $\\Treh$, which is the temperature of the thermal plasma at the beginning of the radiation dominated era when the inflaton field has decayed completely and the energy density of the universe is dominated by the inflaton decay products. The cosmological constraints on $\\nG$ then provide an upper bound on $\\Treh$ of $10^{6-9}$\\gev.  If, as in Refs.~\\cite{nos,krauss,khlopovlinde,EKN,falomkinetal,jss,ENS, ks,kletal,moroi.95,cefo}, one assumes instantaneous reheating after inflation or that $\\Treh$ is the maximum temperature during reheating, then the upper bound on $\\Treh$ makes it difficult to create sufficiently high number densities of GUT gauge and Higgs bosons whose decays could generate the baryon asymmetry of the universe. Similarly, the gravitino bound constrains leptogenesis models in which the lepton asymmetry is generated by the decay of heavy bosons or fermions~\\cite{delepinesarkar}. However, as discussed in Refs.~\\cite{ kolb_book,chungetal}, after the inflationary era the temperature does not rise instantaneously to $\\Treh$ but rises initially to a maximum temperature $\\Tmax$ and then falls to $\\Treh$. In Ref.~\\cite{chungetal} the authors then argue that $\\Tmax$ can be as high as $10^3 \\Treh$ and that sufficient numbers of the lightest heavy Majorana right-handed neutrino of mass $\\sim 10\\Treh$ can be produced during reheating to allow for successful leptogenesis. This issue has also been studied in Ref.~\\cite{delepinesarkar,giudiceetal}.  While the above scenario considers the possible production of heavy neutrinos during reheating it does not consider the possible enhancement in the gravitino production as well. If the gravitino abundance generated post-reheating is proportional to the maximum temperature during that era, namely $\\Treh$, one should ask whether the gravitino abundance generated during reheating is proportional to the maximum temperature, $\\Tmax$, of the reheating era.  If this abundance is larger than the abundance generated in the post-reheating era it could affect the viability of the leptogenesis scenario of Ref.~\\cite{chungetal}. Therefore in this article we explicitly calculate the gravitino abundance generated during the reheating era.  We then compare it to the standard calculation of the gravitino abundance generated after reheating in the radiation dominated era, and discuss its relevance for leptogenesis models.  As argued above, one might expect that the gravitino abundance generated during reheating will be a function of $\\Tmax$. Interestingly, what we find is that by manipulating the relations between $\\Tmax$, $\\Treh$, the inflaton decay rate and the scale of inflation the dependence on $\\Tmax$ cancels out and the gravitino abundance is proportional to $\\Treh$ only. Furthermore, while one would not expect the gravitino abundance generated in the reheating and the post-reheating eras to be similar that is indeed what we find. The resulting contraint on the reheat temperature and on leptogenesis models is then only slightly stronger than before.  Our results are valid for a reheating scenario that does not include preheating.  Gravitino production during preheating has been considered in Ref.~\\cite{greene.99,maroto.00,kallosh.00,tsujikawa.00,nilles.01, nilles&olive.01,podolsky}. ",
        "conclusions": "In conclusion, in this article we have calculated the gravitino abundance generated during reheating. We find that it is linearly proportional to the the reheat temperature $\\Treh$, as in the standard calculation of gravitinos produced in the radiation dominated era after reheating. Further, we find that it is about 1/3 the number density of gravitinos generated in the radiation dominated era. Therefore this lowers the upper bound on $\\Treh$ from cosmological constraints on the gravitino number density by a factor of 4/3. This does not significantly alter the viability of leptogenesis scenarios."
    },
    "0606/astro-ph0606551_arXiv.txt": {
        "abstract": "{Determinations of stellar mass and radius with realistic uncertainties at the level of 1\\% provide important constraints on models of stellar structure and evolution.} {We present a high-precision light curve of the A0\\,IV star \\psicenn, from the star tracker on board the \\wire\\ satellite and the Solar Mass Ejection Imager camera on the Coriolis spacecraft. The data show that \\psicen\\ is an eccentric eclipsing binary system with a relatively long orbital period.} {The \\wire\\ light curve extends over 28.7 nights and contains 41\\,334 observations with 2\\,mmag point-to-point scatter. The eclipse depths are 0.28 and 0.16\\,mag, and show that the two eclipsing components of \\psicen\\ have very different radii. As a consequence, the secondary eclipse is total. We find the eccentricity to be $e=0.55$ with an orbital period of 38.8\\,days from combining the \\wire\\ light curve with data taken over two years from the Solar Mass Ejection Imager camera.} {We have fitted the light curve with \\ebop\\ and have assessed the uncertainties of the resulting parameters using Monte Carlo simulations. The fractional radii of the stars and the inclination of the orbit have random errors of only 0.1\\% and 0.01$^\\circ$, respectively, but the systematic uncertainty in these quantities may be somewhat larger. We have used photometric calibrations to estimate the effective temperatures of the components of \\psicen\\ to be $10\\,450\\pm300$ and $8\\,800\\pm300$\\,K indicating masses of about 3.1 and 2.0\\Msun. There is evidence in the \\wire\\ light curve for $g$-mode pulsations in the primary star.} {} ",
        "introduction": "The study of detached eclipsing binaries is of fundamental importance to stellar astronomy as a means by which we can measure accurately the parameters of normal stars from basic observational data (Andersen \\cite{andersen91}). The masses and radii of the component stars of a detached eclipsing binary (dEB) can be measured to accuracies better than 1\\% (e.g.\\ Southworth et al., \\cite{SMS05}), and the effective temperatures and luminosities can be obtained from spectral analysis or the use of photometric calibrations. An important use of these data is in the calibration of theoretical models of stellar evolution (Pols et al.\\ \\cite{pols+97}; Andersen et al.\\ \\cite{andersen1990}; Ribas et al.\\ \\cite{ribas+00}). Comparing the properties of a dEB to the predictions of theoretical models allows the age and metal abundance of the binary to be estimated (e.g.\\ Southworth et al.\\ \\cite{SMSa}) and the accuracy of the predictions to be checked, particularly if the two components of the dEB are of quite different mass or evolutionary stage (e.g.\\ Andersen et al.\\ \\cite{andersenetal1991}).  Some of the physical effects contained in current theoretical stellar evolutionary models are poorly understood, so are treated using simple parameterisations, for example convective core overshooting and convective efficiency (mixing length). Whilst these parameterisations provide a good way of including such physics, our incomplete knowledge of their parameter values is compromising the predictive power of the models (Chaboyer \\cite{chaboyer95}, Young et al.\\ \\cite{young01}) and must be improved. The amount of convective core overshooting has been studied using dEBs by Andersen et al.\\ (\\cite{andersen1990}) and Ribas et al.\\ (\\cite{ribas+00}), and the mixing lengths in the binary AI~Ph\\oe nicis has been determined by Ludwig \\& Salaris (\\cite{ludwig+99}), but these studies have not been able to provide definitive answers because the age and chemical composition are usually free parameters which can be adjusted when comparing the model predictions to the properties of the dEBs.  Progress in the calibration of convective overshooting and mixing length may be made by increasing the accuracy with which the basic parameters of dEBs are measured, which requires much improved observational data. It is arguably more difficult to significantly improve the quality and quantity of light curves, rather than radial velocity curves, because a large amount of telescope time is needed for each system. Forthcoming space missions such as \\kepler\\ (Basri et al.\\ \\cite{kepler05}) and \\corot\\ (Baglin et al.\\ \\cite{baglin01}) will help to solve this problem by obtaining accurate and extensive light curves of a significant number of dEBs.   We have begun a programme to obtain high-quality photometry of bright dEBs with the Wide field InfraRed Explorer (\\wire) satellite, with the intention of measuring the radii of the component stars to high accuracy. Several of the targets in this program have long orbital periods (\\psicen) or shallow eclipses (AR\\,Cas; \\cite{arcas}) which has made them difficult to observe using ground-based telescopes. \\psicen\\ and AR\\,Cas have secondary components which have much smaller masses and radii than the primary stars, so will be able to provide strict constraints on the predictions of theoretical models. An additional advantage of having components with very different radii is that the eclipses are nearly total, which generally allows the radii to be measured with increased accuracy compared to systems with partial eclipses.    In this paper we report our discovery of eclipses in the bright ($V = 4.1$) A0\\,IV star $\\psi$~Cen made with the \\wire\\ satellite. We present a detailed analysis of the light curve showing that the orbit is eccentric and the secondary eclipse is total. As is turns out, this is a particularly interesting system since its brightest component is located in the region of the Hertzsprung-Russell diagram between the blue edge of the Cepheid instability strip and the region of $g$-mode oscillations seen in slowly pulsating B stars (see, e.g., Pamyatnykh \\cite{pamyat1999}). We present an analysis of the oscillations showing evidence of two low-frequency modes that may be interpreted as global oscillation $g$-modes in the primary star. ",
        "conclusions": "We report the discovery that \\psicen\\ is a bright detached eclipsing binary. Photometry from the star tracker on the \\wire\\ satellite shows that the orbit is eccentric and the period longer than about a month. With the photometry from the \\smei\\ camera we have been able to determine the period of the system to be 38.81252(29)\\,d.  We have fitted the \\wire\\ and \\smei\\ light curves using {\\sc ebop} to find precise photometric parameters for \\psicen. Monte Carlo simulations show that the random error in the fractional radii of the stars is about 0.1\\%. We have attempted to assess the systematic uncertainties in this modelling by obtaining solutions with the {\\sc wink} and Wilson-Devinney codes as well as with {\\sc ebop}, but have so far been unsuccessful. Experiments have shown that this is probably due to convergence problems with the differential correction algorithms in the comparison light curve codes. We have been able to find that the systematic uncertainty in the radii are at most 0.3\\%, and a more accurate assessment can be made once the comparison light curve codes have been modified to use a different optimisation algorithm.  Only a very limited number of radial velocity measurements for \\psicen\\ are available in the literature, and are severely affected by line blending, so new high-resolution spectroscopy will be acquired by our group. With this data we will be able to measure the absolute masses and radii of the component stars of \\psicen\\ with high accuracy, which should allow us to place strong constraints on the predictions of theoretical stellar models for masses around 2--3\\Msun. This will be helped if we are able to determine a precise metal abundance from the spectra of \\psicen, but the high rotational velocity of the primary component may make this difficult.  We found two significant peaks in the amplitude spectrum at 2.0 and 5.1 c/day with amplitudes of only 0.2~mmag. These frequencies are too high to be due to rotational modulation. We tentatively interpret them as low amplitude $g$-mode oscillations in the primary star. A preliminary comparison with pulsation models shows that a mode identification should be possible once \\teff\\ and the mass of primary is in hand. Thus, \\psicen\\ will be an interesting object both as a classical detached eclipsing binary system and as an object for a detailed asteroseismic study. It is very interesting that our results indicate that the region of SPB stars as determined from theory (Pamyatnykh \\cite{pamyat1999}) may have to be extended towards the blue edge of the Cepheid instability strip for main sequence stars. For further discussion of this point we refer to \\cite{aerts05}."
    },
    "0606/astro-ph0606598_arXiv.txt": {
        "abstract": "{New images are presented of the supernova remnant (SNR) HB9 based on 408 MHz and 1420 MHz continuum emission and HI-line emission data from the Canadian Galactic Plane Survey (CGPS). Two different T-T plot methods and new integrated flux densities give spectral index (S$_{\\nu}$$\\propto$$\\nu$$^{-\\alpha}$) for the whole of HB9 of 0.48$\\pm$0.03; 0.49; and 0.47$\\pm$0.06, respectively. These values are lower than the previous spectral index estimated for HB9 ($\\alpha$=0.61). The change is  mainly due to improved 1420 MHz data. No difference in spectral index is detected between strong and weak filaments. A new result is that the spectral index for interior regions is steeper than for the rim. This can be explained by a standard curved interstellar electron energy spectrum combined with lower interior magnetic field compared to that near the outer shock. This results in a larger proportion of steep spectrum emission for lines-of-sight through the central body of the SNR. HI observations show structures probably associated with the SNR in the radial velocity range -3 to -9 km$/$s, suggesting a kinematic distance of 0.8$\\pm$0.4 kpc for the SNR. This is consistent with the distance to the radio pulsar PSR B0458+46, offset from the center of HB9 by 23$^{\\prime}$. However the pulsar spindown and kinematic ages are significantly greater than estimates of the SNR age: the Sedov age for HB9 is 6600 yr and the evaporative cloud model yields ages of 4000-7,000 yr. ",
        "introduction": "The production of high-energy particles in our Galaxy is closely related to shock acceleration in supernova remnants (SNR). The study of the radio spectra produced by high-energy electrons in SNRs allows one to learn about the electron energy spectrum. Due to its large diameter HB9 is one of the best candidates for a study of SNRs' spectral index spatial variations. Variations have been observed to occur in some other large angular sized SNRs (Leahy $\\&$ Tian 2005; Tian $\\&$ Leahy 2005, Uyaniker et al. 2004, Alvarez et al. 2001). Spatial variations in the spectral index across HB9 has been previously noted by Leahy et al. (1998) and Leahy $\\&$ Roger (1991), but a new set of higher resolution and sensitivity observations of the Canadian Galactic Plane Survey (CGPS) combined with new spectral index analysis methods, impel us to carry out a new spectral index study of HB9. A lack of direct estimates of the distance to HB9 in previous research (Leahy $\\&$ Roger 1991, Lozinskaya 1981, Milne 1979) is corrected here by an analysis of the new HI observations of HB9. ",
        "conclusions": "We present new higher sensitivity and higher resolution images of the SNR HB9 at 408 MHz and 1420 MHz, and study its radio spectrum, corrected for compact source flux densities using new improved methods. The T-T plot spectral index for HB9 agrees with the integrated flux-density based 408-1420 MHz spectral index. A study of spatial variations shows no systematic difference in spectral index between weak and strong filaments, supporting the conclusion of Leahy $\\&$ Roger (1991). We find steeper spectral index for interior regions than for rim regions. This can be explained by a standard curved interstellar electron energy spectrum combined with variable magnetic field. Due to lower compression, a lower magnetic field in the interior, compared to that near the outer shock, results in steeper spectrum emission from the interior. Thus a larger proportion of steep spectrum emission for lines-of-sight through the central body of the SNR results in an observed steeper spectral index for interior regions. Based on HI features associated with HB9, we obtain a distance of 0.8$\\pm$0.4 kpc, and give updated Sedov and evaporating cloud model parameters for the SNR. We discuss the possible association between the pulsar and HB9, and conclude that more evidence is necessary to decide this possible association."
    },
    "0606/astro-ph0606284_arXiv.txt": {
        "abstract": "We have identified 41 infrared dark clouds from the 8 $\\mu$m maps of the Midcourse Space Experiment (MSX), selected to be found within one square degree areas centered on known ultracompact HII regions. We have mapped these infrared dark clouds in N$_{2}$H$^{+}$ $1 \\rightarrow 0$, CS $2 \\rightarrow 1$ and C$^{18}$O $1 \\rightarrow 0$ emission using the Five College Radio Astronomy Observatory. The maps of the different species often show striking differences in morphologies, indicating differences in evolutionary state and/or the presence of undetected, deeply embedded protostars. We derive an average mass for these clouds using N$_2$H$^+$ column densities of $\\approx 2500$ M$_{\\odot}$, a value comparable to that found in previous studies of high mass star forming cores using other mass tracers. The linewidths of these clouds are typically $\\sim$2.0 - 2.9 km~s$^{-1}$. Based on the fact that they are dark at 8 $\\mu$m, compact, massive, and have large velocity dispersions, we suggest that these clouds may be the precursor sites of intermediate and high mass star formation. ",
        "introduction": "\\label{intro} The study of star formation has made tremendous strides over the past two decades.  Advances in observational capabilities have allowed a number of phases of the star formation process to be identified and characterized, starting with the centrally concentrated core of molecular gas that collapses to form a star surrounded by a proto-planetary disk.  It has been the isolation of such objects that have not yet formed stars -- pre-stellar cores -- that has allowed us to probe the earliest initial stages of star formation (see Andr\\'{e} et al. 2000; Alves, Lada \\& Lada 2000). Most of the progress has focused on the formation of low-mass stars, predominantly because these objects can form isolated from other nearby stars that reduces confusion, and, conveniently, there is a large sample of such objects located in nearby clouds. However, it has now been recognized that most stars that are more massive than the Sun do not form in isolated fashion, but rather in clusters of $>$100 stars (Zinnecker et al. 1993).  Progress in our understanding of high-mass star formation has been hampered by a number of factors: (1) timescales for massive star formation are short and examples in a given state are intrinsically rare, (2) the large distances to the Giant Molecular Clouds (GMCs) that are the birth sites of massive stars make studies of individual objects difficult, and (3) the very nature of clustered star formation increases confusion (cf. Garay \\& Lizano 1999).  As a result, the objects in the earliest phases -- the ``pre-stellar massive cores''  -- have been difficult to identify.  The traditional method of locating low-mass cores has been to examine optical plates for regions of obscured starlight and then to pursue follow up molecular line observations (Myers \\& Benson 1983; Lee \\& Myers 1999).  A comparison with the IRAS point source catalog then denotes whether these cores are associated with newly formed stars (Beichmann et al. 1986).  This method cannot be applied to massive star forming regions since the greater distances makes isolating individual objects difficult.  Moreover, the size and high column densities of GMCs makes it impossible to use optical plates to find individual objects. An analogous method of searching for molecular cores is to search for obscured regions in galactic mid-infrared background. However, due to atmospheric constraints, ground-based observations in the mid-IR are difficult to obtain. The ISOCAM instrument on the Infrared Space Observatory was used in this fashion, but only with pointed observations towards previously identified cores (e.g. Bacmann et al. 2000).  The Midcourse Space Experiment (MSX) surveyed the galactic plane in mid-infrared bands spanning from 7 to 25~$\\mu$m. This survey revealed a large population of dark clouds, predominantly located toward the inner galaxy (Egan et al. 1998).  Follow-up molecular studies of a few objects confirmed that the obscured regions represent a new population of dense, n(H$_2$) $> 10^5$ cm$^{-3}$, and cold, T $< 20$ K, molecular clouds (Carey et al 1998).  Further comparison to IRAS images demonstrated that most of these clouds are dark from 7--100~$\\mu$m, presumably because these objects either do not contain newly formed stars, or any newly formed stars are very deeply embedded.  As such, a sub-sample of these objects may trace massive pre-stellar cores.  There has been substantial activity in this field recently, with numerous groups analyzing various samples of infrared dark clouds (IRDCs). These studies have shown that it is likely that IRDCs are the birth-sites of high-mass stars and stellar clusters (Sridharan et al. 2005, Menten et al. 2005, Rathborne et al. 2006, Pillai et al. 2006).  We have identified a sample of infrared dark clouds and searched them for emission from the N$_2$H$^+$ $1 \\rightarrow 0$, CS $2 \\rightarrow 1$ and C$^{18}$O $1 \\rightarrow 0$ transitions ($\\S$\\ref{obs}). In most cases, we find that the emission closely corresponds to the MSX dark regions.  Using a gas temperature of T=15~K based on CO $1 \\rightarrow 0$ data, we deduce several properties of the dark clouds including column density and mass ($\\S$\\ref{results}). We summarize our findings and the implications ($\\S$\\ref{summary}). ",
        "conclusions": "\\label{summary} We have identified 41 infrared dark (at 8.8 $\\mu$m) clouds that are opaque, compact, and associated with UC HII regions using MSX survey data. In order to determine some basic characteristics of these dark clouds, we have mapped emission from N$_{2}H^{+}$ $1 \\rightarrow 0$, CS $2 \\rightarrow 1$ and C$^{18}$O $1 \\rightarrow 0$ using the FCRAO. The morphology and relative strengths of these molecular lines varies dramatically, possibly indicating evolutionary differences and/or the presence of undetected embedded protostar(s). Based on the derived kinematic distances and the simplifying assumption that the cores are optically thin, we have determined average properties: diameter $<D> \\approx 0.9$ pc, density $<n> \\approx 5000$ cm$^{-3}$, and mass $<M> \\approx 2500$ M$_{\\odot}$.  The low density estimate likely indicates that the dark clouds are clumpy rather than homogeneous. The derived masses, however, are comparable to those derived for a sample of HMPOs. The linewidths are larger than those seen in low-mass star forming cores and larger than in high-mass star forming cores in Orion and Cepheus.  However, they are narrower than the CS linewidths seen in regions that are actively forming massive stars . These observations, taken together, suggest that the infrared dark clouds may be the relatively quiescent pre-cursors to intermediate or massive star formation, the so-called ``pre-protostellar cores.''  At present, we do not know if these sources are truly starless.  Only by obtaining very sensitive infrared observations can we confirm the starless nature of these cores."
    },
    "0606/astro-ph0606417_arXiv.txt": {
        "abstract": "\\noindent{We present short \\chandra\\ observations of twelve bright ($i$$<$$18$) $z_{em}$$\\sim$$1.5$ quasars from the Sloan Digital Sky Survey chosen to have significantly redder optical colors than most quasars at the same redshift. Of the five quasars with optical properties most consistent with dust reddening at the quasar redshift, four show indirect evidence of moderate X-ray absorption (inferred \\nh\\,$\\sim$\\,10$^{22}$\\,\\cmsq) with a dust-to-gas ratio $<$1\\% of the SMC value. The remaining seven objects show no evidence for even moderate intrinsic X-ray absorption. Thus, while optically red quasars are marginally more likely to show signatures of X-ray absorption than optically selected quasars with normal colors, dust-reddened type 1 AGN (as opposed to fully obscured type 2 AGN) are unlikely to contribute significantly to the remaining unresolved hard X-ray background. The red quasar population includes objects with intrinsically red continua as well as objects with dust-reddened continua. Improved sample selection is thus needed to increase our understanding of either subpopulation. To identify dust-reddened quasars likely to exhibit X-ray absorption, some measure of spectral curvature is preferable to simple cuts in observed or relative broad-band colors.} ",
        "introduction": "It has long been suspected that optical surveys for active galactic nuclei are significantly incomplete due to orientation-dependent (and perhaps evolution-dependent) dust obscuration ({Antonucci} 1993).  This suspicion has been confirmed by recent X-ray and infrared investigations that find higher densities of active galactic nuclei (AGN) than do optical surveys to comparable flux limits (e.g., {Ueda} {et~al.} 2003; {Bauer} {et~al.} 2004; {Stern} {et~al.} 2005).  The exact fraction of quasars missed by optical surveys is not yet clear, as the missing objects appear to be mostly lower-luminosity AGN (e.g., {Ueda} {et~al.} 2003; {Hao} {et~al.} 2005).  Nonetheless, it is clear that a full understanding of quasars and their contribution to the X-ray background requires proper consideration of optically obscured quasars.  Indeed, one of the most active topics in X-ray astronomy today is the origin of the hard X-ray background. With the sensitivity of {\\em Chandra} and \\xmm, $\\sim$90\\% of the X-ray background can be resolved into point sources at $<$6~keV, but only $\\sim$60\\% at 6--8~keV (e.g., {Alexander} {et~al.} 2001; {Worsley} {et~al.} 2005). The spectral shape of the $>10$~keV X-ray background ({Marshall} {et~al.} 1980) clearly indicates that a significant fraction of the very hard X-ray background comprises absorbed sources. The demographics of this absorbed population are uncertain at present, and it is unclear how many are bona-fide type 2 AGN with fully obscured broad emission-line and optical continuum regions as opposed to X-ray-absorbed and optically dust-reddened but otherwise normal type 1 (broad line) AGN extincted beyond the flux limits of optical surveys.  Characterizing this population requires not only statistical arguments demonstrating the existence of obscured quasars, but also focused investigations targeting the properties of individual objects that exhibit evidence of obscuration.  Some recent studies have claimed that quasars with red near-IR colors ($J-K>2$) from the Two-Micron All-Sky Survey (2MASS) represent a population of previously undiscovered quasars that could partially account for much of the ``missing'' hard X-ray background sources ({Cutri} {et~al.} 2002; {Wilkes} {et~al.} 2002, 2005).  In this paper we take a complementary approach and investigate the nature of quasars that have {\\em relative} UV/optical colors that are redder than average (at a given redshift).  This selection has the benefit of being less sensitive to color differences due to emission lines moving in and out of the bands and of exhibiting larger differences due to extinction, since the effects of extinction are strongest at UV wavelengths.  We use {\\em Chandra} observations of these quasars to look for additional signs of extinction, such as a harder X-ray spectrum and lower than expected X-ray flux.  We present the construction of our sample in \\S~2, our X-ray observations and data analysis in \\S~3, the results of our observations in \\S~4, a discussion of our results in \\S~5, and our conclusions in \\S~6. We adopt $H_{\\rm 0}=70$~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_m=0.3$ (e.g., {Spergel} {et~al.} 2006). ",
        "conclusions": "We have used short \\chandra\\ observations to demonstrate that optically red quasars are marginally more likely to exhibit evidence for X-ray absorption than optically blue quasars, but that not all optically red quasars exhibit X-ray absorption (\\S\\,\\ref{sec:res} and \\S\\,\\ref{sec:dis}). Four of the five quasars in our sample with optical spectra most consistent with dust reddening at the quasar redshift (\\S\\,\\ref{sec:sum}) show signatures in \\daox\\ and \\GHR\\ of X-ray absorption from neutral or partially ionized gas with column densities \\nh\\,$\\gtrsim 2\\times 10^{22} {\\rm ~cm}^{-2}$ and dust-to-gas ratios $<$1\\% of the SMC value (\\S\\,\\ref{sec:xo} and \\S\\,\\ref{sec:dis}). However, half of our sample shows no evidence for such absorption. We conclude that dust-reddened type 1 AGN (as opposed to fully obscured type 2 AGN) are unlikely to contribute significantly to the remaining unresolved hard X-ray background, at least for the levels of extinction probed by our sample.  Plausible origins have been identified (see the Appendix) for most of the intrinsic dispersion exhibited by quasars in \\aox\\ (the correlation of \\aox\\ with luminosity, plus flux variability in both bands) and possibly in \\GHR\\ (the possible correlation of \\GHR\\ with \\auv, plus X-ray spectral variability).  However, the existence of rare outliers like PHL 1811 underscores our ignorance of how those parameters interact to produce the observed \\aox\\ and \\GHR\\ values.  Progress in understanding the X-ray properties of the type 1 radio-quiet quasar population will require even more careful work to create samples that are as homogeneous as possible (\\S\\,\\ref{sec:jk2}). Such samples must not only be thoroughly vetted to remove quasars exhibiting strong intervening absorption and BAL troughs, but should also span the observed ranges of luminosity, \\ion{C}{4} blueshift, $\\alpha_{UV}$, optical/UV spectral curvature, estimated $M_{BH}$, and possibly emission-line EW, since all those properties may be interrelated with a quasar's X-ray properties. While such studies can make use of serendipitous archival observations of random quasars (e.g., {Gallagher} {et~al.} 2005; {Strateva} {et~al.} 2005; {Risaliti} \\& {Elvis} 2005), the surface density of relatively bright quasars on the sky is low and the the fraction of the sky covered by hard X-ray observations is small. Serendipitous observations do not yield samples that cleanly span the full range of quasar properties, and so there is a continuing need for targeted X-ray observations of carefully selected quasar samples."
    },
    "0606/astro-ph0606621_arXiv.txt": {
        "abstract": "Values of black hole masses are frequently determined with the help of the reverberation method. This method requires a specific geometrical factor related to the distribution of the orbits of the Broad Line Region clouds. Onken et al. determined the value \\f2~$= 1.37 \\pm 0.45$ from the black hole mass - dispersion relation. In this paper we determine this factor using an independent mass determination from the X-ray variance method for a number of Seyfert 1 galaxies and comparing them with the reverberation results by Peterson et al. We obtain mean value \\f2~$= 1.12\\pm 0.54$, consistent with Onken et al.  Both values are larger than the value 0.75 corresponding to a spherical geometry. It indicates that most probably all values of the black hole masses obtained with the use of the Kaspi et al. formulae should be multiplied by a factor of $\\sim$ 1.7. This also shows that the Broad Line Region is rather flat, and hints for a dependence of the factor \\f2 on a source inclination seem to be present in the data. ",
        "introduction": "The value of the black hole mass is one of the key parameters in the description of the accretion process in an active galactic nucleus (AGN). Much effort was devoted to develop reliable and efficient methods of mass determination from the observational data. Current popular methods include the reverberation mapping (Blandford \\& McKee 1982; Wandel, Peterson \\& Malkan 1999; Kaspi et al. 2000; Onken \\& Peterson 2002; Vestergaard 2002), the stellar and gas kinematics (Kormendy \\& Richstone 1995; Nelson \\& Whittle 1995; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Woo \\& Urry 2002; Tremaine et al 2002; Verolme et al. 2002; Peterson et al. 2004; Onken et al. 2004), determination via water maser emission (Miyoshi et al. 1995; Greenhill et al. 2003), the method based on BH mass -- bulge mass or bulge -- luminosity relation (Magorrian et al. 1998; Kormendy \\& Gebhardt 2001; McLure \\& Dunlop 2002; H\\\" aring \\& Rix 2004), based on the disk luminosity estimation (Tripp, Bechtold \\& Green 1994; Collin et al. 2002) or just the $H_{\\alpha}$ line analysis (Green \\& Ho 2005). These methods were recently supplemented by the methods based on X-ray variability: power spectrum break (Papadakis 2004; McHardy et al. 2005) and excess variance (Nikolajuk et al. 2004; O'Neill et al. 2005).  Each of the methods require the knowledge of a certain normalization factor, and is biased by some systematic errors, so the comparison of two independent methods is extremely important.  In the present paper we compare the results of the reverberation approach with the results of the X-ray excess variance for several Seyfert 1 galaxies. The methods of mass determination and the approach to their comparison is given in Section~\\ref{sect:method}, the results for mean normalization factors and the dependence on the source inclination are shown in Section~\\ref{sect:results}, and we discuss the results in Section~\\ref{sect:discussion}. ",
        "conclusions": "\\begin{itemize}  \\item We determine the squared geometrical factor $f^2$ present in the reverberation method using black hole masses determined from X-ray variance method as a reference.  Our value of the geometrical factor, $f^2 =$ \\four, is consistent with the value \\fOnken\\ obtained by Onken et al. (2004) from the black hole mass - stellar dispersion relation, and applied by Peterson et al.  (2004), and in agreement with $f^2 = 0.96$ of Colin et al. (2006).  Such values are intermediate between the expected values for a spherical and a flat geometry of the BLR.   \\item The values of the black hole masses given by the formulae of Kaspi et al. (2000) are most probably systematically underestimated by a factor of $\\sim$ 1.7-1.8.  \\item complementary use of the reverberation and excess variance methods is a promising tool to constrain the inclination of an AGN  \\end{itemize}"
    },
    "0606/astro-ph0606273.txt": {
        "abstract": "It is sometimes suggested that phenomenological power-law plus cool disc-blackbody models represent the simplest, most robust interpretation of the X-ray spectra of bright ultraluminous X-ray sources (ULXs); this has been taken as evidence for the presence of intermediate-mass black holes (BHs) ($M \\sim 10^3 M_{\\odot}$) in those sources. Here, we assess this claim by comparing the cool disc-blackbody model with a range of other models. For example, we show that the same ULX spectra can be fitted equally well by subtracting a disc-blackbody component from a dominant power-law component, thus turning a soft excess into a soft deficit. Then, we propose a more complex physical model, based on a power-law component slightly modified at various energies by smeared emission and absorption lines from highly-ionized, fast-moving gas. We use the {\\it XMM-Newton}/EPIC spectra of two ULXs in Holmberg II and NGC\\,4559 as examples. Our main conclusion is that the presence %%of a soft excess, a soft deficit, or a combination of both, of a soft excess or a soft deficit depends on the energy range over which we choose to fit the ``true'' power-law continuum; those small deviations from the power-law spectrum are well modelled by disc-blackbody components (either in emission or absorption) simply because they are a versatile fitting tool for most kinds of smooth, broad bumps. Hence, we argue that those components should not be taken as evidence for accretion disc emission, nor used to infer BH masses. %a similar argument can be applied to soft-excess AGN. Finally, we speculate that bright ULXs could be in a spectral state similar to (or an extension of) the steep-power-law state of Galactic BH candidates, in which the disc is now completely comptonized and not directly detectable, and the power-law emission may be modified by the surrounding, fast-moving, ionized gas. ",
        "introduction": "Ultra-luminous X-ray sources (ULXs) are point-like, accreting X-ray sources with apparent isotropic luminosities spanning a range from  $\\sim 10^{39}$ to $\\approx 3 \\times 10^{40}$ erg s$^{-1}$, that is, one or two orders of magnitude greater than the Eddington luminosity ($L_{\\rm Edd}$) of a stellar-mass black hole (BH). The main unsolved issue is whether the accreting sources are more massive than typical Galactic BH candidates (BHCs), perhaps in the intermediate-mass range ($M \\sim 10^3 M_{\\odot}$; Miller, Fabian \\& Miller 2004), or stellar-mass BHs accreting at super-Eddington rates (Begelman 2002); alternatively, their brightness could be due to beaming along the line-of-sight of the observer (King et al.~2001; K\\\"{o}rding, Falcke \\& Markoff~2002; Fabrika \\& Mescheryakov 2001). %% Models for achieving super-Eddington luminosities %% include the 'slim' disc model (e.g. Ebisawa et al. 2003), %% radiation-pressure dominated discs (Begelman 2002), mildly %% anisotropic radiation patterns from sources emitting at or %% below the Eddington limit (King et al. 2001) and relativistic %% beaming (K\u00f6rding, Falcke & Markoff 2002).  The standard, most reliable way to determine the mass of an accreting BH in X-ray binaries is based on phase-resolved spectroscopic and photometric studies of their optical counterparts. By measuring the orbital period and the radial velocity shifts of selected optical lines from the donor star and the accretion disc, one can constrain the mass functions of both components. Further constraints to the inclination angle and the size of the system come from the amplitude of the ellipsoidal variations in the donor star, and the presence/absence of dips and eclipses. Those techniques have been successfully applied to a growing number of BHCs (for a review, see McClintock \\& Remillard 2006). Attempts to apply similar techniques to ULXs have been fruitless or inconclusive, so far, mostly because of their optical faintness. At typical distances of a few Mpc (distance moduli $\\sim 28$--$30$), most candidate optical counterparts are fainter than $V \\sim 24$ mag. In many cases, crowding is also a problem: the X-ray error circle may be consistent with an unresolved group of stars. Pioneering efforts (e.g., Gris\\'{e}, Pakull \\& Motch 2006; Pakull, Gris\\'{e} \\& Motch 2006) may yield results for a few sources in the near future. Meanwhile, though, one has to rely on indirect methods to estimate the BH mass.  One such method is based on X-ray spectral fitting over the ``standard'' $0.3$--$10$ keV band. In Galactic BHCs, the X-ray spectrum consists of essentially two components, power-law and thermal, with varying normalizations and relative contributions in various spectral states. The power-law component is scale-free and without a direct dependence on BH mass. However, its slope and normalization are related to the spectral state and normalized luminosity; for instance, the slope is flatter (photon index $\\Gamma \\sim 1.5$--$2$) in the low/hard state ($L_{\\rm X}/L_{\\rm Edd} \\la 10^{-2}$) and steeper ($\\Gamma \\sim 2.5$--$3$) at $0.1 \\la L_{\\rm X}/L_{\\rm Edd} \\la 1$ (McClintock \\& Remillard 2006). More significantly, the thermal component, interpreted as the spectrum of an optically-thick Shakura-Sunyaev disc (Shakura \\& Sunyaev 1973), contains, in principle, a direct dependence on disc size and BH mass.  In the standard thin-disc approximation, for a non-rotating BH, \\begin{equation} \\sigma T_{\\rm eff}^4(R) \\approx \\frac{3 G M \\dot{M}}{8 \\pi R^3} \\end{equation} (Shakura \\& Sunyaev 1973; Pringle 1981; Frank, King \\& Raine 2002)\\footnote{Strictly speaking, in the standard disc-blackbody approximation the disc temperature quickly drops to zero at the inner edge, due to the zero-torque condition at that boundary, after reaching a maximum at $R \\approx (49/36) R_{\\rm in}$. However, equation (1) corresponds to the simplified version of the disc-blackbody spectrum more commonly used in data fitting ({\\tt diskbb} model: Makishima et al.~1986), in which the temperature is assumed to increase all the way to the innermost stable circular orbit, $R_{\\rm in}$. This is why the fitted peak temperature of a disc-blackbody model is commonly referred to as the inner-disc temperature, $T_{\\rm in}$.}. If the disc extends to the innermost stable circular orbit, $R_{\\rm in} = 6GM/c^2$ in the Schwarzschild geometry, the inner-disc temperature scales as \\begin{equation} T_{\\rm in} \\equiv T(R_{\\rm in}) \\propto M^{-1/2} \\dot{M}^{1/4}. \\end{equation} The total luminosity of a standard disc, mostly emitted in the X-ray band for stellar-mass systems, is (Makishima et al 1986) \\begin{equation} L_{\\rm X} \\approx 4\\pi \\sigma T_{\\rm in}^4 R_{\\rm in}^2 \\approx 9.76 \\times 10^{36} \\left(\\frac{T_{\\rm in}}{{\\rm keV}}\\right)^4 \\left(\\frac{M}{M_{\\odot}}\\right)^2  ~~{\\rm erg~s}^{-1}~ \\end{equation} which implies that $L_{\\rm X} \\propto T_{\\rm in}^4$, varying the accretion rate along lines of constant BH mass. %Mitsuda et al. PASJ, 36, 741, (1984), %Makishima, K., Maejima, Y., Mitsuda, K., Bradt, H. V., %Remillard, R. A., Tuohy, I. R., ApJ 308, 635, (1986). At the Eddington limit, assuming a standard radiative efficiency $\\approx 0.1$, \\begin{equation} \\dot{M}_{\\rm Edd} \\approx \\frac{1.3 \\times 10^{39}}{c^2} \\, \\left(\\frac{M}{M_{\\odot}}\\right) ~~{\\rm g~s}^{-1}, \\end{equation} therefore $T_{\\rm in} \\propto M^{-1/4}$ from Eq.~(2), and \\begin{equation} L_{\\rm Edd} \\approx 1.3 \\times 10^{38} \\left(\\frac{M}{M_{\\odot}}\\right) %\\ \\ {\\rm erg~s}^{-1} \\nonumber\\\\ \\approx 4.5 \\times 10^{39} \\left(\\frac{T_{\\rm in}}{{\\rm keV}}\\right)^{-4} ~~{\\rm erg~s}^{-1}.~ \\end{equation}  There are various caveats attached to those simple relations. For example, the innermost stable circular orbit (assumed to mark the inner-disc boundary at high accretion rates) can vary between $6GM/c^2$ and $GM/c^2$ for a non-rotating or maximally-rotating BH, respectively. The fitted temperature (``color temperature'') may be higher than the effective temperature by a factor as high as $2.6$ (Shrader \\& Titarchuk 2003). On the other hand, this is approximately compensated (Fabian, Ross \\& Miller 2004) by the temperature drop near the inner edge, due to the zero-torque condition. The normalization of the observed spectrum, and therefore also the inferred BH mass, depend on the viewing angle, often poorly constrained.  Nevertheless, equations (3) and (5) provide an overall good scaling and order-of magnitude estimate of the BH mass in Galactic BHCs, and the {\\tt diskbb} model in {\\small XSPEC} (Arnaud 1996) has proved simple and successful. More complex implementations of the disc-blackbody model (e.g., {\\tt diskpn}; Gierlinski et al.~1999), taking into account some of the corrective factors mentioned above, have also been used, at the cost of additional free parameters. It seems reasonable to apply the same simple tools to estimate the mass of the accreting BHs in ULXs, if they are scaled-up versions of Galactic BHCs. ",
        "conclusions": "\\subsection{The weakness of phenomenological models}  Phenomenological models can successfully reproduce the observed X-ray spectral features of ULXs, and may be useful for the classification and comparison of different classes of sources, but are not necessarily more constrained, robust or reliable than complex physical models. In fact, they often introduce a dangerous bias, when we try to attribute them unwarranted physical meaning. The disc-blackbody component often used to model deviations from a pure power-law spectrum in ULXs and AGN is a classical example. We argued that the so-called soft excess can just as easily be described as a soft deficit, depending on the energy range to which we choose to fit the power-law continuum. A small change of the fitted power-law slope can turn an apparent absorption feature into an apparent emission feature. Both the excess and the deficit are well modelled with a disc-blackbody component in emission or absorption, respectively. This is simply because a disc-blackbody component is a versatile tool to model a broad, smooth bump or trough (especially if we can adjust both the temperature and the absorbing column density), regardless of its original meaning of an accretion disc spectrum. Similarly, it could be tempting to classify ULXs as stellar-mass objects, at least for those sources (see Stobbart et al.~2006) where a hot-disc model reproduces the observed curvature around $5$ keV better than a simple power-law. Again, we argued that this is misleading: the curvature is likely to be unrelated to a true disc spectrum, and may instead be the signature of increasing, smeared absorption at energies below $5$ keV.  A more physical model to bright ULX spectra is likely to require an injected power-law spectrum modified by the presence of an ionized medium in our line-of-sight, resulting in a combination of blurred, blended emission and absorption lines, and possibly also reflection bumps. Such kind of models would give a more natural explanation for the reason why the X-ray spectra of many bright ULXs and soft-excess AGN look somewhat similar. We have shown one recent implementation of such complex models, which we have developed thanks to the photoionization code {\\small TITAN}, imported into {\\small {XSPEC}}, and applied to two bright ULXs as an illustrative example for this paper. We are aware of the larger number of free parameters included in this kind of models, but we argue that they are less misleading than {\\tt diskbb} models. Our modelling shows that it is possible to produce broad, smooth emission and absorption features when an injected power-law spectrum is seen through a highly-ionized plasma with midly relativistic motion. This may be an ad hoc condition at this stage, but it is probably less problematic than attributing those features to accretion disc emission. We do not speculate at this stage the geometry of the ionized plasma; %(a forthcoming work, based on larger sample of ultra-soft %objets, will help tackle this issue); we merely point out that such intervening medium would produce an effect consistent with what is observed. In principle, the relative contribution of the outward absorption and emission components to the observed spectrum could help us constrain the geometry of the ionized medium. However, this is not yet possible with the available X-ray data, because the true slope of the injected power-law component is not known a priori and cannot be precisely determined over the small energy range of our detectors. Further work on a more extended sample of ULXs (Soria et al., in preparation) will begin to explore this issue, at least on a statistical level. Future observations with instruments such as the Hard X-ray Telescope on {\\it Constellation-X} will be needed to constrain the slope of the primary continuum over a larger energy range, and thus better determine the relative contribution of each components and the physical origin of the power-law itself.   The uncritical use of a disc-blackbody model, interpreted as robust evidence of cool disc emission, has led to claims of BH masses $\\sim 10^3 M_{\\odot}$, skewing both observational and theoretical studies of ULXs towards the IMBH scenario. The ionized-plasma model does not provide in itself evidence in favour or against IMBHs (in fact, it can also be applied to IMBHs and to even bigger BHs, in AGN); however, it implies that the deviations from a power-law spectrum seen in bright ULXs are not related to disc emission, and therefore are not a measure of their BH masses. Without this piece of information, the remaining evidence in favour of IMBHs is much weakened for the majority of ULXs (except for the brightest source in M\\,82). Time-variability studies may provide a more critical test (e.g., Markowitz et al.~2003; Markowitz \\& Edelson 2004; Fiorito \\& Titarchuk 2004; Uttley \\& McHardy~2005; Done \\& Gierlinski 2005), at least for those few sources for which we have enough counts to investigate high-frequency quasi-periodic oscillations and breaks in the power density spectrum. In particular, the temporal behaviour of Holmberg II X-1 was studied by Goad et al. (2006), who found substantial variability on time-scales of months to years, but very little variability on time-scales of less than a day. Based on a combination of energy spectrum and power spectral density considerations, they concluded that Holmberg II X-1 was not likely to be in the disc-dominated high/soft state, and may instead have been in the steep-power-law state (McClintock \\& Remillard 2006); this would be consistent with a BH mass $\\la 100 M_{\\odot}$. For NGC\\,4559 X-1, a break in the power density spectrum was noted and discussed by Cropper et al.~(2004). However, they concluded that the time-variability data available are not yet sufficient to constrain the mass range significantly: BH masses $\\sim 50 M_{\\odot}$ or $\\sim 1000 M_{\\odot}$ could both be consistent with the observed break frequency, if one takes different assumptions on the interpretation of that break and its scaling with BH mass. For most other ULXs, the strongest constraint to their BH mass remains their X-ray luminosity in comparison with the Eddington limit. This argument suggests an upper limit $\\sim 100$--$200 M_{\\odot}$ if the emission is isotropic (Swartz et al.~2004; Gilfanov, Grimm \\& Sunyaev 2004) and even less if beamed. %Done, C., Gierlinski, M. 2005, MNRAS, 364, 208 %Gilfanov, M., Grimm, H.-J., Sunyaev, R. 2004, Nucl Phys B, 132, 369 %Markowitz, A., Edelson, R. 2004, ApJ, 617, 939 %Markowitz, A. 2003, ApJ, 593, 96 %Uttley, P., McHardy, I. M. 2005, MNRAS, 363, 586 This may still be an order of magnitude higher than the mass of Galactic BHs, but may be accommodated with more ordinary star-formation processes.     \\subsection{ULXs as a new spectral state?}  Complex models based on ionized absorption, emission and reflection assume that the underlying X-ray spectrum is a power-law, only slightly modified. This is also consistent with the detection of some bright ULXs with a pure power-law spectrum. How can the disc be not visible at all, not even as a small bump? We have argued that the brightness of those sources strongly disfavours models in which the disc is simply truncated far from the innermost stable orbit. A possible qualitative alternative is that the disc emission is completely comptonized in a non-thermal corona. This requires that most of the available gravitational power be efficiently released in the corona, or efficiently transferred from the disc to the corona---in this scenario, a magnetic disc/corona coupling could perhaps drain most of the energy from the disc at small radii (e.g., Kuncic \\& Bicknell 2004).  A related problem is to determine the low-energy cut-off to the power-law which, we know for certain, does not extend to the optical band. In comptonization models, the low-energy cut-off identifies the characteristic temperature (or at least the lowest temperature) of the seed photons. In the ionized outflow model, the injected power-law component is assumed to extend without a break even below the {\\it Chandra} or {\\it XMM-Newton} energy band. This is also the case if we adopt a phenomenological disc-blackbody plus power-law model. Moreover, it appears to be the case for those bright ULXs that can be fitted with a simple power-law. Taken at face value, this requires the presence of seed photons at energies $\\la 0.1$ keV: we speculate that UV photons from the outer disc and/or the donor star may also contribute as seed for the comptonization process. The fraction of such photons that may illuminate the comptonizing region and be upscattered depends on unknown parameters such as the radial and vertical size of the corona, the thickness and flaring angle of the disk, the spectral type and radius of the donor star, the binary separation.  %Roberts et al. (2005): %The most severe problem with the non-standard spectral description %is the domination of the power-law continuum at low energies. This %spectral form has not (to the best of our knowledge) previously been %observed for any Galactic BHB. Furthermore, it is physically %inconsistent with the standard accretion disc plus Comptonized-corona %model for BHBs, as one would not expect to see the power-law continuum %extend below the peak emissivity of the MCDBB component, due to a %lack of seed photons available to be Compton-scattered. It is very %unlikely that the corona is simply seeing an alternate source of seed %photons; for instance it would require a radically different disc %emission profile for these photons to originate in the cool outer %regions of the accretion disc. It is also very unlikely that the %secondary star could seed this component; assuming equal primary and %secondary masses of ~10 M, with the B0 Ib secondary overflowing its %Roche lobe (cf. Liu et al. 2004), and a corona extending out to 100 %Schwarzschild radii, then the coronal region intercepts less than %one in 1e6 photons emitted by the star. As the power-law component %emits ~1e49 photon s-1 above 0.3 keV (after correcting for %absorption) this requires implausibly high stellar fluxes. % %So can the dominant soft power-law continuum instead be explained by %an alternative physical process? Foschini et al. (2004) discuss two %possibilities for M33 X-8; mildly relativistic jets, and an %outflowing wind. A mildly relativistic jet is certainly a possibility %in M33 X-8 due to the detection of a radio counterpart. However, %Galactic black holes tend to display spectrally hard power-law %continua (\\Gamma ~ 1.5--2) during episodes of radio jet emission, %predominantly when in the low/hard spectral state (though jets can %also be seen from very high state sources, e.g. GRS 1915+105; %Reig, Belloni & van der Klis 2003). Such a hard power-law continuum %is inconsistent with the X-ray spectral forms seen in these ULXs. %The alternative, an outflowing wind, may be likely if ULXs are %powered by accretion from a high-mass star on to a stellar-mass black %hole. For instance, Podsiadlowski, Rappaport & Han (2003) demonstrate %that the rate of matter transfer from the star could potentially be %far higher than that required to reach the Eddington limit in these %systems. Presumably this excess material must either be accreted on %to the black hole in a radiatively inefficient manner, or more likely %discharged from the system in an outflow (see Begelman 2002; King %2002). This scenario may be particularly appropriate for NGC 5204 %X-1, which possesses a B0 Ib star companion (Liu et %al. 2004). However, whilst such a wind could emit a soft X-ray %spectrum (e.g. King & Pounds 2003), it should have a thermal %blackbody spectrum. Hence this also appears to be unable to provide %an origin for the dominant soft power law.  A similar situation (i.e., dominant power-law component extending to energies $\\la 0.3$ keV and comparatively small disc component) appears to occur in the steep-power-law state of Galactic BHCs (McClintock \\& Remillard 2006), and is not well understood in that case, either. Also, in that state, the disc contribution is already small compared to the power-law. Goad et al. (2006) suggested that the temporal variability of Holmberg II X-1 is similar to that found in the Galactic BHC GRS~1915$+$105 in its steep-power-law state ($\\chi$ class). Thus, we speculate that ULXs represent a further spectral state, contiguous to the steep-power-law state, in which the disc contribution is entirely negligible and, in addition, the dominant power-law component is modified by smeared emission and absorption from the surrounding, highly-ionized, possibly outflowing gas. Interestingly, one of the effects of the broad absorption features at $\\sim 1$ keV is to make the continuum appear flatter than the injected power-law, over the $2$--$10$ keV range, as we noted when comparing positive and negative disc-blackbody models (see also fig.~9 in Chevallier et al.~2006). This may be one reason why many bright ULXs in this class appear to have a flatter power-law slope, when fitted with a CD model, than Galactic BHCs in the steep-power-law state (the latter presumably being less affected by highly-ionized, fast outflowing plasma).  Such a spectral state could be shared by higher-mass accretors such as AGN. Narrow Line Seyfert 1s, in particular, display a soft X-ray excess and characteristic variability which could be associated with a steep-power-law state. It has been shown (Chevallier et al.~2006) that the soft excess in AGN could be fitted with the same relativistically smeared ionized plasma model applied here to ULXs. Thus, our approach offers a possible common explanation to the properties of ULXs, soft-excess AGN and Galactic BHs; it suggests that the main spectral features in this bright state depend on the physical parameters of the outflowing plasma, not on the mass of the accretor. A more detailed discussion of this issue is beyond the scope of this work; we shall address it in a forthcoming paper.  %%Perhaps this spectral state can %%only be reached by BHs at least one order of magnitude %%more massive than the Galactic sample, %%with a correspondingly larger disc and Roche lobe %%(linear size $\\propto M$), and %%a correspondingly higher mass accretion rate %%($\\dot{M} \\propto \\dot{M}_{\\rm Edd} \\propto M$). %%If that is the case, differences between %%the steep-power-law spectral states of lower- %%and higher-mass BHs would have to be found %%in physical parameters that are not %%scale-invariant with respect to BH mass: for example, %%the average gas density and ionization parameter %%inside the Roche lobe. %This is precisely what determines the small %deviations, excesses or deficits, from the injected %power-law spectrum, in our favourite scenario. %%A more detailed discussion of these issues %%is beyond the scope of this work. %($\\ga 10^{-6} M_{\\odot}$/yr). %\\vspace{Observational Test} %Time-variability studies of the hard %and soft X-ray band. Even that may not be enough %to tell the difference between disc emission %and blurred line emission from gas in a fast outflow %or just above the disc. Eg, NGC5408 %seems to have a constant soft component as opposed %to a varying po. %disentangle deviations from the po due to 3 factors. % From your model I would expect %that soft and hard band vary together. On the other %hand, thermal disc models would predict two different %behaviours for soft excess and hard band. This is a strong %observational prediction that can be tested!    %%%%%%%%%%%%%"
    },
    "0606/astro-ph0606399_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc Over the past 10 years abundant evidence has emerged that many (if not all) stars are born with circumstellar disks.   Understanding the evolution of post--accretion disks can provide strong constraints on theories of planet formation and evolution. In this review, we focus on developments in understanding: a) the evolution of the gas and dust content of circumstellar disks based on observational surveys, highlighting new results from the Spitzer Space Telescope; b) the physical properties of specific systems as a means to interpret the survey results; c) theoretical models used to explain the observations; d) an evolutionary model of our own solar system for comparison to the observations of debris disks around other stars; and e) how these new results impact our assessment of whether systems like our own are common or rare compared to the ensemble of normal stars in the disk of the Milky Way. \\\\~\\\\~\\\\~% ",
        "introduction": "At the first Protostars and Planets conference in 1978, the existence of circumstellar disks around sun--like stars was in doubt, with most researchers preferring the hypothesis that young stellar objects were surrounded by spherical shells of material unlike the solar nebula thought to give rise to the solar system ({\\it Rydgren et al.}, 1978). By the time of Protostars and Planets II, experts in the field had accepted that young stars were surrounded by circumstellar disks though the evidence was largely circumstantial ({\\it Harvey}, 1985).  At that meeting, Fred Gillett and members of the IRAS team announced details of newly discovered debris disks, initially observed as part of the calibration program ({\\it Aumann et al.}, 1984).   At PPIII, it was well--established that many stars are born with circumstellar accretion disks ({\\it Strom et al.}, 1993) and at PPIV, it was recognized that many of these disks must give rise to planetary systems ({\\it Marcy et al.}, 2000). Over the last 15 years, debris disks have been recognized as playing an important role in helping us understand the formation and evolution of planetary systems ({\\it Backman and Paresce}, 1993; {\\it Lagrange et al.}, 2000; see also {\\it Zuckerman}, 2001). After PPIV, several questions remained.  How do debris disks evolve around sun--like stars?  When do gas--rich disks transition to debris disks?  Can we infer the presence of extra--solar planets from spectral energy distributions (SEDs) and/or resolved disk morphology?  Is there any connection between debris disks and the radial velocity planets?  Is there evidence for differences in disk evolution as a function of stellar mass?  In answering these questions, our objective is no less than to understand the formation and evolution of planetary systems through observations of the gas and dust content of circumstellar material surrounding stars as a function of stellar age.  By observing how disks dissipate from the post--accretion phase through the planet building phase we can hope to constrain theories of planet formation (cf. chapters by {\\it Durisen et al.} and {\\it Lissauer and Stevenson}).  By observing how debris disks generate dust at late times and comparing those observations with physical models of planetary system dynamics, we can infer the diversity of solar system architectures as well as attempt to understand how they evolve with time.  Today, we marvel at the wealth of results from the Spitzer Space Telescope and high contrast images of spectacular individual systems. Detection statistics that were very uncertain with IRAS and ISO sensitivity now can be compared with models of planetary system evolution, placing our solar system in context.  Advances in planetary system dynamical theory, the discovery and characterization of the Kuiper Belt (see chapter by {\\it Chiang et al.}) have proceeded in parallel and further contribute to our understanding of extrasolar planetary systems.  We attempt to compare observations of disks surrounding other stars to our current understanding of solar system evolution. Our ultimate goal is to learn whether or not solar systems like our own are common or rare among stars in the disk of the Milky Way and what implications this might have on the frequency of terrestrial planets that might give rise to life.  Our plan for this contribution is as follows. In Section 2, we describe recent results from observational surveys for gas and dust surrounding normal stars.  Next we describe detailed studies of individual objects in Section 3. In Section 4, we review modeling approaches used in constraining physical properties of disks from the observations.  Section 5 describes a toy model for the evolution of our solar system which we use to compare to the ensemble of observations.  Finally, in Section 6 we attempt to address whether or not planetary systems like our own are common or rare in the Milky Way galaxy and summarize our conclusions. ",
        "conclusions": ""
    },
    "0606/astro-ph0606485_arXiv.txt": {
        "abstract": "We present a quantitative and homogeneous analysis of the broadband multicolor photometric data sets gathered so far on rapidly pulsating hot B subdwarf stars. This concerns seven distinct data sets related to six different stars. Our analysis is carried out within the theoretical framework developed by Randall et al., which includes full nonadiabatic effects. The goal of this analysis is partial mode identification, i.e., the determination of the degree index $l$ of each of the observed pulsation modes. We assume possible values of $l$ from 0 to 5 in our calculations. For each target star, we compute a specific model atmosphere and a specific pulsation model using estimates of the atmospheric parameters coming from time-averaged optical spectroscopy. For every assumed value of $l$, we use a formal $\\chi^2$ approach to model the observed amplitude-wavelength distribution of each mode, and we compute a quality-of-fit $Q$ probability to quantify the derived fit and to discriminate objectively between the various solutions. We find that no completely convincing and unambiguous $l$ identification is possible on the basis of the available data, although partial mode discrimination has been reached for 25 out of the 41 modes studied. A brief statistical study of these results suggests that a majority of the modes must have $l$ values of 0, 1, and 2, but also that modes with $l =4$ could very well be present while modes with $l =3$ appear to be rarer. This is in line with recent results showing that $l = 4$ modes in rapidly pulsating B subdwarfs have a higher visibility in the optical domain than modes with $l = 3$. Although somewhat disappointing in terms of mode discrimination, our results still suggest that the full potential of multicolor photometry for $l$ identification in pulsating subdwarfs is within reach. It will be a matter of gathering higher S/N ratio observations than has been done up to now. ",
        "introduction": "It is well established that an analysis of multicolor photometric data, whereby one compares oscillation amplitudes and/or phase differences in two or more wavebands, can lead to partial mode identification in pulsating stars (see, e.g., Heynderickx, Waelkens, \\& Smeyers 1994). The technique has been used successfully in the past to infer the values of the degree index $l$ of pulsation modes for many types of oscillating stars. To name just a few, let us mention that it has been applied to $\\delta$ Scuti stars (Garrido, Garc\\'\\i a-Lobo, \\& Rodriguez 1990), $\\beta$ Cepheids (Cugier, Dziembowski, \\& Pamyatnykh 1994), ZZ Ceti white dwarfs (Fontaine et al. 1996), $\\gamma$ Doradus stars (Breger et al. 1997), and Slowly Pulsating main sequence B stars (Dupret et al. 2003). Attempts to understand the observed period spectrum on the basis of multicolor photometry have also been made for EC 14026 stars, first by Koen (1998) and more recently by Jeffery et al. (2004; 2005). The former study gives a qualitative interpretation of the periods observed for KPD 2109+4401 based on the theory of Watson (1988), while the latter assert to have provisionally identified or constrained the $l$ values for some of the modes detected for the fast oscillators KPD 2109+4401, HS 0039+4302, and PG 0014+067. In these cases, the $l$ identification is based on a qualitative comparison with amplitude ratios computed by Ramachandran, Jeffery, \\& Townsend (2004) in the adiabatic approximation for representative models.  A theoretical framework for the quantitative exploitation of multicolor photometry for pulsating sdB stars has been put forward recently by some of us (Randall et al. 2005). In that paper, the potential of the technique as applied to both types of pulsating sdB stars (the short-period $p$-mode pulsators of the EC 14026 type and the slowly oscillating $g$-mode variables of the PG 1716 type) has been explored in detail. The method of Randall et al. (2005) features a full nonadiabatic treatment of the atmospheric layers and uses a designated model atmosphere code which automatically incorporates the wavelength dependence of the limb darkening --- thus avoiding the need for approximate parameterized limb darkening coefficients as used in most other multicolor photometric studies with the notable exception of Ramachandran et al. (2004). Of central importance, the method of Randall et al. (2005) has led to the first (and, so far, only) unambiguous determination of the $l$ index of a pulsation mode in a oscillating sdB star using multicolor photometry, thus demonstrating the feasibility of the approach.  Indeed, at the time of the writing of the Randall et al. paper, the best available amplitude estimates of oscillation modes in a pulsating EC 14026 star were those based on the superb $u'g'r'$ data set gathered by Jeffery et al. (2004) on KPD 2109+4401 ($V$ = 13.38) with {\\sl ULTRACAM} mounted on the 4.2-m William Herschel Telescope (WHT). Considering the largest amplitude mode reported with a period of 182.42 s, and adopting at face value the three amplitudes and their quoted uncertainties ($u' = 8.87 \\pm 0.04$ mmag, $g' = 6.58 \\pm 0.04$ mmag, $r' = 6.15 \\pm 0.04$ mmag), Randall et al. (2005) were able to demonstrate that this mode must be a radial mode ($l$ = 0) as first indicated by Jeffery et al. (2004).  The quality of these unique observations has remained unsurpassed so far, but there are nevertheless other available multicolor data sets that certainly deserve to be analyzed in the same fashion now that there is a proven theoretical framework to do that. While it is known that it is generally difficult to discriminate between modes with $l$ = 0, 1, or 2 on the basis of multicolor photometry for EC 14026 stars (Ramachandran et al. 2004; Randall et al. 2005), one may hope to distinguish between modes of that group from modes with $l$ = 3 or $l$ = 4 or possibly $l$ = 5. Our survey of the published work in the field has revealed the following available data sets 1) the $UBVR$ photometry of Koen (1998; herafter K98) on KPD 2109+4401, 2) the $UBV$ photometry of Silvotti et al. (2000; herafter S00) on PG 1628+563B, 3) the {\\sl BUSCA} photometry of Falter et al. (2003; hereafter F03) on PG 1605+072, 4) the $u'g'r'$/{\\sl ULTRACAM} photometry of Jeffery et al. (2004; hereafter J04) on HS 0039+4309, 5) the $u'g'r'$/{\\sl ULTRACAM} photometry of Jeffery et al. (2004) on KPD 2109+4401, 6) the Str\\\"omgren photometry of Oreiro et al. (2005; hereafter O05) on BAL 090100001, 7) the $UBVR$ photometry of Baran et al. (2005; hereafter B05) on BAL 090100001, 8) the $u'g'r'$/{\\sl ULTRACAM} photometry of Jeffery et al. (2005; hereafter J05) on PG 0014+067, and 9) the $u'g'r'$/{\\sl ULTRACAM} photometry of Aerts et al. (2006) on SDSS J171722.08+58055.8. Except for cases 5) and 9) (see below), we analyze all of the other data sets in the present paper.  Partial mode identification (the identification of the degree index $l$) through multicolor photometry is a worthwhile venture in itself, but it has become particularly important for EC 14026 pulsators now that specific asteroseismic models have been proposed for a few of them. Indeed, by combining the forward method with high S/N ratio spectroscopy, it has been possible to carry out complete asteroseismological analyses for four EC 14026 pulsators so far: PG 0014+067 (Brassard et al. 2001), PG 1047+003 (Charpinet et al. 2003), PG 1219+534 (Charpinet et al. 2005a), and Feige 48 (Charpinet et al. 2005b). The ultimate product of these analyses is the determination of the global structural properties of the stars, but complete mode identification (i.e., a determination of the radial order $k$ and the degree index $l$ of each pulsation mode) is a byproduct of the method. Multicolor photometry can thus be used as testing grounds for the proposed seismic models by checking if the $l$ assignments are correct. Of the four EC 14026 pulsators with a proposed specific seismic model, only PG 0014+067 has been observed so far in several filters simultaneoulsy (J05). We will discuss this interesting case below as part of our analysis of that data set.  The main purpose of the paper is to present a quantitative and homogeneous analysis based on the method of Randall et al. (2005) for each of the available multicolor data sets that we found in the published literature on pulsating EC 14026 stars. Considerable observational efforts went into the gathering of these data sets and we feel that it is most worthwhile to attempt extracting the maximum information from them. Partial mode identification is the goal. We note that the exact same approach has been used recently by Aerts et al. (2006) for analyzing their {\\sl ULTRACAM}/WHT data of the very faint EC 14026 pulsator SDSS J171722.08+58055.8 ($B$ $\\simeq$ 16.7), so it is not necessary to reconsider their data here. Unfortunately, in that case, the identification of $l$ was not possible for the two pulsation modes uncovered because the uncertainties on the observed amplitudes turned out to be too large due to insufficient S/N ratio. In contrast, some interesting mode identifications appear possible for the Jeffery et al. (2004) data set on KPD 2109+4401 (beyond that for the 182.42 s mode discussed in Randall et al. 2005), but we defer our analysis to a separate publication that will include unpublished additional observations that we gathered at the Canada-France-Hawaii 3.6-m Telescope on that star. We note also that two-color ($U$ and $R$) photometry has been gathered and analyzed for two PG 1716 sdB pulsators in order to constrain the $l$ index (see the paper of Randall et al. 2006a on PG 1627+017 and that of Randall et al. 2006b on PG 1338+481). ",
        "conclusions": "We have presented in this paper a quantitative and homogeneous analysis of the broadband multicolor photometric data sets gathered so far on rapidly pulsating sdB (EC 14026) stars. This analysis was carried out within the theoretical framework developed by Randall et al. (2005) with the goal of partial mode identification, i.e., the determination of the degree index $l$ of each of the observed pulsation modes. With the exception of the very high S/N ratio $u'g'r'$/{\\sl ULTRACAM} data set obtained by Jeffery et al. (2004) on KPD 2109+4401 (an analysis of which will be presented in a separate publication), we considered all available data that we could find in the literature. This consists of 7 distinct data sets pertaining to 6 different EC 14016 stars and involving 41 pulsation modes.  Our final results are summarized in Table 10 where we list the name of the target star (1st column), the period of the mode of interest (2nd column), the acceptable values of $l$ for that mode (3rd column), and a qualifier as to the level of mode discrimination achieved (4th column). It can be seen that no completely convincing and unambiguous $l$ identification has been possible on the basis of the available data, although we came close to that goal in the case of two modes in KPD 2109+4401 (K98) and three modes in BAL 090100001 (B05) for which only two possible values of $l$ were found to be consistent with the observed amplitude-wavelength distributions. At the same time, we found that no mode discrimination (between values from $l =0$ to $l =5$) was possible for 16 out of 41 modes. There is no doubt that these disappointing (but perhaps not totally unexpected) results are due to the fact that the available data sets were of insufficient sensitivity. The exception is the data set of B05 which features a remarkably high S/N value for the largest amplitude mode (356.19 s) uncovered in that study. We speculate here that nonlinear effects (observed in the form of the first and second harmonic of that mode) may have undermined our ability to model properly the amplitude-wavelength behavior of that particular mode. The latter shows an exceptionally large amplitude by EC 14026 star standards. We recall, in this context, that the feasibility of our approach has been demonstrated beyond any doubt by Randall et al. (2005) in their analysis of the 182.42 s mode measured by Jeffery et al. (2004) in KPD 2109+4401. Hence, the case of the 356.19 s mode in BAL 090100001 remains enigmatic.  It is interesting to examine the statistics of the numbers shown in Table 10. Considering only those cases where partial mode discrimination has been possible, we find that the $l= 0$ solution comes up 19 times, the $l = 1$ solution 23 times, the $l = 2$ solution 24 times, the $l =3$ solution 3 times, the $l =4$ solution 13 times, and the $l =5$ solution once. While no single $l$ index identification has been achieved in our study, one can argue, from a strictly statistical point of view, that these results are certainly consistent with the view that the pulsations seen in rapidly pulsating sdB stars are low-degree modes, mostly with values of $l =0$, 1, and 2. This is not a great surprise by any means, but it confirms the early interpretation of the EC 14026 phenomenon in terms of low-order, low-degree $p$-mode pulsations as presented by Charpinet et al. (1997). Moreover, the higher frequency of $l =4$ solutions compared to that of $l = 3$ solutions is consistent with the finding of Randall et al. (2005) that the former modes are more visible than the latter (in the optical domain).  In conclusion, we find that in order to exploit the full potential of multicolor photometry for EC 14026 pulsators, it will be necessary in the future to gather higher S/N ratio observations than has been done up to now. This is certainly within reach, however, as our results clearly imply. Furthermore, it will be necessary in future asteroseismological exercises such as those carried out by Brassard et al. (2001) and Charpinet et al. (2005c) to include, as required, theoretical modes with $l = 4$ since those have a higher visibility factor in sdB stars than $l = 3$ modes."
    },
    "0606/astro-ph0606216_arXiv.txt": {
        "abstract": "A proper test of Modified Newtonian Dynamics (MOND) in systems of non-trivial geometries depends on modelling subtle differences in several versions of its postulated theories. This is especially important for lensing and dynamics of barely virialised galaxy clusters with typical gravity of scale $\\sim a_0 \\sim 1\\AA{\\rm s}^{-2}$.  The original MOND formula, the classical single field modification of the Poisson equation, and the multi-field general relativistic theory of Bekenstein (TeVeS) all lead to different predictions as we stray from spherical symmetry.  In this paper, we study a class of analytical MONDian models for a system with a semi-Hernquist baryonic profile.  After presenting the analytical distribution function of the baryons in spherical limits, we develop orbits and gravitational lensing of the models in non-spherical geometries. In particular, we can generate a multi-centred baryonic system with a weak lensing signal resembling that of the merging galaxy cluster 1E 0657-56 with a bullet-like light distribution. We finally present analytical scale-free highly non-spherical models to show the subtle differences between the single field classical MOND theory and the multi-field TeVeS theory. ",
        "introduction": "\\protect\\label{sec:intr} While luminous models for the central parts of galaxies do not usually require a dark matter (DM) component, massive halos of DM must be taken into account if one wants to construct  realistic potential-density pairs for an entire galaxy, in order to reproduce the observed nearly-flat galactic rotation curves.  Curiously, there is a considerable body of evidence that the galactic mass profiles of baryonic and dark matter are not uncorrelated (McGaugh 2005). The correlation between the Newtonian gravity of the baryons ${\\bf g}_N$ and the overall gravity ${\\bf g}$ (baryons plus DM) can be loosely parameterized by Milgrom's (1983) empirical relation \\begin{equation} \\protect\\label{eqn:mueq} {\\mu}(g/a_0) \\, {\\bf g} = {\\bf g}_{{\\rm N}}, \\end{equation} where the interpolating function $\\mu(x)$ is a function which runs smoothly from $\\mu(x)=x$ at $x\\ll 1$ to $\\mu(x)=1$ at $x\\gg 1$ with a dividing gravity scale $a_0 \\sim 1 {\\rm \\AA} {\\rm s}^{-2}$ at the transition.  It is a very serious challenge for cold dark matter simulations to reproduce this empirical relation. On the other hand, this relation can be interpreted as a modification of the Newton-Einstein gravitational law in the ultra-weak field regime, with no actual need for dark matter (at the galaxy scale). This provocative idea was taken as the basis for the MOND theory: however, using Eq.(1) alone to transform ${\\bf g}_{{\\rm N}}$ into ${\\bf g}$ does not provide a respectable gravitational force preserving energy and angular momentum.  Bekenstein \\& Milgrom (1984) then suggested modifying the Poisson Equation in order to produce a non-relativistic MOND potential. This proposal was recently refined by Bekenstein (2004) who presented a Lorentz-covariant theory of gravity, dubbed TeVeS, yielding MONDian behaviour in the appropriate limit. However, there is a subtle difference between the non-relativistic formulation of MOND and the relativistic TeVeS, namely the number of fields involved: in MOND, the potential is modified directly in order to trigger the MOND phenomenology, while in TeVeS, a scalar field is added to the traditional Newtonian potential. These two descriptions are equivalent only in highly symmetric systems (spherical or cylindrical symmetry). For disk galaxies, Eq.(\\ref{eqn:mueq}) is a good approximation of the MOND and TeVeS gravitational theories since the additional curl field (see \\S \\ref{sec:montev}) is small when solving the modified Poisson equation for the potential or the scalar field (Brada \\& Milgrom 1995). For this reason, Eq.(\\ref{eqn:mueq}) has been used with confidence to fit the rotation curves of an impressive list of external galaxies with remarkable accuracy (Sanders \\& McGaugh 2002). However, very little work has been carried out to study the actual effect of the curl field in the non-relativistic MOND theory (Brada \\& Milgrom 1995, 1999; Ciotti, Londrillo \\& Nipoti 2006), while the quantitative difference between the predictions of MOND and TeVeS has not been studied at all. The issue has become urgent as recent evidence against MOND is largely based on multi-centred systems such as satellites of the Milky Way (Zhao 2005) and the bullet cluster of merging galaxies (Clowe et al. 2004).  In this paper, after recalling some basic formulae of MOND and TeVeS (\\S\\ref{sec:montev}), we propose a set of parametric interpolating functions that are physical in TeVeS (\\S \\ref{sec:inter}). Then we present spherical analytical potential-density pairs for the baryon distribution in early type galaxies, galactic bulges and dwarf spheroidals, valid in MOND as well as in TeVeS (\\S \\ref{sec:spherpot}). For this model we work out the density, potential, circular velocity, isotropic and anisotropic distribution functions. We also show a simple analytical expression for the gravitational bending angle. Then we present the combination of such models in multi-centred systems for the classical MOND approximation(\\S \\ref{sec:twocenter}); this provides an indication of the kind of result we could expect from a rigorous modelling of the bullet cluster of galaxies (Clowe et al. 2004). More generally, we then show how to extend those spherical models to oblate ones (\\S \\ref{sec:flatmond}). In those non-spherical situations, the density corresponding to a given potential differs in the classical MOND and in TeVeS. To illustrate this,  we study the effects of the flattening of the potential in the special case of scale-free oblate one-dimensional models (\\S \\ref{sec:scalefree}), and we explicitly show the subtle differences between the different theoretical frameworks. Note that Milgrom (1994) has also suggested that MOND could have a modified inertia basis rather than a modified gravity basis; in that case the original MOND formula is correct for circular orbits in galaxies, the theory would become strongly non-local, the conservation laws would become unusual, and the potential-density approach used hereafter would not apply to that framework. ",
        "conclusions": "\\label{sec:conc} In summary, we have presented analytical models of dynamics and lensing in MOND/TeVeS. \\\\ \\\\ 1. We proposed (\\S \\ref{sec:inter}) a useful set of interpolating functions for MOND, with a physical counterpart in TeVeS, contrary to the standard interpolating function commonly used to fit galactic rotation curves. \\\\ \\\\ 2. Using our interpolating functions, we found a useful family of spherical (\\S \\ref{sec:spherpot}) models in MOND/TeVeS, with moderate $1/r$ cusps. Those models were then extended to non-spherical ones, by considering multi-centred models (\\S \\ref{sec:twocenter}), flattened models (\\S \\ref{sec:flatmond})  and scale-free flattened models (\\S \\ref{sec:scalefree}). \\\\ \\\\ 3. We showed that the lensing and the orbits in these spherical or flattened models are rather similar to the expectation in Einstein-Newton gravity, but still trigger a few surprises in extreme geometries. In multi-centred models, the convergence map does not always reflect the projected matter in the lens plane in MOND. This cautions simple interpretations of the analysis of weak lensing in the bullet cluster 1E 0657-56 (Clowe et al. 2004, see Fig. \\ref{fig:bullet}). \\\\ \\\\ 4. In flattened scale-free models (\\S \\ref{sec:scalefree}) we also found that there are differences in the potential-density pairs between the single-field MOND formulation of Bekenstein \\& Milgrom (1984) and the multi-field TeVeS formulation (Bekenstein 2004). However, the differences are probably not sufficient to solve the MONDian mass discrepancy in galaxy clusters (Sanders 2003) as suggested by Bekenstein (2005), unless other parameters of the TeVeS theory, e.g., $\\Xi$ (see Eq. \\ref{eqn:xi}), are important. \\\\ \\\\ Note that our results are not inconsistent with the well-known result that rotation curves of astronomical disks are insensitive to the detailed formulation of MOND. While the very squashed axisymmetric systems of \\S \\ref{sec:scalefree} are not very realistic by themselves, they seem to suggest that the three versions of MOND are likely to show greater differences in systems of complex multi-centred geometry, which is realistic for systems undergoing mergers. \\footnote{For example, satellites in the tidal field of a galaxy have interesting Roche lobe shapes, which contain detailed information on the different laws of gravity (Zhao \\& Tian, 2006).} Altogether we would argue that the most promising systems to test different versions of MOND are systems of lower symmetry for which MOND was not designed and the least is known. We have shown here that application of this test requires highly non-trivial computation to be done {\\it properly}."
    },
    "0606/astro-ph0606746_arXiv.txt": {
        "abstract": "We present BIMA and IRAM CO(1--0) observations of seven low surface brightness (LSB) galaxies, including three large spirals with faint disks but prominent bulges, and four relatively small LSB galaxies with irregular disks. The giant LSB galaxies are UGC~5709, UGC~6614 and F568-6 (Malin2). The smaller LSB galaxies are NGC~5585, UGC~4115, UGC~5209 and F583-1. The galaxies were selected based on their relatively high metallicity and apparent signs of star formation in their disks. The BIMA maps suggested the presence of molecular gas in 2 of the giant LSB galaxies, F568-6 and UGC~6614. Using the 30m IRAM telescope we detected CO (1--0) emission in the disks of both galaxies and in the nucleus of F568-6. The molecular gas in these galaxies is clearly offset from the nucleus and definitely associated with the LSB disk. In addition we also detected a millimeter continuum source in the center of UGC~6614. When compared with VLA 1.5~GHz observations of the galaxy, the emission was found to have a flat spectrum indicating that the millimeter continuum emission is most likely due to an active galactic nucleus (AGN) in the galaxy. Our results show that giant LSB spirals may contain significant quantities of molecular gas in their disks and also harbor radio bright AGN in their centers. ",
        "introduction": "Low surface brightness (LSB) galaxies have faint stellar disks of central brightness less than 23 B~mag~arcsec$^{-2}$ (Impey \\& Bothun 1997). Although they have diffuse stellar disks they are extremely gas rich with large HI gas disks that extend well beyond their apparent stellar disks. The high gas fraction observed in these galaxies combined with the overall low metallicity suggests that they have had much less star formation compared to regular high surface brightness (HSB) galaxies, probably because their gas surface densities are well below the critical value required for star formation (van der Hulst et al. 1993). They are thus considered to be less evolved compared to HSB galaxies. They also have very large mass to light ratios which indicate that their disks are dominated by massive dark matter halos; the massive halo inhibits the formation of disk instabilities such as bars and spiral perturbations, which further contributes to the overall low star formation rate in these galaxies.  LSB galaxies span a wide range of Hubble type and size but can be roughly divided into two categories. The first category is made up of the very populous dwarf spiral and dwarf irregular galaxies. The majority of LSB galaxies in our local universe are of this type. The second category comprises of late type spiral galaxies that are very large in size; some may have prominent bulges and discernable spiral arms. These so called ``LSB giants\" such as Malin~1 and UGC~6614 are relatively rare. The origin and evolution of these faint galaxies is still not clear. Previous studies suggest that LSB galaxies are relatively isolated compared to other types of galaxies and the lack of galaxy interaction has left them less evolved compared to their HSB counterparts. To obtain a better understanding of their evolution, we need to understand their star formation history. An important step in this direction is to determine the gas distribution in their disks.  Although LSB galaxies are rich in HI, molecular gas (H$_{2}$) has been detected in only a handful of such galaxies (O'Neil, Hofner \\& Schinnerer 2000; Matthews \\& Gao 2001; O'Neil, Schinnerer \\& Hofner 2003; O'Neil \\& Schinnerer 2004; Matthews et al. 2005). The low detection rate is probably a combination of several factors such as a low dust content, low metallicity and a low surface density of neutral gas (de Blok \\& van der Hulst 1998) all of which impede gas cooling, molecule formation and cloud condensation. However, modest amounts of ongoing star formation have been observed in most late type galaxies (McGaugh \\& Bothun 1994; McGaugh, Schombert \\& Bothun 1995; Matthews \\& Gallagher 1997) which must mean that star formation can proceed even in these adverse environments. Very little is known about the physics of star formation in such low density and low metallicity environments. One possibility is that star formation occurs in localized regions over the disk. Molecular gas is then patchy in distribution and hence difficult to detect.  To address some of these issues we searched for molecular gas in seven LSB galaxies. Three of the galaxies are LSB giants and the remaining four are smaller late type galaxies. The parameters of the galaxies are listed in Table~1.  All the galaxies in our sample have relatively high metallicities and show signs of star formation in their disks. We followed up tentative BIMA detections of CO (1--0) emission with IRAM 30m single dish observations. CO (1--0) emission was detected in two galaxies, F568-6 (Malin~2) and UGC~6614. The molecular gas masses were derived from the IRAM fluxes; they represent gas masses for the detected regions only and not for the entire galaxy. We also detected millimeter continuum emission from the nucleus of one galaxy in our sample, UGC~6614. The emission was detected at a mean frequency of 111.2~GHz  and is most likely due to an AGN in the galaxy.  In the following sections we present our observations and discuss the implications of our results. For F568-6 (Malin~2) which has a fairly high redshift (z=0.046), we adopt a luminosity distance of $D_{L}=198.8$~Mpc ($H_{o}=71~kms^{-1}Mpc^{-1}$); this leads to a scale of 0.92~kpc arcsec$^{-1}$. UGC~6614 has a much lower redshift and so we adopt the angular distance of $D=89.5$~Mpc ($V_{sys}=6351~\\kms$), which leads to a distance scale of 0.43~kpc arcsec$^{-1}$. ",
        "conclusions": "We have searched for molecular gas in a sample of seven LSB galaxies using the BIMA interferometer. Molecular gas was detected in the disks of two galaxies, UGC~6614 and F568-6, using the IRAM 30m single dish telescope. Both galaxies have prominent bulges and large, low surface brightness disks. Our results indicate that molecular gas may be present in both the disks of LSB galaxies as well as their nuclei but the distribution is localized over isolated regions and thus difficult to detect in unbiased single dish observations. Overall it appears that molecular gas is rare in these galaxies but not completely absent. We have also detected millimeter continuum emission from the nucleus of one of these galaxies, UGC~6614. Our main results are summarised below.  1. CO (1--0) emission was detected in UGC~6614 at approximately 17~kpc (38\\sec) west of the nucleus. It is associated with a spiral feature in the disk. The mass of molecular gas detected from this location is $\\sim2.8\\times10^{8}~M_{\\odot}$. The BIMA observations did not detect any molecular gas in the nuclear region, probably because the gas surface density averaged over the BIMA beam falls below the detection limits of the observation.  2. CO (1--0) emission in F568-6 was detected at two positions in the inner disk. One position is close to the nucleus and the other is approximately 28~kpc (30\\sec) east of the nucleus. The BIMA intensity maps indicate that the CO (1--0) emission is distributed about the nucleus and extends into the inner disk. The mass of molecular gas associated with the detection is $\\sim2.6\\times10^{9}~M_{\\odot}$.  3. We have also detected a millimeter continuum source in UGC~6614 at the center of the galaxy. Comparing it with VLA FIRST 1.5~GHz maps, we find that it has a flat spectrum between ~ 1 to 110~GHz. The continuum emission is due to the AGN in the galaxy which has been detected at optical and X-ray wavelengths as well. Millimeter continuum emission from Seyferts or radio loud galaxies is rare and so UGC~6614 is fairly unique in its AGN activity compared to most spiral galaxies."
    },
    "0606/astro-ph0606570_arXiv.txt": {
        "abstract": "We show that there are strong local correlations between metallicity, surface brightness, and dynamical mass-to-light ratio within M33, analogous to the fundamental line of dwarf galaxies identified by \\citet{pb02}.  Using near-infrared imaging from 2MASS, the published rotation curve of M33, and literature measurements of the metallicities of \\hii\\ regions and supergiant stars, we demonstrate that these correlations hold for points at radial distances between 140~pc and 6.2~kpc from the center of the galaxy.  At a given metallicity or surface brightness, M33 has a mass-to-light ratio approximately four times as large as the Local Group dwarf galaxies; other than this constant offset, we see broad agreement between the M33 and dwarf galaxy data.  We use analytical arguments to show that at least two of the three fundamental line correlations are basic properties of disk galaxies that can be derived from very general assumptions.  We investigate the effect of supernova feedback on the fundamental line with numerical models and conclude that while feedback clearly controls the scatter in the fundamental line, it is not needed to create the fundamental line itself, in agreement with our analytical calculations.  We also compare the M33 data with measurements of a simulated disk galaxy, finding that the simulation reproduces the trends in the data correctly and matches the fundamental line, although the metallicity of the simulated galaxy is too high, and the surface brightness is lower than that of M33. ",
        "introduction": "\\label{introduction}  \\citet[][hereafter PB02]{pb02} recently found that the dwarf galaxies in the Local Group obey a puzzling correlation between mean metallicity and global dynamical mass-to-light ratio (M/L).  By combining this relationship with the previously known trend of metallicity with central surface brightness \\citep[e.g,][]{ped90,c98}, PB02 demonstrated that all three parameters together form an even stronger relationship, which they labeled the fundamental line of dwarf galaxies.  PB02 showed that the metallicity-M/L ratio correlation can be explained by a simple chemical enrichment model. In this model, galactic winds continuously remove the heavy elements formed in dwarf galaxies from their interstellar medium, and star formation in these galaxies ends when all of the gas has been blown out of the system.  Because more metals are retained by more massive galaxies (which have lower mass-to-light ratios) this model is able to reproduce the observed metallicity-M/L relationship.  However, the model does not make clear why including surface brightness as a third parameter improves the correlation, and it is also surprising that variables such as morphological type and star formation history have no apparent effect on the fundamental line.  PB02 proposed that feedback from star formation or supernovae could be involved in regulating the fundamental line, and \\citet{dw03} and \\citet{tassis03} used analytical arguments and simulations, respectively, to show that supernova feedback can produce relationships like the fundamental line, but additional work is needed to improve this understanding.  Two natural questions arise from these previous investigations of the fundamental line.  First, does the fundamental line hold for massive galaxies as well as dwarfs?  Galactic winds and supernova feedback should operate less effectively in larger galaxies as a result of their deeper potential wells, so if either of these processes is the primary mechanism responsible for setting the fundamental line correlations, the fundamental line should not be the same for massive systems.  Second, is the fundamental line a local law as well as a global one?  If not, it would be puzzling that all dwarf galaxies know to follow the relationship globally without having any corresponding correlations on smaller scales.  The data to test the fundamental line locally in dwarfs do not yet exist, because spatially resolved metallicity measurements have not been obtained, and few dSph and dE galaxies have well-resolved kinematics either.  Answering either of these questions with existing data therefore requires us to move to more massive galaxies.  In order to explore further the nature of the fundamental line, we consider the case of a well-resolved disk galaxy with known variations of its metallicity, surface brightness, and mass-to-light ratio with radius.  The best example of a galaxy with these necessary characteristics is M33.  We will address the following questions with this study.  1) Does the fundamental line hold locally, at every point within a galaxy?  2) Do the parameters of the global and local (assuming one exists) fundamental lines agree?  3) What is the origin of the fundamental line?  4) Can simulations of galaxy formation reproduce the fundamental line?  In this paper, we use several recently published data sets to investigate the local fundamental line in M33.  In the next section, we combine a deep Two Micron All Sky Survey (2MASS) mosaic of M33 \\citep{block04}, high resolution CO and \\hi\\ rotation curves \\citep{c03,cs00}, and literature metallicity measurements to construct the necessary observational quantities.  In \\S \\ref{relationships} we perform fits to the various relationships between mass-to-light ratio, metallicity, and surface brightness, and in \\S \\ref{discussion} we compare the locally determined fundamental line of disk galaxies to the fundamental line of dwarf galaxies, give derivations of two of the fundamental line relationships, and discuss the implications of our findings.  We also use the recent $\\Lambda$CDM disk galaxy simulation by \\citet{robertson04} to test how accurately current simulations can reproduce the properties of galaxies and we employ additional toy models of galaxy formation to understand the impact of feedback on the fundamental line.  We briefly summarize our results and conclude in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have shown that there are strong correlations between dynamical mass-to-light ratio and metallicity, dynamical mass-to-light ratio and surface brightness, and metallicity and surface brightness that appear to hold \\emph{locally}, at every point within M33.  The trends of these relationships are that low metallicity, low surface brightness regions have high M/L, and high metallicity, high surface brightness regions have low M/L.  By analogy with the fundamental line of dwarf galaxies identified by PB02, we combine the three correlations into the fundamental line of disk galaxies (Figure \\ref{flplot}), using local properties instead of global ones.  The M33 data follow the same trends as the Local Group dwarf galaxies, but with zero-point offsets, such that at a given metallicity, the mass-to-light ratio in M33 is approximately a factor of $5-10$ higher than that of the dwarfs, and at a given surface brightness M33 has a mass-to-light ratio that is 3 times higher than the dwarf galaxies.  These conclusions are largely independent of the particular mass model chosen for M33.  We demonstrated using analytical calculations that the metallicity-surface brightness relationship in galaxy disks is a generic prediction of ordinary chemical evolution models.  Similarly, the M/L-surface brightness correlation can be derived straightforwardly under the assumptions of an exponential stellar disk and a dark matter halo that follows any of the proposed structures in the literature.  It is not yet clear whether the M/L metallicity relationship shares a similar explanation or is merely a residual of the other two correlations.  These results indicate that at least two of the three relationships that make up the local fundamental line originate from the basic properties of galaxy disks, and should therefore be universally applicable to other galaxies.  Feedback, which was proposed by PB02, \\citet{dw03}, and \\citet{tassis03} as the physical mechanism responsible for the fundamental line of dwarf galaxies, does not seem to play such an important role in the fundamental line of disk galaxies.  Finally, we compared the fundamental line relationships in M33 with their equivalents in a simulated disk galaxy from \\citet{robertson04}. We found good agreement in the slopes of each relationship, although the metallicity of the simulated galaxy is too high.  Additional toy model simulations showed that feedback controls the tightness of the fundamental line, supporting the conclusion of our analytical models that feedback is not necessary to create the local fundamental line. Forcing simulated disks to match the fundamental line and its scatter will improve current numerical recipes for star formation and feedback."
    },
    "0606/astro-ph0606093_arXiv.txt": {
        "abstract": "With central densities way above the density of atomic nuclei, neutron stars contain matter in one of the densest forms found in the universe. Depending of the density reached in the cores of neutron stars, they may contain stable phases of exotic matter found nowhere else in space.  This article gives a brief overview of the phases of ultra-dense matter predicted to exist deep inside neutron stars and discusses the equation of state (\\eosp) associated with such matter. ",
        "introduction": "Neutron stars are dense, neutron-packed remnants of stars that blew apart in supernova explosions. Many neutron stars form radio pulsars, emitting radio waves that appear from the Earth to pulse on and off like a lighthouse beacon as the star rotates at very high speeds. Neutron stars in x-ray binaries accrete material from a companion star and flare to life with a burst of x-rays. Measurements of radio pulsars and neutron stars in x-ray binaries comprise most of the neutron star observations. Improved data on isolated neutron stars (e.g. \\rxj1856, \\psrotwoofive) are now becoming available, and future investigations at gravitational wave observatories focus on neutron stars as major potential sources of gravitational waves (see Ref.\\ \\refcite{villain06:a} for a recent overview). Depending on star mass and rotational frequency, the matter in the core regions of neutron stars may be compressed to densities that are up to an order of magnitude greater than the density of ordinary atomic nuclei. This extreme compression provides a high-pressure environment in which numerous subatomic particle processes are believed to compete with each other\\cite{glen97:book,weber99:book}. The most spectacular ones stretch from the generation of hyperons and baryon resonances ($\\Sigma, \\Lambda, \\Xi, \\Delta$) to quark ($u, d, s$) deconfinement to the formation of boson condensates ($\\pi^-$, $K^-$, H-matter)\\cite{glen97:book,weber99:book,heiselberg00:a,lattimer01:a,weber05:a,sedrakian06:a}. It has also been suggested (strange matter hypothesis) that strange quark matter may be more stable than ordinary atomic nuclei. In the latter event, neutron stars could in fact be made of absolutely stable strange quark matter rather than ordinary hadronic matter\\cite{alcock86:a,alcock88:a,madsen98:b}. Another striking implication of the strange matter hypothesis is the possible existence of a new class of white-dwarfs-like strange stars (strange dwarfs)\\cite{glen94:a}. The quark matter in neutron stars, strange stars, or strange dwarfs ought to be in a color superconducting state\\cite{rajagopal01:a,alford01:a,alford98:a,rapp98+99:a}. This fascinating possibility has renewed tremendous interest in the physics of neutron stars and the physics and astrophysics of (strange) quark matter\\cite{weber05:a,rajagopal01:a,alford01:a}.  This paper reviews the possible phases of ultra-dense nuclear matter expected to exist deep inside neutron stars and its associated equation of state\\cite{glen97:book,weber99:book,heiselberg00:a,lattimer01:a,weber05:a,sedrakian06:a}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606600_arXiv.txt": {
        "abstract": "The Seyfert 1 nucleus of NGC\\,4395 is energized by a black hole of mass $3.6 \\times 10^5 M_\\odot$ (Peterson et al.), making it one of only two nuclear black holes of intermediate mass, $10^3 - 10^6 M_\\odot$, detected in the radio regime.  Building upon UV and X-ray evidence for outflows from this Seyfert nucleus, the VLBI High Sensitivity Array was used at 1.4~GHz to search for extended structure on scales greater than 5~mas (0.1~pc).  Elongated emission was discovered, extending over 15~mas (0.3~pc) and suggesting an outflow on sub-parsec scales from this intermediate-mass black hole.  The Seyfert nucleus is located at the center of an elliptical star cluster, and the elongation position angle of the sub-parsec radio structure is only 19\\arcdeg\\, from the star cluster's minor axis. ",
        "introduction": "Dynamical studies have established that supermassive black holes, with masses $10^6 - 10^9 M_\\odot$, occur in the nuclei of most nearby galaxies with stellar bulges.  The focus has now turned to searches for intermediate-mass black holes (IMBHs), with masses $10^3 - 10^6 M_\\odot$, in nearby galactic nuclei.  Finding IMBHs will help to define the nature of the seed black holes that figure in models for the growth of black holes over cosmic time and to predict the gravitational radiation background expected for {\\em LISA\\/} due to mergers of IMBHs.  In this mass regime, dynamical searches currently fail beyond the Local Group and are being replaced by surveys for signatures of active galactic nuclei \\citep[AGN;][]{gre04}.  Such AGN searches are strongly guided by the discovery of a candidate IMBH in \\object{NGC 4395} \\citep{fil03}, recently established as a bona fide nuclear IMBH with mass (3.6$\\pm$1.1)$\\times 10^5 M_\\odot$ \\citep{pet05}.  NGC\\,4395 is a bulgeless Sdm galaxy at a Cepheid distance of 4.3$\\pm$0.3~Mpc with a scale 0.021~pc~mas$^{-1}$ \\citep{thi04}.  Its nucleus has the emission properties of a Seyfert 1, with broad permitted optical and UV emission lines \\citep{fil93,fil03,pet05} and a pointlike, hard X-ray source that is highly time-variable \\citep{shi03,mor05,vau05}.  The reverberation study by \\citet{pet05} reports a characteristic diameter for the UV broad-line-region (BLR) of 7.0$\\times 10^{-5}$~pc or 0.0033~mas.  There are hints of an outflowing UV absorber close to the nucleus of this AGN \\citep{cre04}. Also, in luminous AGNs, warm X-ray absorbers close to the nucleus are often associated with outflows; in NGC\\,4395 such absorbers have been inferred from variations in their ionization \\citep{shi03} and/or column densities \\citep{mor05}.  While these X-ray and UV hints of outflows are intriguing, seeking structural evidence for an outflow is an essential next step.  NGC\\,4395 is one of only two nuclear IMBHs that have been detected in the radio regime \\citep{gre06}, important because VLBI on mas scales can then be used to search for structural evidence for an outflow on sub-parsec scales.  At 1.4~GHz, the Very Long Baseline Array \\citep[VLBA;][]{nap94} detected one source, with flux density 0.53$\\pm$0.13~mJy, spectral power $1.2 \\times 10^{18}$~W~Hz$^{-1}$, diameter less than 11~mas (0.2~pc), and brightness temperature more than 2 million K \\citep{wro01}.  While these traits are consistent with the VLBA detection being accretion-powered by the IMBH, the resolution and sensitivity of the VLBA image failed to provide quantitative structural information.  This Letter reports improved imaging of the Seyfert nucleus of NGC\\,4395 at 1.4~GHz, using a VLBI High Sensitivity Array (HSA) consisting of the VLBA, the Very Large Array \\citep[VLA;][]{tho80} operating in its phased mode, and the Robert C. Byrd Green Bank Telescope \\citep[GBT;][]{jew04}.\\footnote{The VLBA, the VLA, and the GBT are operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.}  The HSA imaging is described in Section 2.  The implications for the Seyfert nucleus are explored in Section 3, with an emphasis on the discovery of elongated structure on sub-parsec scales. ",
        "conclusions": ""
    },
    "0606/hep-th0606223_arXiv.txt": {
        "abstract": "We present a model for quintessential inflation using a string modulus for the inflaton - quintessence field. The scalar potential of our model is based on generic non-perturbative potentials arising in flux compactifications. We assume an enhanced symmetry point (ESP), which fixes the initial conditions for slow-roll inflation. When crossing the ESP the modulus becomes temporarily trapped, which leads to a brief stage of trapped inflation. This is followed by enough slow roll inflation to solve the flatness and horizon problems. After inflation, the field rolls down the potential and eventually freezes to a certain value because of cosmological friction. The latter is due to the thermal bath of the hot big bang, which is produced by the decay of a curvaton field. The modulus remains frozen until the present, when it becomes quintessence. ",
        "introduction": "Introduction} An array of recent observations has ascertained that we live in a Universe which is engaging into accelerated expansion at present \\cite{all}. The simplest explanation for this observation is to assume a non-zero cosmological constant, corresponding to vacuum energy density comparable to the density of matter today. However, such a value for the cosmological constant is extremely fine-tuned compared to theoretical expectations \\cite{L}. As a result, a number of alternatives has been suggested in the literature. Many of the alternative solutions postulate the existence of an unknown exotic substance, called Dark Energy, whose properties (e.g. equation of state) are such that it would drive the Universe to accelerated expansion if it dominates the Universe content (for a recent review see \\cite{ed}). A particular type of such substance, which fulfils the requirements of dark energy is a potentially dominated scalar field. Under this hypothesis, the cosmological constant may be set to zero, as once assumed, so that the extreme fine-tuning of its value can be avoided. This means that the scalar field either lies at a metastable minimum of the scalar potential or is rolling down a slope in the scalar potential leading to the minimum, which corresponds to zero density. This latter case has been envisaged some time ago in models of `dynamical cosmological constant' \\cite{early}. With respect to accounting for dark energy, the scalar field has been named `quintessence' because it is the fifth element after baryons, photons, CDM and neutrinos \\cite{Q}.  The idea of using a rolling scalar field in order to achieve a phase of accelerated expansion in the Universe history is, of course, not new. In fact, it is the basis of the inflationary paradigm, where the scalar field is the so-called `inflaton' \\cite{Guth:1980zm}. Hence, quintessence corresponds to nothing more than a late time inflationary period, taking place at present. In this respect, the credibility of the quintessence idea has been enhanced by the fact that the generic predictions of the inflationary paradigm in the Early Universe are very much in agreement with the observations.  Since they are based on the same idea, it was natural to attempt to unify early Universe inflation with quintessence. Quintessential inflation was thus born \\cite{quinf,QI,jose,eta}. The advantages of such unified models are many. Firstly, quintessential inflation models allow the treatment of both inflation and quintessence within a single theoretical framework, with hopefully fewer and more tightly constrained parameters. Another practical gain is the fact that quintessential inflation dispenses with a tuning problem of quintessence models; that of the initial conditions for the quintessence field. It is true that attractor - tracker quintessence alleviates this tuning but the problem never goes away completely. In contrast, in quintessential inflation the initial conditions for the late time accelerated expansion are fixed in a deterministic manner at the end of inflation. Finally, a further advantage of unified models for inflation and quintessence is the economy of avoiding to introduce yet again another unobserved scalar field to explain the late accelerated expansion.  For quintessential inflation to work one needs a scalar field with a runaway potential, such that the minimum has not been reached until today and, therefore, there is residual potential density, which, if dominant, can cause the observed accelerated expansion. String moduli fields appear suitable for this role because they are typically characterised by such runaway potentials. The problem with such fields, however, is how to stabilise them temporarily, in order to use them as inflatons in the early Universe.  In this paper (see also Ref.~\\cite{BuenoSanchez:2006eq}) we achieve this by considering that, during its early evolution our modulus crosses an enhanced symmetry point (ESP) in field space. As shown in Ref.~\\cite{Kofman:2004yc}, when this occurs the modulus may be trapped at the ESP for some time. This can lead to a period of inflation, typically comprised by many sub-periods of different types of inflation such as trapped, eternal, old, slow-roll, fast-roll and so on. After the end of inflation the modulus picks up speed again in field space resulting into a period of kinetic density domination, called `kination' \\cite{kination}. Kination is terminated when the thermal bath of the hot big bang (HBB) takes over. During the HBB, due to cosmological friction \\cite{cosmofric}, the modulus freezes asymptotically at some large value and remains there until the present, when its potential density becomes dominant and drives the late-time accelerated expansion \\cite{eta}.  Is is evident that, in order for the scalar field to play the role of quintessence, it should not decay after the end of inflation. Reheating, therefore should be achieved by other means. In this paper we assume that the thermal bath of the HBB is due to the decay of some curvaton field \\cite{curv} as suggested in Refs.~\\cite{curvreh}. Note that the curvaton can be a realistic field, already present in simple extensions of the standard model (for example it can be a right-handed sneutrino \\cite{sneu}, a flat direction of the MSSM and its extensions \\cite{mssm,nmssm} or a pseudo Nambu-Goldstone boson \\cite{pngb,orth} possibly associated with the Peccei-Quinn symmetry \\cite{PQ} etc.).  Thus, by considering a curvaton we do not necessarily add an ad~hoc degree of freedom. The importance of the curvaton lies also in the fact that the energy scale of inflation can be much lower than the grand unified scale \\cite{liber}. In fact, in certain curvaton models, the Hubble scale during inflation can be as low as the electroweak scale \\cite{orth,low}.  Throughout our paper we use natural units, where \\mbox{$c=\\hbar=1$} and Newton's gravitational constant is \\mbox{$8\\pi G=m_P^{-2}$}, with \\mbox{$m_P=2.4\\times 10^{18}$GeV} being the reduced Planck mass. ",
        "conclusions": "\\label{sec:SaC}   We have investigated in detail a realisation of quintessential inflation within the framework of string theory. Our inflaton - quintessence field is a string modulus with a characteristic runaway scalar potential. In our treatment we have avoided to specify in full detail our string inspired theoretical basis in order for our considerations to remain as generic as possible. In that sense, our quintessential inflation realisation is more like a paradigm than a specific model.  In this spirit we considered a broadly accepted form of the non-perturbative superpotential (which can be due to, for example, gaugino condensation or $D$-brane instantons) and of the tree-level K\\\"{a}hler potential shown in Eqs.~(\\ref{W}) and (\\ref{eq:treeKahler}) respectively. The resulting non-perturbative potential appears in Eq.~(\\ref{scalarpot0}) and is of the form \\mbox{$V_{\\rm np}(\\sigma)\\propto \\sigma^{-\\ell}e^{-\\mu c\\sigma}$}, where \\mbox{$\\sigma=$ Re $T$} and $\\ell,\\mu=1,2$ depending which term dominates. Due to Eq.~(\\ref{eq:treeKahler}) the canonically normalised modulus $\\phi$ is in fact related to $\\sigma$ as \\mbox{$\\sigma(\\phi)=e^{\\sqrt{\\frac{2}{3}}\\phi/m_P}$} [c.f. Eq.~(\\ref{canonfield})]. Hence, the potential $V_{\\rm np}(\\phi)$ is, in fact, of the form of a double exponential.  Such a potential is too steep to support inflation. This is why, we have assumed the existence of an enhanced symmetry point (ESP) at some value $\\phi_0$, corresponding to potential density \\mbox{$V_0=3H_*^2m_P^2$}. The most salient effect of the ESP is that it can trap the rolling modulus and effectively stop it from rolling. Indeed, as the modulus crosses the ESP, particle production can generate a contribution $V_{\\rm int}(\\phi)$ to the scalar potential able to stop the roll of the modulus and trap the latter at the ESP. However, apart from the phenomenon of particle production described in \\cite{Kofman:2004yc}, in this paper we have also taken into account that the ESP generates a `flat patch' on the scalar potential because of the condition \\mbox{$V'(\\phi_0)=0$}. Hence, all things considered, the trapping mechanism that operates at the ESP may set the field at a locally flat region in field space. In this paper we show how, after the field is released from its trapping, this flatness may result in enough inflation to solve the flatness and horizon problems.  After being trapped, the modulus oscillates around the ESP under the influence of $V_{\\rm int}(\\phi)$. These oscillations deplete its initial kinetic density until the latter decreases below $V_0$, in which case a phase of so-called trapped inflation begins \\cite{Kofman:2004yc}. Trapped inflation dilutes $V_{\\rm int}$ until the latter is unable to restrain the field, at which time the modulus is released and continues rolling over $V(\\phi)$. As we have shown, trapped inflation is brief and cannot suffice for the solution of the horizon and flatness problems. The duration of trapped inflation is independent from the value of the initial kinetic density of the modulus, because all excess of the latter is depleted before the onset of trapped inflation. Hence, our framework is largely independent of initial conditions, provided trapping occurs.  To study the dynamics after the field is released from its trapping, we have followed a phenomenological approach. In it, the scalar potential $V(\\sigma)$ consists of a `background' potential $V_{\\rm np}(\\sigma)$ of non-perturbative origin (from gaugino condensation, for example), and a phenomenological potential $V_{\\rm ph}(\\sigma)$ chosen so that the scalar potential $V(\\sigma)=V_{\\rm np}(\\sigma)+V_{\\rm ph}(\\sigma)$ satisfies $V^{\\prime}(\\phi_0)=0$. The guiding line to pick out $V_{\\rm ph}$ is to demand that it decreases faster than $V_{\\rm np}$, so that the former is only important when the field is close to the ESP. Given the characteristic form of $V_{\\rm np}(\\sigma)$, the simplest choice for $V_{\\rm ph}$ is to take $V_{\\rm ph}\\propto e^{-q\\sigma}$ far from the ESP with $q>\\mu c$ a constant. A particular realisation of $V_{\\rm ph}$ is given in Eq.~(\\ref{eq:vph}). In this setting, we show in Appendix \\ref{app:minapp} that the terms of order forth and higher in a series expansion around the ESP become important only after the end of inflation. Therefore, the inflationary dynamics can be fully accounted for through the scalar potential $V(\\phi)=V_0+\\frac{1}{2}m_{\\phi}^2\\bar{\\phi}^2+\\frac{1}{3!}\\vppp \\bar{\\phi}^3$, where $\\bar{\\phi}\\equiv\\phi-\\phi_0$.  In this paper we are concerned with the observable amount of inflation. We thus pay special attention to the last $N\\approx60$ $e$-foldings of inflation, and compute the curvature perturbation, $\\pphi$, generated by the inflaton field when the observable Universe leaves the horizon. We show how in our model $\\vppp$ is fully determined by $H_*$ and ${\\cal P}^{1/2}_{\\phi}$, Eq.~(\\ref{eq:pert}). This offers two significant advantages. First, the dynamics becomes described in terms of simple and meaningful physical quantities relevant to inflation. Second, the parameter space of the model becomes automatically constrained by the observed value of the curvature perturbation, which $\\pphi$ cannot exceed, i.e. $\\pphi\\lesssim4.8\\times10^{-5}$. Note that the observed curvature perturbation in our model is generated by a curvaton field which is also employed to reheat the Universe.  A side effect of our model is that it imposes a lower bound in $\\pphi$ [c.f. Eq.~(\\ref{eq:intpert})], which in turn translates into an upper bound on $H_*$, Eq.~(\\ref{eq:upperm0}). We recognise these bounds as an artifact of the model, that appears when we parametrise our ignorance about the behaviour of $V(\\phi)$ around the ESP. A simple way around these unphysical bounds is to let $q$ have a running.  We find that we can attain enough inflation to solve the flatness and horizon problems if $|m_{\\phi}|\\lesssim H_*$ and $\\vppp\\ll H_*$. In particular, we discover that the field may have two different inflationary histories, or regimes, ending in a phase of slow-roll during which the observable Universe leaves the horizon. Whereas one of the inflationary regimes consists only of slow-roll inflation, the other regime involves a previous phase of eternal inflation. Eqs.~(\\ref{eq:intsecsr2}) and (\\ref{eq:inteternal}) [c.f. Eq.~(\\ref{eq:f(g)})] determine which of these inflationary regimes will be driven by the field. Of course, consistency requires that the regimes must be compatible with the trapping condition, Eq.~(\\ref{eq:trp}). The values of $g$ corresponding to the different inflationary regimes, and to the trapping condition, are pictured in Figs.~\\ref{fig:infmap1} and \\ref{fig:infmap2}, corresponding to $m_{\\phi}^2<0$ and $m_{\\phi}^2\\geq0$ respectively. The results in both cases are qualitatively the same: the inflationary regime involving eternal inflation is always compatible with trapping. However, the inflationary regime where only slow-roll inflation occurs, is compatible with trapping only if $|\\eta|$ when sufficiently large, Eq.~(\\ref{eq:eta}). Finally, an acceptable inflationary cosmology is also achievable through fast-roll inflation. This possibility arises only when the ESP results in a maximum with $|m_{\\phi}^2|\\sim H_*^2$. In this case, $H_*$ must satisfy an upper bound, Eq.~(\\ref{eq:Hfr}), so that the field does not generate an excessive amount of perturbations.  After inflation is terminated the field becomes dominated by its kinetic density and the Universe enters a period of so-called kination \\cite{kination}. During this period the modulus is oblivious of the scalar potential and its rapid roll can send it very far in field space. To reheat the Universe we have assumed a curvaton field, whose decay products in time dominate the kinetic density of the rolling modulus \\cite{curvreh}. The existence of a suitable curvaton, apart from reheating the Universe, can also provide the required density perturbations \\cite{curv} and allow for a low inflation scale \\cite{orth,low}. After reheating, the modulus roll is inhibited by cosmological friction as discussed in Refs.~\\cite{eta,cosmofric}. Consequently, in little time, the modulus freezes at a value $\\sigma_F$ shown in Eq.~(\\ref{eq:sigmafrozen}). The modulus remains frozen until the present, which guarantees the non-variation of the fundamental constants through-out the hot big bang despite the fact that the modulus is not stabilised in a local minimum. Eq.~(\\ref{eq:sigmafrozen}) shows that $\\sigma_F$ depends on the reheating temperature, which, in turn, depends on the details of the curvaton model. Therefore, in our framework curvaton physics can determine the value of the string modulus in our Universe.  After the end of kination the modulus is expected to freeze at a large value at the tail of the scalar potential. At such values, we assume that the latter may be dominated by a contribution of the form \\mbox{$V(\\sigma)=C_n/\\sigma^n$} [c.f. Eq.~(\\ref{exponentials})], where $C_n$ is a density scale and the value of $n$ is model dependent. The value of $C_n$ is determined by the coincidence requirement, if the modulus is to become quintessence at present. This requirement results in the constraint in Eq.~(\\ref{eq:reheatingtemp}), where we see that $C_n$ depends also on the reheating temperature and, therefore, it is determined by curvaton physics. The requirement that reheating occurs before Big Bang Nucleosynthesis (BBN) sets an upper bound on $C_n$, shown in Eq.~(\\ref{eq:upperbound}). Crucially, the bound is strongly $n$-dependent.  This type of uplifting potential is again broadly accepted in string flux compactifications.We have discussed some possibilities for the origin of the uplifting potential, such as RR and NS fluxes on $\\overline{D3}$-branes (where $n=2$ and $C_2^{1/4}\\lesssim 10$~MeV), gauge field fluxes on $D7$-branes (where $n=3$ and $C_3^{1/4}\\lesssim 1$~TeV) and $\\alpha'$-corrections (where $n=9/2$ and $C_{9/2}^{1/4}\\lesssim 10^{11}$~GeV). $C_n$ may be further constrained in order to protect BBN from heavy KK modes. It turns out that this constraint is important only in the $n=9/2$ case where the bound on $C_{9/2}$ is strengthened to $C_{9/2}^{1/4}\\lesssim 10^6$~GeV.\\footnote{The uplifting potential considered may even have a perturbative origin, due to duality transformations $\\sigma\\rightarrow 1/\\sigma$ of polynomial terms at large values. KD wishes to thank P.M.~Petropoulos for pointing this out.} Note that the above bounds are much more reasonable than the fine-tuning of $\\Lambda$ required in $\\Lambda$CDM.  According to Eq.~(\\ref{exponentials}), the scalar potential for the canonically normalised field $\\phi$ is of exponential form. For the above discussed values of $n$ and according to Ref.~\\cite{Cline:2001nq} the modulus can cause a brief period of accelerating expansion while it unfreezes at present \\cite{jose}. Thus, our model does not lead to eternal acceleration and therefore does not suffer from the existence of future horizons \\cite{Hellerman:2001yi}. In the future, the modulus eventually approaches $+\\infty$ which corresponds to a supersymmetric ground state. If $\\sigma$ is the volume modulus then this final state leads to decompactification of the extra dimensions.  Some numerical simulations more recent than Ref.~\\cite{Cline:2001nq} suggest that the parameter space for brief acceleration is somewhat reduced. In particular, in Ref.~\\cite{LopesFranca:2002ek} the authors demonstrate that brief acceleration can indeed take place in the interval \\mbox{$\\sqrt 2<b\\leq\\sqrt 3$} corresponding to the range \\mbox{$\\sqrt 3<n\\leq 3/\\sqrt 2$}, which includes the case $n=2$. This is confirmed by Refs.~\\cite{Kallosh:2002gf} and \\cite{Kehagias:2004bd}, where it is implied that brief acceleration, which is terminated at present, can be attained even for larger values of $b$. In Ref.~\\cite{Blais:2004vt} it is indeed shown that one may have brief acceleration up to the value \\mbox{$b=\\frac{3}{4}\\sqrt 6$} corresponding to \\mbox{$n=9/4$}. The authors of Ref.~\\cite{Blais:2004vt}  acknowledge the fact that the parameter space determined by observations is affected by assumed priors and may be extended further. Hence, the case \\mbox{$n=3$} might still be successful. However, in view of the above, the case of $\\alpha'$-corrections to the K\\\"{a}hler potential (Sec.~\\ref{sec:Qspt}) is probably excluded.  In conclusion, we have investigated a paradigm of quintessential inflation using a sting modulus as the inflaton - quintessence field. We have shown that there exists ample parameter space for successful quintessential inflation to occur with reasonable values of the model parameters (with possibly a mild tuning of gravitational couplings between the modulus and the standard model fields, required to overcome the 5th force problem). The possibility of attaining an acceptable inflationary cosmology depends crucially on the values of $\\vppp$ and $m_{\\phi}$ relative to the Hubble scale $H_*$ during inflation. The latter is determined by the location of the ESP and can be much different to $m_\\phi$ or $\\vppp$. Furthermore, a crucial role is played by the value of the reheating temperature, which controls both the value of the frozen modulus and the required value of the density scale in the uplift potential. The reheating temperature is determined, in turn, by the particulars of the assumed curvaton field. Because the modulus remains frozen until the present, the model may conceivably be linked to the intriguing possibility that some of the fundamental constants, such as the fine-structure constant, may begin to vary at present."
    },
    "0606/astro-ph0606108_arXiv.txt": {
        "abstract": "{\\sl Aims.} The purpose of this study is to examine the prediction that the deuterated H$_3^+$ ion, H$_2$D$^+$, can be found exclusively in the coldest regions of molecular cloud cores.  This is also a feasibility study for the detection of the ground-state line of {\\sl ortho}-H$_2$D$^+$ at 372 GHz with APEX.  {\\sl Methods.} The $(1_{10} \\rightarrow 1_{11})$ transition of H$_2$D$^+$ at 372 GHz was searched towards selected positions in the massive star forming cloud OriB9, in the dark cloud L183, and in the low- to intermediate mass star-forming cloud R CrA.  {\\sl Results.} The line was detected in cold, prestellar cores in the regions of OriB9 and L183, but only upper limits were obtained towards other locations which either have elevated temperatures or contain a newly born star. The H$_2$D$^+$ detection towards OriB9 is the first one in a massive star-forming region. The fractional {\\sl ortho}-H$_2$D$^+$ abundances (relative to H$_2$) are estimated to be $\\sim 1\\,10^{-10}$ in two cold cores in OriB9, and $3\\,10^{-10}$ in the cold core of L183.  {\\sl Conclusions.} The H$_2$D$^+$ detection in OriB9 shows that also massive star forming regions contain very cold prestellar cores which probably have reached matured chemical composition characterized, e.g., by a high degree of deuterium fractionation.  Besides as a tracer of the interior parts of prestellar cores, H$_2$D$^+$ may therefore be used to put contraints on the timescales related to massive star formation. ",
        "introduction": "The trihydrogen ion, H$_3^+$, is supposed to become the principal carrier of positive charge in the centres of cold, dense cores when 'heavy' elements like C, O and N, are nearly completely depleted (\\cite{walmsley2004}). Because deuterium fractionation reactions are favoured in cold gas, relatively large abundances of the isotopoloques H$_2$D$^+$ and D$_2$H$^+$ are to be expected in these objects. This has been also confirmed by observations (\\cite{caselli2003}; \\cite{vastel2004}). While infrared absorption spectroscopy of H$_3^+$ can be used to extract vital information on the gas columns towards infrared sources (e.g. \\cite{mccall1999}), the rotational lines of H$_2$D$^+$ and D$_2$H$^+$ probe the physical conditions of cold, prestellar cores. The abundance of H$_3^+$ and its deuterated forms depend on the cosmic ray ionization rate of H$_2$, and on the abundances of destructing agents: electrons, gaseous neutral species like CO and N$_2$, and negatively charged grains (e.g. \\cite{caselli2003}; \\cite{walmsley2004}). Furhermore, the H$_3^+$ abundance and the H$_2$D$^+$/H$_3^+$ abundance ratio depend heavily on the {\\sl ortho:para} ratio of H$_2$, which in turn is a function of time and density (\\cite{pineaudesforets1991}; \\cite{flower2006}).  The ground-state $(1_{10} \\rightarrow 1_{11})$ transition of {\\sl ortho}-H$_2$D$^+$ lies between adjacent atmospheric O$_2$ and H$_2$O absorption lines and its observation requires extremely good conditions.  A reasonable limit is that the precipitable water vapour content of the atmosphere, PWV, is less than 0.5 mm, which can be achieved at high-altitude observatories only. H$_2$D$^+$ has been previously detected from Mauna Kea towards a protostellar core (\\cite{stark1999}), and in a small number of prestellar dark cloud cores (\\cite{caselli2003}; \\cite{stark2004}; \\cite{vastel2006}).   In this {\\em Letter} we report on the first H$_2$D$^+$ observations with the Atacama Pathfinder Experiment, APEX, during its Science Verification periods in July, October and November 2005. The main goal of this series of observations was to test the feasibility of the 372 GHz line observations with this instrument. In the course of these measurements H$_2$D$^+$ was detected in a core belonging to a high-mass star forming region. This may open new vistas to the chemical evolution preceding the collapse of massive stars. ",
        "conclusions": "The performance of APEX and its equipment, and the atmospheric transmission on Chajnantor are found to meet very high standards. The telescope is therefore likely to become a very important tool for studies of molecular cloud interiors and star formation using H$_2$D$^+$, and other 'difficult' molecules.  The present observations towards a small sample of dense cores with some diversity of physical characteristics suggest that either an elevated temperature (as in R CrA NW) or the presence of an embedded star (as towards IRAS 05405-0117), even if there is little evidence for star-cloud interaction, decreases the changes to find H$_2$D$^+$.  The kinetic temperature, velocity dispersion and the fractional H$_2$D$^+$ abundance in OriB9 are similar to those in the prestellar dark cloud cores L1544, 16293E, and L183, where strong emission of this line has been detected previously. The masses and the central densities of these nearby cores are $2-3\\,M_\\odot$ and $\\sim 10^6\\,{\\rm cm}^{-3}$, respectively (\\cite{vastel2006} and references therein).  The total mass of the OriB9 core estimated from ammonia is of the order of $100 \\, M_\\odot$ (\\cite{harju1993}), and the subcondensations seen in the N$_2$H$^+$ map (\\cite{caselli1994}) are likely to be an order of magnitude more massive than L1544, 16293E, or L183.  A (sub)millimetre continuum map is needed to confirm this. Nevertheless, OriB9 seems to be capable of forming a massive star or a dense cluster of low- to intermediate-mass stars, and is therefore exceptional among sources detected in H$_2$D$^+$ so far.  This detection confirms the existence of very cold, quiescent, dense cores in massive star forming regions. The previous evidence for such objects is scarce.  Some infrared dark clouds (IRDCs) have gas temperatures approaching 10~K and ammonia linewidths slightly below 1 km\\,s$^{-1}$ (\\cite{pillai2006}). Yet another massive core with these characteristics has been recently found in the region of ISOSS J18364-0221 (\\cite{birkmann2006}).  The linewidths in these objects are, however, clearly larger than in OriB9. On the other hand, only a small fraction of IRDCs have been observed in spectral lines to date.  Because compression leads to an intensified cooling by molecules and dust, the collapse of all dense cores should be preceded by a cold, quiescent phase. Very cold GMC cores may have remained indiscernible because of their short life-time or the fact that large-scale surveys are usually biased towards the presence of a certain molecular species which might be depleted in the coldest regions.  As discussed recently by \\cite{flower2006}, the H$_2$D$^+$ abundance depends inversely on the {\\sl ortho}:{\\sl para} ratio of H$_2$, which is largest at early stages of core evolution. Consequently, a high degree of deuterium fractionation is a sign of matured chemistry characterized by a low {\\sl o}/{\\sl p} H$_2$ and probably a high degree of molecular depletion. The {\\sl ortho}-H$_2$D$^+$ detection towards OriB9 suggests an evolved chemical stage and tells of a longlasting prestellar phase. However, estimates of the {\\sl o}/{\\sl p} ratio of H$_2$ and the degree of deuterium fractionation using other tracers are needed to confirm this. Despite these obstacles, it seems viable to use the 372 GHz H$_2$D$^+$ line together with chemistry models to estimate timescales related to the early evolution of massive cores.  \\vspace{3mm}  We thank Dr. J.R. Pardo for making the ATM program available, and Malcolm Walmsley and the anonymous referee for helpful comments on the manuscript.  The Helsinki group acknowledges support from the Academy of Finland through grants 1201269, 1210518 and 206049."
    },
    "0606/astro-ph0606422_arXiv.txt": {
        "abstract": "We present \\aco\\ observations of the high--redshift quasi--stellar objects (QSOs) BR\\,1202-0725 ($z=4.69$), PSS\\,J2322+1944 ($z=4.12$), and APM\\,08279+5255 ($z=3.91$) using the NRAO Green Bank Telescope (GBT) and the MPIfR Effelsberg 100\\,m telescope.  We detect, for the first time, the CO ground--level transition in BR\\,1202-0725.  For PSS\\,J2322+1944 and APM\\,08279+5255, our observations result in line fluxes that are consistent with previous NRAO Very Large Array (VLA) observations, but they reveal the full line profiles. We report a typical lensing--corrected velocity--integrated intrinsic \\aco\\ line luminosity of $L'_{\\rm CO} = 5 \\times 10^{10}\\,$K \\kms pc$^2$ and a typical total H$_2$ mass of $M({\\rm H_2}) = 4 \\times 10^{10}$\\,M$_{\\odot}$ for the sources in our sample.  The CO/FIR luminosity ratios of these high--$z$ sources follow the same trend as seen for low--$z$ galaxies, leading to a combined solution of log$(L_{\\rm FIR}) = (1.39 \\pm 0.05) \\times {\\rm log} (L_{\\rm CO})-1.76$.  It has previously been suggested that the molecular gas reservoirs in some quasar host galaxies may exhibit luminous, extended \\aco\\ components that are not observed in the higher--$J$ CO transitions.  Utilizing the line profiles and the total intensities of our observations and large velocity gradient (LVG) models based on previous results for higher--$J$ CO transitions, we derive that emission from all CO transitions is described well by a single gas component where all molecular gas is concentrated in a compact nuclear region.  Thus, our observations and models show no indication of a luminous extended, low surface brightness molecular gas component in any of the high--redshift QSOs in our sample.  If such extended components exist, their contribution to the overall luminosity is limited to at most 30\\%. ",
        "introduction": "Understanding when and how galaxies form is one of the primary objectives in both observational and theoretical astrophysics. The mere fact that active galaxies, such as radio galaxies, QSOs, and emission line galaxies are observed up to redshifts of $z = 6.6$ (e.g., Fan et al.\\ 2001; Rhoads et al.\\ 2001; Hu et al.\\ 2002; Kodaira et al.\\ 2003; Taniguchi et al.\\ 2005), less than one Giga--year after recombination, implies that galactic scale ($\\sim$10\\,kpc), gravitationally bound structures exist at this early epoch. The study of the masses and dynamical state of these young systems serves as a direct constraint to the models describing the growth of large scale structures since the epoch of recombination.  Studies of the molecular and dusty interstellar medium (ISM) in these galaxies are of fundamental importance, since it is this medium out of which stars form; accurate determination of the molecular gas mass may therefore serve as an indicator of the evolutionary state of a galaxy. The detection of carbon monoxide (CO) is also a strong confirmation that star formation is going on in some of the highest redshift systems. In fact, the combination of molecular gas and dust detections with large far--infrared (FIR) luminosities provides the strongest evidence that a significant fraction of high--$z$ galaxies is undergoing starbursts at prodigious rates ($>10^3\\,$M$_\\odot {\\rm yr}^{-1}$), consistent with the formation of a large elliptical galaxy on a dynamical timescale of $\\sim$10$^7$--10$^8$\\,yr.  Over the past decade, more than 30 galaxies at $z>2$ have been detected in CO emission (Solomon \\& Vanden Bout 2005) out to a redshift of $z=6.42$ (Walter et al.\\ 2003, 2004; Bertoldi et al.\\ 2003), confirming the presence of intense starbursts in numerous high--$z$ galaxies. As most of these observations were obtained using millimeter interferometers, these detections were typically achieved by observing high--$J$ CO($J \\to J-1$) transitions ($J \\geq 3$). Although these high--$J$ lines exhibit in general higher peak flux densities than the ground--state ($J=1$) transition, it is possible that the higher order transitions are biased to the excited gas close to a central starburst and do not necessarily trace the entire molecular gas reservoir seen in \\aco.  Also, the conversion factor ('$\\alpha$') to derive molecular (H$_2$) masses from measured CO luminosities has mostly been estimated for the \\aco\\ line (e.g., Downes \\& Solomon 1998; Weiss et al.\\ 2001).  Observing \\aco\\ has the additional advantage that properties of the highly redshifted sources can be {\\em directly} compared to the molecular gas properties of nearby (starburst) galaxies which are predominantly mapped in the \\aco\\ transition.  These are main motivations for observing the \\aco\\ ground--state transition.  However, due to the faintness of the line and the bandwidth limitations of current radio telescopes, such high--$z$ \\aco\\ observations only exist for two QSOs and two radio galaxies to date.  All of these observations have been obtained with radio interferometers operating at centimeter wavelengths -- the NRAO Very Large Array (VLA; Papadopoulos et al.\\ 2001; Carilli et al.\\ 2002a; Greve et al.\\ 2004) and the Australia Telescope Compact Array (ATCA; Klamer et al.\\ 2005). In particular, due to the bandwidth limitations of the VLA, obtaining better constraints on the spectral line shape and total flux of the \\aco\\ transition is desirable even for already detected sources. Today's largest single dish telescopes, such as the NRAO Green Bank Telescope (GBT)\\footnote{The Green Bank Telescope is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.}  and the MPIfR Effelsberg telescope\\footnote{The Effelsberg telescope is a facility of the Max-Planck-Gesellschaft (MPG), operated by the Max-Planck-Institut f{\\\"u}r Radioastronomie (MPIfR).}, can eliminate some of those issues due to their larger spectral bandwidths.  Here we report on first observations of \\aco\\ in three CO--bright high--$z$ QSOs, which also exhibit ultra--luminous IR emission, using the GBT and the Effelsberg telescope. In Section 2, we describe our observations. Section 3 summarizes our results on the individual objects (BR\\,1202-0725 at $z=4.69$, PSS\\,J2322+1944 at $z=4.12$, and APM\\,08279+5255 at $z=3.91$).  Section 4 provides an analysis and discussion, and Section 5 closes with a summary of our results.  We assume a standard $\\Lambda$ cosmology throughout, with $H_0 = 71\\,$\\kms\\ Mpc$^{-1}$, $\\Omega_M = 0.27$, and $\\Omega_{\\Lambda} = 0.73$ (Spergel et al.\\ 2003). ",
        "conclusions": "\\subsection{LVG modeling} \\label{LVG}  To investigate how much of the \\aco\\ emission in our target QSOs is associated with the molecular gas reservoirs detected in the higher--$J$ CO transitions, we have used spherical, one--component large velocity gradient models (LVG).  All these LVG calculations use the collision rates from Flower \\& Pineau des Forets (2001) with an ortho/para H$_2$ ratio of 3 and a CO abundance per velocity gradient of [CO]/$\\nabla v = 1 \\times 10^{-5}\\,{\\rm pc}\\,$(\\kms)$^{-1}$ (e.g.\\ Weiss et al.\\ 2005). Models were fitted to those lines {\\em above} the \\aco\\ transition listed in Table \\ref{tab-01} for each source. The turnover of the CO line SED (and therefore the slope beyond the turnover) is not well determined. Thus, a large degeneracy exists between the kinetic gas temperature $T_{\\rm kin}$ and density $\\rho_{\\rm gas}({\\rm H_2})$, the two main free parameters in our study.  As an example, Figure \\ref{fig6} shows data of all transitions and three representative models for BR\\,1202-0725: Model 1 (solid line) assumes $T_{\\rm kin} = 60\\,$K and $\\rho_{\\rm gas}({\\rm H_2}) = 10^{4.1}\\,$cm$^{-3}$, and gives the overall best fit to all transitions. Models 2 (dashed line) and 3 (dotted line) are shown as a representation of the parameter space allowed by the data within the error bars. Model 2 assumes $T_{\\rm kin} = 120\\,$K and $\\rho_{\\rm gas}({\\rm H_2}) = 10^{3.7}\\,$cm$^{-3}$, while model 3 assumes $T_{\\rm kin} = 30\\,$K and $\\rho_{\\rm gas}({\\rm H_2}) = 10^{4.6}\\,$cm$^{-3}$. The LVG predicted \\aco\\ flux of the different models based on the $J>1$ CO transitions is fairly well constrained by the solutions. Most of our calculated models suggest that the CO emission is close to thermalized up to the \\dco\\ transition and optically thick ($\\tau \\gtrsim 5$), which implies that the LVG predicted \\aco\\ line fluxes are similar to those we would derive by assuming a $\\nu^2$ scaling of the line flux densities from the mid--$J$ CO transitions.  The predicted LVG integrated flux ranges of the \\aco\\ transition are 0.10--0.12\\,Jy \\kms\\ (BR\\,1202-0725), 0.20--0.23\\,Jy \\kms\\ (PSS\\,J2322+1944), and 0.13--0.20\\,Jy \\kms\\ (APM\\,08279+5255) and are in good agreement with our observations for all sources.  As described in Sect.\\ 3.2.2, the observed \\aco\\ linewidth in our new observations and in the VLA spectrum (Carilli et al.\\ 2002a) of PSS\\,J2322+1944 is lower than that of the higher--$J$ transitions, thus the integrated model fluxes are a bit higher than our result. If we compare the peak fluxes only, we obtain a good agreement.  Our observations and models are in agreement with the assertion that {\\em all} observed \\aco\\ flux density is associated with the highly excited molecular gas seen in the high--$J$ CO lines. We thus find no evidence for an additional luminous, more extended low surface brightness gas component surrounding the central region of our target QSOs, in contrast to what has been suggested previously for APM\\,08279+5255 (Papadopoulos et al.\\ 2001).  Given the accuracy of our measurements, we conclude that at most 20--30\\% of the \\aco\\ luminosity may be associated with such a diffuse component.  We note however that, if the $L'_{\\rm CO}$ to $M_{\\rm gas}({\\rm H_2})$ conversion factor ('$\\alpha$') for a faint extended component were higher (e.g.\\ Galactic), a higher H$_2$ mass may be hidden in such an extended component.  \\subsection{Correlations of high and low redshift galaxies} \\label{gaosolomon}  CO results for our three high--redshift IR--luminous QSOs are summarized in Tab.~\\ref{tab-03} and \\ref{tab-04}.  Our sample consists of all three high--$z$ QSOs for which the CO ground--state transition has been detected to date, and covers $\\sim$20\\% of the CO--detected $z>2$ quasars.  As discussed above, the \\aco\\ transition provides the best information of the total amount of molecular gas in a system, quantified by $L'_{\\rm CO(1-0)}$. For most QSOs/galaxies detected in CO at high $z$, \\cco\\ and \\dco\\ are the lowest transitions that have been detected to date. Based on our results in the previous paragraph, however, we can now estimate their $L'_{\\rm CO(1-0)}$ by assuming constant brightness temperature for all transitions from \\aco\\ to \\dco\\ (i.e., $L'_{\\rm CO}$ is the same for those transitions).  This assumption is also in agreement with observations towards the radio galaxy 4C60.07 (Greve et al.\\ 2004). However, it is important to note that all high--$z$ sources were selected via higher--$J$ CO transitions, which could, in principle, introduce a bias towards highly excitated starburst environments.  Following Sanders et al.\\ (1991), Gao \\& Solomon (2004) found a non--linear relation between the logarithms of the FIR luminosity $L_{\\rm FIR}$ and the \\aco\\ line luminosity $L'_{\\rm CO}$ for a sample of local spiral galaxies, luminous, and ultra--luminous infrared galaxies. Their sample consists mostly of galaxies whose FIR luminosity is powered by star formation only.  Our \\aco\\ observations and LVG models suggest that $L'_{\\rm CO}$ can be estimated for all CO--detected high--$z$ sources with some degree of confidence, even if only observations of higher--$J$ CO transitions exist for most of these sources. To put this result into context, we now aim to discuss the $L'_{\\rm CO}$--$L_{\\rm FIR}$ relation and its implications for some selected samples of galaxies at low and high redshift.  Figure \\ref{fig7} shows the relationship between log($L_{\\rm FIR}$) and log($L'_{\\rm CO}$) for our three targets, all other high--$z$ CO detections (except TN\\,J0924-2201, which does not have a measured FIR luminosity), the $z < 0.2$ Palomar--Green (PG) QSOs from Alloin et al.\\ (1992), Evans et al.\\ (2001), and Scoville et al.\\ (2003) including PDS\\,456 (Yun et al.\\ 2004), local ($z < 0.3$) ULIRGs from Solomon et al.\\ (1997), and the local galaxy sample ($z < 0.1$) of Gao \\& Solomon (2004). All FIR luminosities for the high--$z$ sources are re--derived as described by Carilli et al.\\ (2005) unless stated otherwise. Lensing magnification factors were taken into account. In addition, data for the Milky Way disk (Fixsen et al.\\ 1999) and an (extrapolated) relation for Galactic molecular clouds (Mooney \\& Solomon 1988) are given for comparison. The solid line is a straight--line least squares fit to the Gao \\& Solomon (2004) sample, corresponding to log$(L_{\\rm FIR}) = (1.26 \\pm 0.08) \\times {\\rm log} (L_{\\rm CO})-0.81$ (the power law index is 1.25 $\\pm$ 0.08 in the original publication).  Fitting the high--$z$ sources, the Gao \\& Solomon (2004) data, the Solomon et al.\\ (1997) ULIRGs, and the PG QSOs together, we find the relation log$(L_{\\rm FIR}) = (1.39 \\pm 0.05) \\times {\\rm log} (L_{\\rm CO})-1.76$. Both relations are clearly non--linear, as their power--law indexes are significantly larger than unity (see also the corresponding discussion in Gao \\& Solomon 2004). We recover a Schmidt--Kennicutt law (power law index of 1.4, Kennicutt 1998a,b) for the larger, more heterogenous sample.  More importantly, the high--$z$ sources seem to follow the same slope (within the errors) as seen at low $z$, even though most high--$z$ sources harbor central AGNs.  In addition, the high--$z$ QSOs and high--$z$ sub--millimeter and radio galaxies (likely lacking a luminous AGN) statistically occupy the same area in the plot. The same behaviour is found for the PG QSOs in comparison to the low--$z$ spirals/LIRGs/ULIRGs.  It is in general not clear that the FIR emission comes mostly from star--formation, i.e.\\ that it is not strongly biased by active nuclei or the interstellar radiation field. E.g., SED modeling of the $z=2.6$ quasar H1413+117 (the Cloverleaf) reveals that only $\\sim$20\\% of its FIR luminosity is powered by star formation (Solomon et al.\\ 2003), but it is not an outlier in the $L'_{\\rm CO}$--$L_{\\rm FIR}$ diagram.  In this context, the elevated $L_{\\rm FIR}$/$L'_{\\rm CO}$ ratio in APM\\,08279+5255 may be explained by a combination of differential lensing and a relatively high contribution of the AGN to $L_{\\rm FIR}$ (see Rowan--Robinson 2000 for models of the IR SED).  Outliers like APM\\,08279+5255 are also found at low $z$ (e.g.\\ the QSO PDS\\,456 and the warm ULIRG IRAS\\,08572+3915, see Fig.\\ \\ref{fig7}).  As the relative number of such outliers is very small, this may also be explained by a brief FIR--bright AGN phase (Yun et al.\\ 2004).  However, it is likely that the dominant energy source in most ULIRGs is an extreme starburst rather than heating by a dust--enshrouded AGN (e.g., Solomon et al.\\ 1997; Downes \\& Solomon 1998).  This picture is supported by the finding that ULIRGs harbor large quantities of {\\em dense} molecular gas, which is more intimately involved with star formation than the major fraction of the mostly diffuse CO (Gao \\& Solomon 2004).  It has to be kept in mind that the high--$z$ sources are highly selected, and probably fulfill the Malmquist bias (i.e., the apparent increase in the average $L'_{\\rm CO}$ and $L_{\\rm FIR}$ towards high redshift is probably a consequence of the flux limitation in the sample, e.g.\\ Sandage 1994). As an additional consequence of flux limitation in the sample, lensing influences the results (i.e., stronger lensing magnification allows to probe deeper, and therefore the intrinsic CO luminosity drops).  We can also set an approximate lower limit on the duration of the intense starburst phase. Considering the total H$_2$ masses and star--formation rates in Table \\ref{tab-04}, the depletion timescale of the molecular gas is of order 10$^7$\\,yr for all three QSOs under the assumption of a constant SFR and 100\\% star--forming efficiency. This implies that the starburst itself can be relatively short--lived, (and compact, as the dynamical time must be less than the starburst lifetime) unless the molecular gas in which the star--formation occurs can be re--supplied on timescales of $\\sim 10^7$\\,yr."
    },
    "0606/astro-ph0606087_arXiv.txt": {
        "abstract": "We report the most sensitive water maser survey towards Bok globules to date, using NASA's 70m antenna in Robledo de Chavela (Spain). We observed 207 positions within the \\citet{cb} catalog with a higher probability of harboring a young star, using as selection criteria the presence of radio continuum emission (from submillimeter to centimeter wavelengths), geometrical centers of molecular outflows, peaks in maps of high-density gas tracers (NH$_3$ or CS), and IRAS point sources. We have obtained 7 maser detections, 6 of which (in CB 34, CB 54, CB 65, CB 101, CB 199, and CB 232) are reported for the first time here. Most of the water masers we detected are likely to be associated with young stellar objects (YSOs), except for CB 101 (probably an evolved object) and CB 65 (uncertain nature). The water maser in CB 199 shows a relatively high shift ($\\simeq 30$ km s$^{-1}$) of its velocity centroid with respect to the cloud velocity, which is unusual for low-mass YSOs. We speculate that high-velocity masers in this kind of object could be related with episodes of energetic mass-loss in close binaries. Alternatively, the maser in CB 199 could be pumped by a protoplanetary or a young planetary nebula. CB 232 is the smallest Bok globule ($\\simeq 0.6$ pc) known to be associated with water maser emission, although it would be superseded by the cases of CB 65 ($\\simeq 0.3$ pc) and CB 199 ($\\simeq 0.5$ pc) if their association with YSOs is confirmed. All our selection criteria have statistically compatible detection rates, except for IRAS sources, which tend to be a somewhat worse predictor for the presence of maser emission. ",
        "introduction": "Bok globules \\citep*{bok47} are small (projected size $\\la 20'$), isolated, and relatively simple  molecular clouds, with typical masses of $\\simeq 5-50$ M$_\\odot$ \\citep*{mar78, cle91}. These globules are usually identified and cataloged as dark patches in optical images (e.g., \\citealt*{cb}, hereafter CB). The reference catalogs of Bok globules compiled by CB and  \\citet*[hereafter BHR]{bhr} in the north and south hemispheres, respectively, focused on the group of smaller Bok globules, using a size of $<10'$ as a selection criterion. Under the assumption that they are nearby clouds ($d\\la 500$ pc), as suggested by the small number of foreground stars, the globules in the CB and BHR catalogs would in general have linear sizes $\\la 1$ pc.    Some Bok globules are sites of low- and intermediate-mass star  formation (e.g., \\citealt*{yun91}, \\citeyear{out-yc92}; \\citealt*{rei92}). Since they are identified from optical images, without any characterization of their possible star-forming activity, catalogs of Bok globules may span a wide range of evolutionary stages, from quiescent dark cores to clusters of Herbig Ae/Be or T-Tauri stars. Therefore, systematic studies of these globules are potentially useful to study the evolution of phenomena related to  star formation, like collapse, fragmentation, mass-loss, or formation of protoplanetary disks. Moreover, being small and relatively simple molecular clouds, one can study these phenomena with a lower chance of contamination from multiple generations of young stellar objects (YSOs) within the same region. For these reasons, they are interesting laboratories in the study of the star-formation processes and its evolution.  Another important reason to study star formation in Bok globules is that field stars in the solar neighborhood, and the Sun itself, may have formed in this kind of cloud. Moreover, some T Tauri stars apparently not related to known molecular clouds, may have originated in Bok globules which have dispersed, leaving isolated pre-main-sequence objects \\citep*{mm-lh97}.   An important phenomenon related to star formation is the occurrence of water maser emission at 22 GHz. This emission is a powerful tool to study the morphology and kinematics of the environment of YSOs at high angular and spectral resolution \\citep*{tor02,cla05}. The physical conditions needed for maser excitation can be reached in both the inner parts of circumstellar disks surrounding the YSOs and in shocked gas related to winds. This  dichotomy has been suggested to be an evolutionary effect by \\citet*{tor97,tor98}, in which masers may trace bound motions (e.g., disks) in the youngest protostars, and  outflows in  more evolved objects.  Although water masers are more intense and widespread in high-mass star forming regions, they are also present around low-mass YSOs \\citep*{wil87,ter92}. In the case of low-mass star formation, there seems to be an evolutionary trend, with water masers tracing the youngest YSOs, i.e., Class 0 protostars. In their survey of water maser emission towards low-mass YSOs, \\citet*{fur01} obtained a  detection rate of 40$\\%$ on Class 0 sources, 4$\\%$ on Class I, and 0$\\%$ on Class II.  Since there are several evolutionary aspects related to water maser emission, it is interesting to study this emission in Bok globules, where we can find YSOs in different evolutionary stages. However, there are very few studies of water masers towards Bok globules. Only \\citet*{sca91} conducted a search for water masers specifically in Bok globules. These authors observed 80 globules, and obtained a detection towards CB 3 (LBN 594), with a peak flux density of $\\simeq 77.6$ Jy. However, their detection threshold ($\\sim 4.4-15.7$ Jy) would miss a high fraction of masers around low-mass YSOs.  Other surveys included among their targets some objects within the globules in the CB catalog \\citep*{fel92,pal93,wou93,per94,cod95}. To our knowledge, only three Bok globules of the CB catalog have been reported to harbor a water maser: the mentioned one in CB 3 \\citep{sca91}, a detection of $\\simeq 1-2$ Jy towards the Herbig Be star HD 250550 in CB 39 \\citep*{sch75}, and another one of $\\simeq 7.5$ Jy in CB 205 (LDN 810) obtained by \\citet*{nec85}. We are not aware of any water masers in the southern Bok globules of the BHR catalog. However, all these three globules (CB 3, CB 39, and CB 205) are somewhat different from most of the clouds in the CB catalog, since they are located at distances $\\geq 1$ kpc, and thus are larger and more massive than the globules that CB intended to compile.   In this paper, we present the most extensive and sensitive survey of water masers in Bok globules to date. We have searched for this emission towards the positions within the globules in the CB catalog with the highest probability of harboring a YSO (see Sec. \\ref{sec:sample}). This survey aims to be nearly complete for possible star-forming regions in Bok globules north of declination $\\simeq -36^\\circ$, to study the conditions and characteristics of maser emission in these isolated globules. This work is complemented by high-resolution interferometric studies of the detections in this survey, to accurately determine their spatial and velocity distribution, and to pinpoint their exciting source  (de Gregorio-Monsalvo et al. 2006, hereafter DG06). The present paper is structured as follows: In Sec. \\ref{sec:observations} we describe the technical details of the observations, in Sec. \\ref{sec:sample} we explain the criteria used to select the target sources, and we present our results in Sec. \\ref{sec:results}, which are further discussed in Sec. \\ref{sec:discussion}. We summarize our conclusions in Sec. \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In this paper, we present the most sensitive survey for water maser emission towards Bok globules to date, using NASA's 70m antenna in Robledo de Chavela (Spain). A total of 203 target positions within the clouds of the \\citet{cb} catalogs were observed. Our main results are as follow: \\begin{itemize} \\item We have obtained 7 maser detections, 6 of which (in CB 34, CB 54, CB 65, CB 101, CB 199, and CB 232) are new. Of the previously known masers in the CB catalog, only the one in CB 3 was detected by us. No emission was seen towards the previously reported masers in either CB 39 or CB 205. \\item Of our detections, the ones in CB 3, CB 34, CB 54, and CB 232 are most likely associated with YSOs. We suggest that the maser in CB 101 is associated with an evolved object. The nature of CB 65 is uncertain. \\item In the case of CB 199, the relatively large shift ($\\simeq 30$ km s$^{-1}$) of the centroid velocity of the maser emission with respect to the cloud velocity is unusual for YSOs, specially for low-mass ones. We speculate that the presence of high-velocity masers in low-mass YSOs might be related with episodes of energetic mass loss in close binaries. Alternatively, the maser in CB 199 could be related to a protoplanetary or a young planetary nebula. \\item CB 232 is the smallest Bok globule (size $\\simeq 0.6$ pc) known to be associated with water maser emission. However, if their association with YSOs is confirmed, CB 65 (size $\\simeq 0.3$ pc) and CB 199 ($\\simeq 0.5$ pc) are even smaller globules. These objects are good candidates for the study of relatively isolated star formation with high spatial resolution. \\item Of our selection criteria, the more restrictive ones for the identification of YSOs (radio continuum emission, peak of a high-density tracer, and geometrical center of a molecular outflow) show statistically compatible rates of water maser detection. Only the presence of IRAS sources tends to be a somewhat worse predictor for the presence of masers. \\end{itemize}"
    },
    "0606/astro-ph0606752_arXiv.txt": {
        "abstract": "We report on the spectral analysis results of the neutron star, atoll type, low mass X-ray Binary \\source observed by \\integral and \\sax satellites. Spectral behavior in three different epochs corresponding to three different spectral states has been deeply investigated. Two data set spectra show a continuum well described by one or two soft blackbody plus a Comptonized components with changes in the Comptonizing electrons and black body temperature and the accretion rates, which are typical of the spectral transitions from high to low state. In one occasion \\integral spectrum shows, for first time in this source, a hard tail dominating the emission above 30 keV. The total spectrum is fitted as the sum of a Comptonized component similar to soft state and a power-law component ($\\Gamma=2.76$), indicating the presence of a non thermal electron distribution of velocities. In this case, a comparison with hard tails detected in soft states from neutron stars systems and some black hole binaries suggests that a similar mechanism could originate these components in both cases. ",
        "introduction": "\\source is a neutron star low mass X-ray binary (LMXB) classified as a atoll source (Hasinger \\& van der Klis 1989), with an orbital period of 3.8 hr derived from the optical variability of its companion V801 Arae (Pedersen, van Paradijs \\& Lewin 1984) and at distance of 3.7--6.5 kpc (Fujimoto et al. 1988, Smale \\& Lochner 1992, Augusteijn et al. 1998).  While the X-ray burst properties and timing signatures have been analyzed extensively (see Jonker et al. 2005, Belloni et al. 2005 and references therein) the spectral characteristics have been studied only at low energy with \\einstein\\/, \\exosat\\/, \\temna\\/ and \\asca\\/. In general, the spectrum was acceptably fitted by a Comptonization model plus a black body component. A coronal temperature of 2.3-\u00ad2.6 keV, an optical depth of 13, and a soft black body temperature of 0.55 keV were the best fit parameter for the Einstein data (Christian and Swank 1997). \\exosat\\/ data reveal two different source intensity: for the maximum intensity the coronal temperature was $\\sim 3.5$ keV and the optical depth $\\sim 10.4$, while for the minimum intensity the temperature was $\\sim 6.1$ keV and the optical depth $\\sim 7.2$. White et al. (1988) found the coronal temperature to be 1.8 keV and the optical depth to be 16.7. A model consisting of a black body, a multi-color disk and a broad Fe line, was used by Asai et al. (1998) to fit the \\asca data, confirming previous \\temna results (Waki et al. 1984). We report here a broad band spectral analysis performed on data from \\sax and \\integral satellites, which allowed us to better constrain the spectral parameters and to detect the presence of a high energy tail dominating the spectrum above $\\sim 30$ keV. A similar feature has been observed in other LMXBs namely GX~17+2 (Di Salvo et al. 2000), GX~349+2 (Di Salvo et al. 2001), Sco~X-1 (D'Amico et al. 2001), 4U~1608-522 (Zhang et al. 1996), XB~1254-690 (Iaria et al. 2001), Cir~X-1 (Iaria et al. 2002) and 4U~0614+091 (Piraino et al. 1999). ",
        "conclusions": "\\label{discussion} According to our present understanding, the black-body component in the soft states could originate at both the neutron star surface and the surface of an optically-thick accretion disk. In our observations the two black body components seems to originate from two different parts of the disk, corrisponding to two different temperature. The high optical thickness of the plasma rules out the black body component as originates from the neutron star cap regions. Anyway, independently from the optical thickness of the plasma, it is very improbable to have emission from the polar caps of the neutron star for this source. Infact the magnetic field of the LMXBs is $\\sim$ $10^8$ Gauss implying that the accretion onto the neutron  star is not magnetically dominated and the matter accretes onto the whole NS surface. The Comptonization component may arise from a corona above the disk and/or between the disk and the neutron star surface. In the hard states, accretion probably assumes the form of a truncated outer accretion disk and a hot inner flow, joining the disk and the stellar surface as previously reported by Barret \\& Olive (2002) for LMXB 4U~1705-44. The spectral transitions are generally, but not necessarily, coupled with changes in luminosity, indicating they are driven by variability of the accretion rate or change of the geometry of the system as for Black Hole hard/soft transition at constant luminosity (Belloni et al. 2005). For \\source the accretion rate is lower in the hard state than in the soft state: for an accretion efficiency of $\\eta = 0.2$ (corresponding, e.g., to $M_{\\rm NS}=1.4 {\\rm M}_{\\sun}$ and $R_{\\rm NS}=10$ km ) and using our model luminosities, $L_{0.1-200\\,\\rm keV}$, we get $\\dot{M}_{soft}\\simeq 7.3 \\times 10^{-9} {\\rm M}_{\\sun}$ yr$^{-1}$ and $\\dot{M}_{hard}\\simeq 4.0 \\times 10^{-9} {\\rm M}_{\\sun}$ yr$^{-1}$. Transitions are apparently accompanied by changing in the geometry of the flow, and in the relative contribution of the blackbody-like and Comptonization components. In the peculiar state the accretion rate becames very high ($\\dot{M}_{peculiar}\\simeq 2.1 \\times 10^{-8} {\\rm M}_{\\sun}$ yr$^{-1}$), but this value can be influenced by the steep power law that dominates the energy spectrum at low energies. The emission is best described by Comptonization from a complex electron distribution due to a low temperature ($\\sim 3-4$keV) thermal electron distribution together with non-thermal power-law electrons. This two-component electron distribution could arise from non-thermal electron acceleration regions powered by magnetic reconnections above a disc. Low-energy electrons cool preferentially by Coulomb collisions leading to a thermal distribution while the high-energy electrons cool by Compton scattering, preserving a non-thermal distribution (Coppi 1999). Alternatively, the thermal and non-thermal electrons could be spatially distinct, e.g.\\ magnetic reconnection above the disc can produce a non-thermal electron distribution, while overheating of the inner disc produces the thermal Comptonization (Kubota et al.~2001).  LMXBs exhibit hard X-ray states which are very similar to the Black Hole Candidates ones and investigating the possible disc-jet coupling is very interesting and still matter of discussion. The power-law component could be produced by Componization by synchrotron emission in the jet (Bosch-Ramon et al. 2005, Fender 2004) as proposed for others LMXBs: the power law component in Cir~X-1 is explained by Iaria et al. (2002) as jet emission that has been resolved using radio interferometry by Fender \\& Kuulkers (2001); a significant correlation between radio and X-ray flux has been detected for 4U~1728-34 (Migliari et al. 2003), indicating a clear signature of disc-jet coupling. In our source, this hypothesis is strengthened by radio detection: Sydney University Molonglo Sky Survey catalog gives a flux of 7.5 mJy at 843 GHz (Mauch et al. 2003). Assuming a radio spectral index of -0.5 to estimate the flux density at 5 GHz based on other observations, we found a radio loudness $P_R/P_X$ (as defined in Fender \\& Kuulkers 2001) in the range 1.1-2.2 Jy/Crab, where $P_R$ is the radio flux density at 5 GHz and $P_X$ is the peak X-ray flux measured in the soft band ($<12~keV$), even though the two data set are not simultaneous. The radio loudness value is consistent with the one for Cir X-1 (Fender \\& Kuulkers 2001) and it is very high respect to other LMXBs and quite similar to that of Back Hole Candidates. These founding are in favour of the jet hypothesis as origin of the power-law observed for \\source similarly to GX~5-1, Cir~X-1 and GX~17+2 being all of them associated with variable radio sources (Fender \\& Hendry 2000, and reference therein).  Simultaneous high energy and radio observations during spectral transition are crucial to disentangle the power law emission in the LMXBs containing a neutron star.  Finally, Laurent \\& Titarchuk (1999) suggest the detection of a power-law component at high energy to be a signature of presence black hole in an X-ray binary system. Our data are supporting result on GX~17+2 by Di Salvo et al. (2000), clearly showing this criterion proposed to distinguish black hole versus neutron star binaries is inadequate."
    },
    "0606/astro-ph0606034_arXiv.txt": {
        "abstract": "Mass-varying neutrino (MaVaN) models propose a source of dark energy in a new scalar field called the acceleron.  Recent work has shown that nonrelativistic neutrino fields in these theories are unstable to inhomogeneous fluctuations, and form structures that no longer behave as dark energy.  One might expect that in multiple-neutrino models, the lighter species could continue to act as a source for the acceleron, generating dark energy without the help of heavier species. This paper shows that by considering the evolution of the acceleron field for a large class of models, the result of any component becoming unstable is that all components become unstable within a short time on cosmological scales.  An alternate model employing a second scalar field in a hybrid potential is shown to have stable MaVaN dark energy even in the presence of unstable heavier components. ",
        "introduction": "Recent results in precision cosmology have created a number of puzzles.  Among the unexplained phenomena are several apparent coincidences in which the energy density of two components are comparable despite having different redshift properties.  Another is the existence of a negative pressure cosmological fluid that accelerates the rate of expansion of the universe.  The model of mass-varying neutrinos (MaVaNs), introduced by Fardon et al. in~\\cite{Fardon:2003eh}, suggests that neutrino and dark energy densities have tracked each other throughout the lifetime of the universe through a new scalar field called the acceleron.  The energy density of the scalar potential of this new field contributes to the dark energy.  The authors showed that this new field is capable of explaining the present cosmological expansion.  Since the original MaVaN proposal, several authors have investigated the stability of the \\emph{dark sector}, composed of neutrinos and the scalar field, under perturbations to the neutrino density~\\cite{Afshordi:2005ym,Takahashi:2006jt}.  Both groups subject the model to a hydrodynamic analysis.  In this approximate picture, the speed of sound squared in the cosmological fluid, given by~\\cite{Mukhanov:1990me}, \\beq c_s^2=\\frac{\\dot{P}}{\\dot{\\rho}}=w+\\frac{\\dot{w}\\rho}{\\dot{\\rho}}=w-\\frac{\\dot{w}}{3H(1+w)} \\eeq is positive only if \\beq \\frac{\\partial w}{\\partial z} \\geq -\\frac{3w(1+w)}{1+z} \\eeq where $z$ is the redshift.  However, for nonrelativistic neutrinos, $\\frac{\\partial w}{\\partial z}$ is a negative quantity while the right hand side is positive for $w$ close to -1.  Since at least one neutrino must be nonrelativistic today, the dark sector appears unstable.  Additionally, Afshordi, et al. perform a stability analysis using kinetic theory to account for neutrino streaming~\\cite{Afshordi:2005ym}.  This analysis also shows that perturbations in the neutrino field become unstable when the mass of the neutrino is of order the neutrino temperature.  The ratio of mass to temperature when the instability occurs is a function of the acceleron potential, but is approximately 7 for the potentials considered in this paper.  Afshordi, et al., also examine the result of undergoing a phase transition in a neutrino component, and suggest that the unstable neutrino field may rapidly form nonlinear structures termed \\emph{neutrino nuggets}, which redshift as dark matter and cannot drive the cosmic expansion.  In this paper, we will assume that the neutrinos are initially relativistic.  As the universe expands and cools, the neutrinos become less relativistic until their mass and temperature are approximately equal. At this point, we assume that they decouple from the scalar field into some structure such as the one described by Afshordi et al.  The resulting dark matter will not provide a large contribution to the energy density, and we will not include the contribution in the dark sector energy.  Note that in this paper dark sector will always refer to the energy density contributions of the stable neutrinos and the scalar field.  In a model with multiple neutrinos, the dark energy may still be driven by relativistic species even after the heavier components have become unstable. However, the coupling between the neutrinos and acceleron create a feedback mechanism. The shift in the acceleron expectation value when a neutrino becomes unstable can be sufficient to necessarily change the mass of another species so that it goes from relativistic to nonrelativistic, causing it to become unstable as well.  This \\emph{cascaded instability} can make all neutrinos unstable at about the same time that the heaviest neutrino becomes unstable. This paper will examine a particular class of models, and will show that see-saw MaVaN models with flat scalar potentials suffer from precisely this problem. The timing between the instability in successive neutrinos is strongly dependent on the flatness of the scalar potential, but a flat potential is also required to generate the observed dark energy.  This result may point toward models that do no suffer from a cascaded instability, and this paper concludes with a simple example. ",
        "conclusions": "Several authors have found that a nonrelativistic MaVaN neutrino field is unstable to inhomogeneous fluctuations.  By considering the evolution of the acceleron expectation value when a neutrino field becomes unstable, we have argued above that all neutrino fields in the theory are susceptible to a cascaded instability in which they all become unstable at nearly the same time.  This occurs as long as the scalar potential has a nearly flat dependence on the acceleron, and the neutrino masses vary inversely with the acceleron.  Including a very light neutrino is not sufficient to avoid this problem. Since there are at least three neutrino species, and the atmospheric neutrino deficit requires at least one mass scale above 1eV, the instability poses a constraint on all physical MaVaN models.  There are a number of possible resolutions.  One is to increase the curvature of the scalar potential.  In models with a single scalar field, this makes it difficult for the scalar field potential to form dark energy.  However, this is easily remedied by including a second scalar field.  A simple example is illustrated above in the hybrid MaVaN model.  A similar potential arises naturally in supersymmetric models, such as those in~\\cite{Fardon:2005wc} and~\\cite{Takahashi:2005kw}.  Another solution is to modify the theory so that the dark sector never reaches a state where the adiabatic condition applies.  Such models do not suffer from the instability described above.  One such theory is presented in~\\cite{Brookfield:2005td}.  Reducing the dependence of the acceleron on the heavy neutrino components also forms a class of possible solutions.  If the acceleron is decoupled from each heavy component before it becomes unstable, the acceleron expectation value is not quickly driven to a new scale.  The mass of the lighter neutrinos would be largely unaffected, avoiding instability."
    },
    "0606/astro-ph0606202_arXiv.txt": {
        "abstract": "We present the first observations of magnetic cataclysmic variables with the Spitzer Space Telescope. We used the Infrared Array Camera to obtain photometry of the polars EF Eri, GG Leo, V347 Pav, and RX J0154.0-5947 at 3.6, 4.5, 5.8, and 8.0 $\\mu$m.  In all of our targets, we detect excess mid-infrared emission over that expected from the component stars alone.  We explore the origin of this IR excess by examining bremsstrahlung, cyclotron emission, circumbinary dust, and L/T brown dwarf secondary stars. Bremsstrahlung and cyclotron emission appear unlikely to be significant contributors to the observed fluxes.  At present, the most likely candidate for the excess emission is dust that is probably located in a circumbinary disk with an inner temperature near 800 K. However, a simple dust disk plus any reasonable low mass or brown dwarf-like secondary star is unable to fully explain the observed flux densities in the 3--8 $\\mu$m region. ",
        "introduction": "Cataclysmic variables (CVs) are interacting binary stars containing a white dwarf (WD) primary and a low mass secondary (Warner 1995). CVs that contain a highly magnetic WD, with surface fields ranging from 10--250 MG, are called polars. In a polar, accreted material flows along magnetic field lines, does not form a viscous disk, and accretes directly onto the WD at/near its magnetic pole(s).  Polars are well known to stop and (re)start their mass transfer, causing changes of $\\sim$2--3 (or more) optical magnitudes at apparently random intervals. During low states, the accretion flux disappears and the two component stars can often be cleanly observed. For the past few decades, theoretical models of CV evolution (Kolb 1993; Howell et al.\\ 2001) have predicted that the very shortest orbital period systems should contain low mass, brown dwarf-like secondary stars. More recently, near-IR ($JHK$) observations of the shortest period CVs (e.g., Harrison et al.\\ 2003; Howell et al.\\ 2004) have provided evidence that they likely do contain secondaries similar to L or T dwarfs.  We have accomplished an initial survey of four polars (EF Eridani, GG Leonis, V347 Pavonis, and RX J0154.0-5947) using the Infrared Array Camera (IRAC) on the Spitzer Space Telescope. We specifically chose ultra-short period polars ($P_{\\rm orb} \\leq 90$ min) to facilitate studying the brown dwarf-like secondary stars predicted to be present. We found that all four of the polars show emission in the 3--8 $\\mu$m region in excess of that produced solely by the WD plus an M or L dwarf. In this Letter, we briefly discuss our observations and possible sources of this excess emission.  The best explanation is cool circumbinary dust that would likely dynamically settle into a disk, and/or possibly a T dwarf-like secondary star. ",
        "conclusions": "Figure 3 shows an IRAC two-color diagram for our polars, along with several brown dwarfs from the Patten et al.\\ sample (L6 = SDSS 1331-0116; T5 = 2MASS 0559-1404; T8 = 2MASS 0415-0935). We have also plotted GD1400, a detached WD+L5/6 binary with no dust (Farihi et al.\\ 2006), and our simple dust disk model (based on GD362), as shown in Figure 2. The effect of changing the GD362 model disk temperature from $T_{\\rm inner}=1200$ K to 800 K is shown in the lower right of the figure. For comparison, the IRAC colors of T Leonis and VY Aquarii are also plotted; these are not-yet-published CVs that we have observed with Spitzer. They have similar orbital periods and distances as our polars, but contain non-magnetic WDs surrounded by hot ($\\sim 8000$ K) accretion disks.  The SEDs of T Leo and VY Aqr are dominated by the RJ tails of the secondary star and accretion disk well into the IRAC region. Our observed polars occupy a region of color-space intermediate between WD+brown dwarf binaries without dust, single brown dwarfs, and dust disk-dominated systems.   \\begin{figure} \\includegraphics[angle=-90, scale=0.335]{spitzer_polars4.ps} \\caption{IRAC color-color diagram showing our four polars (red points with error bars). The WD+L5 green point is GD1400 (Farihi et al. 2006) and the two green points at 800K and 1200K, represent our scaled GD362 (WD+L8) dust disk model as described in \\S3.3. The white points are observed single L6, T5, and T8 dwarfs, and the two blue points are observed accretion disk dominated CVs (T Leo and VY Aqr). The F(4.5$-$5.8) value for the T8 dwarf is off scale at 1384 $\\mu$Jy as indicated by the arrow in the figure. See the text for details.} \\end{figure}  Bremsstrahlung, while often producing a flat spectral continuum, likely contributes essentially nothing to our observed IRAC fluxes. Cyclotron emission also appears unlikely to be a major contributor to the SEDs of our polars, since it should be present only as a low level falling RJ tail. However, our understanding of the complex nature of cyclotron emission is far from perfect, so we hesitate to say with complete certainty that none of the excess emission is due to the cyclotron process. At present, the best candidates for the origin of the observed IR excess are cool dust (in or near the CV) and/or secondary stars that are similar to late T dwarfs.  If the dust is produced by aeons of mass outflow from the CV (in winds or nova outbursts) or is possibly remnant material from the common envelope stage ejected with some orbital angular momentum (e.g., Hinkle et al.\\ 2006), then dynamic considerations lead us to expect that this material would settle into a circumbinary disk. If a dust disk is present and is similar to that seen around GD362, then we expect the inner disk temperature to be around 800 K. The emitting volume of such a disk must be large compared to the component stars (which are Earth-sized and Jupiter-sized), as it outshines even the cool brown dwarf-like secondary in the IRAC bands. If the 3--8 $\\mu$m SED originates primarily from the secondary star, then it would have to be similar to a late T dwarf and these CVs would have to be far closer than observations of their WDs suggest.   The likely presence of a circumbinary dust disk in polars naturally leads us to postulate that such disks might be common in other CVs. The 3--8 $\\mu$m region still contains a strong spectral contribution from the RJ tail of the accretion disk in CVs with weakly- or non-magnetic WDs. In longer period systems, the larger, brighter G, K, and early M secondary stars also contribute in the IRAC bandpass. Thus, the contrast from the weak IR emission of a dust disk against that produced by the stellar and other components in all but the polars may prevent detection shortward of 8 $\\mu$m. Some support for this is provided by Dubus et al.\\ (2004), who detected SS Cygni and AE Aquarii at 11 and 17 $\\mu$m, and by Abada-Simon et al.\\ (2005), who detected AE Aqr at wavelengths out to 170 $\\mu$m\\footnote{AE Aqr is a special case, however, as most of its far-IR emission is attributed to synchrotron emission.}.  Neither a simple dust disk, as described here, nor a brown dwarf-like secondary star alone can fully explain the observed IR excess in the polars, especially redward of $\\approx 6$ $\\mu$m. Spectral observations would help distinguish between the possibilities discussed above and provide details regarding the cause of the observed IR excess."
    },
    "0606/astro-ph0606728_arXiv.txt": {
        "abstract": "We present first results from a Keck/DEIMOS spectroscopic survey of red giant branch (RGB) stars in M33. The radial velocity distributions of the stars in our fields are well described by three Gaussian components, corresponding to a candidate halo component with an uncorrected radial velocity dispersion of $\\sigma \\simeq 50$\\,km\\,s$^{-1}$, a candidate disk component with a dispersion $\\sigma \\simeq 16$\\,km\\,s$^{-1}$, and a third component offset from the disk by $\\sim 50$\\,km\\,s$^{-1}$, but for which the dispersion is not well constrained. By comparing our data to a model of M33 based upon its HI rotation curve, we find that the stellar disk is offset in velocity by $\\sim 25$\\,km\\,s$^{-1}$ from the HI disk, consistent with the warping which exists between these components.  The spectroscopic metallicity of the halo component is [Fe/H] $\\simeq -1.5$, significantly more metal-poor than the implied metallicity of the disk population ([Fe/H] $\\simeq -0.9$), which also has a broader colour dispersion than the halo population. These data represent the first detections of individual stars in the halo of M33 and, despite being $\\sim 10$ times less massive than M31 or the Milky Way, all three of these disk galaxies have stellar halo components with a similar metallicity. The color distribution of the third component is different to the disk and the halo, but is similar to that expected for a single, coeval, stellar population, and could represent a stellar stream. More observations are required to determine the true nature of this intriguing third kinematic component in M33. ",
        "introduction": "M33, the Triangulum Galaxy ($1^h 33^m 51^s, 30^\\circ 39^\\prime 36^{\\prime\\prime}$), is the third brightest Local Group galaxy, with an integrated luminosity of $M_V = -18.9$. It is a late-type spiral, Sc II-III, with two open spiral arms and no clear evidence of any bulge component (\\citealt{bothun1992,minniti1993,mclean1996}). The overwhelming majority of light in M33 is distributed in an exponential disk component (\\citealt{devaucouleurs1959}), and the optical disk is tilted at nearly $30^\\circ$ to the strongly warped HI envelope (\\citealt{rogstad1976}). The mass-to-light ratio of its nucleus is small ($M/L < 0.4$), ruling out the presence of a supermassive black hole (\\citealt{kormendy1993,lauer1998}) like that implied to exist in the Milky Way (MW; \\citealt{genzel1997}) and M31 (\\citealt{dressler1988,kormendy1988}).  \\cite{chandar2002} analyse the kinematics of stars clusters in M33, and find two distinct kinematic populations, which they argue represent a disk and halo population, although individual stars belonging to a halo component in M33 had not been directly observed. In a companion paper by A. Ferguson et al. {\\it in preparation} (see also \\citealt{ferguson2006}), we present our wide field photometry of this galaxy taken with the Isaac Newton Telescope Wide Field Camera (INT~WFC), and show that a low level stellar component dominates the radial profile of M33 at large radius, implying the presence of an extended stellar halo in M33. Unlike its massive neighbor, M31, there is no evidence of substructure in M33 to a surface brightness threshold of $\\mu_V \\simeq 30$\\,mags\\,sq.arcsec.  In this {\\it Letter}, we present first results from a spectroscopic survey of red giant branch (RGB) stars in M33 using the DEep Imaging Multi-Object Spectrograph (DEIMOS) on Keck~II (\\citealt{faber2003}). Section~2 describes our observations. Section 3 analyzes the radial velocity distributions in our fields. Section 4 analyzes the metallicities of our target stars, and Section 5 discusses our results and concludes. We assume a distance to M33 of $809 \\pm 24$\\,kpc (\\citealt{mcconnachie2004a,mcconnachie2005a}). ",
        "conclusions": "\\label{final}  M33 is considered an example of the quintessential disk galaxy, with no evidence for any bulge or spheroid component. However, the radial profile for M33 presented in Ferguson et al. (2006), based upon our INT-WFC observations, reveals for the first time the presence of a spatially extended, low density component. Using Keck/DEIMOS, we have identified a population of RGB stars which have velocities centered on the systemic velocity of M33 and which appear to have a large dispersion ($\\sigma \\sim 50$\\,km\\,s$^{-1}$). These stars are bluer than other RGB stars in M33, with a small color dispersion. Their average metallicity is [Fe/H] $\\sim -1.5$, and we attribute them to the newly discovered halo component of M33. The haloes of the MW and M31 have similar metallicities to this (\\citealt{eggen1962,chapman2006,kalirai2006b}), even although both these galaxies are $\\sim 10$\\,times more massive than M33. It is important to emphasise that there is no evidence of a trend linking halo metallicity to galaxy mass or luminosity for the Local group galaxies, as has been suggested from recent Hubble Space Telescope studies of galaxies within $\\sim 10$\\,Mpc (\\citealt{mouhcine2005a,mouhcine2005b}). Instead, all the large Local Group spiral galaxies appear to have very similar halo metallicities.  We also observe stars in the disk of M33. Their $v_h$ distribution is well described by a Gaussian with an uncorrected dispersion of $\\sigma \\simeq 16$\\,km\\,s$^{-1}$, corresponding to a corrected dispersion of $\\simeq 12.5$\\,km\\,s$^{-1}$. These stars are redder and have a broader color dispersion than the halo stars. The stellar disk is offset in $v_h$ from the HI disk in our simple model by $\\sim 20 - 25$\\,km\\,s$^{-1}$, which is fully consistent with the extreme warping of the HI disk implied by \\cite{rogstad1976}.  Finally, we present evidence for a third kinematic feature in M33, revealed by the $v_h$ distribution of the southern fields. The velocity dispersion of this component is not well constrained by our data. A comparison to our simple model of the disk of M33 shows that this feature does not arise as a result of our line-of-sight through M33 sampling stars at different circular velocities in the main disk. The locus of these stars in the CMD has a remarkably small color dispersion, reminiscent of a single stellar population and inconsistent with belonging to a foreground population. We believe the most likely explanation for this feature is that it is a stellar stream. However, more spectroscopic observations of fields at different locations around the disk of M33 are required before the true nature of this intriguing feature is revealed."
    },
    "0606/astro-ph0606491_arXiv.txt": {
        "abstract": "We describe the results of fitting simple spherically symmetric models to the first overtone CO and AlO A--X (2--0) bands in V838~Mon. We find that the temperature and column of both CO and AlO systematically decline over the period 2002 October -- 2005 February and that an additional, hotter and denser, component is present from 2005 January. We also describe the results of an observation of the star and its light echo at 850\\micron. We do not detect the `infrared' echo at these wavelength, and place an upper limit of $3\\times10^7$~cm$^{-2}$ on the column of grains. ",
        "introduction": "We began a programme of regular infrared (1--5\\micron) spectroscopic monitoring of V838~Mon soon after the 2002 eruption was reported by \\cite{brown}; the results are described in \\cite{evans,rushton-mega,rushton-sio,phd}. All of the data reported here cover the first overtone CO bands (bandhead at 2.29\\micron) and the A--X (2--0) AlO bands at 1.67\\micron. In this contribution we discuss our attempts to model the complex evolution of these bands, as revealed by our low resolution obervations of CO and AlO.  Details of the infrared data we obtained at high resolution, and what the data tell us about the complex velocity systems in the environment of V838~Mon, is discussed by Geballe elsewhere in these proceedings.  We also describe here the results of an observation at 850\\micron\\ with the James Clerk Maxwell Telescope (JCMT). ",
        "conclusions": "We have fitted a spherically symmetric, expanding shell of CO and AlO to the near-infrared spectra of V838~Mon. We find a declining column of CO and AlO, and a declining $T_{\\rm ex}$ for both species. However there is an additional hot CO component from 2005 Jan.  We find an upper limit of $3\\times10^7$~cm$^{-2}$, and corresponding upper mass limit of $0.1$~M$_\\odot$, for the dust in light echo material; this is significantly lower than that implied by the Spitzer observations.  Further details of this work will be published elsewhere."
    },
    "0606/astro-ph0606172_arXiv.txt": {
        "abstract": "Using cosmological hydrodynamical simulations, we investigate the prospects of the thermal Sunyaev-Zel'dovich (tSZ) effect to detect the missing baryons in the local universe. We find that at least 80\\% of the tSZ luminosity is generated in collapsed structures, and that $\\sim$ 70\\% of the remaining diffuse tSZ luminosity (i.e., $\\sim 15$\\% of the total) comes from overdense regions with $\\delta_{gas}>$10, such as filaments and superclusters. The gas present in slightly overdense and underdense regions with $\\delta_{gas} < 10$, despite making up 50\\% of the total baryon budget, leaves very little tSZ signature: it gives rise to only $\\sim$ 5\\% of the total tSZ luminosity. Thus, future Cosmic Microwave Background (CMB) observations will be sensitive to, at best, one half of the missing baryons, improving the current observational status, but still leaving one half unobserved.  Since most of the tSZ is generated in haloes, we find a tight correlation between gas pressure and galaxy number density. This allows us to predict the CMB Comptonization from existing galaxy surveys and to forecast the tSZ effect from the local structures probed by the Two Micron All Sky Survey (2MASS) galaxy catalog. ",
        "introduction": "Recent cosmological observations like galaxy surveys (e.g.,2dFGRS, \\citet{twodf}, SDSS, \\citet{sdss}), Supernovae (e.g., SNLS, \\citet{SNLS}) or Cosmic Microwave Background (CMB) (e.g., WMAP, \\citet{wmap1,wmap2}) give strong support to the LCDM model of the universe, in which only $\\sim$ 18\\% of matter is in the form of baryons.  Baryons are mostly probed by CMB, Lyman--$\\alpha$ forest (e.g., \\citet{croft,McDonald}), and X-rays observations. The first two probes show that the Universe, at much younger epochs, contained an amount of baryons that exceeds by a factor of $\\sim$ 9 the baryons seen in today's universe, \\citep{fukpeebles}. This is the missing baryons problem. \\begin{figure*} \\centering \\plotancho{./f1.eps} \\caption[fig:fig1]{ {\\it (a)} Relative contribution to the total tSZ luminosity of regions with gas density contrast bigger than $\\delta_{gas}$. The vertical solid line indicates the average gas overdensity in haloes. {\\it (b)} Relative contribution to the total baryon mass of regions with gas density contrast bigger than $\\delta_{gas}$. The vertical solid line indicates the average gas overdensity in haloes. Note the asymmetry with respect to the previous (left hand side) panel.  {\\it (c)} Ratio of gas pressure produced in galaxy groups (small haloes) over pressure produced by diffuse gas {\\em and} galaxy groups, versus total pressure (including galaxy clusters), averaged on scales of 12 h$^{-1}$Mpc. The solid vertical line corresponds to regions with average matter density. In overdense regions, the pressure of haloes always dominates the diffuse pressure. } \\label{fig:fig1} \\end{figure*} Indeed, according to \\citet{cenost1,cenost2}, about half of the baryons should be in the form of a Warm Hot Intergalactic Medium (WHIM), at relatively high temperature ($\\sim 10^5-10^7$ K) but low densities. Only a small fraction of the total gas is hot enough ($^>_{\\sim} 10^8$K) to enable X-ray detection via {\\it bremmstrahlung}: this mechanism led in the nineties to the identification of massive galaxy clusters from ROSAT X-ray observations (e.g., \\citet{rosat}), and more recently has motivated the search (and tentative detections) of diffuse X-ray emission from filaments and superclusters, \\citep{xr1,xr2}. Although further pieces of evidence for local gas have been found in quasar spectra, in the form of Oxygen absorption features at $z\\simeq 0$ \\citep{qspec1,qspec2}, it is clear that we do not have a complete map of the baryon distribution in the local Universe.  With the advent of high resolution and high sensitivity CMB experiments it has been proposed to trace the local gas using the thermal Sunyaev-Zel'dovich effect (tSZ, \\citet{tSZ}). The tSZ effect describes the Compton scattering of CMB photons off hot electrons, in which a net transfer of energy from electrons to radiation distorts the Black Body spectrum of the CMB. The relative % intensity fluctuations introduced in this scattering are redshift independent, and their amplitude is proportional to the integrated electron pressure along the line of sight.  Observational efforts by \\citet{tenef1} and \\citet{mito} have detected strong decrements in Corona Borealis, and interpreted them as Compton scattering being generated by extended gas in a filament aligned with the line of sight. Other studies \\citep{hansen05, dolag05, atrio06} addressed this problem from a theoretical point of view, trying to characterize the tSZ signal generated by non-collapsed baryons.  Here we use a hydrodynamical numerical simulation to explore the prospects of future tSZ observations to search for the missing baryons.  We find that, as the tSZ effect is sensitive to gas pressure, most of the tSZ luminosity (around 80\\%) is generated in collapsed structures, which host only $\\sim$30\\% of the total baryon mass.  Most baryons are in environments with low overdensities, giving rise to very little tSZ signal: more than 60\\% of baryons are in regions with $\\delta_{gas} < 20$, which contribute only $\\sim$10\\% of the total tSZ luminosity. More than two thirds of the diffuse tSZ luminosity (not generated in collapsed structures) is coming from overdense environments with $\\delta_{gas} > 10$ such as filaments or superclusters. The significant amount of baryons present in underdense regions ($\\sim$ 18\\% where $\\delta_{gas} <0$) is practically invisible to the tSZ effect ($ \\sim 0.5$ \\% contribution to tSZ luminosity). While the tSZ will be, at best, sensitive to $ \\delta_{\\rm gas} > 10$ (i.e. superclusters and filaments), which corresponds to $\\sim 50\\%$ of all baryons, X-rays are only sensitive to the core centers of galaxy clusters in our neighborhood which corresponds to less than 5\\% of all baryons, \\citep{fukpeebles}.  Because in overdense regions most of the tSZ is generated in haloes, we find a tight correlation between the gas pressure giving rise to the tSZ and the number galaxy density field, $p_e \\propto n_{gal}^{2.2}$. This allows us to build tSZ templates and to make predictions for the tSZ signal from existing galaxy surveys. ",
        "conclusions": ""
    },
    "0606/astro-ph0606344_arXiv.txt": {
        "abstract": "We discuss the panchromatic properties of 99,088 galaxies selected from the Sloan Digital Sky Survey Data Release 1 ``main'' spectroscopic sample (a flux-limited sample for 1360 deg$^2$). These galaxies are positionally matched to sources detected by ROSAT, GALEX, 2MASS, IRAS, GB6, FIRST, NVSS and WENSS. The matching fraction varies from $<1$\\% for ROSAT and GB6 to $\\sim$40\\% for GALEX and 2MASS. In addition to its size, the advantages of this sample are well controlled selection effects, faint flux limits and the wealth of measured parameters, including accurate X-ray to radio photometry, angular sizes, and optical spectra. We find strong correlations between the detection fraction at other wavelengths and optical properties such as flux, colors, and emission-line strengths. For example, $\\sim$2/3 of SDSS ``main'' galaxies classified as AGN using emission-line strengths are detected by 2MASS, while the corresponding fraction for star-forming galaxies is only $\\sim$1/10. Similarly, over 90\\% of galaxies detected by IRAS display strong emission lines in their optical spectra, compared to $\\sim$50\\% for the whole SDSS sample. Using GALEX, SDSS, and 2MASS data, we construct the UV-IR broad-band spectral energy distributions for various types of galaxies, and find that they form a nearly one-parameter family. For example, the SDSS $u$- and $r$- band data, supplemented with redshift, can be used to ``predict'' $K$-band magnitudes measured by 2MASS with an rms scatter of only 0.2 mag. When a dust content estimate determined from SDSS spectra with the aid of models is also utilized, this scatter decreases to 0.1 mag and can be fully accounted for by measurement uncertainties. We demonstrate that this interstellar dust content, inferred from optical SDSS spectra by Kauffmann et al. (2003a), is indeed higher for galaxies detected by IRAS and that it can be used to ``predict'' measured IRAS 60 $\\mu$m flux density within a factor of two using only SDSS data. We also show that the position of a galaxy in the emission-line-based Baldwin-Phillips-Terlevich diagram is correlated with the optical light concentration index and $u-r$ color determined from the SDSS broad-band imaging data, and discuss changes in the morphology of this diagram induced by requiring detections at other wavelengths. Notably, we find that SDSS ``main'' galaxies detected by GALEX include a non-negligible fraction (10-30\\%) of AGNs, and hence do not represent a clean sample of starburst galaxies. We study the IR-radio correlation and find evidence that its slope may be different for AGN and star-forming galaxies and related to the $H_\\alpha/H_\\beta$ line strength ratio. ",
        "introduction": "The study of global galaxy properties has been recently invigorated by modern sensitive large-area surveys across a wide wavelength range. The Sloan Digital Sky Survey (SDSS, York et al. 2000, for more details see Appendix A1) stands out because it has already provided near-UV to near-IR five-color imaging data and high-quality spectra (R$\\sim$1800) for over 500,000 galaxies. The ``main'' spectroscopic galaxy sample is defined by a simple r-band flux limit (Strauss et al. 2002), and will include close to 1,000,000 galaxies.  A number of detailed galaxy studies based on SDSS data have already been published. Strateva et al. (2001) and Shimasaku et al. (2001) demonstrated a tight correlation between the $u-r$ color, concentration of the galaxy's light profile, and morphology. Blanton et al. (2003) presented the SDSS galaxy luminosity function, and Kauffmann et al. (2003ab) determined and analyzed stellar masses and star-formation histories for 100,000 SDSS galaxies.  In addition to ``stand-alone'' studies based on only SDSS data, SDSS can be used as a cornerstone for panchromatic studies of galaxies aided by recent surveys at wavelengths outside the optical range (0.3-1 \\mic). The special role of SDSS in such studies is due to its rich optical information, in particular high-quality spectra and photometry. Nevertheless, galaxies emit a substantial fraction of their bolometric flux outside the wavelength range accessible to SDSS. For example, in starburst and Seyfert 2 galaxies the mid/far-IR wavelength range is the most important contributor to the bolometric flux, e.g. Schmitt et al. (1997). Information obtained by other surveys offers important observational constraints for models of galaxy formation and evolution.  Numerous studies that utilize SDSS and surveys at other wavelengths have already been published. For example, Finlator et al. (2000) analyzed the properties of point sources detected by SDSS and Two Micron All Sky Survey (2MASS), and Ivezi\\'{c} et al. (2001a) discussed the colors and counts of SDSS sources detected by SDSS, 2MASS, and FIRST surveys. Ivezi\\'{c} et al. (2002) cross-correlated SDSS and the survey Faint Images of the Radio Sky at Twenty Centimeters (FIRST), and analyzed the optical and radio properties of quasars and galaxies. Bell et al. (2003) used SDSS and 2MASS data to estimate the baryonic mass functions of galaxies, and Anderson et al. (2003) studied the properties of AGN galaxies detected by SDSS and ROSAT. Best et al. (2005ab) studied radio galaxies, Chang et al. (2005) analyzed the SDSS-2MASS colors of elliptical galaxies, and Goto (2005) and Pasquali, Kauffmann \\& Heckman (2005) studied the optical properties of SDSS galaxies detected by IRAS. A detailed analysis of rest-frame colors in the Str\\\"{o}mgren system synthesized from SDSS spectra was presented by Smol\\v{c}i\\'{c} et al. (2006). They found that the galaxy distribution in the resulting color-color diagrams forms a very narrow locus with a width of only 0.03 mag. This finding agrees well with the conclusion by Yip et al. (2004), based on a principal component analysis of SDSS spectra, that galaxy spectra can be described by a small number of eigenspectra.  Here we cross-correlate the catalog of galaxies from SDSS Data Release 1 (Abazajian et al. 2003) with catalogs of sources detected by ROSAT (X ray), GALEX (UV), 2MASS (near-IR), IRAS (mid/far-IR), GB6 (6 cm), FIRST (20 cm), NVSS (20 cm), and WENSS (92 cm). References and a description of each survey are listed in Appendix A. The panchromatic galaxy samples discussed here are $\\sim$10-100 times larger than those used in older pre-SDSS studies. In addition, they are selected by simple flux limits, and benefit from a wealth of accurately measured parameters including X-ray to radio photometry, angular sizes, and optical spectra. The main aim of this paper is to quantify the fraction and basic properties of SDSS ``main'' galaxies detected by other surveys using a uniform approach for all analyzed surveys. However, due to the size and quality of the resulting samples, even a simple, preliminary analysis presented here is sufficient to yield a wealth of additional results.  We describe our matching and analysis methods in Section 2. In Section 3 we discuss the detection fraction of SDSS galaxies by other surveys, and in Section 4 we present a preliminary analysis of some panchromatic properties of galaxies in our sample. We discuss and summarize our results in Section 5. ",
        "conclusions": "GALEX, SDSS, 2MASS, and IRAS photometry, as well as the interstellar dust content inferred from SDSS optical spectra with  the aid of sophisticated models.  A large sample of galaxies that have SDSS spectra and are detected by IRAS and NVSS allowed us to study the IR-radio correlation separately for star-forming and AGN galaxies. We confirm that both galaxy types follow a tight correlation, and find that a large fraction (80\\%) of optically classified SDSS-NVSS AGN galaxies show significantly more radio emission than expected from their IR flux (technically, we show that IR fluxes predicted from observed radio emission for galaxies that are not detected by IRAS are higher than the corresponding IRAS upper flux limits, see Section~\\ref{radioIR}). We also find marginal evidence for different slopes of IR-radio correlation for AGN and star-forming subsamples, an effect that seems to be related to the $H_\\alpha/H_\\beta$ line strength ratio.   Perhaps the most important conclusion of our study is that little more than a single datum can be learned about galaxies from photometric data whose accuracy is not demonstrably better than $\\sim$0.1 mag. Fortunately, all three major modern galaxy surveys, GALEX, SDSS, and 2MASS, appear to have achieved this goal and, together with surveys such as FIRST and NVSS, opened unprecedented opportunities for detailed studies of galaxies."
    },
    "0606/astro-ph0606303_arXiv.txt": {
        "abstract": "{  In this paper we use neutral hydrogen (\\HI) and optical spectroscopic observations to compare the timescales of a merger event, starburst episode and radio-AGN activity in the radio galaxy \\object{B2 0648+27}. We detect a large ring-like structure of \\HI\\ in emission around the early-type host galaxy of \\object{B2 0648+27} ($M_{\\rm \\HI} = 8.5 \\times 10^{9} M_{\\odot}$, diameter = 190 kpc). We interpret this as the result of a major merger that occurred $\\gtrsim$ 1.5 Gyr ago. From modelling optical long-slit spectra we find that a young stellar population of 0.3 Gyr, indicative of a past starburst event, dominates the stellar light throughout the galaxy. The off-set in time between the merger event and the starburst activity in \\object{B2 0648+27} suggests that the starburst was triggered in an advanced stage of the merger, which can be explained if the gas-rich progenitor galaxies contained a bulge. Although the exact age of the radio source remains uncertain, there appears to be a significant time-delay between the merger/starburst event and the current episode of radio-AGN activity. We also observe an outflow of emission-line gas in this system, which is likely related to superwinds driven by the stars that formed during the starburst event. We argue that the radio galaxy \\object{B2 0648+27} is a link in the evolutionary sequence between Ultra-Luminous Infrared Galaxies (ULIRGs) and genuine early-type galaxies. ",
        "introduction": "\\label{sec:intro}  Merger events are often invoked as the trigger of the activity in galaxies. Extreme examples are major merger systems such as Ultra-Luminous Infra-Red Galaxies (ULIRGs), which owe their extreme infra-red colors to a massive starburst, often in combination with a dust-enshrouded AGN. In the hierarchical model of galaxy formation, early-type galaxies such as E's and S0's form the end products of merging systems. The overwhelming majority of bright, low-$z$ radio sources are hosted by these early-type galaxies, which often still show optical signatures of the merger event \\citep[like optical tails, bridges, shells;][]{smi89:auth}. In addition, a growing number of radio galaxies is found to contain a young or intermediate age stellar population, indicating that they are post-starbust systems \\citep[e.g.][]{are01:auth,wil02:auth,wil04:auth,tad02:auth,tad05:auth,rai05:auth}. The same holds for optically selected AGN from the Sloan Digital Sky Survey (SDSS), of which a significant fraction experienced bursts of star formation in the recent past \\citep{kau03:auth}.  A connection between mergers and galactic scale starburst events has been well established and modelled \\citep[e.g.][]{bar91:auth,bar96:auth,mih94:auth,mih96:auth,spr05:auth,kap05:auth}. Although the models of \\citet{spr05:auth} and \\citet{kap05:auth} predict a connection also between mergers and AGN activity, observationally there remain considerable uncertainties about this, and in particular about the timing of the events. While some studies do find trends between merger/interaction events and AGN activity \\citep[e.g.][]{can01:auth,wu98:auth,hec86:auth}, others find no such trends  \\cite[e.g.][]{gro05:auth,dun03:auth,lut98:auth}. To investigate this further it is worth looking at the ``order-of-events'' in individual nearby galaxies that show signs of both merger and starburst activity as well as AGN activity.  An excellent object to do this in detail is the nearby ($z$ = 0.0412)\\footnote{$H_{\\circ} = 71$ km~s$^{-1}$~Mpc$^{-1}$ used throughout this paper. This puts \\object{B2 0648+27} at a distance of 174 Mpc and 1 arcsec = 0.84 kpc.} radio galaxy \\object{B2 0648+27}.\\footnote{In this paper we use the name \\object{B2 0648+27} for both the radio source as well as the galaxy that hosts the source.} This galaxy contains a compact radio AGN (log $P_{\\rm 1.4 GHz}$ = 23.75 W Hz$^{-1}$). The early-type host galaxy has an elliptical morphology, but deep optical colour images reveal a low surface brightness envelope and faint plume- or tail-like structures \\citep{hei94:auth}. The {\\sl HST} image from \\citet{cap00:auth} shows a patchy distribution of dust, as well as what appear to be regions of star formation. \\citet[][hereafter \\paper]{mor03:auth} found a ring-like structure of \\HI\\ gas that surrounds the host-galaxy, suggesting a major merger happened in this system. In this paper we present the results from new \\HI\\ data (Sect. \\ref{sec:hi}) as well as a stellar population analysis from new optical spectra (Sect. \\ref{sec:spectra}). We use these to study the timescales between the merger event, starburst activity and onset of radio-AGN activity in \\object{B2 0648+27} and to determine the evolutionary stage of this galaxy. ",
        "conclusions": "\\label{sec:conclusions}  We detect 8.5$\\times$10$^{9}$ $M_{\\odot}$ of \\HI\\ in a large ring-like structure of 190 kpc around the radio galaxy \\object{B2 0648+27}. We also find that the light from the host galaxy is dominated by a post-starburst stellar population. The extended \\HI-structure and post-starburst stellar population are the result of a major merger event. There appears to be a significant time-delay between the merger/starburst event and the current episode of radio-AGN activity.  From the derived properties of the merger and starburst event we conclude that \\object{B2 0648+27} represents an important link in the evolutionary sequence between ULIRGs and normal early-type galaxies. More research is important to further investigate the role of the AGN activity in this respect. The relative proximity of \\object{B2 0648+27} allows a detailed study of the physical processes that occur in this system, which is important for high-$z$ studies, where major mergers are much more common."
    },
    "0606/astro-ph0606629_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are sources of energetic, highly variable fluxes of $\\gamma$ rays, which demonstrates that they are powerful particle accelerators. Besides relativistic electrons, GRBs should also accelerate high-energy hadrons, some of which could escape cooling to produce ultra-high energy cosmic rays (UHECRs).  Acceleration of high-energy hadrons in GRB blast waves will be established if high-energy neutrinos produced through photopion interactions in the blast wave are detected from GRBs.  Limitations on the energy in nonthermal hadrons and the number of expected neutrinos are imposed by the fluxes from pair-photon cascades initiated in the same processes that produce neutrinos.  Only the most powerful bursts at fluence levels $\\gtrsim 3\\times 10^{-4}\\,\\rm erg \\ cm^{-2}$ offer a realistic prospect for detection of $\\gg$ TeV neutrinos.  Detection of high-energy neutrinos is likely if GRB blast waves have large baryon loads and Doppler factors $\\lesssim 200$. Cascade $\\gamma$ rays will accompany neutrino production and might already have been detected as anomalous emission components in the spectra of some GRBs.  Prospects for detection of GRBs in the Milky Way are also considered. ",
        "introduction": "There is general consensus that acceleration of cosmic rays by supernova remnants (SNRs) is the main source of galactic cosmic rays at energies below $\\sim 100$ TeV (e.\\ g., \\cite{drury}). It is also generally thought that all cosmic rays with energies up to at least the second knee in the cosmic-ray spectrum at $E_2\\sim 3\\times 10^{17}\\,\\rm eV $ (e.g. \\cite{som01}) or even up to the ankle at $E_{\\rm ank} \\simeq 3\\times 10^{18}$ eV are produced in our Galaxy (see, e.g., \\cite{nw00}).  Meanwhile, cosmic-ray acceleration to energies significantly exceeding $\\approx 10^{14}$ eV with the conventional mechanism of nonrelativistic first-order shock acceleration by SNRs from typical (Type Ia and II) supernovae (SNe) is problematic \\cite{lc83}. The origin of the knee in the cosmic-ray spectrum in the form of a spectral-index break in the power-law all-particle spectrum by $\\approx 0.3$ units at $E_1 \\simeq 3\\times 10^{15}\\,\\rm eV$, accompanied by a change in the cosmic-ray composition, seems to suggest a new contribution to cosmic rays at these energies.  GRBs have been proposed as effective accelerators of cosmic rays in the universe \\cite{vie95} and as probable sources of cosmic rays up to ultra-high energies in our Galaxy \\cite{der02}.  Cosmic rays with energies below $\\sim 100$ TeV are produced in the conventional scenario of continuous injection due to nonrelativistic shock acceleration by SNRs formed in all types of SNe, with subsequent modification of the source spectrum through energy-dependent propagation (see \\cite{drury,ber90}).  GRBs in our own Galaxy are very rare and probably take place at a rate $\\lesssim 10^{-4}$ yr$^{-1}$. Our best opportunity to establish cosmic-ray acceleration in GRB blast waves is by looking for hadronic emission signatures in the spectra of GRBs, or characteristic signatures from remnants of past GRBs. Conclusive evidence for cosmic-ray acceleration in GRBs will come from the detection of high-energy neutrinos.  In the relativistic blast-wave model for GRBs, photohadronic production is the most important mechanism for producing high-energy neutrinos. In Section 2, the photohadron process is described and an approximation used for computing the spectra of secondaries formed in photohadronic processes, which is also useful for analytical studies, is given. Calculations of secondary neutrino and cascade gamma-ray production are presented in Section 3, and the requirements for detection of high-energy neutrinos from GRBs are derived. Models for UHECRs from GRBs are described in Section 4. A discussion of signatures of GRBs in the Galaxy is found in Section 5, where we briefly consider recent reports that GRBs favor sites of low metallicity.  Summary and conclusions are given in Section 6. ",
        "conclusions": "The possibility that GRBs accelerate relativistic hadrons could solve the problem of the origin of the high-energy cosmic rays. There are several ways to establish whether GRBs accelerate cosmic rays, from:  high-energy neutrino detection from GRBs; anomalous radiation signatures associated with hadronic acceleration and energy losses in GRB blast waves; radiations made by neutrons that escape from the blast wave; and evidence for cosmic-ray production from GRB remnants in our own and nearby galaxies.  In the case of neutrino and $\\gamma$-ray production, at most only a few high-energy $\\nu_\\mu$ can be detected with km-scale neutrino detectors even from bright GRBs at the fluence level $\\Phi_{tot} \\gtrsim 3\\times 10^{-4} \\,\\rm erg \\; cm^{-2}$, and only when the baryon loading is high \\cite{da03,wda04}. This is because the detection of a single $\\nu_\\mu$ requires a $\\nu_\\mu$ fluence $ \\gtrsim 10^{-4} \\,\\rm erg \\; cm^{-2}$ above 1 TeV. Since the energy release in high-energy neutrinos and electromagnetic secondaries is about equal, this energy will be reprocessed in the pair-photon cascade and emerge in the form of observable radiation at $\\gamma$-ray energies, and this radiation cannot exceed the measured fluence in this regime. Detection of high-energy neutrinos and hadronic $\\gamma$-ray emission components from GRBs will mean that GRBs have a large load in relativistic baryons, which is required for UHECR production.  The advent of the IceCube neutrino telescope at the South Pole and the launch of GLAST in late 2007 will provide the important instruments to establish whether GRBs accelerate ultra-relativistic hadrons. In the meantime, searches for GRB remnants at other wavelengths will indicate the importance of GRBs in the local universes. With these developments, the longstanding puzzle of the origin of the cosmic rays---now nearly 100 years since their discovery---should finally be answered.  \\vskip0.2in \\noindent We thank A.\\ Reimer for permission to reproduce Fig.\\ 1 from Ref.\\ \\cite{muc99}, and for useful comments on the manuscript. The work of C.\\ D.\\ D.\\ is supported by the Office of Naval Research. We also acknowledge support of GLAST Science Investigation DPR-S-1563-Y."
    },
    "0606/astro-ph0606135_arXiv.txt": {
        "abstract": "As part of an ongoing program to map better the empirical instability strip of pulsating ZZ Ceti white dwarfs, we present a brief progress report based on our last observing season. We discuss here high-speed photometric measurements for 6 new pulsators. These stars were selected on the basis of preliminary measurements of their effective temperature and surface gravity that placed them inside or near the known ZZ Ceti instability strip. We also report detection limits for a number of DA white dwarfs that showed no sign of variability. Finally, we revisit the ZZ Ceti star G232-38 for which we obtained improved high-speed photometry. ",
        "introduction": "The ZZ Ceti stars are pulsating white dwarfs whose optical spectra are dominated by hydrogen lines (DA stars). They are found in a very well defined region in the $\\Te$-$\\logg$ plane, known as the ZZ Ceti instability strip. They exhibit pulsation periods in the range 100-1400 s, corresponding to low degree gravity-mode oscillations. The detected pulsation modes have amplitudes that are found in the range from a few millimagnitudes in the low-amplitude pulsators to a fraction of a magnitude in the largest amplitude ones. The ZZ Ceti stars show some remarkable trends, although there are a few exceptions. For instance, the cooler pulsators tend to show longer periods and larger amplitudes. Also, for a given effective temperature, the observed periods tend to be longer for lower gravity objects.  Since the early 90's, our group in Montr\\'eal has been carrying out a systematic study aimed at defining the empirical boundaries of the ZZ Ceti instability strip. We use quantitative time-averaged optical spectroscopy to pin down the locations of candidate stars in the $\\Te$-$\\logg$ diagram, and we follow up on these candidates with ``white light'' fast photometry to determine if a target pulsates or not. The determination of these boundaries is essential if we are to understand better the ZZ Ceti phenomenon and to establish on a more quantitative footing the trends already alluded to above. The question of the purity of the instability strip (i.e., the absence of photometrically constant stars within) is also an important issue, and this for several reasons. First, a pure strip implies that the ZZ Ceti phenomenon is an evolutionary phase through which all DA white dwarfs must pass. It is then possible to apply what is gleaned from asteroseismological studies of these stars, such as information on their internal structure, to the entire class of DA white dwarfs. Second, knowing that the strip is pure allows us to predict the variability of a star simply through a measure of its atmospheric parameters, $\\Te$ and $\\logg$.  The paper of \\citet{bergeron95} presented, for the first time, a complete and homogeneous picture of the ZZ Ceti stars in terms of their time-averaged properties. It was based on the sample of 22 pulsators known at the time. In the meantime, we pursued our survey to improve statistics, and we reported several discoveries and new findings in a series of papers, the last of which is \\citet{gianninas05} where references to our previous work can be found. One of the key results of \\citet{gianninas05}, based on our homogeneous spectroscopic sample of 39 pulsators and 121 constant stars, is the demonstration that the ZZ Ceti strip is indeed pure. Furthermore, that paper incorporated the first results of an extended ongoing survey whose primary aim is to define more accurately the empirical boundaries of the instability strip. That survey is based on spectroscopic measurements of a subsample of DA white dwarfs from the Catalog of Spectroscopically Identified White Dwarfs of \\citet{mccook99}. We thus obtain optical spectra for each star and then fit the Balmer line profiles using a grid of synthetic spectra generated from detailed model atmospheres. This allows us to accurately measure $\\Te$ and $\\logg$ for each star. It is then possible to identify white dwarfs whose atmospheric parameters place them near or within the ZZ Ceti instability strip. In this paper, we report on the latest results of this ongoing survey after the 2005-2006 observing season.  \\begin{table*}[!t] \\caption{Journal of Observations} \\begin{center} \\begin{tabular}{clc@{}cccc} \\hline \\hline &      & \\multicolumn{2}{c}{Run}    & Date & Start Time & Total Number of \\\\ WD & Name & \\multicolumn{2}{c}{Number} &  UT  &     UT     &   Data Points   \\\\ \\hline 0016$-$258 & MCT 0016$-$2553 & kit-013 & & 2005 Oct 27 & 05:43 &  867 \\\\ 0036+312   & G132-12         & kit-005 & & 2005 Oct 25 & 05:25 & 1042 \\\\ &                 & kit-017 & & 2005 Oct 28 & 04:42 & 1348 \\\\ 0136+768   & GD 420          & kit-009 & & 2005 Oct 26 & 05:55 &  410 \\\\ 0145+234   & MK 362          & kit-002 & & 2005 Oct 24 & 07:28 &  736 \\\\ 0302+621   & GD 426          & kit-010 & & 2005 Oct 26 & 07:18 &  532 \\\\ 0344+073   & KUV 03442+0719  & kit-003 & & 2005 Oct 24 & 10:04 &  899 \\\\ &                 & kit-011 & & 2005 Oct 26 & 09:07 & 1188 \\\\ &                 & kit-015 & & 2005 Oct 27 & 11:25 &  366 \\\\ 0347$-$137 & GD 51           & kit-006 & & 2005 Oct 25 & 08:49 &  505 \\\\ 0416+701   & GD 429          & kit-014 & & 2005 Oct 27 & 08:54 &  801 \\\\ 0741+248   & LP 366-3        & kit-007 & & 2005 Oct 25 & 10:31 &  737 \\\\ 1150$-$153 & EC 11507$-$1519 & kit-021 & & 2006 Mar 23 & 08:47 &  723 \\\\ 1211+320   & CBS 54          & kit-018 & & 2006 Mar 23 & 03:43 &  582 \\\\ 1503$-$093 & EC 15036$-$0918 & kit-022 & & 2006 Mar 23 & 11:11 &  439 \\\\ 1820+709   & GD 530          & cfh-111 & & 2005 May 01 & 13:42 &  328 \\\\ &                 & cfh-119 & $^{\\rm a}$ & 2005 Aug 14 & 05:52 &  468 \\\\ &                 & kit-012 & & 2005 Oct 27 & 02:21 & 1048 \\\\ &                 & kit-016 & & 2005 Oct 28 & 01:54 &  918 \\\\ 2148+539   & G232-38         & kit-004 & & 2005 Oct 25 & 02:04 & 1082 \\\\ 2148$-$291 & MCT 2148$-$2911 & cfh-117 & $^{\\rm a}$ & 2005 Aug 13 & 07:33 &  382 \\\\ 2311+552   & GD 556          & kit-001 & & 2005 Oct 24 & 02:17 & 1724 \\\\ 2336$-$079 & GD 1212         & kit-008 & & 2005 Oct 26 & 02:23 & 1083 \\\\ \\hline \\multicolumn{7}{l}{$^{\\rm a}$ Observed through a B-band filter.} \\\\ \\multicolumn{7}{l}{Notes. -- The sampling time is 10 s for all observations.} \\\\ \\end{tabular} \\end{center} \\end{table*}  We note that with the recent releases of data sets from the SDSS, the number of known ZZ Ceti stars has grown in a spectacular way, by more than a factor of 2.5 \\citep[see, e.g.,][]{mukadam04b,mullally05,kepler05,castanheira06}. The vast majority of the new pulsators, however, are significantly fainter than those previously known and, therefore, have a much less interesting asteroseismological potential. Moreover, the relatively low S/N ratio spectroscopic observations reached in the SDSS limits the usefulness of statistical studies based on them. Of particular interest here, \\citet{mukadam04a} claimed that the ZZ Ceti instability strip contains several nonvariable stars on the basis of the SDSS sample of \\citet{mukadam04b}. We conclusively showed in \\citet{gianninas05}, however, that this incorrect result could be explained in terms of insufficient S/N ratio in the SDSS spectra. Thus, we believe that improvements in our empirical description of the ZZ Ceti instability strip necessarily rests on high S/N ($>$ 80) time-averaged spectroscopy such as in the survey we are currently pursuing. This limits the potential targets to relatively bright stars ($V < 17$), but as we are finding out as well as others \\citep[see, e.g.,][]{silvotti05,voss06}, there are still bright ZZ Ceti pulsators to be discovered. ",
        "conclusions": ""
    },
    "0606/astro-ph0606245_arXiv.txt": {
        "abstract": "{} {We present the results of spectroscopic-- and orbit--sampled photometric observations of the faint UV-excess object PG~2200+085.} {The optical CCD photometry observations of this object were performed by the Russian-Turkish 1.5-meter telescope RTT150 at the TUBITAK National Observatory (Turkey). The long-slit optical spectroscopy observations with 2.6 \\AA ~resolution were carried out by 6-meter telescope BTA at the Special Astrophysical Observatory (Russia).} {The photometric variations over two nights are almost sinusoidal with an amplitude $\\Delta m_V = 0.^{m}04$ and a period of $P = 0.3186$d. Such a light curve is typical of a detached close binary with an illumination effect or the ellipsoidal deformation of a secondary star. The observed spectrum clearly displays a featureless blue continuum of a hot component and a rich absorption--line and molecular band K--star spectrum. The CaII line profiles with strong emission cores are remarkably similar to those of V471 Tau.} {We tentatively classify PG~2200+085 as a pre-cataclysmic binary of the V471 Tau type.} ",
        "introduction": "PG~2200+085 = USNO-A2.0 0975-20754051 ($\\alpha_{2000}=22^{h}03^{m}19^{s}.72$, $\\delta_{2000}=08^{\\circ}45'36''.5$) is one of the faint ultraviolet sources, $m_V =14.^{m}21$ included in the Palomar-Green (PG) catalogue  \\cite{gr}.  A finding chart for the object is shown in Fig. 1. The spectrum of this star was classified as composite, and its colour index indicates radiation excess in the red spectral region.  Later, Schultz et al. \\cite{schu} found an absorption feature in the spectrum of PG~2200+085 in the region of the $H_{\\alpha}$ line. Apart from that, PG~2200+085 has not been studied in any further detail.  During the preliminary photometric observations of  PG catalogue objects that we performed  in 2002,  we found indications of light variations of PG~2200+085 with an amplitude of $\\Delta m_V=0.^{m}02$. Therefore, in 2003, we performed a more focused  investigation of PG~2200+085, including longtime photometry with a high-quality CCD--detector over several nights and spectroscopy with the 6--meter telescope of the Special Astrophysical Observatory (Russia).  \\section[]{Observations} \\subsection{Photometry}  Photometric observations of PG~2200+085 were performed on September 18-20 2003, with the  1.5--meter Russian-Turkish telescope RTT-150 at the TUBITAK National Observatory (Turkey) on Bakirlitepe Mountain (see Table 1). The thermoelectrically cooled CCD camera ANDOR (model DW436, $2048 \\times 2048$ with $13.5 \\times 13.5~ \\mu$m pixels) with an operating temperature of -60$^{\\circ}$ C installed at the Cassegrain focus was used. All observations were carried out in the V-band with an exposure time of 30 sec and a full frame readout time of 37 sec (binning 2 x 2). The total observation time during each night was about 6 hours at an average seeing of 2.5 arcsec. Four stars of similar magnitudes in the field of the CCD ($8' \\times 8'$), which are identified in Fig. 1, have been used as comparison stars (see Table 2). Five Landolt standard stars have been observed over both nights, and they have been used to determine the V-magnitudes of the reference stars in the field around PG~2200+085. The brightness of these stars was constant within the error of the differential photometry of the two objects ($\\Delta m_V=0.^{m}01$).  \\begin{table} \\begin{center} \\caption{Log of photometric observations.} \\label{tab1} \\begin{tabular}{cccc} \\hline Data & Start  & Duration & Number \\\\ &  JD    &  hours   & of points \\\\ \\hline 19.09.2003 & 2452901.29470  & 6.25 & 331 \\\\ 20.09.2003 & 2452902.29269  & 6.16 & 327 \\\\  \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\label{tab1a} \\caption{Variable and comparison stars. Coordinates and magnitudes are from the USNO-2A catalogue (Monet et al. 1998). The stars are also identified in Fig.1.} \\begin{tabular}{lccccc} \\hline  Star & V(RTT) &  B & V &  $\\alpha_{2000}$ & $\\delta_{2000}$ \\\\ \\hline PG2200 & 14.36 & 14.8 & 14.4 &  22 03 19.7 & +08 45 36  \\\\ Ref.1 & 14.72 & 14.8 & 14.6 &  22 03 24.4 & +08 47 22 \\\\ Ref.2  & 14.61 & 15.2 & 14.7 &  22 03 09.2 & +08 45 12 \\\\ Ref.3  & 14.29 & 14.5 & 14.3 &  22 03 29.3 & +08 43 35 \\\\ Ref.4  & 14.64 & 15.4 & 14.8 &  22 03 32.6 & +08 46 14 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure} \\rotatebox{0}{\\includegraphics[scale=0.32]{5146fig1.eps}} \\label{fig1a} \\caption{Finding chart of PG~2200+085. North is at the top  and east to the left. The size is $8' \\times 8'$.} \\end{figure}  \\begin{figure} \\rotatebox{0}{\\includegraphics[width=\\columnwidth]{5146fig2.eps}} \\label{fig1} \\caption{The light curve of PG~2200+085 on September 19-20, 2003. The solid curve is the sine function best fitted to the observed data.} \\end{figure}  \\subsection{Spectroscopy}  Spectroscopic observations of PG~2200+085 were carried out at the 6--meter telescope BTA of the Special Astrophysical Observatory with a long--slit spectrograph \\cite{Af} and with a nitrogen--cooled CCD ($1024 \\times 1024$  with $24 \\times 24~ \\mu$m  pixels) installed at the prime focus. Observations were performed on August 3, 2003, under excellent weather conditions, with an average seeing of about 1.$''$3 (see Table 3). We used a 1302 lines mm$^{-1}$ grating, which gives a  $\\Delta\\lambda$ = 2.6 \\AA~ resolution in the wavelength region  $\\Delta \\lambda$ 3860--5100 \\AA. Three subsequent spectra  with the same exposure time of 600 s and with a signal-to noise ratio of $S/N \\approx 100$ were taken. For the calibration of wavelengths, fluxes, and radial velocities $Ar$--$Ne$--$He$ lamp  spectra and spectra of the spectrophotometric standard star BD +28 2106 from the review by Bohlin \\cite{bo} were taken.  \\begin{table} \\begin{center} \\caption{Log of spectroscopic observations.  N: the spectrum number, $S/N$ : signal-to-noise ratio, $\\varphi_1$ and $\\varphi_2$: orbital phases with respect to ephemeris 1 and 2.} \\label{tab1c} \\begin{tabular}{ccccccc} \\hline N & UT & HJD & S/N & $\\varphi_1$ & $\\varphi_2$ \\\\ \\hline 1 & 22$^h$ 16$^m$ & 2452855.4326 & 91 & 0.354 & 0.177 \\\\ 2 & 22$^h$ 26$^m$ & 2452855.4395 & 103 & 0.376 & 0.188 \\\\ 3 & 22$^h$ 37$^m$ & 2452855.4468 & 98 & 0.398 & 0.199 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The photometric and spectroscopic data reduction was performed using the $MIDAS$ \\cite{ba} program package.  \\begin{figure} \\includegraphics[width=\\columnwidth]{5146fig3.eps} \\caption{\\label{fig2}The light curves of PG~2200+085 folded with the period $0.^{d}31858$ (a) and $0.^{d}63716$ (b).} \\end{figure}  \\section[]{Data analysis}  \\subsection{The light curves and ephemeris}  The dependence of the magnitude $m_V$ of PG~2200+085 versus the Julian date of our observations on September 19 and 20, is shown in Fig. 2. It is obvious that the light curves during both nights are almost sinusoidal, with a similar amplitude of $\\Delta m_V = 0.^{m}04$.  The fitting of the light curve  with a simple sine  by the $\\chi^2$ technique resulted in the following ephemeris: \\begin{equation}\\label{equ2} HJD = 2452900.558 (\\pm 0.001) + 0.31858 (\\pm 0.00022) \\varphi, \\end{equation} where $\\varphi = 0.0$ corresponds to the minimum of $m_V$. The light curve folded with this period is shown in  Fig. \\ref{fig2}a. Such a light curve is typical of a close binary with a reflection effect where the ultraviolet radiation of the hot primary component is absorbed and reemitted in the optical region from the facing hemisphere of the secondary.  Fitting the light curve by the three--parameter expression \\begin{equation}\\label{equ3} m_V = m_0 + \\frac{m_1}{2} * cos(2\\pi\\varphi) + \\frac{m_2}{2} * cos(4\\pi\\varphi) \\end{equation} leads to $m_1 = 0.^{m}045 \\pm 0.^{m}0003$ and $m_2 << 0.^{m}002$. The small value of $m_1$ means that either the reflection effect or the orbital inclination is small, and the very small value of $m_2$ indicates that any superposed ellipsoidal variations in this light curve are very small.  However, based on the available photometric data, we cannot exclude the possiblity that the orbital period is twice the value given in (\\ref{equ2}) and that the light curve is double-humped. The indirect argument in favour of this possibility is the $10\\%$ difference in the amplitude of the variation of $m_V$ during the nights September 19-20. Correspondingly, the light curve could also be described (see Fig. \\ref{fig2}b) by the ephemeris: \\begin{equation}\\label{equ4} HJD = 2452900.558 (\\pm 0.001) + 0.63716 (\\pm 0.00051) \\varphi. \\end{equation} Fitting the light curve with Eq. (\\ref{equ3}) yields  the parameters  $m_1 = 0.^{m}004 \\pm 0.^{m}0003$ and $m_2 = 0.^{m}044 \\pm 0.^{m}0003$. Therefore, if ephemeris (\\ref{equ4}) applies, then  the main light variations ($0.^{m}044$) arise due to  ellipsoidal deformation of the secondary.  \\begin{figure} \\includegraphics[width=\\columnwidth]{5146fig4.eps}  \\caption{\\label{fig3} Observed spectrum of PG~2200+085, theoretical spectrum of the K--star and blue continuum of the hot component (a); observed profiles of \\rm{CaII$~\\lambda \\lambda$ 3933, 3968 \\AA~} lines (b).} \\end{figure}  \\subsection{The spectra}  Cross-correlation analysis of the Doppler line shifts in the three spectra of PG~2200+085 showed the constancy of the radial velocities with a measuring accuracy of $\\Delta V$ = 3 km s$^{-1}$ during the 30 min of observations. At the same time, all important features in these spectra had similar intensity.  Therefore, we added the three spectrograms. This yielded the final spectrum with $S/N > 170$ shown in Fig. \\ref{fig3}a. The presence of strong absorption lines from atoms and ions of heavy elements and weak molecular bands allows us to classify one of the components as a main sequence K--star.  To determine the temperature of this star, we computed synthetic spectra for different values of $T_\\mathrm{eff}$ and with $\\log g = 4.5$, $[A] = 0.0$ and compared them to the observed spectrum. The computations were performed by means of the program $SPECTR$ \\cite{sa}, whereby we ignored the possible reflection effect in the binary. The observed ratios of the intensities of the metal lines  are well modelled (see Fig. \\ref{fig3}a) with a temperature $T_\\mathrm{eff} = 4600 \\pm 200K$. However, the high level of the observed continuum and the low line intensites indicate the presence of additional flux from the hot component of PG~2200+085. To extract the contribution of this component, we subtracted the theoretical spectrum of the K--star (scaled to the best agreement of observed and theoretical line intensities) from the observed spectrum of PG~2200+085. The residual spectrum shown in  Fig. \\ref{fig3}a does not show any significant absorption or emission lines. The strong absorption feature near $\\Delta \\lambda$ 4750--4800 \\AA~ is the result of neglecting the molecular bands in the synthetic K--star spectrum. Therefore, this feature is not part of the spectrum of the hot component. We note that the real flux distribution of the hot component is possibly even steeper than shown in  Fig. \\ref{fig3}a because we have ignored interstellar absorption. As a result, we can conclude that the observed spectrum of PG~2200+085 is composite (as mentioned by Green et al. (1986), with contributions of a hot white dwarf and a K--star.  At the time of our observations, the radial velocity of the K--star  was  $V_r = 57 \\pm 3 $ km s$^{-1}$. According to our ephemeris (\\ref{equ2}), the spectrograms were taken at the phases $\\varphi$ = 0.34, 0.36 and  0.38. The absence of radial velocity variations of the K--star between these phases allows us to set an upper limit on the orbital radial velocity amplitude  of $V_r$ $\\le$ 30 km s$^{-1}$. On the other hand, if we use (\\ref{equ4}), the corresponding  upper limit for the  radial velocity amplitude of the K--star is 110 km s$^{-1}$.  In the spectrum of PG~2200+085 emission lines of ionized calcium $\\lambda \\lambda $ 3933, 3968 \\AA~ can be seen ( Fig. \\ref{fig3}b). The measured equivalent widths of these lines are $W_\\lambda(3933)$ =1.1 \\AA~and $W_\\lambda(3968)$=0.9 \\AA. These widths are too small for a typical K--star without flares. Similar intensities of \\rm{CaII} emission lines are, however, seen in published spectra \\cite{sk} of the pre-cataclysmic binary V471 Tau at an orbital phase $\\varphi = 0.25$. Moreover, the intensities of all absorption lines in the spectra of V471 Tau and PG~2200+085 are very similar. ",
        "conclusions": "The observations presented in this paper allow us to conclude that PG~2200+085 is a pre-cataclysmic binary consisting of a hot ( pre-) white dwarf and a K--dwarf companion, and with an orbital period of either $P = 0.3186$d or $P= 0.6372$d. When comparing PG~2200+085 to the pre-cataclysmic binaries listed by Schreiber \\& G\\\"ansicke \\cite{sg}, or to the most recent version (release 7.4) of the Ritter \\& Kolb \\cite{ri} catalogue, we note that PG~2200+085 is only the second such system consisting of a white dwarf and a K--star companion (the other one being V471~Tau). All other known short--period pre-cataclysmic binaries either consist of a white dwarf and an M--dwarf companion, or, if the companion is a K--star, the hot component is an extremely hot sdO star. From the evolutionary point of view, systems like V471~Tau and PG~2200+085 are intersting because they are the direct progenitors of cataclysmic variables with a comparatively long orbital period, i.e., $P \\ge 6$hr. We cannot exclude the other possibility, namely, that this object is a wide binary consisting of a hot, highly magnetized white dwarf and a K--star companion. But, according to the statistical arguments presented above, this is much less probable. Nonetheless, if it were true, PG~2200+085 is an even more interesting object than a pre-CV. For this reason alone, PG~2200+085 deserves further and detailed study.  In the framework of the pre-CV hypothesis, we have presented theoretical and observational arguments that resolve the orbital period ambiguity in favour of the shorter period $P = 0.3186$d, i.e., we favour the interpretation that the optical modulation is due to a reflection effect on the facing hemisphere of the K--star, which, in turn, is caused by irradiation from the moderately hot ($T_{\\rm eff,1} \\ge 38500$ K) white dwarf. However, for a final decision about which of the two proposed periods is the correct one, a detailed radial velocity study of PG~2200+085 is required."
    },
    "0606/astro-ph0606590_arXiv.txt": {
        "abstract": "The Galactic high-mass X-ray binary and jet source (microquasar) LSI~+61~303 has recently been detected at TeV $\\gamma$-ray energies by the MAGIC telescope. We have applied a time-dependent leptonic jet model to the broadband spectral energy distribution and suggested (though not unambiguously detected) orbital modulation of the very high energy $\\gamma$-ray emission of this source. Our model takes into account time dependent electron injection and acceleration, and the adiabatic and radiative cooling of non-thermal electrons. It includes synchrotron, synchrotron self-Compton and external inverse Compton (with seed photons from the companion star), as well as $\\gamma\\gamma$ absorption of $\\gamma$-rays by starlight photons. The model can successfully reproduce the available multiwavelength observational data. Our best fit to the SED indicates that a magnetic field of $B_0 \\sim 5 \\times 10^3$~G at $\\sim 10^3 \\, R_g$ is required, and electrons need to be accelerated out to TeV energies ($\\gamma_2 = 10^6$) with a nonthermal injection spectrum with a spectral index of $q = 1.7$, indicating the operation of acceleration mechanisms beyond the standard first-order Fermi mechanism at relativistic or non-relativistic shocks. The orbital modulation of the VHE $\\gamma$-ray emission can be explained solely by the geometrical effect of changes in the relative orientation of the stellar companion with respect to the compact object and jet as it impacts the position and depth of the $\\gamma\\gamma$ absorption trough. Such a scenario predicts a trend of spectral hardening during VHE $\\gamma$-ray low orbital phases. ",
        "introduction": "X-ray binaries with relativistic jets, or microquasars, have recently been established as a new class of $\\gamma$-ray emitting sources. Two sources detected by EGRET on board the {\\it Compton Gamma-Ray Observatory} were found to be spatially consistent with the locations of microquasars, namely LS 5039 and LSI +61 303 \\citep{hartman99}. Observations of $\\gtrsim$ 250 GeV $\\gamma$-rays from LS 5039 with the High Energy Stereoscopic System (HESS) have recently shown this high-mass source to be a TeV emitter \\citep{aharonian05}, also strengthening its earlier tentative identification with the EGRET source 3EG J1824-1514 \\citep{paredes2000}. The second source, LSI +61 303, associated with the COS-B source 2CG 135+01 \\citep{hermsen77, gt78} and the EGRET source 3EG J0241+6103 \\citep{kniffen97}, was recently observed with the Major Atmospheric Gamma-ray Imaging Chrenkov (MAGIC) telescope by \\cite{albert06}. They detected variable $\\gamma$-ray emission above 100 GeV over six orbital cycles, suggesting a periodic modulation on the time scale of the orbital period of $29.496$~d, though further observations are necessary to firmy establish the correlation between the VHE variability and the orbital period. The strongest detections were not at the periastron, suggesting that the modulation might not be related to a modulation of the accretion rate \\citep[as suggested by, e.g.,][]{romero03,bosch04b}, but rather to geometrical effects. Geometrical effects causing an orbital modulation of the high-energy emission include the azimuthal-angle dependence of the $\\gamma\\gamma$ absorption of high-energy emission by starlight photons \\citep{boettcher05,dubus06} or of the Compton upscattering of starlight photons by relativistic electrons in the microquasar jet \\citep{dermer06}. \\cite{dubus06} has performed a detailed analysis of the $\\gamma\\gamma$ absorption of VHE photons in the photon field of the stellar companion, taking into account the finite size of the star and the eccentricity of the orbit for the case of VHE emission from the surface of a putative neutron star associated with LSI +61 303. Note that he used a definition of the phase such that phase 0 corresponds to the periastron passage, so that his results should be shifted by $\\Delta\\psi = 0.23$ (see below) when compared to the definition commonly used throughout the literature on this object, including this {\\it Letter}. He found that the $\\gamma\\gamma$ absorption depth is expected to show a pronounced maximum at phases just before the periastron passage, when the compact object is located behind the stellar companion. The recent VHE detections, along with the observation of X-ray jet structures in several microquasars using $Chandra$ and $XMM-Newton$ \\citep[e.g.,]{corbel02, tomsick02}, have re-ignited interest in jet models for the high energy emission from microquasars, analogous to the commonly adopted models for blazars \\citep[for a recent review see, e.g.][]{boettcher02}.  In the leptonic model of microquasars, Very High Energy (VHE) emission might most likely originate near the base of the mildly relativistic jet. The soft photons from the companion star, the accretion disk as well as from jet synchrotron radiation can be Compton upscattered by the ultra-relativistic electrons in the jet. Steady-state leptonic jet models of the $\\gamma$-ray emission from microquasars have been presented, e.g., by \\cite{bosch04a,bosch04b}, and \\citep{dermer06}. A time-dependent, broadband leptonic jet model was presented in \\cite{gupta06}, (henceforth, Paper 1) where an analytical solution to the electron kinetic equation was presented, restricting the analysis to Compton scattering in the Thomson regime. This $Letter$ follows up on the analysis of Paper 1, now incorporating a full Klein-Nishina treatment of the Compton scattering as well as the angle dependence of the stellar radiation field. We then apply the new model to broadband observations of LSI +61 303, iincluding the effect of the orbital modulation on the $\\gamma\\gamma$ absorption in the $>$100GeV range. In \\S \\ref{model}, we present a general outline of the model geometry and the radiative processes involved. \\S \\ref{application} shows the application to LSI +61 303, and we conclude with a brief summary in \\S \\ref{summary}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606559_arXiv.txt": {
        "abstract": "In a multiply connected space, the two twins of the special relativity twin paradox move with constant relative speed and meet a second time without acceleration.  The twins' situations appear to be symmetrical despite the need for one to be younger due to time dilation. Here, the suggestion that the apparent symmetry is broken by homotopy classes of the twins' worldlines is reexamined using space-time diagrams. It is found that each twin finds her own spatial path to have zero winding index and that of the other twin to have unity winding index, i.e. the twins' worldlines' relative homotopy classes are symmetrical. Although the twins' apparent symmetry is in fact broken by the need for the non-favoured twin to non-simultaneously identify spatial domain boundaries, the non-favoured twin {\\em cannot} detect her disfavoured state if she only measures the homotopy classes of the two twins' projected worldlines, contrary to what was previously suggested. We also note that for the non-favoured twin, the fundamental domain can be chosen by identifying time boundaries (with a spatial offset) instead of space boundaries (with a temporal offset). ",
        "introduction": "Since the confirmation by Wilkinson Microwave Anisotropy Probe (WMAP) observations of the cosmic microwave background that the large scale ($r > 10 {\\hGpc}$ in comoving units) auto-correlation function of temperature fluctuations is close to zero [fig.~16, \\citealt{WMAPSpergel} (orthogonally projected spatial scale); fig.~1 of \\citealt{RBSG08} (non-projected spatial scale)] and difficult to reconcile with a standard cosmic concordance model ({\\SSS}7 of \\citealt{WMAPSpergel}; though see also \\citealt{EfstNoProb03b}) there has recently been considerable observational interest in understanding multiply connected models of the Universe. The Poincar\\'e dodecahedral space (PDS) has been proposed as a better model of the 3-manifold of comoving space rather than an ``infinite'' flat space (e.g. \\citealt{LumNat03,RLCMB04,Aurich2005a,Aurich2005b,Gundermann2005,RBSG08}; though see also \\citealt{NJ07}). Curvature estimates are consistent with the model: analysis of the 3-year WMAP data \\citep{WMAPSpergel06} results in best estimates of the total density parameter $\\Omtot = 1.010_{-0.009}^{+0.016}$ (when combined with HST key project on $H_0$ data) and $\\Omtot = 1.015_{-0.016}^{+0.020}$ (when combined with Supernova Legacy Survey data), or together with the third, fourth and fifth acoustic peaks, as estimated from their observations near the South Galactic Pole using the Arcminute Cosmology Bolometer Array Receiver (ACBAR), $\\Omtot = 1.03^{+0.04}_{-0.06}$ is obtained \\citep{ACBAR08}, consistently with that expected from the PDS analyses, which require positive curvature in this range of $\\Omtot$ values.  It has also recently become noticed that global topology in a universe containing density perturbations can, in principle, have at least some effect on the metric, even though the effect is expected to be small. At the present epoch, in the case of a $\\mathbb{T}^3$ model of fundamental lengths which are slightly unequal by a small fraction $\\delta$, a new term appears in the acceleration equation, causing the scale factor to accelerate or decelerate in the directions of the three fundamental lengths in a way that tends to equalise them \\citep{RBBSJ06}.   In this context, the properties and implications of the twin paradox of special relativity in a multiply connected space need to be correctly understood. It has already been shown \\citep{BransStewart73,Peters83,Peters86,Low90,UzanTwins02,LevinTwins01} that resolving the twin paradox of special relativity in a multiply-connected ``Minkowski'' space-time implies new understanding of the paradox relative to the case in simply connected Minkowski space.  Moreover, it is known that, at least in the case of a static space with zero Levi-Civita connection, multiple connectedness implies a favoured space-time splitting. This should correspond to the comoving reference frame \\citep{UzanTwins02,LevinTwins01}. This could be of considerable importance to the standard cosmological model, since it would provide a novel physical (geometrical) motivation for the existence of a favoured space-time foliation, i.e. the comoving coordinate system.    The difference between the twin paradox of special relativity in a multiply connected space relative to that in a simply connected space is that in a multiply connected space, the two twins can move with constant relative speed and meet each other a second time, {\\em without requiring any acceleration}. The paradox is the apparent symmetry of the twins' situations despite the time dilation effect expected due to their non-zero relative speed. It is difficult to understand how one twin can be younger than the other --- why should moving to the left or to the right be somehow favoured? Does the time dilation fail to occur?   As shown by several authors \\citep{Peters83,Low90,UzanTwins02,LevinTwins01}, the apparent symmetry is violated by the fact that (at least) one twin must identify the faces of the fundamental domain of the spatial 3-manifold {\\em non-simultaneously}, and has problems in clock synchronisation.   \\cite{UzanTwins02} suggested that the apparent symmetry is also violated by an asymmetry between the homotopy classes of the worldlines of the two twins.  Here, we reexamine this suggestion.   The multiply connected space version of the twins paradox is briefly reviewed in \\SSS\\ref{s-paradox}, the question of clarifying the asymmetry is described in \\SSS\\ref{s-asymmetry}, and a multiply connected space-time, with a standard Minkowski covering space-time, and the corresponding space-time diagrams, are introduced in \\SSS\\ref{s-st-diagrams}. In \\SSS\\ref{s-worldlines-etc}, the worldlines of the two twins, projections from space-time to space, homotopy classes and winding numbers are expressed algebraically.  In \\SSS\\ref{s-results}, the resulting projections of the twins' paths from space-time to space are given and their homotopy classes inferred. We also calculate whether or not the generator of the multiply connected space and the Lorentz transformation commute with one another, in \\SSS\\ref{s-lorentz}.  Discussion suggesting why this result differs from that of \\cite{UzanTwins02} is presented in \\SSS\\ref{s-discuss}, including a brief description in \\SSS\\ref{s-tperiodic} of how the non-favoured twin can choose a time-limited instead of a space-limited fundamental domain of space-time. Conclusions are summarised in \\SSS\\ref{s-conclu}. A thought experiment to develop physical intuition of the non-favoured twin's understanding of the spatial homotopy classes is given in Appendix~\\ref{s-stretchable-cord}.  For a short, concise review of the terminology, geometry and relativistic context of cosmic topology (multiply connected universes in the context of modern, physical cosmology), see \\cite{Rouk00BASI} (slightly outdated, but sufficient for beginners). There are several in-depth review papers available \\citep{LaLu95,Lum98,Stark98,LR99} and workshop proceedings are in \\cite{Stark98} and the following articles, and \\cite{BR99}. Comparisons and classifications of different {\\em observational} strategies, have also been made \\citep{ULL99b,LR99,Rouk02topclass,RG04}.  We set the conversion factor between space and time units to unity throughout: $c_{\\mbox{space-time}}=1$. ",
        "conclusions": "\\label{s-conclu}  Finding an asymmetry in the twin paradox of special relativity in a multiply connected space is less obvious than in a simply connected space, since neither twin accelerates. It was already known that the asymmetry required is the fact that (at least) one twin must identify space-time events non-simultaneously and has problems in clock synchronisation.  Here, space-time diagrams and the corresponding  algebra have been presented in order to examine whether or not the homotopy classes of the twins' worldlines provide another asymmetry. They show that homotopy classes (numbers of windings) do {\\em not} show which of the two twins of the twin paradox has a preferred status, contrary to what was previously suggested: each twin finds her own spatial path to have zero winding index and that of the other twin to have unity winding index (in the direction of travel of the other twin).  Although the twins' apparent symmetry is broken by the need for the non-favoured twin to non-simultaneously identify spatial domain boundaries, and by the non-favoured twin's problems in clock synchronisation (provided that she has precise clocks), the non-favoured twin {\\em cannot}  detect her disfavoured state by measuring the topological properties of the two twins' worldlines in the absence of precise metric measurements with clocks or rods.  On the other hand, a non-favoured twin capable of making precise metric measurements will notice many surprising properties of space-time. For example, a non-favoured twin could define the fundamental domain of space-time to be $\\mathbb{R} \\times [0,\\beta\\gamma {L})$ where the first dimension is space and the second is time. In words, time would be ``periodic with period $\\beta\\gamma L$ and a spatial offset when matching time boundaries'' as an alternative explanation to having space be ``periodic with period $\\gamma L$ and a time offset when matching space boundaries''."
    },
    "0606/cond-mat0606514_arXiv.txt": {
        "abstract": "Binding of few-electron systems in two-dimensional potential cavities in the presence of an external magnetic field is studied with the exact diagonalization approach. We demonstrate that for shallow cavities the few-electron system becomes bound only under the application of a strong magnetic field. The critical value of the depth of the cavity allowing the formation of a bound state decreases with magnetic field in a non-smooth fashion, due to the increasing angular momentum of the first bound state. In the high magnetic field limit the binding energies and the critical values for the depth of the potential cavity allowing the formation of a bound system tend to the classical values. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606715_arXiv.txt": {
        "abstract": "Results are presented on component masses and system angular momenta for over a hundred low-temperature contact binaries. It is found that the secondary components in close binary systems are very similar in mass. Our observational evidence strongly supports the argument that the evolutionary process goes from near-contact binaries to A-type contact binaries, without any need of mass loss from the system. Furthermore, the evolutionary direction of A-type into W-type systems with a simultaneous mass and angular momentum loss is also discussed. The opposite direction of evolution seems to be unlikely, since it requires an increase of the total mass and the angular momentum of the system. ",
        "introduction": "\\label{sect:intro} In his two seminal papers, \\citet{Lucy1968a,Lucy1968b} not only showed that two stars can exist in an envelope of a common equipotential and thus resolved the overall hitherto unexplainable properties of W~UMa-type stars, but also very clearly indicated that such contact binaries must have very dissimilar components. As pointed out by \\citet{Hazle1970}, evolution can create this dissimilarity.  The evolutionary state of contact binary stars remains unclear because their spectra cannot be analyzed for abundances due to the extreme broadening and blending of spectral lines. Indirect information though, such as their progressively increasing numbers with age in old open \\citep{Rci1998} and globular \\citep{Rci2000} clusters, as well as the kinematic characteristics \\citep{GB1988, Bilir2005}, very strongly suggest an advanced age of $>2$ Gyr.  Recently, \\citet{Stepien2004} has developed a model with the currently less massive component being the more evolved one. Such a model is conceptually very close to that used to explain the semi-detached Algols. In his model, the current secondary (less-massive) components must be in some cases very low in mass to explain systems like AW~UMa or SX~Crv \\citep{Rasio1995}.  In this Letter we present a summary of results on the component masses (for 112 systems) and system angular momenta (for 93 systems) for low-temperature contact binaries. The sample used was collected mainly from the list of contact binaries defined by \\citet{Kreiner2003}. Half of the systems have solutions published in the frame of the W~UMa project (papers I-VI) \\citep{Kreiner2003, Baran2004, Zola2004, Gazeas2005, Zola2005, Gazeas2006}. All the rest were collected from the literature, the physical parameters of which have been determined accurately, using photometric light curves and radial velocity measurements for both components. We also included the near-contact binaries (NCBs) and the detached close binaries (DCBs) listed in Tables 2 and 3 of \\citet{YE2005}, in order to compare their physical parameters with those of our sample.  We had to exclude the cases, where the third light and the low-amplitude light variation give unreliable solutions, such as V2150~Cyg, V899~Her, HT~Vir, BL~Eri and GO~Cyg. In some cases, close binaries in triple systems have led to spurious solutions and for this reason they were excluded from our sample. Other cases with third light contribution do not seem to produce any problem. Such systems, which are members in multiple systems, have better geometrical configuration and usually have very good spectroscopic determination of the third light contribution, allowing accurate determination of the orbital and physical parameters. Two systems, V351 Peg and V402 Aur were very probably incorrectly classified as W-type binaries, although the shape of their light curves does not support such a classification. Both systems have equal minima in their light curves, making it difficult to distinguish them from each other. Since their physical parameters (masses and periods) were closer to those of A-type binaries, they were classified as A-type systems.  Only recently, good spectroscopic data has become available for more than a hundred contact binaries. Since even small-mass secondaries can be observed spectroscopically in contact systems, most objects have been analyzed as double-lined binaries bypassing any need of inferences based on single-lined data or solely on sometimes highly unreliable photometric elements (particularly mass ratios).  In our study we consider component \"1\" as the more massive one. Our assumption is based on the double-lined spectroscopic observations, where the mass ratio is taken as $q=M_{\\rmn{2}}/M_{\\rmn{1}} <1$. ",
        "conclusions": "\\label{sect:disc}  Do both A and W-type contact binaries have the same origin? Is the one type progenitor of the other? The above questions are still open for investigation.  The main result of the present study, taking into account the masses and angular momenta of contact binaries, is that the W-type systems cannot produce A-type binaries, since angular momentum and mass cannot be added to a system, but only lost from it. It seems more reasonable that A-type systems evolve to W-type systems by loosing mass and angular momentum. A similar evolutionary direction from long to short period binaries \\citep{Bilir2005} showed that systems with longer periods are kinematically young (age 2 Gyr) in contrast to those with shorter periods (age 8 Gys).  An interesting result extracted from our plots is that the secondary components in close binary systems are very similar in mass. Scatter of the points in our plots is mainly due to inaccurate photometric and spectroscopic solutions and/or a possible undetected third light contribution. Since the best of the available data is used in our sample, only the third light could produce a problem and this is why some solutions are excluded from our sample. A very recent study of \\citet{PR2006} about the formation of contact binaries in multiple systems, suggests that a large percentage of close binaries is formed into triple and multiple systems. In such a case, a small amount of unreliability is added to all the solutions, if they are affected from an undetected third light.  Our observational evidence (Figures 1-4) strongly supports the argument that the evolutionary process is from NCBs to A-type contact binaries, without any need of mass loss from the system. The next step in this scenario may lead either to a transformation of A-type to W-type systems with a simultaneous mass and angular momentum loss, or to A-type systems with extremely low mass ratio. These systems will eventually merge into a single, fast-rotating object. The opposite direction of evolution seems to be impossible, since it requires an increase of the total mass and angular momentum of the system."
    },
    "0606/astro-ph0606009_arXiv.txt": {
        "abstract": "{ We examine the distributions of eccentricity and host star metallicity of exoplanets as a function of their mass. Planets with $M \\sin i \\gtrsim 4$ M$_{\\rm J}$ have an eccentricity distribution consistent with that of binary stars, while planets with $M \\sin i \\lesssim 4$ M$_{\\rm J}$ are less eccentric than binary stars and more massive planets. In addition, host star metallicities decrease with planet mass. The statistical significance of both of these trends is only marginal with the present sample of exoplanets. To account for these trends, we hypothesize that there are two populations of gaseous planets: the low-mass population forms by gas accretion onto a rock-ice core in a circumstellar disk and is more abundant at high metallicities, and the high-mass population forms directly by fragmentation of a pre-stellar cloud. Planets of the first population form in initially circular orbits and grow their eccentricities later, and may have a mass upper limit from the total mass of the disk that can be accreted by the core. The second population may have a mass lower limit resulting from opacity-limited fragmentation. This would roughly divide the two populations in mass, although they would likely overlap over some mass range. If most objects in the second population form before the pre-stellar cloud becomes highly opaque, they would have to be initially located in orbits larger than $\\sim 30$ AU, and would need to migrate to the much smaller orbits in which they are observed. The higher mean orbital eccentricity of the second population might be caused by the larger required intervals of radial migration, and the brown dwarf desert might be due to the inability of high-mass brown dwarfs to migrate inwards sufficiently in radius. } ",
        "introduction": "The discovery of the first exoplanets around main-sequence stars a decade ago (Mayor \\& Queloz 1995; Butler \\& Marcy 1996; Marcy \\& Butler 1996; Butler et al. 1997) has provided us with a first glimpse at the fascinating diversity of planetary systems in the universe. At present, with about 200 exoplanets known (most of them having been found by the method of radial velocities around stars similar to the Sun, and a few by the methods of transits and microlensing events), we have learned that: {\\em (a)} Jovian planets are found over a wide range of orbital radii $a$, from the smallest orbits where they can survive evaporation ($a \\sim 0.03$ AU) to the largest ones at which they can be detected ($a\\sim 5$ AU); {\\em (b)} Jovian planets that have not been tidally circularized show a wide range of eccentricities, with a median value of $\\sim 0.3$; {\\em (c)} orbital resonances are commonly found in systems when more than one planet is detected; {\\em (d)} there is a power-law distribution of planet masses that is roughly uniform in $\\log M$ from $\\sim 0.3$ to $\\sim10$ Jupiter masses (e.g., Tabachnik \\& Tremaine 2002), although with a strong deficit of objects with mass greater than $\\sim0.015$ M$_\\odot$ (designated as brown dwarfs), which is known as the ``brown dwarf desert'', and probably an increase in the planet abundance at lower masses, as suggested by the recent detections of Neptune-like planets with the method of microlensing (Beaulieu et al. 2006; Gould et al. 2006). All of these findings have come as surprises, and none of them was predicted or expected from theories of planetary formation.  With the increasing sample of exoplanets available for statistical studies, the fundamental question of the planet formation mechanism and their subsequent orbital evolution that results in the observed distribution may start to be addressed. The only information we have so far that relates to the formation process of planets is the distribution of their orbital periods, eccentricities, and planet masses, which can also be Compared to the same distributions for binary stars and brown dwarfs. In addition, properties of the host stars such as their metallicity can be included. Jovian planets are thought to form in circumstellar disks from the coalescence of planetesimals and gravitational accretion of gaseous material (Pollack et al.\\ 1996). This process should lead to orbits that are initially circular, but they could subsequently be perturbed by several mechanisms causing an increase of the eccentricity, such as disk-planet interactions (Kley \\& Dirksen 2006), the Kozai mechanism (Kozai 1962; Holman et al. 1997; Takeda \\& Rasio 2005), and close encounters or resonant interactions between planets (Chiang et al. 2002; Ford et al. 2005; see also the review on eccentricity growth mechanisms by Tremaine \\& Zakamska 2004). These mechanisms need to be effective for a large majority of planets formed in disks since most of the detected exoplanets are found to possess eccentricities much larger than solar system planets (and hence, the solar system must be an oddity among planetary systems).  It has been noted before that the distribution of eccentricities of the exoplanets seems to be remarkably similar to that of binary systems (Heacox 1999; Stepinski \\& Black 2000, 2001), with a slight tendency for the eccentricities to increase with planet mass (see Tremaine \\& Zakamska 2004; Marcy et al. 2005; Papaloizou \\& Terquem 2006). This result is surprising because it is not clear how a dynamical process can result in higher eccentricities acquired by more massive planets. Although interactions with a gaseous disk can generate eccentricities more easily for the most massive planets, it is doubtful that this process alone can excite the eccentricities up to the observed distribution. While the Kozai mechanism predicts an eccentricity distribution independent of planetary mass, perturbations among planets would likely tend to leave lower mass planets with higher eccentricities, basically from energy equipartition arguments; it is possible, however, that if one observes the most massive planet that has survived in every planetary system, the systems with greater total mass compared to the host star have been more strongly perturbed.  Contrary to planetary objects, the most favored hypothesis to explain the formation of binary stars and brown dwarf companions is by fragmentation of the parent molecular cloud during the gravitational collapse process, or as a result of gravitational instability or fission of a rapidly rotating pre-stellar cloud (Tohline 2002). If the fragmentation process resulted in an orbit randomly selected from phase space for a fixed orbital energy, the distribution of eccentricities should be uniform in $e^2$, with a median eccentricity of $0.7$ in a binary star sample. However, this is not observed (e.g., Abt 2005), and even after removing systems that may have been affected by tidal circularization the typical eccentricities are substantially lower. It is particularly striking that the distribution of eccentricities of exoplanets and binary stars are remarkably similar, since the two types of systems are thought to form by different mechanisms.  In this paper, we consider the possibility that some planetary mass objects have also formed by the process of direct cloud fragmentation. The known exoplanets would therefore be part of two populations (as proposed earlier by, e.g., Black 1997; Mayor et al. 1998; Papaloizou \\& Terquem 2001; Udry et al. 2002), stemming from two different mechanisms of mass growth: either gas accretion from a disk onto a seed planetesimal, or fragmentation during the collapse of the gas cloud. Even though exoplanets and brown dwarfs are usually considered as separate classes of objects (with the separation chosen at a mass of $0.013$ M$_\\odot$, the minimum mass required for deuterium burning), there is no fundamental reason why these objects should belong to distinct classes from the point of view of their formation. In general, the different formation processes mentioned above (as well as disk fragmentation by gravitational instability) may be relevant for objects over wide ranges of mass that may overlap, and should be unrelated to the minimum mass for deuterium burning or for hydrogen burning. We shall search for possible signatures of the presence of two populations of objects among the known exoplanets and brown dwarfs with the improved statistics available today, and also examine the correlation with the metallicity of the host star. ",
        "conclusions": "\\label{sec_disc}  \\subsection{A hypothesis for the physical origin of two planet populations}  The statistical evidence for the two trends we have examined, of increasing orbital eccentricity and decreasing host star metallicity with planet mass, should be considered as marginal at this point. However, the presence of the two independent trends reinforce each other in suggesting that there may be different planetary formation mechanisms that give rise to planets of different masses. In this section, we discuss a possible physical explanation for these trends, assuming that they are correct and bearing in mind that they will need to be confirmed (or refuted) by improved statistical evidence as the number of known exoplanets increases.  The metallicity of the host star should not be affected by the fact that a planet has formed and has stayed in orbit around the star. Therefore, it seems reasonable to conclude that the majority of the high-mass planets form in a way that is less influenced by the metallicity than low-mass planets. The most natural interpretation is that many of the high-mass planets form through the same mechanism postulated for brown dwarfs, by direct fragmentation of the pre-stellar cloud (possibly combined with the interruption of gas accretion by radiation feedback from nearby stars) or of a circumstellar gaseous disk that is gravitationally unstable (e.g., Whitworth et al. 2006), while the low-mass planets form by the initial coagulation of a core of rock and ice and subsequent gas accretion onto the core (Pollack et al. 1996). There is indeed no reason to suppose that substellar objects formed by these two mechanisms would not overlap in mass. The formation of a planet from a core of rock and ice has a natural mass upper limit determined by the total mass that can be accreted from a circumstellar disk, at the time and the radius where the planetary core reaches a critical mass allowing gas accretion to start. At the same time, a natural lower limit to the mass of objects formed by fragmentation and direct gravitational collapse of gas is the opacity limit, which assumes that fragmentation will not be efficient once a pre-stellar cloud becomes opaque to its own cooling radiation, and is given by (Rees 1976): \\begin{equation} M_{\\rm op}= M_{\\rm Ch}\\, \\mu^{-9/4}\\, f^{-1/2}\\, \\left( { kT \\over m_p c^2 } \\right)^{1/4} ~, \\label{opam} \\end{equation} where $M_{\\rm Ch}$ is the Chandrasekhar mass, $\\mu$ is the mean molecular weight in units of the proton mass, $f$ is the emissivity of the cloud at the moment it becomes opaque (expressed as the fraction of the blackbody radiation that is emitted), and $T\\sim 10$ K is the cloud temperature. This opacity limit mass is likely to represent the smallest possible mass of an object formed by direct gravitational collapse of gas: even if fragmentation is still possible by gravitational instability in a rapidly rotating, highly opaque disk, it is unlikely that the gas temperature in a circumstellar disk around a young star would be low enough to bring the Jeans mass below the opacity limit value in the conditions of the protostellar nebula. The opacity limit mass is close to a Jupiter mass, suggesting that objects formed by gas fragmentation may extend into the high-mass planet regime and be found orbiting other stars with an abundance that is independent of the host star metallicity. On the other hand, gas planets formed by core accretion may not usually grow to more than a few Jupiter masses, form preferentially around metal-rich host stars starting from nearly circular orbits, and grow their orbital eccentricities to an average value lower than the more massive objects formed by fragmentation.  \\subsection{Orbital migration of the high-mass population}  However, this simple idea cannot account for the observations without including an additional ingredient. If the lower limit to the mass of substellar objects formed by fragmentation of gas clouds is to have anything to do with the opacity limit, then these objects should form at very large distances from their stars because the density at which opacity sets in is $n_H \\sim 10^{10}$ cm$^{-3}$ (e.g., Whitworth et al. 2006), and the size of a region containing the mass in Eq. (\\ref{opam}) at this density is $\\sim 30$ AU. Therefore, these objects would have to be formed at large distances and then migrate to the much smaller orbits at which the known exoplanets have been detected.  It may not be unreasonable for the population of high-mass planets to have undergone a large radial migration. After all, we know that it is necessary for most of the planets to have experienced radial migration because it is believed that they can form only at radii large enough to allow for the presence of ice particles, and many planets are found within the ice-line. If planets always start on circular orbits when forming from a rock-ice core, their eccentricities would have to increase as they experience radial migration to account for the observed high eccentricities. The need for a large interval of radial migration for objects formed by fragmentation might then be the reason behind the increasing eccentricity with mass: while most of the high-mass planets would have migrated over large intervals, some of the low-mass planets, which could form from solid cores in the protoplanetary disk at distances much closer than the high-mass population, may have migrated over radial intervals that are too small to grow the eccentricity by the same amount.  Possible mechanisms for the migration of planets formed at large distances include the dynamical interaction with low-mass planets formed in the circumstellar disk, a hydrodynamic and/or gravitational interaction with the gaseous circumstellar disk, or interaction with a disk of planetesimals. The orbit of a distant, high-mass planet could be perturbed to a small pericenter by external torques, interactions with other distant planets, or a Kozai mechanism, and then its semimajor axis could be greatly reduced by exchanging energy with planets formed at short distances, or by dissipative hydrodynamic or gravitational interactions with the disk.  If brown dwarfs and high-mass planets in orbit around solar-type stars are indeed formed only at large distances, this might account also for the presence of the brown dwarf desert (note that brown dwarfs are abundant around M-type stars at typical distances of $\\sim$ 1 AU, where they have probably formed by a mechanism similar to binary stars; Close et al. 2003). As the brown dwarf mass increases, a more massive disk needs to be present around the star in order that a sufficient amount of angular momentum can be absorbed by the disk to allow for the migration of the brown dwarf. It is possible that planets or brown dwarfs above some critical mass are not typically able to migrate to within a few AU of their host star, because their angular momentum is too large to be exchanged with the circumstellar disk.  Clearly, a model where high-mass planets and brown dwarfs form only at large distances predicts that many of these objects should still be found at late times in orbits of large radius around stars. If $\\sim$ 1\\% of solar-type stars contain high-mass planets within 5 AU, then a similar number of stars should contain the same planets and brown dwarfs at distances $>$ 30 AU, even if the migration process of these distant planets can be highly efficient (that is, a high fraction of the distant planets end up in orbits within 5 AU). In general, brown dwarfs and planets with $M> 5 {\\rm M}_{\\rm J}$ are abundant in young clusters, accounting for $\\sim$ 10\\% of all the objects (see Gonz\\'alez-Garc\\'\\i a et al. 2006, although in some cases the abundance may be lower, see Lyo et al. 2006). These same objects in young clusters are found only rarely as companions to solar-type stars on distant orbits (in $\\sim$ 1\\% of the stars; R. Rebolo, priv. communication), although the abundance of these companions could be lower in clusters containing many stars compared to the field, because of tidal destruction of the binary systems. The case of the star $\\epsilon$ Indi is noteworthy: at a distance of 3.6 parsecs, it is the 20th nearest star to the Sun and it has a pair of brown dwarfs at a separation of 1500 AU that was discovered only very recently (McCaughrean et al. 2004). In summary, it seems plausible that there is a large enough reservoir of high-mass planets and brown dwarfs formed at large distances, some of which may undergo radial migration and account for the high-mass planets found within 5 AU of their stars. The alternative possibility to the migration scenario would be that the high-mass planets can form at a radius near their final orbit at which they are observed, by direct gas collapse through gravitational instability of a rotating, opaque disk. However, if this mode of formation is to account for both binary stars and high-mass planets, the presence of the brown dwarf desert has no clear explanation.  \\subsection{The eccentricity distribution of planets and binary stars}  A possible test for the idea that the large eccentricities of the high-mass planet population originate in a large interval of radial migration is related to the fact that all planets on small orbits should have migrated by large intervals, because planets formed by the core accretion process are thought to have formed in orbits outside the ice-line, so they can reach small orbits only by migrating. Therefore, at a small semimajor axis the eccentricity distribution should become more alike for low- and high-mass planets. However, it is hard to tell from Fig. 2 if there is any tendency of increasing eccentricity with decreasing period, and as mentioned in Sect. \\ref{sec_ecdis} any such dependence might be influenced by selection effects arising from the methods by which planets are found in Doppler surveys.  The hypothesis we have examined for the origin of the population of high-mass planets provides no satisfying explanation for why the eccentricity distributions of spectroscopic binaries and high-mass planets are so similar. Clearly, binary stars would have even greater difficulty for migrating inwards from a large orbit than brown dwarf companions. Binary stars are likely to form on an orbit of similar size to their final one, perhaps by gravitational instabilities in massive circumstellar disks or fission of a rapidly rotating pre-stellar cloud (see Tohline 2002 for a review of binary formation theories). In any case, it is clear that binary stars must form by very different mechanisms than any high-mass planets collapsing directly out of gas. Even if the high-mass planets formed on orbits of similar size as the binary stars, the need to account for the brown dwarf desert strongly suggests totally different formation mechanisms. Binary stars are likely to form in different ways in any case, in view of the wide range of orbital sizes, the known excess of twin binaries (with very similar component masses; Pinsonneault \\& Stanek 2006), and the excess of small eccentricity orbits we mentioned at the end of Sect. \\ref{sec_ecdis}. The similar eccentricity distributions of binary stars and high-mass planets may more likely be related to processes that occur after their formation; for example, migration of a high-mass planet through a disk and the formation of a star from gravitational instability in the disk might give rise to similar eccentricity distributions at the end of the process. In addition, a series of perturbations that change orbital eccentricities in a random way may typically produce a characteristic eccentricity distribution (see Juric \\& Tremaine 2006)."
    },
    "0606/astro-ph0606523_arXiv.txt": {
        "abstract": "The conversion of neutron matter to strange matter in a neutron star have been studied as a two step process. In the first step, the nuclear matter gets converted to two flavour quark matter. The conversion of two flavour to three flavour strange matter takes place in the second step. The first process is analysed with the help of equations of state and hydrodynamical equations, whereas, in the second process, non-leptonic weak interaction plays the main role. Velocities and the time of travel through the star of these two conversion fronts have been analysed and compared. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606379_arXiv.txt": {
        "abstract": "{} {We address the issue of the multiplicity of the three brightest early-type stars of the young open cluster IC\\,1805, namely HD\\,15570, HD\\,15629 and HD\\,15558.} {For the three stars, we measured the radial velocity by fitting Gaussian curves to line profiles in the optical domain. In the case of the massive binary HD\\,15558, we also used a spectral disentangling method to separate the spectra of the primary and of the secondary in order to derive the radial velocities of the two components. These measurements were used to compute orbital solutions for HD\\,15558.} {For HD\\,15570 and HD\\,15629, the radial velocities do not present any significant trend attributable to a binary motion on time scales of a few days, nor from one year to the next. In the case of HD\\,15558 we obtained an improved SB1 orbital solution with a period of about 442 days, and we report for the first time on the detection of the spectral signature of its secondary star. We derive spectral types O5.5III(f) and O7V for the primary and the secondary of HD\\,15558. We tentatively compute a first SB2 orbital solution although the radial velocities from the secondary star should be considered with caution. The mass ratio is rather high, i.e. about 3, and leads to very extreme minimum masses, in particular for the primary object. Minimum masses of the order of 150\\,$\\pm$\\,50 and 50\\,$\\pm$\\,15\\,M$_\\odot$ are found respectively for the primary and the secondary.} {We propose that HD\\,15558 could be a triple system. This scenario could help to reconcile the very large minimum mass derived for the primary object with its spectral type. In addition, considering new and previously published results, we find that the binary frequency among O-stars in IC\\,1805 has a lower limit of 20\\,\\%, and that previously published values (80\\,\\%) are probably overestimated.} ",
        "introduction": "\\label{intro}  The study of stellar populations of young open clusters has been the purpose of many works in the last years (e.g. \\citealt{Sag}; \\citealt{RM}). One of the crucial questions addressed in this context concerns the massive star content and the binary frequency of the earliest stars harboured by these clusters. Recent numerical simulations of the star formation in open clusters predict a high level of mass segregation, along with the formation of the most massive stars through simultaneous accretion and stellar collisions, resulting in the presence of the most massive stars in binary systems \\citep{BB2002}. The investigation of the multiplicity of the massive star population of young open clusters is consequently of particular interest. For instance, \\citet{GM1805} presented radial velocity measurements for 37 O- and B-stars in the open cluster NGC\\,6231, and proposed binary frequencies for a series of open clusters, among which is IC\\,1805.  IC\\,1805 is a young rich open cluster in the core of the Cas\\,OB6 association, in the molecular cloud W\\,4 in the Perseus spiral arm of our Galaxy. \\citet{massey1805} inferred an age of 1-3\\,Myr for the cluster, in agreement with other previous estimates (see \\citealt{fein1805}, and references therein, although these latter authors proposed an age $<$\\,1\\,Myr). The spectroscopic study of \\citet{SH1805} revealed that about 80 of the members of IC\\,1805 are O or B stars. \\citet{GM1805} (see also \\citealt{Ish}) estimated a binary frequency of 80\\,\\% among the 10 O-stars in IC\\,1805. However, in a previous paper \\citep[][Paper I]{ic1805_1}, we showed that two suspected binaries, BD\\,+60$^{\\circ}$\\,501 (O7V((f))) and BD\\,+60$^{\\circ}$\\,513 (O7.5V((f))), were most probably single stars, suggesting that the binary frequency proposed by \\citet{GM1805} might be overestimated. In this paper, we investigate the multiplicity of the three earliest O-type stars of the cluster: HD\\,15570, HD\\,15629, and HD\\,15558.\\\\  HD\\,15570 (O4If$^+$, V\\,=\\,8.10) was proposed to be the most massive member of IC\\,1805, and incidentally one of the most massive and most luminous stars known in our Galaxy, with a present evolutionary mass of about 80\\,M$_\\odot$ following \\citet{HPV}. HD\\,15629 (O5V((f)), V\\,=\\,8.42) has also been proposed to be a very massive star. According to \\citet{HPV}, this star could have evolved from an initial mass of about 70\\,M$_\\odot$, and presently be on the way towards evolutionary stages close to those of HD\\,14947 (O5If$^+$) and HD\\,210839 (O6I(n)fp, $\\lambda$\\,Cep). The same authors inferred a present-day evolutionary mass of about 61\\,M$_\\odot$. For both stars, no clear evidence of binarity has been reported in the literature, even though some radial velocities quoted in the WEBDA data base\\footnote{Available at http://obswww.unige.ch/webda} suggest the occurrence of variations.  HD\\,15558 (m$_{V}$\\,=\\,8.0) was classified as an O5III(f) star by \\citet{Mat2}. It was reported to be a spectroscopic binary with a period of about 420\\,d for the first time by Trumpler (according to \\citealt{und67}). Up to now, the only orbital solution proposed for this system is that of \\citet{GM}, revealing a highly eccentric binary (e\\,=\\,0.54) with a period of about 440\\,d, i.e. the longest of any O-type spectroscopic binary known at that time. This star is also believed to be very massive. \\citet{HPV} derived spectroscopic and evolutionary masses of about 90\\,M$_\\odot$ for HD\\,15558. Finally, this massive binary is classified as a non-thermal radio emitter. In this context, the study of its binarity is of crucial interest \\citep[see][ for a detailed discussion]{thesis}.\\\\  This paper is organized as follows. Section\\,\\ref{spec} presents the optical spectrum of the stars. The results of the radial velocity study for HD\\,15570 and HD\\,15629 are given in Sect.\\,\\ref{single}. The case of HD\\,15558 is discussed in detail in Sect.\\,\\ref{hd15558}. Section\\,\\ref{disc} consists of a discussion of our results. A summary of the main results and the conclusions are provided in Sect.\\,\\ref{conc}. ",
        "conclusions": "}  We have presented the results of an intensive spectroscopic study of the brightest massive stars in the young open cluster IC\\,1805: HD\\,15570, HD\\,15629 and HD\\,15558. For HD\\,15570 and HD\\,15629, the RVs do not present any significant trend attributable to binary motion on a time scale of a few days, nor from one year to the next. This is in line with the results obtained by \\citet{hillwigcasob6} who performed a search for RV variation on time-scales of a few days for some O-type stars of IC\\,1805 including HD\\,15629 and HD\\,15570.  In the case of HD\\,15558, we have derived an SB1 orbital solution with significantly refined parameters as compared to those obtained by \\citet{GM}. The system appears to be eccentric ($e$\\,$\\sim$\\,0.4) and we obtain a period of 442\\,$\\pm$\\,12\\,d. A careful inspection of the spectra obtained close to the extrema of the radial velocity curve reveals the presence of the companion in the profiles of the \\ion{He}{i} $\\lambda$ 4471, \\ion{He}{ii} $\\lambda$ 4542, \\ion{C}{iv} $\\lambda$ 5812 and \\ion{He}{i} $\\lambda$ 5876 lines. We have simultaneously fitted Gaussians to the profiles of the \\ion{He}{i} $\\lambda$ 4471 and \\ion{He}{ii} $\\lambda$ 4542 lines in order to separate the primary and secondary components. The determination of the equivalent width of these two lines allowed us to derive O5.5 and O7 spectral types respectively for both stars. Considering in addition that the \\ion{He}{ii} $\\lambda$ 4686 line of the primary has a P-Cygni profile whilst the same line is in absorption in the case of the secondary, along with the fact that we do not clearly observe the \\ion{N}{iii} $\\lambda$$\\lambda$ 4634-4641 lines in emission for the secondary, we propose that HD\\,15558 is an O5.5III(f) + O7V binary.\\\\  We estimated the radial velocities of the secondary of HD\\,15558 following two approaches: (1) a simultaneous fit of line profiles and (2) a disentangling method. Both techniques allowed us to determine minimum masses of the order of 150\\,$\\pm$\\,50 and 50\\,$\\pm$\\,15\\,M$_\\odot$ respectively for the primary and the secondary. We also obtained the first SB2 orbital solution for HD\\,15558. Although we note that the quality of this SB2 solution is rather poor, our data point clearly to a rather high mass ratio (about 3), leading to an extreme minimum mass for the primary. Our results require independent validation using an improved disentangling procedure.  The main problem in considering our results is to reconcile the very extreme mass of the primary with its spectral type. It is indeed unlikely that a very massive main-sequence star could cool down enough during its evolution to become an O5.5 giant. A possible scenario can however be considered where HD\\,15558 is not a binary but a triple system. The primary may be a yet unrevealed close binary system. In this case, as the mass of the primary is estimated on the basis of the motion of the secondary, the primary object -- i.e.\\,the hypothetical close binary -- would appear to be a massive object whose mass is the sum of the masses of two stars. Even though our data do not provide any evidence for this scenario, we estimate that at this stage it should not be rejected, and that it could constitute a valuable working hypothesis for future investigations concerning this system.  From our new and previously published results, we briefly address the question of the multiplicity of the early-type stars in IC\\,1805. The binary frequency among O-stars should be of at least 20\\,\\%, since out of 10 O-stars only 2 are confirmed binaries, and should not exceed 60\\,\\%. We therefore conclude that the previously claimed binary frequency of 80\\,\\% was overestimated."
    },
    "0606/astro-ph0606186_arXiv.txt": {
        "abstract": "We use {\\it Hubble Space Telescope} ({\\it HST}) NICMOS continuum and Pa$\\alpha$ observations  to study the near-infrared and star-formation properties of a representative sample of 30 local ($d\\sim 35-75\\,$Mpc)  luminous infrared galaxies (LIRGs, infrared [$8-1000\\,\\mu$m] luminosities of $\\log L_{\\rm IR}=11-11.9\\,[{\\rm L}_\\odot]$). The data provide spatial resolutions of $25-50$\\,pc and cover the central $\\sim 3.3-7.1\\,$kpc regions of these galaxies.  About half of the LIRGs show compact ($\\sim 1-2\\,$kpc) Pa$\\alpha$ emission with a high surface brightness in the form of nuclear emission, rings, and mini-spirals. The rest of the sample show Pa$\\alpha$ emission along the disk and the spiral arms extending over scales of $3-7\\,$kpc and larger. About half of the sample contains HII regions with H$\\alpha$ luminosities significantly higher than those observed in normal galaxies.  There is a linear empirical relationship between the mid-IR $24\\,\\mu$m and hydrogen recombination (extinction-corrected Pa$\\alpha$) luminosity for these LIRGs, and the HII regions in the central part of M51. This relation holds over more than four decades in luminosity suggesting that the mid-IR emission is a good tracer of the star formation rate (SFR). Analogous to the widely used relation between the SFR  and total IR luminosity of Kennicutt (1998), we derive an empirical calibration of the SFR in terms of the monochromatic $24\\,\\mu$m luminosity that can be used for luminous, dusty galaxies. ",
        "introduction": "The importance of infrared (IR) bright galaxies has been recognized since their discovery more than 30 years ago (Rieke  \\& Low 1972) and the detection of large numbers by the {\\it IRAS} satellite (Soifer et al. 1987). Specifically, there has been controversy over the extent to which luminous and ultraluminous IR galaxies (LIRGs, $L_{\\rm IR}=L[8-1000\\mu {\\rm m}] = 10^{11}-10^{12}\\,{\\rm L}_\\odot$, and ULIRGs $L_{\\rm IR} > 10^{12}\\,{\\rm L}_\\odot$, respectively) in the local universe are powered by intense star formation as opposed to active galactic nuclei (see the review by Sanders \\& Mirabel 1996). The process of merging with accompanying super-starbursts that produces LIRGs and ULIRGs appears to be an important stage in galaxy evolution, possibly even converting spiral galaxies into ellipticals (e.g., Genzel et al. 2001;  Colina et al. 2001, and references therein) and giving rise to quasars (Sanders et al. 1988).  In recent years with the advent of the new generation of IR satellites, there has been a considerable effort in understanding the properties of distant IR-selected galaxies. The majority of these galaxies at $z<1$ are in the LIRG class, and they  make a significant contribution to the star formation rate density at $0.5<z<2$ (Elbaz et al. 2002; P\\'erez-Gonz\\'alez et al. 2005; Le Floc'h et al. 2005). At $z\\sim 0.7 - 1$ a significant fraction of the LIRG population appears to be morphologically disturbed spiral galaxies (Bell et al. 2005; Shi et al. 2006), rather than interacting and merging systems. These high-$z$ IR bright galaxies show morphologies and properties (see e.g., Zheng et al. 2004; Papovich et al. 2005; Melbourne, Koo, \\& Le Floc'h 2005; Shi et al. 2006) typical of low-$z$ LIRGs and ULIRGs (among others, Murphy et al. 1996;  Surace, Sanders, \\& Evans 2000, Scoville et al. 2000; Alonso-Herrero et al. 2002; Colina et al. 2001;  Colina, Arribas, \\& Clements 2004; Arribas et al. 2004). Lagache, Puget, \\& Dole (2005) have recently reviewed the properties of the high-$z$ IR galaxy population.   The local density of LIRGs is two orders of magnitude higher than that of ULIRGs (e.g., Soifer et al. 1987). Moreover, Murphy et al. (2001) have argued that in the local universe interacting galaxies may spend a significant fraction of their lifetime as LIRGs, whereas the ULIRG phase may just be short and recurrent during the merging process. However, the majority of the detailed studies of the local universe IR galaxies have been focused on the objects with the largest IR luminosities (i.e., ULIRGs). The LIRGs have been generally neglected, except for detailed studies of a few famous examples (e.g., Gehrz, Sramek, \\& Weedman 1983; Doyon, Joseph, \\& Wright 1994; Genzel et al. 1995; Satyapal et al. 1999; Sugai et al. 1999; L\\'{\\i}pari et al. 2000, 2004; Alonso-Herrero et al. 2000, 2001). Thus, understanding  low-$z$ LIRG samples is critical for interpreting the properties of IR-selected high-$z$ galaxy populations; in addition, they represent the link between ULIRGs and the population of field galaxies as a whole.    \\begin{deluxetable*}{llcclcccc}  \\tiny \\tablewidth{18cm} \\tablecaption{The sample of local universe LIRGs.}   \\tablehead{Galaxy Name  & IRAS Name & $v_{\\rm hel}$ & Dist & $\\log L_{\\rm IR}$ & Spect. & Refs  & Morphology\\\\ & & (km s$^{-1}$) & (Mpc) & (L$_\\odot$) & class\\\\ (1) & (2) & (3) &(4) & (5) & (6) & (7) & (8)} \\startdata  NGC~23 & IRASF~00073+2538 & 4566  & 59.6 & 11.05 &H\\,{\\sc ii}   &1 & paired with NGC~26\\\\ MCG~+12-02-001      & IRASF~00506+7248 & 4706  & 64.3 & 11.44  &--    &-- & in group?\\\\ NGC~633            & IRASF~01341$-$3735 & 5137  & 67.9 & 11.09: &H\\,{\\sc ii} &2,3 & paired with ESO~297-G012 \\\\ UGC~1845           & IRASF~02208+4744 & 4679  & 62.0 & 11.07  &--   &-- & isolated\\\\ NGC~1614           & IRASF~04315$-$0840 & 4746  & 62.6 & 11.60  &H\\,{\\sc ii}   &1,2,3 & merger\\\\ UGC~3351           & IRASF~05414+5840 & 4455  & 60.9 & 11.22  &Sy2   &4,5 & paired with UGC~3350\\\\ NGC~2369           & IRASF~07160$-$6215 & 3237  & 44.0 & 11.10  &--    &--  & isolated\\\\ NGC~2388           & IRASF~07256+3355 & 4134  & 57.8 & 11.23: &H\\,{\\sc ii} &1  & paired with NGC~2389\\\\ MCG~+02-20-003      & IRASF~07329+1149 & 4873  & 67.6 & 11.08: &--    &--  & in group\\\\ NGC~3110           & IRASF~10015$-$0614 & 5034  & 73.5 & 11.31: &H\\,{\\sc ii} &2 & paired with MCG~$-$01-26-013\\\\ NGC~3256           & IRASF~10257$-$4339 & 2814  & 35.4 & 11.56  &H\\,{\\sc ii} & 6 & merger\\\\ NGC~3690/IC~694     & IRASF~11257+5850 & 3121  & 47.7 & 11.88  &Sy2-H\\,{\\sc ii} & 7 & close interacting pair\\\\ ESO~320-G030       & IRASF~11506$-$3851 & 3232  & 37.7 & 11.10  &H\\,{\\sc ii}   & 8 & isolated\\\\ MCG~$-$02-33-098E/W   & IRASF~12596$-$1529 & 4773  & 72.5 & 11.11  &H\\,{\\sc ii} (both)  &1,3 & close interacting pair\\\\ IC~860             & IRASF~13126+2453 & 3347  & 59.1 & 11.17: &H\\,{\\sc ii}     &5 & isolated\\\\ NGC~5135           & IRASF~13229$-$2934 & 4112  & 52.2 & 11.17  &Sy2   &2,3 & in group\\\\ NGC~5653           & IRASF~14280+3126 & 3562  & 54.9 & 11.06  &H\\,{\\sc ii}   &1 & isolated\\\\ NGC~5734           & IRASF~14423$-$2039 & 4074  & 59.3 & 11.06: &NO     &1 & paired with NGC~5743\\\\ IC~4518E/W         & IRASF~14544$-$4255 & 4715  & 69.9 & 11.13  &Sy2 (W)  &2 & close interacting pair \\\\ Zw~049.057         & IRASF~15107+0724 & 3897  & 59.1 & 11.27: &H\\,{\\sc ii} &1,5 & isolated\\\\ NGC~5936           & IRASF~15276+1309 & 4004  & 60.8 & 11.07  &H\\,{\\sc ii}   &1 & isolated\\\\ ---                & IRASF~17138$-$1017 & 5197  & 75.8 & 11.42  &H\\,{\\sc ii}   &2,3 & isolated?\\\\ IC~4687/IC~4686         & IRASF~18093$-$5744 & 5200/4948  & 74.1 & 11.55:  &H\\,{\\sc ii} (both)  &3,9 & close interacting pair\\\\ IC~4734            & IRASF~18341$-$5732 & 4680  & 68.6 & 11.30  &H\\,{\\sc ii}/L &2 & in group\\\\ NGC~6701           & IRASF~18425+6036 & 3965  & 56.6 & 11.05  &H\\,{\\sc ii}:    &1 & isolated\\\\ NGC~7130           & IRASF~21453$-$3511 & 4842  & 66.0 & 11.35  &L/Sy  &1,2 & peculiar\\\\ IC~5179            & IRASF~22132$-$3705 & 3422  & 46.7 & 11.16  &H\\,{\\sc ii}   &1 & isolated\\\\ NGC~7469           & IRASF~23007+0836 & 4892  & 65.2 & 11.59  &Sy1   &1 & paired with IC~5283\\\\ NGC~7591           & IRASF~23157+0618 & 4956  & 65.5 & 11.05  &L     &1 & isolated\\\\ NGC~7771           & IRASF~23488+1949 & 4277  & 57.1 & 11.34   &H\\,{\\sc ii}   &1& paired with NGC~7770\\\\ \\enddata \\tablecomments{Column~(1): Galaxy name. Column~(2): {\\it IRAS} denomination from Sanders et al. (2003). Column~(3): Heliocentric velocity from NED. Column~(4): Distance taken from Sanders et al. (2003). Column~(5): $8-1000\\,\\mu$m IR luminosity taken from Sanders et al. (2003), where the suffix ``:'' means large uncertainty (see Sanders et al. 2003 for details). NGC~633 was identified as the {\\it IRAS} source in Sanders et al. (1995). Interacting pairs for which Surace et al. (2004) obtained {\\it IRAS} fluxes for the individual components: NGC~633 $\\log L_{\\rm IR}=10.56$, and for NGC~2388, NGC~7469, and NGC~7771 with values of $\\log L_{\\rm IR}$ similar to those given in Sanders et al. (2003). Column~(6): Nuclear activity class from spectroscopic data (``Sy''=Seyfert, ``L''=LINER, ``H\\,{\\sc ii}''=H\\,{\\sc ii} region-like, ``NO''=no classification was possible, ``--''=no data available). Column~(7): References for the spectroscopic data: 1. Veilleux et al. (1995); 2. Corbett et al. (2003); 3. Kewley et al. (2001); 4. V\\'eron-Cetty \\& V\\'eron  (2001); 5. Baan, Salzer, \\& LeWinter (1998); 6. L\\'{\\i}pari et al. (2000); 7. Garc\\'{\\i}a-Mar\\'{\\i}n et al. (2006). The Sy2 classification is for B1 in NGC~3690; 8. van den Broek  et al. (1991); 9. Sekiguchi \\& Wolstencroft (1992). Column~(8): Morphological description.}  \\end{deluxetable*}     In this paper we present {\\it Hubble Space Telescope} ({\\it HST}) NICMOS continuum (1.1, 1.6, and $1.876\\,\\mu$m) and Pa$\\alpha$ emission line observations of a volume limited sample (30 systems, 34 individual galaxies) of local universe ($z<0.017$) LIRGs. The paper is organized as follows. \\S2 gives details on the sample selection and properties. In \\S3 we describe the {\\it  HST}/NICMOS observations, data reduction, and continuum and H\\,{\\sc ii} region photometry. In \\S4 we discuss the near-IR continuum properties of the central regions of local LIRGs. In \\S5, and \\S6,we describe the Pa$\\alpha$ morphology, and the properties of individual H\\,{\\sc ii} regions, respectively. In \\S7 we estimate the extinction to the Pa$\\alpha$ emitting regions. In \\S8  we estimate the extended H$\\alpha$ emission, and in \\S9 we discuss the star formation rates of local LIRGs. Our conclusions are given in \\S10. Throughout this paper we use $H_0=75\\,$km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have analyzed {\\it HST}/NICMOS $1.1-1.89\\,\\mu$m continuum and Pa$\\alpha$ emission line observations of a volume limited sample of 30 local universe LIRGs ($\\log L_{\\rm IR}=11-11.9\\,[{\\rm L}_\\odot]$).  The galaxies have been selected so that the Pa$\\alpha$ emission line could be observed with the NICMOS F190N narrow-band filter ($2800 <v < 5200\\,$km s$^{-1}$). The sample comprises approximately 80\\% of all the LIRGs in the RBGS in the above velocity interval, and is representative of local LIRGs in general. The NICMOS observations cover the central $\\simeq 3.3-7.2\\,$kpc of the galaxies. We find the following:  \\begin{enumerate}  \\item The most common morphological continuum features in the central regions are bright star clusters, and large-scale spiral arms, sometimes extending down to the inner kpc (mini-spirals). In highly inclined systems there are dust lanes crossing the disks of the galaxies and in some cases, even hiding the nucleus of the galaxy. The  ``extreme'' (perturbed) morphologies commonly seen in ULIRGs are only observed in some of the most IR luminous examples in our sample.  \\item Approximately half of the LIRGs show compact ($\\sim 1-2\\,$kpc) Pa$\\alpha$ morphologies in the form of nuclear star formation rings, mini-spiral structure, and emission associated with the nucleus of the galaxy. The typical observed (not corrected for extinction) H$\\alpha$ surface brightnesses are between $2-10 \\times 10^{41}\\,{\\rm erg \\, s}^ {-1}$ kpc$^{-2}$, although the most extreme cases, for instance the ring of star formation of NGC~1614, can reach $\\simeq 60 \\times 10^{41}\\,{\\rm erg \\, s}^ {-1}$ kpc$^{-2}$.  \\item The remaining half of the galaxies show Pa$\\alpha$ emission extending over scales of $\\simeq 3.3-7.2\\,$kpc and larger, with bright H\\,{\\sc ii} regions in the spiral arms (face-on systems) or along the disks of the galaxies (edge-on systems), with or without bright nuclear emission. The H$\\alpha$ surface brightnesses are about one order of magnitude fainter than in galaxies with compact Pa$\\alpha$ emission.  \\item Most of the LIRGs show a population of numerous bright H\\,{\\sc ii} regions, and in about half of them the typical H\\,{\\sc ii} region (the median H$\\alpha$ luminosity) is as bright as or brighter than the giant H\\,{\\sc ii} region 30 Dor. The most IR  (as well as in H$\\alpha$) luminous galaxies tend to host the brightest median and first-ranked H\\,{\\sc ii} regions.   \\item The extinctions to the gas (derived from the H$\\alpha$/Pa$\\alpha$ and Pa$\\alpha$/Br$\\gamma$ line ratios) and to the stars (derived from the $m_{\\rm F110W}-m_{\\rm F160W}$ colors) are on average $A_V\\sim 2-6\\,$mag over the Pa$\\alpha$ emitting regions covered by the {\\it HST}/NICMOS images.  \\item There exists a good correlation between the mid-IR $24\\,\\mu$m and the extinction-corrected Pa$\\alpha$ luminosities of LIRGs, ULIRGs, normal galaxies, and H\\,{\\sc ii} regions of M51, covering nearly five decades in luminosity.  This suggests that the mid-IR luminosity of galaxies undergoing dusty, intense star formation is a good indicator of the SFR. The increasing $L(24\\,\\mu{\\rm m})/L({\\rm Pa}\\alpha)_{\\rm corr}$ ratio for more luminous Pa$\\alpha$ emitters (that is, galaxies with  higher SFRs) may be due to the presence of deeply embedded sources in the youngest star-forming regions for which we have underestimated the extinction, even at near-IR wavelengths.  \\item Analogous to the widely used relation between SFR and total IR luminosity (Kennicutt 1998), we show that  SFRs can be determined accurately in luminous, dusty galaxies as:  SFR(M$_\\odot$ yr$^{-1}) = 8.45 \\times 10^{-38}$ (L(24$\\mu$m)/(erg s$^{-1}$))$^{0.871}$   \\end{enumerate}   $\\,$\\\\  Based on observations with the NASA/ESA Hubble Space Telescope 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. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration (NASA).  We thank an anonymous referee for a careful reading of the manuscript and useful suggestions. We would like to thank D. Calzetti and R. Kennicutt for enlightening discussions. We are grateful to T. Hattori for providing us with H$\\alpha$ images of the galaxies. We thank T. D\\'{\\i}az-Santos for helping us with the {\\it Spitzer}/MIPS images of the sample of LIRGs. AAH and LC acknowledge support from the Spanish Programa Nacional de Astronom\\'{\\i}a y Astrof\\'{\\i}sica under grant AYA2002-01055 and Plan Nacional del Espacio ESP2005-01480, and PGPG from the Spanish Programa Nacional de Astronom\\'{\\i}a y Astrof\\'{\\i}sica under grant AYA 2004-01676. This work has been funded by NASA grant HST-GO-10169 and by NASA through contract 1255094 issued by JPL/Caltech."
    },
    "0606/astro-ph0606465_arXiv.txt": {
        "abstract": "We report the discovery of two highly dispersed pulsars in the direction of the Galactic Centre made during a survey at 3.1~GHz with the Parkes radio telescope. Both PSRs~J1745$-$2912 and J1746$-$2856 have an angular separation from the Galactic Centre of less than 0.3\\degr\\ and dispersion measures in excess of 1100~cm$^{-3}$pc, placing them in the top 10 pulsars when ranked on this value. The frequency dependence of the scatter-broadening in PSR~J1746$-$2856 is much shallower than expected from simple theory. We believe it likely that the pulsars are located between 150 and 500~pc from the Galactic Centre on the near side, and are part of an excess population of neutron stars associated with the Centre itself. A second survey made at 8.4~GHz did not detect any pulsars. This implies either that there are not many bright, long-period pulsars at the Galactic Centre or that the scattering is more severe at high frequencies than current models would suggest. ",
        "introduction": "Surveys for radio pulsars have been extremely successful over the past decade \\cite{lor05}. All-sky surveys and deep surveys of the Galactic plane have doubled the total number of known radio pulsars to above 1700. At the same time, targeted surveys have discovered very young pulsars in supernova remnants, old millisecond pulsars in globular clusters and a population in the Small and Large Magellanic Clouds. In spite of these successes, no pulsars are known within $\\sim$1\\degr\\ of the Galactic Centre (GC).  Following the Johnston et al. (1995)\\nocite{jwv+95} targeted search of the GC at 1500~MHz (no detections), the most sensitive search at 1500 MHz is the Parkes multibeam survey which integrated for 35~mins over a large fraction of the Galactic plane including the GC. That survey discovered PSR~J1747$-$2802, which has a period of 2.8~s, the smallest angular separation from the GC of all the known pulsars ($\\sim$1\\degr) and a high dispersion measure (DM) of 835~cm$^{-3}$pc \\cite{mhl+02}. At higher frequencies, the most sensitive survey was undertaken with the Effelsberg telescope which surveyed a small region (0.2 degrees or 32 pc radius) around the GC at 4850~MHz \\cite{kkl+00,kle04}. In spite of a sensitivity below 100~$\\mu$Jy and the high observing frequency the survey failed to discover any new pulsars.  Sgr A* is at the heart of our Galaxy and is believed to be a black hole with a mass of $3\\times 10^6 M_{\\odot}$ \\cite{egos02} at a distance of $\\sim$8~kpc. A nearby stellar cluster appears to contain early-type stars with masses of $10-20 M_{\\odot}$. Two of the stars in this cluster are in highly eccentric orbits about the black hole with periods of 15 and 30 years \\nocite{gsh+05}. It is likely that shorter orbits exist, however confusion in the infra-red has so far prevented their detection. The presence of these high mass stars and young supernova remnants is a good indicator that active radio pulsars are also likely to exist close to the GC. Pfahl \\& Loeb (2004)\\nocite{pl04} estimate that there are 1000 radio pulsars within a few pc of the GC, a small fraction of which are potentially detectable using current instruments provided that scatter broadening does not render them invisible at any sensible observing frequency (see also Cordes \\& Lazio 1997\\nocite{cl97}). However, it is likely that the entire GC volume out to $\\sim$200~pc contains an overabundance of pulsars generally. Estimates show that about 10\\% of all the high mass stars (the pulsar progenitors) in the Galaxy are contained within this volume \\cite{frk+04}.  One can therefore expect to detect a population of pulsars which would be extremely useful in probing the GC and its conditions: their number and age distribution would probe the past star formation history \\cite{har95}; their period derivatives can constrain the gravitational potential in the GC; pulsar timing would enable us to probe the space-time around the super-massive black hole in the GC due to a variety of relativistic effects \\cite{wk99}. The high stellar density of the GC makes it, like the globular clusters, a possible site of a millisecond pulsar orbiting a stellar-mass black hole though these will be extremely difficult to detect. Pulsar timing can also probe the plasma density in the central regions through measurements of the (variable) dispersion measure of a pulsar as it orbits the central black hole.  We therefore embarked on two surveys of the GC at frequencies of 3.1 and 8.4~GHz, frequencies which bracket that of the Effelsberg 5~GHz survey. The lower frequency reduces the effects of scatter broadening by a factor of $\\sim$16 compared to surveys at 1.5~GHz, but at the same time the telescope beam size is large enough to survey a significant area of sky in a reasonable amount of time. At the higher frequency, the small beam restricts the sky coverage and the flux densities of pulsars are smaller but the scatter broadening is lower by a factor of more than 1000 compared to 1.4~GHz. Cordes \\& Lazio (1997)\\nocite{cl97} identified frequencies near 8~GHz as the ideal for this type of search. ",
        "conclusions": "We have discovered two pulsars within 0.3\\degr\\ of the GC during a targeted survey at 3.1~GHz. Both pulsars have very high DMs and scatter broadening times. It seems unlikely that the survey has penetrated through the scattering screen which surrounds the GC at a distance of $\\sim$150~pc. However, the pulsars are likely to have originated in an excess population associated with the GC and be located within a few hundred pc of it. Both the DM and scattering are then higher than expected in the current electron density models perhaps favouring a filled centre model rather than the existing annular one. No detections were made at 8.4~GHz with the implication being that either there are less than 3000 pulsars beaming in our direction in the inner pc of the Galaxy or that the scattering is more severe than previously thought."
    },
    "0606/hep-ph0606199_arXiv.txt": {
        "abstract": "We present a theory of non-solar cosmic rays (CRs) in which the bulk of their observed flux is due to a single type of CR source at all energies. The total luminosity of the Galaxy, the broken power-law spectra with their observed slopes, the position of the `knee(s)'  and `ankle', and the CR composition and its variation with energy are all predicted in terms of very simple and completely `standard' physics.  The source of CRs is extremely `economical': it has only one parameter to be fitted to the ensemble of all of the mentioned data. All other inputs are `priors', that is, theoretical or observational items of information independent of the properties of the source of CRs, and chosen to lie in their pre-established ranges. The theory is part of a `unified view of high-energy astrophysics' ---based on the `Cannonball' model of the relativistic ejecta of accreting black holes and neutron stars. The model has been extremely successful in predicting all the novel properties of Gamma Ray Bursts recently observed with help of the Swift satellite. If correct, this model is only lacking a satisfactory theoretical understanding of the `cannon' that emits the cannonballs in catastrophic processes of accretion. ",
        "introduction": "\\label{s1}  The field of cosmic-ray (CR) physics was born as a lucky failure. The 1912 attempt by Victor Hess to measure the {\\it decrease} of the Earth's radioactivity in an ascending balloon gave an opposite result: there was an extra-terrestrial source of what are now known to be high-energy nuclei and electrons. Almost a century later, the origin of non-solar CRs is still a subject of intense research and little consensus \\cite{Reviews}. We shall refer throughout to non-solar cosmic rays simply as CRs.  Over almost a century, an impressive set of CR data have been gathered, e.g.~the {\\it all-particle} spectrum (of nuclei, without distinction of charge and mass) has been measured over some 13 orders of magnitude in energy and more than 30 orders of magnitude in flux (perhaps only Coulomb's law is measured over an even wider range). It has become standard practice to present the spectral data as the flux $dF/dE$ times a power of energy, which emphasizes the spectral `features' and the discrepancies between experiments, while de-emphasizing the pervasive systematic errors in energy. The all-particle spectrum $E^3\\,dF/dE$ is shown in Fig.~\\ref{features} for energies $E>10^{11}$ eV. The figure shows that the spectrum is roughly describable as a broken power law \\cite{Reviews,Kampert,HIRES1}: \\begin{eqnarray} &&\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;{dF / dE} \\propto E^{-\\beta},\\nonumber\\\\ \\beta&\\simeq& 2.7;\\;\\;\\; E<E[{\\rm knee}]\\sim 3\\times10^{15}\\rm\\;eV, \\nonumber\\\\ \\beta&\\simeq& 3.0;\\;\\;\\; E[{\\rm knee}]\\,<\\,E<\\,E[{\\rm knee2}] \\sim 2\\times10^{17}\\rm\\;eV, \\nonumber\\\\ \\beta&\\sim& 3.1;\\;\\;\\; E[{\\rm knee2}]\\,<\\,E<\\,E[{\\rm ankle}], \\nonumber\\\\ \\beta& \\sim& 2.7;\\;\\;\\; E>E[{\\rm ankle}]\\sim 3\\times10^{18}\\rm\\;eV. \\label{flux} \\end{eqnarray}  \\begin{figure}[] \\centering \\epsfig{file=DDAllFeaturesbrief.eps,height=5.3cm,width=8.8cm} \\vspace*{5pt} \\caption{The all-particle  CR spectrum wheighted by $E^3$, with its three changes of spectral index: the `knee', the `second knee', and the `ankle' \\cite{Kampert}.} \\label{features} \\end{figure}  Below ${E[{\\rm knee]}}$, protons constitute  $\\sim 96\\%$ of the CRs at fixed energy {\\it per nucleon}. Their flux above ${E_p\\sim 10}$ GeV is~\\cite{Haino}: \\begin{equation} {dF_p\\over dE}\\simeq {1.37\\pm 0.13\\over\\rm cm^{2}~s~sr~GeV}\\;  \\left [E \\over \\rm GeV \\right]^{-2.73\\pm 0.03}\\, . \\label{protons} \\end{equation}  The conventional theory of CRs \\cite{GP} posits that supernova remnants are the site of acceleration of (non-solar) CRs for energies up to $E[{\\rm knee}]$. No consensus on a preferred accelerator site or mechanism exists for energies between $E[{\\rm knee}]$ and $E[{\\rm ankle}]$. It has long being argued that CRs of energy above $E[{\\rm ankle}]$ are {\\it extragalactic} in origin \\cite{Cocconi,Morrison}: they cannot be isotropized by the Galactic magnetic fields, but their observed arrival directions are  isotropic \\cite{HIRES2}. We refer to CRs with  $E>E[{\\rm ankle}]$ as ultra-high energy cosmic rays (UHECRs). They are the subject of great interest, considerable controversy and imaginative model-building; see, for instance, the review by J. Cronin in \\cite{Reviews}. In this paper we use many abbreviations. They are listed in Table \\ref{table0}.   \\begin{table} \\caption{Frequently used abbreviations \\label{table0}} \\begin{tabular}{lclc} \\colrule Afterglow(s) & AG(s) \\\\ Active Galactic Nucleus(i) & AGN(s) \\\\ Cosmic Background Radiation & CBR \\\\ Cosmic Microwave Background & CMB \\\\ $[$Non-solar$]$ Cosmic Ray(s) & CR(s)  \\\\ Cosmic-Ray Electron(s) & CRE(s) \\\\ Gamma Background Radiation & GBR \\\\ $[$Long-duration$]$ Gamma-Ray Burst(s)& GRB(s) \\\\ Inter-Galactic Medium & IGM \\\\ Inter-Stellar Medium & ISM \\\\ Inverse Compton Scattering & ICS \\\\ Greisen, Zatsepin \\& Kuzmin & GZK \\\\ Lorentz Factor(s) & LF(s) \\\\ Magnetic Field(s) & MF(s) \\\\ Short Hard [$\\gamma$-ray] Burst(s) & SHBs \\\\ Starlight & SL \\\\ Superbubble(s) & SB(s) \\\\ Supernova(e) & SN(e) \\\\ Supernova Remnant(s) & SNR(s) \\\\ Synchrotron Radiation & SR \\\\ Ultra-High Energy Cosmic Ray(s) & UHECR(s) \\\\ X-Ray Flash(es) & XRF(s) \\\\ \\end{tabular} \\end{table}  Radio, X-ray and $\\gamma$-ray observations of supernova remnants (SNRs) provide clear evidence that electrons are accelerated to high energies in these sites. So far, they have not provided unambiguous evidence that SNRs accelerate CR nuclei and are their main source in any energy range \\cite{Plaga}. Moreover, SNRs cannot accelerate CRs to energies as large as $E[{\\rm knee}]$ \\cite{Hillas}, though this point is still debated. A direct proof ---such as a localized source--- of an extragalactic origin of the UHECRs was lacking until very recently \\cite{Reviews,HIRES2}. The precise interpretation of the recent results of Auger \\cite{PAC} on the correlation of directions between UHECRs and Active Galactic Nuclei (AGNs), which we discuss in detail in Sections \\ref{Auger} and \\ref{AugerAGNcorr}, may still be debatable.  There is mounting observational evidence that, in addition to the ejection of a non-relativistic spherical shell, the explosion of a core-collapse supernova (SN) results in the emission of highly relativistic bipolar jets of plasmoids of ordinary matter, {\\it Cannonballs}. Evidence for the ejection of such jets in SN explosions is not limited to GRBs but comes also from optical observations of SN 1987A \\cite{NP}, from X-ray \\cite{UH} and infrared \\cite{OK} observations of Cassiopeia A and, perhaps, from the morphology of radio SNRs \\cite{RNM}. These jets may be the main source of CR nuclei at all energies \\cite{DP,DAD,Florence}. They also explain long-duration GRBs \\cite{DETAL}, as advocated in the CB model \\cite{GRB1,DD}.  In this paper we elaborate on a previous theory of CRs \\cite{DP}, which is very different from the conventionally accepted theories \\cite{Reviews}. For much of the required input, we exploit the subsequently acquired information provided by the CB-model analysis of long-duration $\\gamma$-ray bursts (GRBs) and X-ray flashes (XRFs). The jets of CBs responsible for GRBs are akin to the jets of CBs emitted by quasars and microquasars. The former jets, we shall argue, are also responsible for the generation of CRs.  The essence of our considerations may be pictorially conveyed. The quasar Pictor A is shown in Fig.~\\ref{Pictor}. The X-ray picture in the top panel shows one of its extremely narrow jets, which we interpret as X-ray emission from a jet of CBs. The lower panel shows the two opposite jets, and contour plots of their radio-emission fluence. We interpret the radio signal as the synchrotron radiation of `cosmic-ray' electrons. Electrons and nuclei were scattered by the CBs, which encountered them at rest in the intergalactic medium (IGM), kicking them up to high energies. Thereafter, these particles diffuse in the ambient magnetic fields (MFs) and the electrons efficiently emit synchrotron radiation. In applying this picture to the CRs in our Galaxy, we will simply replace the quasar for all past Galactic and extragalactic core-collapse SNe, and fill in the details.   \\begin{figure}[] \\centering \\vbox{% \\epsfig{file=PictorAjet_mod.eps,height=4.4cm,width=7.7cm} } \\vbox{% \\epsfig{file=PictorRadio1_mod.eps,height=8.4cm,width=7.35cm,angle=-90} } \\vspace*{-5pt} \\caption{Top: X-ray image of the  galaxy Pictor A. A non-expanding jet extends across 360000 light years towards a hot spot at least 800000 light years away from where the jet originates \\cite{PG}. Bottom: XMM/p-n image of Pictor A in the 0.2--12 keV energy interval, centred at the position of the leftmost spot in the upper panel, superimposed on the radio contours of a 1.4 GHz radio VLA map \\cite{WYS}.} \\label{Pictor} \\end{figure}  As a CB from a core-collapse SN travels through the interstellar medium (ISM), it encounters ISM matter that has been previously ionized by the passage of the GRB's $\\gamma$ radiation. The CB's density is low enough for individual interactions between its ionized plasma constituents and those of the ISM to be irrelevant. The ISM ions and electrons are only deflected by the collective effects of the CB's inner MFs, generated by the very same ions and electrons. We shall see in detail that this makes a CB act as a formidably efficient {\\it relativistic magnetic-racket} accelerator, which loses essentially all of its energy to the recoiling particles: the newly born CRs. We argue that this very simple concept explains all observed properties of non-solar CRs at all observed energies.  Cosmic-ray sources other than high-energy jets ---such as the traditional expanding SN envelopes, novae, stellar flares, stellar winds and non-relativistic jets--- may be relevant at low energies. Galactic high-energy CRs are also emitted by ordinary pulsars, by soft $\\gamma$-ray repeaters, by microquasars, and probably in the final merger of neutron stars and black holes in binary systems.  The total CR luminosity of these objects is smaller than the observed one by more than two orders of magnitude. Similar considerations lead us to neglect, or to discuss {\\it cum grano salis,} the extragalactic contribution of relativistic jets from massive black holes in AGNs, perhaps the most luminous potentially competitive sources. These topics are discussed in Section \\ref{AugerAGNcorr} and Appendix \\ref{Accelerators}.    Our predictions for the $E^3$-weighted fluxes of the most abundant nuclei and `groups' of nuclei are shown in Fig.~\\ref{Groups}, which previews and summarizes our results. The source spectra are shown in Fig.~\\ref{Groups}a; two of their features are: `knees' at energies proportional to the atomic number $A$; and `Larmor' cutoffs (proportional to the nuclear charge $Z$) beyond which our CR acceleration mechanism is no longer operative. The CR spectra arriving to our planet are shown in Fig.~\\ref{Groups}b. The differences between these two figures ---which are significant and will be discussed in minute detail--- are due to the many {\\it `tribulations'} a CR suffers in travelling to Earth from the location of its source. Two examples of tribulations are:  (1) Below a certain momentum (some $3\\times 10^9\\; Z$ GeV/c) the local flux of CRs of Galactic origin is enhanced by a factor proportional to their momentum-dependent Galactic `confinement' time \\cite{Swordy}: \\begin{equation} \\tau_{\\rm conf} \\propto \\left({Z\\,{\\rm GeV/c}\\over p}\\right)^{\\beta_{\\rm conf}}; \\,\\,\\,\\,\\,\\,\\,{\\beta_{\\rm conf}}\\sim 0.6\\pm 0.10. \\label{cintrod} \\end{equation} This is the origin of the differing slopes of the lower-energy fluxes in Figs.~\\ref{Groups}a,b (note their different scales).  (2) Extragalactic CRs other than protons are efficiently photo-dissociated by the cosmic background infrared radiation on their way to our Galaxy. This is part of the explanation for the very different relative abundances of the elements at the higher energies in Figs.~\\ref{Groups}a and \\ref{Groups}b.  One can see in Fig.~\\ref{Groups}b how H and He dominate up to their knees, which add up to the knee feature of the all-particle spectrum; how the composition thereafter becomes `heavier' and dominated by Fe up to its knee, which is the second knee of the all-particle spectrum; and how the flux becomes once again `lighter'  above this feature. The UHECRs are entirely extragalactic in origin and are dominantly protons.  \\begin{figure} \\begin{center} \\epsfig{file=DDGroups4.eps, width=8.5cm,angle=-0} \\end{center} \\caption{Predicted CR spectra for the most abundant elements and groups. The vertical scales are $E^3\\,dF/dE$. (a): The source spectra, with a common arbitrary normalization. (b): The CR spectra at the location of the Earth. Notice that both the horizontal and vertical scales are different.} \\label{Groups} \\end{figure}  Ours is also a theory of CR electrons and of their radiative contribution to the diffuse $\\gamma$ background radiation (GBR). The diffuse GBR at low Galactic latitudes originates mainly from $\\pi^0$-generating collisions of CR nuclei with the ISM, followed by $\\pi^0\\to 2\\gamma$ decay. At high Galactic latitudes the diffuse GBR, we contend \\cite{DDGBR}, is dominated by inverse-Compton radiation from CR electrons in the ISM and the halo of galaxies (including ours) and in the IGM. In this sense, CR nuclei, CR electrons and a good fraction of the diffuse GBR have the same origin \\cite{Anton}, the latter radiation being a CR `secondary'.  More specific results, to be derived in detail and in agreement with the data, are: \\begin{itemize}  \\item The slope of the CR source spectra below the knees is predicted to be $\\beta_s\\simeq 2.17$, as derived in Section \\ref{belowtheknee}. Modified by Galactic confinement, this corresponds to observed spectra with a slope $\\beta_s+{\\beta_{\\rm conf}}\\approx 2.77$, with ${\\beta_{\\rm conf}}$ as in Eq.~(\\ref{cintrod}).  \\item A very slight composition dependence of the slope of the source spectra is expected. This can be discerned in Fig.~\\ref{Groups}. It is discussed in Section \\ref{composlopes} and shown in Fig. \\ref{ArnonSlopes}.  \\item The CR spectra at the lowest energies are affected by solar effects. The predictions agree with the data at minimum solar activity; see Section \\ref{lowenergies} and Fig.~\\ref{VeryLowEnergy}.  \\item The spectra of the individual CR nuclei are predicted to have `knees', scaling as the atomic weight, at energies around $ E[{\\rm knee}]\\!\\sim\\! 3\\!\\times\\! 10^{15}\\,A\\,{\\rm eV}$; see Eq.~(\\ref{Eknee}) and Figs.~\\ref{Groups},~\\ref{ProtonGEG}. The observed and predicted spectra for the individual elements H, He and Fe, at energies up to their knees, are shown in Fig.~\\ref{KASKADE}.  \\item The CR spectrum is predicted to change rather abruptly in slope, dominant composition (Fe to H) and dominant origin (Galactic to extragalactic) at the `ankle' energy, $E[{\\rm ankle}]\\sim 3\\times 10^{18}\\;{\\rm eV}$, see Section \\ref{ankleregion}.   \\item Our CR acceleration mechanism has a cutoff at the energies of Eq.~(\\ref{toe}), proportional to atomic charge and roughly coincident with the conventional Greisen--Zatsepi--Kuzmin (GZK) cutoff \\cite{GZK}.  These cutoffs do not seem to be present in the AGASA data \\cite{AGASA}, but are compatible with the  HIRES \\cite{HIRES2}, Fly's Eyes' and Auger \\cite{Mello,Unger,PAC} data, which agree well with our theory: see Fig. \\ref{UHECR}.  \\item The predicted normalization of the UHECR flux is approximate but `absolute', i.e.~parameter-free; see Section \\ref{UHECRtext}.  \\item The prediction for the  all-particle spectrum is compared with the data in Fig.~\\ref{AllPart}.  \\item Detailed predictions for the `primary' CR abundances relative to hydrogen are discussed in Section  \\ref{relabundances}. They are compared with data  at 1 TeV in Table \\ref{table1} and  illustrated in Fig.~\\ref{f1}. \\item Data on two rough indicators of the evolution of CR composition with energy, $\\langle \\ln A \\rangle$ and the depth of shower maximum $X_{\\rm max}$, are compared with predictions in  Figs.~\\ref{lnAknee} and \\ref{Xmax}. \\item The confinement volume and the confinement time of CRs in the Galaxy can be estimated theoretically. They agree with the estimates extracted from observations, as discussed in Section \\ref{timevolume}. \\item Below their respective knees, the source spectra of CR nuclei and electrons are predicted to have the same slope: $\\beta_s=13/6$. For relatively high-energy electrons, radiation cooling steepens the slope to $\\beta_e=\\beta_s+1\\sim 3.17$. The observed slope is $\\simeq 3.2$, as shown in Fig.~\\ref{CREspectrum}. The normalization of the electron spectrum, we cannot predict. \\item The slope of the diffuse GBR is predicted to be $(\\beta_e+1)/2\\simeq 2.08$. The observation is $\\simeq 2.1$, as shown in Fig.~\\ref{GBRspectrum}. \\end{itemize}   Admittedly, the `predictions' we have referred to in the above items are `postdictions' of existing data. Yet, the theory on which they are based is very `predictive': only one parameter specific to the CR source will be fitted to the hadronic CR data.  Otherwise, only {\\it priors} (items of information independent of the CR source) have been used as inputs, and kept at their `central' values, or within their error brackets~\\cite{cosmopriors}.  The study of GRBs, some 40 years old, is in its infancy, if compared with the century-old study of CRs. In the GRB realm, novel and very precise data, in particular at X-ray energies and mainly thanks to the Swift satellite, are being gathered. The predictions of the CB model have been precisely verified by the data having appeared since our first posting of this paper in June 2006. This subject is very briefly summarized in Section \\ref{Swift}.  A posteriori the distinction between post- and pre-dictions, or parameters and priors, is somewhat artificial. But there are other assets of the CR theory presented here: it works simply and very well, and it is based on {\\it a single source of CR acceleration at all energies}. Moreover, the underlying theory ---originally inspired by an analogy with the relativistic ejecta of quasars and microquasars--- is part of  a unified model of high-energy astrophysical phenomena \\cite{Florence}, which also offers simple and successful explanations of the origin and properties of `long-duration' GRBs and X-ray flashes and their respective afterglows (AGs) \\cite{DD,AGoptical,AGradio,DDDXRF}, the natal kicks of neutron stars \\cite{DP}, the MFs and radio emission from within and near galaxy clusters  \\cite{DDMF}, and the X-ray emission from  galactic clusters allegedly harbouring `cooling flows' \\cite{CDD}.  The many titles and subtitles in this paper should suffice to convey its organization. We discuss in detail or summarize in Appendices some of the relevant background information: how a CB expands, photo-dissociation, the least debatable `priors' common to all theories of CRs, jets in astrophysics, the CB model, the evidence for the ejection of relativistic jets in SN explosions, the supernova--GRB association and the power supply by CR accelerators other than the one we propose.   Our main point is our proposed mechanism of CR acceleration. A reader primarily interested in it may choose to read first Section \\ref{collisionless} on `Collisionless magnetic rackets'. A reader primarily intrigued by the results may choose to start with Chapter \\ref{CBresults}. ",
        "conclusions": "We have sketched a theory wherein cosmic rays are ions of the interstellar medium, encountered by relativistic CBs and magnetically kicked up to higher energies ---either  elastically, or `inelastically' (i.e., after a succession of accelerating encounters with the CBs' inner turbulent magnetic fields). The elastic component is entirely analogous to the mechanism which ---we contend--- generates the prompt $\\gamma$-rays of a GRB:  `inverse' Compton scattering, by the electrons comoving with a CB, of the `ambient light' they encounter around their parent exploding star. The inelastic CR component is also analogous to the high-energy tail of the spectra of GRBs \\cite{DD}, originating from a small fraction of electrons that have been accelerated within a CB. In this sense, our theory of CRs is but a straightforward generalization to cosmic rays of the very successful CB model of GRBs. All we have done is to substitute the scattering of `ambient' light by the scattering of `ambient' ions and electrons.  Our theory agrees with the classic proposal of Baade and Zwicky \\cite{BZ} that SN explosions are at the origin of CRs. But our mechanism is different from that of the generally-accepted CR theory, in which it is the non-relativistically-expanding SN shells ---as opposed to their relativistic jets--- that accelerate relatively low-energy CRs. We have demonstrated how our simple {\\it and single} accelerators ---cannonballs--- are effective at all observed energies.  Our intention in this paper was not to reproduce the CR observations in minute detail, or to refine at maximum those of the conventional inputs that could be refined ---such as the details of the photo-dissociation of UHECRs. Thus, we have chosen to minimize the number of fit parameters to a grand total of one. The rest of the required input has been gathered from information  independent of the theory of the source of CRs, or fixed from the simplest choices we could make at each point.  Most of our results are `robust' in that ---within very large brackets--- they do not depend on the specific choices of parameters and priors: \\begin{itemize} \\item{} An all-particle power-law spectrum with four successive features: two steepenings at the knee and the second knee, a softening at the ankle, and an end-point at the GZK and proton-acceleration cutoffs, which are roughly coincident. \\item{} An UHECR flux above the ankle, which is predicted ---to within a factor of a few--- and otherwise parameter-free. \\item{} A composition dependence at 1 TeV with the observed trend, so different from that of the ISM for the relative abundances of H and He versus those of the heavier elements. \\item{} A very low-energy flux whose spectral shape is independent of any CB-model `prior' parameters. \\item{} Individual-element knees that scale like $A$ and occur at the predicted energies. \\item{} A non-trivial shape of the individual knees: an abrupt decrease in flux, followed by a spectrum steeper than that below the knee. \\item{} An ankle with the observed shape. The dominantly Galactic-Fe flux below it and the dominantly extragalactic-proton flux above it are comparable in magnitude at the estimated escape energy of Galactic protons. \\item{} A composition dependence that is almost energy-independent below the knee, becomes `heavier' from the knee to the second knee, `lighter' again above it, and finally heavier at yet-unmeasured ultra-high energies. \\item{} An `extended' distribution of CR sources along CB trajectories that emerge from the central realms of the Galaxy, where most SN explosions take place, implying a CR flux at the Earth's position with a much smaller and less energy-dependent anisotropy than that of conventional SNR models of CRs. \\item{} Predictions for the values of a consistently related set of observables: the CR luminosity, confinement time and volume of the Galaxy, the spectral indices of CR electrons and of the diffuse GBR. \\end{itemize} Our results describe the observed properties of hadronic non-solar CRs very well from the lowest energies  to $\\sim 10^{10}$ GeV. Above that energy and up to the highest observed energies, $\\sim 10^{11}$ GeV, our theory opts for the data gathered with fluorescence detectors, corroborated by hybrid detectors such as Auger. Overall, the energy range for which the theory is successful covers ten decades and the flux extends over  three times as many.  The CR theory we have discussed is currently incomplete in various respects. The confinement of CRs, either in the Galaxy or in the CBs themselves, is not well understood, but our assumptions on the subject ---the simplest--- appear to work very well. There is insufficient information on the Galactic CR wind to model its effects with confidence; we have had to experiment with various limiting possibilities. We have argued that CR diffusion need not be explicitly considered, but we have not proved that to be the case, by considering it in detail. Moreover, our theory is based on a two-stage acceleration: that of CBs by core-collapse SNe, and that of CRs by CBs. For the former, we have relied on observations, rather than on a deeper understanding.  In spite of the above limitations, our claims are supported by the simplicity of the theory, its extreme economy of free parameters, and by the good quality of its description of the fairly elaborate ensemble of CR data. The theory is an item in a more general understanding of high-energy astrophysical phenomena, including cooling flows, large-scale magnetic fields, GRBs and XRFs. On the two later topics, there has been great observational progress since the first posting of this article, which very precisely corroborated the CB model of GRBs and XRFs.  The astronomy of non-thermal light sources, from radio frequencies to TeV energies, as well as high-energy neutrino astronomy, are the studies of the interactions of CR nuclei and electrons with ambient matter and magnetic fields. We have illustrated this point by discussing the Gamma `Background' Radiation as a CR `secondary', and by briefly commenting on high-energy $\\gamma$-ray and neutrino astronomy.  We contend that we have identified the acceleration mechanisms promoting the constituents of the interstellar medium of the Milky Way, and other galaxies, to become the bulk of non-solar cosmic rays of all energies, and the `magnetic-racket' accelerators themselves: the cannon balls emitted by a large fraction of ordinary core-collapse supernovae. \\\\ \\\\ {\\bf Aknowledgements} We are indebted to Giuseppe Cocconi, Andy Cohen, Shlomo Dado, Friedrich Dydak, Shelly Glashow, Karl-Heinz Kampert, Etienne Parizot and Rainer Plaga for discussions,   and to the last of them for his comments on the manuscript. This research was supported in part by the Helen Asher Space Research Fund at the Technion Institute. AD thanks the TH division at CERN for its hospitality. ADR thanks the Technion Institute at Haifa for the same reason.  \\appendix"
    },
    "0606/astro-ph0606471_arXiv.txt": {
        "abstract": "Medium-resolution optical spectroscopy of the candidate YSOs associated with the small, nearby molecular cloud Lynds~1333 revealed four previously unknown classical T~Tauri stars, two of which are components of a visual double, and a Class~I source, \\textit{IRAS}~02086+7600. The spectroscopic data, together with new $V,R_{C},I_{C}$ photometric and 2MASS $J$, $H$, and $K_s$ data allowed us to estimate the masses and ages of the new T~Tauri stars. We touch on the possible scenario of star formation in the region. L\\,1333 is one of the smallest and nearest known star forming clouds, therefore it may be a suitable target for studying in detail the small scale structure of a star forming environment. ",
        "introduction": "\\label{Sect_1}  Filamentary molecular clouds with embedded dense cores form a remarkable subset of star forming clouds in our galactic environment \\citep[e.g.][]{Onishi96,NC05,HTPG05}. Young stellar objects (YSOs) are associated with several cores along the filaments. The formation scenario of the filaments and stars within them, however, are not well understood. The filaments may be parts of shells, swept up by powerful stellar winds or supernovae \\citep*[e.g.][]{KMT04}, or may result from fragmentation of sheet-like structures \\citep{Hartmann02}, or may be shaped by large-scale flows like the galactic rotation \\citep{Koda}. Detailed studies of their density and velocity structures, as well as the properties of the YSOs born in them may help understand  their formation and evolution.  Lynds~1333, a small dark cloud of opacity class 6 (Lynds 1962) in Cassiopeia, at (l,b)=(128\\fdg88,+13\\fdg71) is part of a filamentary complex. According to the available observations L\\,1333 is starless, and thus has been included in several studies of starless cores \\citep*[e.g.][]{LMT99,LMT01,Lee04}.  \\citet*[][hereinafter referred to as Paper~I]{OKS}, studied first this cloud. They derived a distance of $180\\pm20$\\,pc from the Sun using Wolf diagram method. Their $^{13}$CO and C$^{18}$O observations have shown L\\,1333 to be part of a long, filamentary molecular structure, stretching from $l \\sim 126\\degr$ to $133\\degr$ and from $b \\sim +13\\degr$ to  $+15\\degr$, and referred to this molecular complex as {\\em L\\,1333 molecular cloud\\/}. The angular extent of the molecular complex corresponds to a length of some 30\\,pc at a distance of 180\\,pc. \\citet{KMT04}  found that the L\\,1333 complex is part of a giant far infrared loop GIRL~G\\,126+10.  Recent star formation in the L\\,1333 molecular cloud complex has been indicated by the presence of the \\textit{IRAS}~source \\textit{IRAS}~02086+7600, whose  \\textit{IRAS} colour indices are indicative of a Class~I protostar, nevertheless it coincides with a faint star in the  \\textit{Digitized Sky Survey} image. Due to its appearance as an optically visible star with large far-infrared excess several authors considered this object as a possible evolved star. \\citet*{FNP} included \\textit{IRAS}~02086+7600 in a multiband photometric survey for candidate post-AGB stars. They could not confirm the post-AGB nature of the star, and noted that it may be an ultracompact \\mbox{H\\,{\\sc ii}} region, or a post-AGB star, or a YSO. \\textit{IRAS}~02086+7600 appeared as a possible planetary nebula in the target lists of \\citet{PM88} and \\citet{VSP95}.  Based on its \\textit{IRAS} colours, \\citet{Slysh} included this object, as a candidate ultracompact \\mbox{H\\,{\\sc ii}} region, in their search for OH maser emission. They detected it as a thermal OH source at the velocity of 3.1km\\,s$^{-1}$. The molecular maps presented in Paper~I revealed that this \\textit{IRAS} source is projected on a dense C$^{18}$O core of a nearby molecular cloud whose radial velocity is +3.0\\,km\\,s$^{-1}$, same as that of the OH source, suggesting that \\textit{IRAS}~02086+7600 most probably is a low-mass YSO. The C$^{18}$O spectrum observed at its position exhibited a wing-like feature, indicative of molecular outflow (Paper~I). No known Herbig--Haro object is associated with this source.  In addition to \\textit{IRAS}~02086+7600, 18 H$\\alpha$ emission stars have been detected in objective prism Schmidt plates in the region of L\\,1333 (Paper~I). Three of these stars are associated with the \\textit{IRAS} point sources \\textit{IRAS}~F02084+7605, 02103+7621, and 02368+7453.  The aim of our present study is to establish an elementary data base on the star forming activity of the L\\,1333 complex. We observed the optical spectra of the candidate YSOs in order to establish their pre-main-sequence nature and their spectral types. We also performed optical photometry of the objects in order to determine their luminosities and positions in the HRD. We describe our observational data in Sect.~\\ref{Sect_2}. Our results on the properties of the observed stars, a short description of the large-scale environment of the cloud, and the possible star formation scenario are presented in Sect.~\\ref{Sect_3}.  Sect.~\\ref{Sect_4} gives a short summary. ",
        "conclusions": "\\label{Sect_4}  We identified five low-mass YSOs in the small filamentary molecular complex associated with the dark cloud Lynds~1333. Their masses are in the interval 0.15--0.8\\,M$_{\\sun}$, and they are 1--5 million years old. We confirmed that \\textit{IRAS}~02086+7600 is a Class~I YSO associated with the L\\,1333 complex, and found its age to be comparable to those of the CTTSs born in the same cloud. The relative distribution of YSOs and clouds suggests that the star formation might have been triggered by the collision of high velocity gas with Loop~III."
    },
    "0606/astro-ph0606192_arXiv.txt": {
        "abstract": "Motivated by the claimed discovery of a very massive galaxy (HUDF-JD2; $M \\simeq 5 \\times 10^{11}\\, {\\rm M_{\\odot}}$) at extreme redshift ($z = 6.5$) within the Hubble Ultra Deep Field (HUDF) (Mobasher et al. 2005), we have completed a systematic search for comparably massive galaxies with $z > 4$ among the 2688 galaxies in our $K_S < 23.5\\, (AB)$ catalogue within the CDFS/GOODS-South field. This search was conducted using redshift estimates based on the recently-completed, uniquely-deep 11-band ($B,V,i,z,J,H,K_S,3.6\\mu m,4.5\\mu m,5.8\\mu m,8.0\\mu  m$) imaging in this 125 square arcmin field, $\\simeq 25$ times larger than the NICMOS HUDF. To ensure completeness, our approach places no special emphasis on the standard $V$-drop, $i$-drop or $z$-drop criteria commonly used to pre-select candidate high-redshift galaxies.  Initial spectral fitting, based on published catalogue SExtractor photometry, led us to conclude that at least 2669 of the galaxies in our sample lie at $z < 4$. This list includes several galaxies for which redshifts $z > 4$ have been previously proposed. We carried out a detailed investigation of the 19 remaining $z > 4$ candidates, performing aperture photometry on all images, and including marginal detections and formal non-detections in the fitting process. This led to the rejection of a further 13 galaxies to lower redshift. Moreover, subjecting HUDF-JD2 to the same analysis, we find that it lies at $z \\simeq 2.2$, rather than the extreme redshift favoured by Mobasher et al. (2005).  The 6 remaining candidates appear to be credible examples of galaxies in the redshift range $z = 4 - 6$, with plausible stellar ages. However, refitting with allowance for extreme values of extinction we find that, even for these objects, statistically acceptable solutions can be found at $z < 3$. In fact only 2 galaxies retain formally preferred high-redshift solutions. Moreover, the recently-released Spitzer MIPS imaging in GOODS-South has revealed that 5 of our 6 final $z > 4$ candidates are detected at 24$\\mu m$. This was also the case for HUDF-JD2 (Mobasher et al. 2005), and provides further circumstantial evidence in favour of the moderate-redshift dusty solutions. We conclude that there is no convincing evidence for any galaxy with $M > 3 \\times 10^{11}\\, {\\rm M_{\\odot}}$ and $z > 4$ within the 125 square-arcmin GOODS-South field. We briefly discuss the implications of this null result, and revised expectations for the much larger (0.8 sq. degree), and deeper near-infrared UKIDSS Ultra Deep Survey now underway with WFCAM on the UKIRT. ",
        "introduction": "Several hundred convincing galaxy candidates have now been uncovered at $z \\simeq 5-6.5$ ({\\it e.g.} Bunker \\& Stanway 2004; Taniguchi et al. 2005; Ouchi et al. 2005; Shioya et al. 2005; Yan et al. 2005; Bouwens et al. 2006), with a few (somewhat less convincing) galaxy candidates even reported at $z > 7$ (Bouwens et al. 2004). The discovery of such objects has been used to set interesting new constraints on the cosmic history of star-formation density ({\\it e.g.} Stark \\& Ellis 2005, Bouwens \\& Illingworth 2006), but has not, as yet, presented a serious challenge to current theories of galaxy formation. This is because the masses of essentially all of these objects are relatively modest ($M \\simeq 10^{10} {\\rm M_{\\odot}}$), and the observed large numbers of such objects are consistent with the predictions of at least some current galaxy-formation models (e.g. Nagamine et al. 2006).  By contrast, the discovery of even a small number of very massive galaxies at these extreme redshifts can present a stern challenge  for both semi-analytic and hydrodynamic galaxy-formation models. Indeed, given the steep decline in the predicted number density of high-mass halos at $z > 4$ (e.g. Somerville 2004), the discovery of a significant number density of very massive objects at such redshifts has the potential to provide an interesting test of the now well-established paradigm of hierarchical structure growth within $\\Lambda CDM$.  For this reason, two recent studies have generated a lot of interest. First, Eyles et al. (2005), in their detailed study of 3 spectroscopically confirmed Lyman-break galaxies at $z \\simeq 5.5 - 6$, reported that these objects already contained a substantial, evolved mass of stars, apparently formed at redshifts as high as $z \\simeq 7.5 - 13.5$. Second, Mobasher et al. (2005) presented apparently convincing evidence for the existence of an extremely massive galaxy (HUDF-JD2; $M \\simeq 5 \\times 10^{11}\\, {\\rm M_{\\odot}}$) lying within the NICMOS Hubble Ultra Deep Field (HUDF) at the extreme redshift of $z = 6.5$. Although this is just one object, its discovery within such a very small-area survey such as the HUDF is undoubtedly surprising and has already generated considerable interest ({\\it e.g.} Panagia et al. 2005).  Motivated by the discovery of HUDF-JD2, we decided to revisit our existing redshift determinations for $K_S$-band selected galaxies in GOODS-South (Caputi et al. 2004; 2005; 2006), and to conduct a systematic search for the existence of {\\it any} galaxies in this sample at very high redshift ($z > 4$). The key point here is that HUDF-JD2 is sufficiently massive that it is relatively bright in the near-infrared, with $K_S = 23.9$. Thus it should be possible to detect comparable objects (less than a factor of 1.5 more massive at $z \\simeq 6.5$) within our complete $K_S < 23.5$ sample of 2898 objects (galaxies and stars) in the GOODS-South field, which covers a solid angle $\\simeq 25$ times larger than that subtended by the NICMOS HUDF. It is also clear that, in this extreme mass domain, the discovery of even one object within the entire GOODS-South field would be extremely important.  Given the apparent evidence for some moderately evolved stellar populations in high-redshift galaxies, and because we wished to search for any galaxies over a relatively wide redshift range $4 < z < 8$, we decided not to base our search on strict colour criteria. Such criteria are frequently adopted in the selection of high-redshift galaxies, in order to ensure minimal contamination from low-redshift interlopers (e.g. Bunker \\& Stanway 2004). However, especially at $z = 4-5$, such clean selection inevitably comes at the expense of completeness, and introduces an inevitable bias in favour of the youngest, and hence lowest mass:light ratio galaxies. Instead, we updated the dataset ultilised by Caputi et al. (2006) with the addition of the recently released ISAAC $H$-band and complete Spitzer 4-band IRAC data, and then derived new redshift estimates for all 2688 galaxies in the $K_S < 23.5$ sample by fitting a range of single and double component spectrophotometric models to the full 11-band optical-infrared photometry.  The full results of this process are described in Cirasuolo et al. (2006) in which we have utilised this new dataset to explore the cosmological evolution of the galaxy mass function out to $z = 4$. In the current paper we use this work simply as a starting point for the careful study of the relatively small subsample of galaxies which were found to have even marginally convincing solutions at $z > 4$.  The paper is structured as follows. In Section 2 we briefly describe how we have utilised the updated public dataset within the GOODS-South field to create a revised, complete sample of 2688 galaxies with $K_S < 23.5$ and full, aperture-matched 11-band HST/ACS+VLT/ISAAC+Spitzer/IRAC photometry. In Section 3 we describe the redshift-estimation technique, and demonstrate that it yields robust solutions at the correct redshifts for known, spectroscopically-confirmed galaxies at $z \\simeq 5-6$. Then, in Section 4, we describe the results of applying this technique to isolate a maximal sample of 32 potential $z > 4$ galaxies within our GOODS-South sample, (using 2.8-arcsec diameter aperture photometry) which was then refined down to a subsample of 19 serious candidates using smaller aperture SExtractor (Bertin \\& Arnout 1996) measurements. Section 5 then presents the results of a detailed investigation of the multi-frequency data for these remaining 19 candidates, with manual aperture photometry leading to the rejection of a further 13 objects to lower redshift. In this section we also show that, by subjecting HUDF-JD2 to the same treatment, we find that it does not lie at $z = 6.5$ (as concluded by Mobasher et al. 2005), but is in fact a dusty, evolved galaxy at $z \\simeq 2.15$. The final interrogation of our remaining six candidate $z > 4$ massive galaxies in presented in Section 6. In particular, we explore the robustness of the derived redshifts when the assumed optical extinction is allowed to float up to values as large as $A_V = 6$. Finally, in Section 7 we briefly discuss the implications of our results in the context of theoretical models of structure formation.  All optical and infrared magnitudes are given in the $AB$ system (Oke \\& Gunn 1983). Masses and ages have been calculated assuming a cosmological model with $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70\\,{\\rm km s^{-1} Mpc^{-1}}$.  \\begin{figure*} \\vspace*{15cm} \\special{psfile=fig1a.ps hoffset=20 voffset=170 hscale=35 vscale=35 angle=0} \\special{psfile=fig1b.ps hoffset=300 voffset=170 hscale=35 vscale=35 angle=0} \\special{psfile=fig1c.ps hoffset=20 voffset=-50 hscale=35 vscale=35 angle=0} \\special{psfile=fig1d.ps hoffset=300 voffset=-50 hscale=35 vscale=35 angle=0} \\caption{\\small Spectral fits and $\\chi^2$ versus redshift $z$ for the four galaxies in the HUDF and GOODS-South field with measured spectroscopic redshifts $z > 5$, and published 7-band photometry as tabulated by Eyles et al. (2005) and Stark et al. (2006). Details of the model fits, and a comparison of estimated and spectroscopic redshifts are given in Table 1.} \\end{figure*} ",
        "conclusions": "To illustrate the implications of this failed search for high-mass galaxies at $z > 4$ we show, in Figure 6, a plot of the comoving number density of galaxies with mass $M > 3 \\times 10^{11} {\\rm M_{\\odot}}$ within the GOODS-South field, as a function of redshift. This limit simply corresponds to the lowest mass found for the high-redshift solutions to the 6 galaxies listed in Table 6.  The data points for the redshift bins $0 < z < 1$, $1 < z < 2$, $2 < z < 3$, and $3 < z < 4$ are derived from the analysis of Cirasuolo et al. (2006) (updated from Caputi et al. 2006), while the reference point at $z = 0$ has been derived by the appropriate integration of the galaxy mass function provided by Cole et al. (2001) (assuming a Salpeter IMF). The result of the unsuccessful search at $z > 4$ described here is illustrated by the upper limit plotted at $z = 5$, which represents the comoving number density in the redshift bin $4 < z < 6$ if {\\it one} of the 6 final high-redshift candidates listed in Table 6 really does lie within this high-redshift bin. Also plotted in this figure are the comoving number density of virialized dark-matter halos above 3 different mass thresholds, as derived from modified Press-Schechter theory (Percival, priv. comm.).  Several features of this diagram are worthy of comment. First, at no redshift does the number density of high-mass galaxies present a fundamental problem for $\\Lambda$CDM; the number density of potential dark-matter halos (10-20 times more massive than the stellar masses of the galaxies) clearly exceeds the inferred number density of the massive galaxies.  Nevertheless, out to redshift $z \\simeq 3$, the number density of these most massive galaxies changes only slowly (compared to the dark matter curves) and hence, for this class of galaxies, the inferred dark-matter:stellar mass ratio appears to evolve from $\\simeq 80$ at $z \\simeq 0$ to $\\simeq 20$ at $z \\simeq 3.5$. Such inferred values do not appear unreasonable (see, for example, Mandelbaum et al. 2006), and this apparent evolution of the dark-matter:stellar mass ratio can be viewed as yet another manifestation of `downsizing' or apparently `anti-hierarchical' galaxy formation (e.g. Heavens et al. 2004).  However, the upper limit derived here indicates that this `anti-heirarchical' behaviour does not persist beyond $z \\simeq 4$, and that, at higher redshift, the number density of massive galaxies drops off as rapidly (or possibly more rapidly) than the number density of potential host dark-matter halos.  An improved measurement of the comoving number density of these rare high-mass galaxies at high-redshift clearly requires a substantially larger survey than the 125 square arcmin covered by the GOODS-South multi-wavelength imaging. The first such survey of the necessary depth and area is now underway. This is the UKIDSS Ultra Deep Survey (UDS), which covers 0.8 square degrees, and is designed to ultimately reach a $K$-band 5-$\\sigma$ detection limit of $K = 25$ (AB). This survey will therefore cover 25 times the area of GOODS-South, to a magnitude limit substantially deeper than the $K_S$-limit of the GOODS-South survey analysed here. A failure to find any galaxies with $M > 3 \\times 10^{11} {\\rm M_{\\odot}}$ and $z > 4$ within the UKIDSS UDS would move the upper limit shown at $z \\simeq 5$ in Figure 6 down by over an order of magnitude. A first analysis, based on the UKIDSS early data release, has provided evidence for a number of moderately massive galaxies at $ z > 5$, but as yet has not revealed any as massive as $M > 3 \\times 10^{11} {\\rm M_{\\odot}}$ (McLure et al. 2006).  \\begin{figure} \\vspace*{7.5cm} \\special{psfile=fig6.ps hoffset=-40 voffset=-28 hscale=55 vscale=55 angle=0} \\caption{\\small A comparison of our best estimate of the evolution of the comoving number density of galaxies with stellar mass $M > 3 \\times 10^{11} {\\rm M_{\\odot}}$ with the predicted evolution of the comoving number density of virialized dark-matter halos above 3 different mass thresholds (as derived from modified Press-Schechter theory; Percival, priv. comm.). The data points for the redshift bins $0 < z < 1$, $1 < z < 2$, $2 < z < 3$, and $3 < z < 4$ are derived from the analysis of Cirasuolo et al. (2006) (updated from Caputi et al. 2006), while the reference point at $z = 0$ has been derived by the appropriate integration of the galaxy mass function provided by Cole et al. (2001) (assuming a Salpeter IMF). The result of the unsuccessful search at $z > 4$ described here is illustrated by the upper limit plotted at $z = 5$, which represents the comoving number density in the redshift bin $4 < z < 6$ if {\\it one} of the 6 final high-redshift candidates listed in Table 6 really does lie within this high-redshift bin. The curves show the comoving number density of virialized dark-matter halos with $M > 5 \\times 10^{12} {\\rm M_{\\odot}}$ (solid line), $M > 1 \\times 10^{13} {\\rm M_{\\odot}}$ (dashed line), and $M > 2.5 \\times 10^{13} {\\rm M_{\\odot}}$ (dotted line). The last (most massive) of these curves was deliberately chosen to coincide with the data-point at $z = 0$, and implies a dark-matter:stellar mass ratio of $\\simeq 80$ for these most massive galaxies in the present-day universe.} \\end{figure}"
    },
    "0606/astro-ph0606647_arXiv.txt": {
        "abstract": "We attempt to document complete energetic transactions of stars in their life.  We calculate photon and neutrino energies that are produced from stars in their each phase of evolution from 1 to 8${\\rm M}_\\odot$, using the state-of-the-art stellar evolution code, tracing the evolution continuously from pre-main sequence gravitational contraction to white dwarfs. We also catalogue gravitational and thermal energies and helium, and heavier elements that are stored in stars and those ejected into interstellar space in each evolutionary phase. ",
        "introduction": "Stars are taken as an engine that uses hydrogen as a fuel, producing energy mostly in photons and leaving helium and heavier elements as ashes. Baryons in stars are only 6\\% of those in the Universe, yet stars cause varieties of most spectacular phenomena, and also give a key to understanding the energetics of the Universe (\\citealp{fp04}; hereafter FP). FP presented an accounting of energies that are relevant to stars using rather rough estimates in view of the difficulty in assembling the relevant astrophysical quantities. There are many calculations of stellar evolution done to date, some of the most recent examples being the Geneva tracks \\citep{gen92,char96} and the Padova tracks \\citep{gir00} among others, but quantities like integrated luminosities and gravitational binding energies stored in stars are not explicitly presented.  The purpose of this paper is to present evolution of stars mainly from the energetics point of view. For example, how much radiation energy is produced in main sequence phase, how much in red giant branch and asymptotic branch phases, and how much energy is stored in heavy elements as well as in gravitational binding, and so forth? The energies in radiation observed in the extragalactic background light must balance those produced by stars besides a small fraction contributed by active galactic nuclei, and they are documented in the nuclear binding energy as a fossil record. We also want to know how much fraction of helium and heavy element produced in stars are liberated into interstellar space. Then, what fraction of energies go to neutrinos, which account for the difference between the nuclear binding energy and the energy in radiation?  We focus our calculations on low and intermediate mass stars from 1 to 8 solar masses that eventually end up as white dwarfs, and we follow the evolution from pre-main sequence contraction to the cool white dwarf with luminosity $\\log(L/L_\\odot)=-4.5$. This enables us to compute integrated output of stars in a consistent way.  We also add the case for 0.8${\\rm M}_\\odot$, calculated to the beginning of helium ignition to allow a continuation to lower mass stars. We consider stars with solar metallicity.  We remark that white dwarfs are the most important reservoir of heavy elements today, which stores about 80\\% of heavy element in the Universe. The gravitational binding energy stored in white dwarfs is 5 times that in main sequence stars in the cosmic mean. Stellar winds in the thermal-pulse (TP) AGB phase are also the important source of chemical enrichment besides supernovae. Thus, white dwarfs are one of the most prominent fossil records of energy transactions in the Universe (FP). ",
        "conclusions": "The central result of this paper is Table 1, which documents products of stars integrated over time for each phase. We consider stars from 1 to 8 ${\\rm M}_\\odot$. This can be used to link baryon mass locked into stars to extragalactic background light, and also to predict chemical enrichment.  The table also tells us the relative importance of each phase for the light emitted by stars.  Some discussion is made to understand the trend of numbers we present in the table.  When the numbers in the table are integrated over the initial mass function, they give the output of photon and neutrino energies from respective phases of stellar evolution and helium and CNO enrichment per given mass of cooled baryons, although the calculation still must be supplemented by a similar table for more massive stars. When they are further integrated over star formation history, we can estimate differential contributions to the energetics of the galaxies and Universe. Such procedures will make the work similar to that of population sysnthesis model, the most recent example of which is Bruzual \\& Charlot (2003). The difference lies in the fact that the latter is specifically oriented for the calculation of the detailed spectral output of the stellar radiation, while our aim is to calculate all outputs in a gross form from the stellar population divided into each stage."
    },
    "0606/astro-ph0606537_arXiv.txt": {
        "abstract": "Low state optical photometry of EF Eri during an extended low accretion state combined with GALEX near and far UV time-resolved photometry reveals a source of UV flux that is much larger than the underlying 9500K white dwarf, and that is highly modulated on the orbital period. The near UV and optical light curves can be modeled with a 20,000K spot but no spot model can explain both the large amplitude FUV variations and the SED. The limitations of limb darkening, cyclotron and magnetic white dwarf models in explaining the observations are discussed. ",
        "introduction": "Accreting close binaries containing white dwarfs with high (10-200 MG) magnetic fields are termed AM Herculis variables or simply polars (review in Wickramasinghe \\& Ferrario 2000). They tend to have lower mass transfer rates than disk accreting systems of similar orbital period and spend perhaps the majority of their time (Ramsay et al. 2004) in states where the mass transfer can drop by orders of magnitude or stop completely (termed low states). During these times, most of the effects associated with active accretion (X-rays, strong optical cyclotron continuum, strong He and Balmer emission lines) disappear. This provides the opportunity to directly view and study the underlying stars (white dwarf and late type secondary). The polar cataclysmic variable EF Eri received notoriety as it remained in one of these inactive states for 9 years, starting in 1997 (Wheatley \\& Ramsay 1998; WR98) and only becoming active again at the beginning of 2006.  In the high state, EF Eri is 14th mag and shows a highly variable light curve over the 81 min orbital period (Bailey et al. 1982). In the infrared, the J light curve revealed a narrow dip which Bailey et al. used to define  phase 0.0. At this phase, the $UBV$ light curves showed a broad minimum, the X-ray had a sharp minimum and circular polarization was a maximum. These features were interpreted as an eclipse of the accretion column by the accretion stream at phase 0. Cyclotron humps led to estimates of 16.5 and 21 MG for two cyclotron regions.  When EF Eri entered its prolonged low state in 1997, the V mag dropped to 18, broad Balmer absorption lines from the white dwarf were apparent, and Zeeman splitting indicated a magnetic field of 14 MG (WR98). While Euchner et al. (2003) fit multipole models (5 components with fields up to 100MG) to spectra obtained in 2000, the fit was not definitive. WR98 concluded the similarity of the low state field to that measured in the small accretion regions during the high state indicated either a uniform field strength over the white dwarf, or continued accretion heated spots. Beuermann et al. (2000) modeled the optical spectrum with a 9500K white dwarf with a 15,000K hot spot covering 6\\% of the surface. $BVRI$ photometry obtained by Harrison et al. (2003; H03) in 2001, showed sinusoidal light curves with a peak near phase 0.9 and a minimum near phase 0.4. They modeled the light curves with a 12,000K spot covering 6\\% of the area of a 9500K white dwarf, with a best fit obtained for an inclination of 35$^\\circ$ and an angle of 35$^\\circ$ between the spin and magnetic pole. The $H$ and $K$ light curves were anti-phased from the optical, with a peak near phase 0.5 and a minimum near phase 0. IR spectra obtained in 2002 (Harrison et al. 2004; H04) showed different cyclotron harmonics were present at phase 0 than at phase 0.5, indicating that cyclotron emission was still present from both accretion poles.  Given that the optical spectrum of EF Eri at the low state shows no obvious sign of an accretion stream, it is intriguing that a hot spot model still fits the optical light curves and cyclotron is still present in the IR. During a study of polars at low states, Araujo-Betancar (2005) found that 30,000-70,000K hot spots contributing 20-30\\% of the white dwarf flux were needed to model the FUV fluxes. To further explore the heating effects at extremely low accretion rates, we obtained UV light curves with GALEX. Our results show the value of GALEX time-resolved photometry in providing information confirming that some area of the white dwarf contributes a large amount of UV flux even after 7 years of extremely low mass transfer. ",
        "conclusions": "Our results show that GALEX time-resolved photometry is a very useful means of studying the hot components in close binaries. Our near and far UV light curves obtained during the extended low accretion state of EF Eri reveal the presence of a symmetrical, large and highly variable source of UV light. If this is a hot spot on the 9500K white dwarf, this spot must have a temperature $>$20,000K to fit the SED. However, the amplitude of the FUV variation is far too large to be fit with a spot of this type, unless large limb darkening is invoked just for the FUV. Since the UV light curves are in phase with the optical, and opposite in phase from IR variations that are known to originate from cyclotron harmonics, it is natural to invoke changing optical depth of harmonics from optically thick in the IR to optically thin in the optical/UV. If there was a harmonic present in the far UV band, it could explain the excess FUV flux and amplitude, but this would require extremely high fields and harmonics should be visible in the optical as well. It is intriguing that small deviations in the UV and optical light curves match the phases of peak IR flux and that the excess UV luminosity over that of the 9500K white dwarf is comparable to the cyclotron luminosity in the IR. However, the lack of good models for cyclotron effects on the spectrum of a white dwarf prevents us from reaching any definite conclusion. Until we can correctly model the observed excess UV flux in the cases of high magnetic field white dwarfs, we cannot ascertain its cause as a hot spot or cyclotron emission or atmospheric effects due to high fields and low accretion."
    },
    "0606/astro-ph0606701_arXiv.txt": {
        "abstract": "In this paper we present the first $BVI$ CCD photometry of six overlooked old open clusters (Berkeley~44, NGC~6827, Berkeley~52, Berkeley~56, Skiff~1 and Berkeley~5) and derive estimates of their fundamental parameters by using isochrones from the Padova library (Girardi et al. 2000). We found that all the clusters are older than the Hyades, with ages ranging from 0.8 (NGC~6827 and Berkeley~5) to 4.0 (Berkeley~56) Gyr. This latter is one of the old open clusters with the largest heliocentric distance.\\\\ In the field of Skiff~1 we recognize a faint blue Main Sequence identical to the one found in the background of open clusters in the Second and Third Galactic Quadrant, and routinely attributed to the Canis Major accretion event. We use the synthetic Color Magnitude Diagram method and a Galactic model to show that this population can be easily interpreted as Thick Disk and Halo population toward Skiff~1. We finally revise the old open clusters age distribution, showing that the previously suggested peak at 5 Gyr looses importance as additional old clusters are discovered. ",
        "introduction": "The present day age distribution of open star clusters in the Galactic disk is the result of the two competing processes: the star formation history of the Galactic disk and the dissolution rate of star clusters (de la Fuente Marcos \\& de la Fuente Marcos 2004).\\\\ The dissolution of star clusters is particularly important for the older clusters, the typical open cluster lifetime being of the order of 200 Myr (Wielen 1971). This way  to trace back the cluster formation history in the Galactic disk is a challenging task. Recent compilations (Friel 1995, Ortolani et al. 2005) show that the age distribution of old open cluster has an e-folding shape with a possible peak at 5 Gyrs. The reality of this peak is however quite difficult to assess, and indeed a more recent analysis (Carraro et al. 2005, Fig.~10) shows that the inclusion of a few overlooked clusters significantly weakens the reality of this peak, and illustrates the importance of a careful hunting of old clusters before drawing definitive conclusions. The recent study of the old cluster Auner~1 (Carraro et al. 2006) with an age of 3.5 Gyr again stresses the fact that we are still missing several old clusters.\\\\  Beginning with the paper of Phelps et al. (1994) several attempts have been made to enlarge the sample of studied old open clusters (Hasegawa et al. 2004, Carraro et al. 2005 and references therein).  In an attempt to further contribute to this interesting field, in this paper we present the first photometric study of 6 overlooked, old open clusters having $53^o\\leq l \\leq 130^o$ and $-5^o.2\\leq b \\leq +5^o.6$ (see Table~1) and provide homogeneous derivation of basic parameters using the Padova (Girardi et al. 2000) family of isochrones.\\\\ These clusters are NGC~6827, Berkeley 5, 52, 44, and 56 (Setteducati \\& Weaver 1960) and Skiff 1 (Luginbuhl \\& Skiff 1990).   \\noindent The plan of the paper is as follows. Sect.~2 describes the observation strategy and reduction technique. Sect.~3 deals with star counts and radius determination. The Color-Magnitude Diagrams (CMD) are described in Section~4, while Section~5  illustrates the derivation of the clusters' fundamental parameters. Section~6 concentrates on the star cluster Skiff~1. Finally, Sect.~7 provides a detailed discussion of the results.     \\begin{table} \\caption{Basic parameters of the clusters under investigation. Coordinates are for J2000.0 equinox and have been visually re-determined by us.} \\begin{tabular}{ccccc} \\hline \\hline \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{$RA$}  & \\multicolumn{1}{c}{$DEC$}  & \\multicolumn{1}{c}{$l$} & \\multicolumn{1}{c}{$b$} \\\\ \\hline & {\\rm $hh:mm:ss$} & {\\rm $^{o}$~:~$^{\\prime}$~:~$^{\\prime\\prime}$} & [deg] & [deg]\\\\ \\hline Berkeley 44    & 19:17:12 & +19:28:00 &  53.21 & +3.35\\\\ NGC 6827       & 19:48:54 & +21:12:00 &  58.25 & -2.35\\\\ Berkeley 52    & 20:14:18 & +28:58:00 &  67.89 & -3.13\\\\ Berkeley 56    & 21:17:42 & +41:54:00 &  86.04 & -5.17\\\\ Skiff 1        & 00:58:24 & +68:28:00 & 123.57 & +5.60\\\\ Berkeley 5     & 01:47:48 & +62:56:00 & 129.29 & +0.76\\\\ \\hline\\hline \\end{tabular} \\end{table}  \\begin{figure} \\centerline{\\psfig{file=fig1.eps,width=\\columnwidth}} \\caption{V  60 secs image centered on Berkeley 44. North is up, east on the left. The field is 10 arcmin on a side.} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=fig2.eps,width=\\columnwidth}} \\caption{V  60 secs image centered on NGC 6827. North is up, east on the left.} \\end{figure}   \\begin{figure} \\centerline{\\psfig{file=fig3.eps,width=\\columnwidth} } \\caption{V  60 secs image centered on Berkeley 52. North is up, east on the left.} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=fig4.eps,width=\\columnwidth}} \\caption{V  60 secs image centered on Berkeley 56. North is up, east on the left.} \\end{figure}   \\begin{figure} \\centerline{\\psfig{file=fig5.eps,width=\\columnwidth} } \\caption{V  60 secs image centered on Skiff 1. North is up, east on the left.} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=fig6.eps,width=\\columnwidth} } \\caption{V  60 secs image centered on Berkeley 5. North is up, east on the left.} \\end{figure} ",
        "conclusions": "We have presented CCD BVI photometry for 6 previously unstudied possibly old open clusters, namely Berkeley~44, NGC~6827, Berkeley~52, Berkeley~56, Skiff~1 and Berkeley~5.\\\\  \\noindent We have found that all the clusters are actually old, and the ages range from 0.8 to 4 Gyr. This sample of clusters represents an important contribution to the poorly populated old open clusters family in the Galactic disk. In Fig.~18 we show an updated age distribution of the old open cluster (older than 500 Myr) so far known.\\\\ This comes from Carraro et al. (2005), where we added the new clusters  studied in this paper and Auner~1 (3.5 Gyr, Carraro et al. 2006).\\\\  The new age distribution can be easily fitted with an exponential relation having an e-folding time of 2 Gyr. This means than on the average the oldest clusters in the Milky Way do not survive more than 2 Gyr. This estimate is an order of magnitude larger than the typical life-time of an open cluster (200 Myr), and suggests that old open clusters survive longer possibly due to particular situations, like birth-places high onto the Galactic plane, or the preferentially high total mass at birth. It might also possible that some open clusters, especially in the anti-centre could have entered the Milky Way in the past together with cannibalized satellites (Frinchaboy et al. 2004)\\\\  Much firmer conclusions might be drawn as additional old clusters are discovered and studied.   \\begin{table*} \\caption{Parameters of the studied clusters. The coordinate system is such that the Y axis connects the Sun to the Galactic Center, while the X axis is positive in the direction of galactic rotation. Y is positive toward the Galactic anti-center, and X is positive in the first and second Galactic quadrants (Lynga 1982).} \\fontsize{8} {10pt}\\selectfont \\begin{tabular}{ccccccccccc} \\hline \\multicolumn{1}{c} {$Name$} & \\multicolumn{1}{c} {$Radius$} & \\multicolumn{1}{c} {$E(B-V)$}  & \\multicolumn{1}{c} {$(m-M)$} & \\multicolumn{1}{c} {$d_{\\odot}$} & \\multicolumn{1}{c} {$X_{\\odot}$} & \\multicolumn{1}{c} {$Y_{\\odot}$} & \\multicolumn{1}{c} {$Z_{\\odot}$} & \\multicolumn{1}{c} {$R_{GC}$} & \\multicolumn{1}{c} {$Age$} \\\\ \\hline & ${\\prime}$& mag & mag& kpc & kpc & kpc & pc & kpc & Myr \\\\ \\hline Berkeley~44  & $\\geq 5.0$ & 1.40$\\pm$0.10  & 15.6$\\pm$0.2 &  1.8 &  1.4 &  -1.1 &   100 &  7.6 & 1300$\\pm$200 \\\\ NGC~6827     & 1.5        & 1.05$\\pm$0.05  & 16.3$\\pm$0.2 &  4.1 &  3.5 &  -2.1 &  -170 &  7.3 &  800$\\pm$100 \\\\ Berkeley~52  & 1.5        & 1.50$\\pm$0.10  & 18.1$\\pm$0.2 &  4.9 &  4.5 &  -1.8 &  -270 &  8.1 & 2000$\\pm$200 \\\\ Berkeley~56  & 1.0        & 0.40$\\pm$0.05  & 16.6$\\pm$0.2 & 12.1 & 12.0 &  -0.8 & -1100 & 14.3 & 4000$\\pm$400 \\\\ Skiff~1      & $\\geq$5.0  & 0.85$\\pm$0.05  & 13.7$\\pm$0.2 &  1.6 &  1.3 &   0.9 &   160 &  9.5 & 1200$\\pm$100 \\\\ Berkeley~5   & 1.0        & 1.30$\\pm$0.10  & 18.0$\\pm$0.2 &  6.2 &  4.8 &   3.9 &    80 & 13.3 &  800$\\pm$100 \\\\ \\hline \\end{tabular} \\end{table*}"
    },
    "0606/hep-ph0606287_arXiv.txt": {
        "abstract": "We investigate cosmological evolution and implications of cosmic strings with time-dependent tension. We derive basic equations of time development of the correlation length and the velocity of such strings, based on the one scale model. Then, we find that, in the case  where the tension depends on some power of the cosmic time, cosmic strings with time-dependent tension goes into the scaling solution if the power is lower than a critical value. We also discuss  cosmic microwave background anisotropy and matter power spectra produced by these strings.  The constraints on their tensions from the Wilkinson microwave anisotropy probe (WMAP) three year data and Sloan digital sky survey (SDSS) data are also given. ",
        "introduction": "Topological defects can be produced as a result of thermal \\cite{Kibble} or nonthermal \\cite{KVY} phase transitions at the symmetry breaking in the early universe. Among them, cosmic strings have been extensively investigated mainly because they could act as the seeds for large scale structure. Recent observations of the cosmic microwave background (CMB) anisotropy \\cite{WMAP} reveal that cosmic strings cannot be the primary source of primordial density fluctuations, and give the bound on the tension of cosmic strings $G\\mu < 2.7 \\times 10^{-7}$ at $95\\%$ confidence level \\cite{PWW}, which is obtained by considering a hybrid scenario of adiabatic fluctuations generated during inflation plus isocurvature fluctuations produced by cosmic strings. The constraint given in other analysis such as Refs.~\\cite{Fraisse} and \\cite{Seljak} are also in good agreement with theirs.  However, cosmic strings can still affect various astrophysics \\cite{VSHK} such as the early reionization \\cite{ER}, gravitational radiation \\cite{GR}, gravitational lensing effects \\cite{GL}, the neutrino masses \\cite{BMY}, and so on. Furthermore, an interesting possibility is revived that fundamental strings of the superstring theory can be expanded to cosmological sizes and act as cosmic strings \\cite{CSS}. In fact, F-strings and D-strings are produced at the end of brane inflation \\cite{BI}.  In almost all research so far, the tensions of cosmic strings have been assumed to be constant. Recently however, one of the authors (M.Y.) pointed out that tensions of cosmic strings can depend on the cosmic time \\cite{MY}. For example, a complex scalar field $\\phi$, which allows a string solution, couples to another field $\\chi$, \\beq V(\\phi,\\chi) = \\frac{\\lambda}{4}(|\\phi|^2 - \\chi^2)^2 + \\frac12 m_{\\chi}^2 \\chi^2. \\eeq When a coupling constant $\\lambda$ is small enough, the backreaction to the oscillation of $\\chi$ is negligible. Then, tension of a string $\\mu$ is determined by the root mean square of the expectation value of $\\chi$ and given by $\\mu \\propto a^{-3}$ ($a$:the scale factor), which is proportional to $t^{-3/2}$ in the radiation dominated era and $t^{-2}$ in the matter dominated era. Hence, the tension $\\mu$ depends on the power of the cosmic time. Furthermore, in the warped geometry, the tension depends on the position of the brane in the bulk, which can depend on the cosmic time before the radion is fixed completely. Thus, it is quite natural to consider cosmic strings with time-dependent tension.  The key property of cosmological evolution of cosmic strings is scaling, which is confirmed by extensive investigations of cosmic strings with constant tension \\cite{VSHK}. It has been shown that the cosmic string network goes into the scaling regime, in which the typical length of the cosmic string network grows with the horizon scale. Then, the number of infinite strings per horizon volume is a constant irrespective of time and hence the ratio of the energy density of infinite strings to that of the background universe is constant. Thus, cosmic strings can generate scale invariant density fluctuations. Scaling property of the cosmic string network is confirmed both analytically \\cite{Kibble} and numerically \\cite{AT,BB,AS} by using the Nambu-Goto action \\cite{NG}, which can be obtained after integrating out heavy modes of particles and neglecting high curvature of the geometry.\\footnote{Scaling property of the evolution is known not only for cosmic (local) strings but also for global strings \\cite{gs} and global monopoles \\cite{gm}.}  As for cosmic strings with time-dependent tension, the followings are shown in Ref.~\\cite{MY}.  First of all, it is shown that the effective action of a cosmic string with time-dependent tension is given by the Nambu-Goto action with an additional factor for the time-dependent tension. By making use of such a Nambu-Goto-like action, the equation of motion in an expanding universe is derived and the evolution of cosmic strings with time-dependent tension is investigated. Then, it is confirmed that, in the case where the tension changes as the power $q$ of time, the string network goes into the scaling regime, in which the characteristic scale of the string network grows in proportion to the cosmic time. One should notice that the ratio of the energy density of infinite strings to that of the background universe is {\\it not} necessarily constant due to the time dependence of the tension, different from the case of cosmic strings with constant tension. However, in Ref.~\\cite{MY}, the constancy of string velocity is implicitly assumed to derive the scaling solution. In this paper, we also derive the equation of time development of string velocity based on the velocity-dependent one scale model \\cite{MS} and find a critical power of time dependence of the tension, below which the scaling property is realized and string velocity becomes almost constant.  When the tension of cosmic strings depends on time, its effects on astrophysics and cosmology can be significantly different from those of the conventional cosmic strings.  Thus we should consider the implications of such cosmic strings on various astrophysical and cosmological issues.  Among them, we study its effects on CMB and large scale structure in this paper.  Since the time dependence of the tension strongly depends on models, we investigate it in some general settings.  Here we assume the time dependence of the tension as $\\mu \\propto a^n$ and $\\mu \\propto \\tau^n$ where $\\tau$ is the conformal time and $n$ parametrize the dependence on the power.  With this assumption, we study the CMB and matter power spectrum in models with the time dependent tension. For this purpose, we modified the CMBACT code developed and made publicly available by Pogosian and Vachaspati \\cite{PV,PWW} to calculate the CMB anisotropy and the matter power spectrum induced by cosmic strings with time dependent tension. We will also give the constraint on the cosmic strings by comparing the predictions with recent cosmological observations, especially, the WMAP three-year data \\cite{WMAP} and the SDSS data \\cite{SDSS}.  We would like to stress that these constraints are important in discussing astrophysical applications such as the early reionization, gravitational radiation, gravitational lensing effects, the neutrino masses, and so on.  This paper is organized as follows. In the next section, we derive basic equations to investigate cosmological evolution of cosmic strings with time-dependent tension and obtain the condition under which the string network goes into the scaling solution. In section III, we calculate the predictions of the CMB anisotropy and the matter power spectrum induced by cosmic strings with time-dependent tension and give the constraints on their tensions from the WMAP three year results and SDSS data. Since the time dependence of the tension strongly depends on a model, we concentrate on the case that the tension changes in proportion to the power of the conformal time or the scale factor. The final section is devoted to conclusions and discussion. ",
        "conclusions": "In this paper, we have discussed cosmological evolution and implications of cosmic strings with time-dependent tension. First of all, we have derived equations of motions based on the velocity-dependent one scale model in the expanding universe. By using these equations, we investigate whether cosmic strings with time-dependent tension go into the scaling solution when the tension depends on some power of the cosmic time $\\mu \\propto t^q$. Then, we find that such strings relax into the scaling solution only when $q<1$ in the radiation domination and $q<2/3$ in the matter domination.  Then we showed the CMB and matter power spectra sourced by cosmic strings with time-dependent tension. As shown in the previous section, the spectra can be different significantly from those produced by the conventional cosmic strings with constant tension.  We have also discussed the constraints on the time-dependent tension from the Wilkinson microwave anisotropy probe (WMAP) three year data and Sloan digital sky survey (SDSS) data.  Since the time dependence of the tension strongly depends on models, here we considered it in some general settings.  For the time dependence of string tensions, we assumed that the tension $\\mu$ changes as some power of the scale factor $a$, $\\mu \\propto a^n$, and the conformal time $\\tau$, $\\mu \\propto \\tau^n$.  The constraints on $\\mu$ for various values of $n$ were given. For the cases with negative powers, the bounds on string tensions at present time become much stronger than those for the case with constant tension.  One should, however, notice that the tension can be large in the past in this case, which may have implications on the structure formation at very small scales. On the other hand, the bounds are weakened for positive powers, which may lead to large amount of gravitational radiation background. We will explorer these possibilities in future work."
    },
    "0606/gr-qc0606020_arXiv.txt": {
        "abstract": "Spatially averaged inhomogeneous cosmologies in classical general relativity can be written in the form of effective Friedmann equations with sources that include backreaction terms. In this paper we propose to describe these backreaction terms with the help of a homogeneous scalar field evolving in a potential; we call it the `morphon field'. This new field links classical inhomogeneous cosmologies to scalar field cosmologies, allowing to reinterpret, e.g., quintessence scenarios by routing the physical origin of the scalar field source to inhomogeneities in the Universe. We investigate a one--parameter family of scaling solutions to the backreaction problem. Subcases of these solutions (all without an assumed cosmological constant) include scale--dependent models with Friedmannian kinematics that can mimic the presence of a cosmological constant or a time--dependent cosmological term. We explicitly reconstruct the scalar field potential for the scaling solutions, and discuss those cases that provide a solution to the Dark Energy and coincidence problems. In this approach, Dark Energy emerges from morphon fields, a mechanism that can be understood through the proposed correspondence: the averaged cosmology is characterized by a weak decay (quintessence) or growth (phantom quintessence) of kinematical fluctuations, fed by `curvature energy' that is stored in the averaged 3--Ricci curvature. We find that the late--time trajectories of those models approach attractors that lie in the future of a state that is predicted by observational constraints. ",
        "introduction": "The fact that the spatially averaged inhomogeneous Universe does not evolve as the standard model of cosmology, furnished by a homogeneous--isotropic solution of Einstein's laws of gravity, has recently become a major topic aiming at a possible solution to the Dark Energy problem  \\cite{wetterich}, \\cite{beanetal:darkenergy}, \\cite{straumann}. While the averaging problem in relativistic cosmology has a long history, initiated by George Ellis \\cite{ellis} (references may be found in \\cite{ellisbuchert} and \\cite{buchert:static}), the backreaction terms, i.e. those averaged contributions that lead to deviations from a Friedmannian cosmology, have been considered quantitatively unimportant. Recent work that conjectures a large backreaction effect (e.g. \\cite{kolbetal}) has been accompanied both by counter conjectures (e.g. \\cite{wald}) and special solutions of the averaged Einstein equations \\cite{buchert:grgdust} that support the claim that backreaction could be made responsible for the Dark Energy gap (e.g. \\cite{buchert:darkenergy}, \\cite{nambu}, \\cite{alnes:LTB}, \\cite{rasanen:constraints,rasanen:model}, \\cite{curvatureLTB} and references and discussion in \\cite{ellisbuchert}).  In parallel to this discussion, other models have been extensively investigated, most notably quintessence models that invoke the presence of a scalar field source, this scalar field being a standard one or a phantom field (with a negative kinetic energy) \\cite{Wetterich88,Peebles2003,Faraoni2005,Jassal2006} and references in the reviews \\cite{sahnistarobinskii} and \\cite{copeland}. Those models imply the possibility of a scenario featuring a time--dependent cosmological term. Particular properties of the scalar field potential are discussed on phenomenological grounds, which can lead to an acceleration of the regional Universe and so to an explanation of Dark Energy. A further problem in conjunction with modeling a repulsive component in universe models is known as the {\\it coincidence problem}, i.e. a recent domination of this component is favoured, possibly occuring around the epoch of structure formation. Its solution also motivates the construction of models with a time--dependent cosmological term that needs no `fine--tuning'.  In this paper we also consider scalar field cosmologies, but we shift the perspective from the usual interpretation of a scalar field {\\it source} in Einstein's equations to a mean field description of averaged inhomogeneities. We see a number of advantages entailed by such a description that motivated the present investigation. First, a (homogeneous) scalar field as a model of spatially averaged (i.e. effectively homogeneous) geometrical degrees of freedom that are physically present does not need to be justified as an additional source arising from fundamental field theories; it has a well--defined physical status and, as such, does not suffer from a phenomenological parameterization, since it is constrained by Einstein's field equations. Second, inhomogeneities encoded in backreaction terms {\\it give rise to} the scalar field cosmology, fixing  its potential, its equation of state, etc.; they not only influence the evolution of the scalar field, but they determine it. Third, the proposed correspondence allows a realistic reinterpretation of quintessence models and other phenomenological approaches involving scalar fields. These approaches  must not be considered as independent alternatives, they here describe the same physics that also underlies the backreaction approach, and so can be confronted with physical constraints; in other words, if backreaction can indeed be made responsible for the late--time acceleration of the Universe, then the effort spent on quintessence models and other approaches \\cite{copeland} can be fruitfully exploited in terms of the proposed rephrasing. Fourth, since the underlying effective equations are general and do not invoke perturbative assumptions, the scalar field cosmology provides access to the whole solution space, notably to the non--perturbative regime. Finally,  subcases of exact solutions include a cosmological constant, which therefore can or cannot be included in Einstein's equations; the cosmological constant is a particular solution of the averaged inhomogeneous model without necessarily being present in Einstein's equations.  Since the scalar field introduced in this way stems from averaged geometrical degrees of freedom, i.e. kinematical backreaction terms encoded in the extrinsic curvature of spatial hypersurfaces, as well as their scalar curvature, we call it the {\\it morphon field}: it effectively models the  inhomogeneities {\\it shaping} the structure in the Universe and capturing the total effect of kinematical backreaction, i.e. the effect leading to a deviation from the kinematics of the standard model.  Speaking in terms of geometry, we may roughly say that the morphon models the geometrical vacuum degrees of freedom in Einstein's theory, which are neglected when modeling the Universe with the help of averaged matter sources only. We call the scalar field a {\\it morphon} only if it arises from the proposed mean field description; we can of course still consider other scalar fields as sources of Einstein's equations. We expect, and we shall demonstrate this for quintessence fields, that the {\\it morphon} is capable of acting like any other scalar field model,  e.g., assigned to an inflaton, a curvaton, a dilaton, etc., depending on our capability to exploit the proposed correspondence, i.e. the possibility to find  solutions of the averaged Einstein equations and to reconstruct the scalar field potential for this solution.  We here provide a first step that should help to appreciate the various possibilities, but this investigation does not claim to be exhaustive. In this line we are going to choose a minimal parameterization of the scalar field cosmology in the following sense: first, we do not include the cosmological constant and the constant curvature term in the `Friedmannian part' of the averaged equations, since both arise as subcases of special solutions for the backreaction terms; second, we choose the simplest foliation (constant lapse and vanishing shift in the ADM (Arnowitt Deser Misner) formulation of Einstein's equations) and the simplest matter model `irrotational dust' in order to exemplify the correspondence; third, we realize the correspondence for a particular class of scaling solutions with single--power laws. In this framework we show that the correspondence holds with a standard minimally coupled scalar field that can play the role of a quintessence field. If we include the full coordinate degrees of freedom as well as an inhomogeneous pressure source in a perfect fluid energy momentum tensor (or even imperfect fluid sources), more sophisticated scalar field cosmologies would arise. We comment on this possibility at the end of the paper.  The proposed correspondence can also be inverted, e.g. we may start with a known model of quintessence and try to recover the corresponding solution of the averaged Einstein equations. We shall  not exemplify this inversion in the present paper, we just note that free parameters in a given quintessence model can be determined through this correspondence.  We proceed as follows. Section 2 recalls the basic equations and relations needed for the present work. Section 3 sets up the correspondence between kinematical backreaction and scalar field models. Section 4 investigates a family of  scaling solutions to the backreaction problem. Section 5 exploits the proposed correspondence. Here we explicitly reconstruct the potential for the scaling solutions, explore the solution space of inhomogeneous cosmologies with the help of the scaling solutions, and discuss some particular cases, which have been advanced before in the literature, and which are now interpreted with the help of a morphon field. We also give a concrete model and confront it with observational constraints. Section 6 contains a summary and an outlook on possible generalizations. ",
        "conclusions": "Our proposal of a mean field description of backreaction effects through a minimally coupled `morphon field', does not only provide a rephrasing of the kinematics of backreaction in terms of a scalar field cosmology, but it also justifies existence of the latter due to the fact that we identify averaged inhomogeneities in Einstein gravity as the underlying fundamental physics. Thus far, such a well--defined link was missing; alternatively, it is thought and there exist some plausible hints that the low--energy scale in some contemporary  particle physics models would provide this link (e.g. \\cite{nilles}, \\cite{polchinski}, \\cite{karthauser}), i.e. that the scalar field emerges from the dynamical degrees of freedom stemming from extra dimensions (`moduli fields').  The set of spatially averaged equations together with the mean field description of kinematical backreaction by a morphon field shows, in particular, that the framework of Friedmann--type equations is very robust. We understand now that an effective and scale--dependent Friedmannian framework is applicable with the surprising new input that, besides averaged matter sources, also a scalar field component emerges and has not to be invented.  \\subsection{Summary}  We have exploited the proposed correspondence by investigating a family of scaling solutions of the backreaction terms. The study of scaling solutions is well--advanced in research work on scalar field models, and our study allows to translate those results into the backreaction context. Here, our discussion was not exhaustive. We concentrated on exploring the solution space of inhomogeneous cosmologies with the help of scaling solutions with a `minimal' parameterization. Therefore, as a next step, this parameterization could be expanded and congruences with other work on scalar field models could be worked out. Note that an obvious such expansion would analyze superimposed scaling solutions that could be modeled either again by a single effective scalar field, or else by multiple scalar fields \\cite{copeland:multiple}, \\cite{collinucci}. One example of a superposition of scaling solutions is provided in Appendix A. Furthermore, observational constraints set on quintessence or phantom quintessence models have direct relevance to constraints that have to be set on morphon fields \\cite{maartens:phantom}. However, observations have to be reinterpreted within the inhomogeneous cosmology underlying the morphon dynamics.  In making the correspondence concrete we have not touched the question whether realistic dynamical models of the inhomogeneous Universe would comply with the class of averaged models that we singled out as `realistic case studies'. For the scaling solutions the scalar field correspondence revealed, that the constancy of the fraction of kinetic to potential energy of the morphon field relates to averaged models that are driven by strong coupling between averaged 3--Ricci curvature and kinematical fluctuations, the constancy implying their direct proportionality. While in the scalar field context this assumption is often employed, here we get an interesting picture of what this represents. The key for a physical interpretation of the proposed correspondence is the equivalence of the integrability condition (\\ref{integrability}) and the Klein--Gordon equation (\\ref{kleingordon}): $$ a_\\CD^{-6}\\partial_{t}\\left(a_\\CD^{6}{\\cal Q}_\\CD\\right)+ a_\\CD^{-2}\\partial_{t}\\left(a_\\CD^{2}\\average{\\CR}\\right)\\,=\\,0\\;\\;, \\;\\;\\Leftrightarrow\\;\\; {\\ddot\\Phi}_\\CD + 3 H_{\\cal D}{\\dot\\Phi}_\\CD + \\epsilon\\frac{\\partial}{\\partial \\Phi_\\CD}U(\\Phi_\\CD)\\;=\\;0\\;\\;. $$ Although the resulting picture may appear complex, it can be structured by looking at a single player: the averaged 3--Ricci curvature. If it is positive, it represents  an energy reservoir, stored in a negative potential in the correspondence. If this reservoir exceeds the `virial energy', then the excess energy is converted (in the scaling solutions directly) into an excess of kinetic energy, i.e. kinematical backreaction. The averaged curvature then decays faster than the `Friedmannian' curvature. On the other hand, if it is negative, then the same logic applies: the averaged curvature consequently becomes stronger negative compared with the `Friedmannian' curvature. Solutions of the Dark Energy problem, where a large positive value of backreaction is needed, will have to make strong use of this conversion, while `near--Friedmannian' models don't. A Newtonian or post--Newtonian approach suppresses these degrees of freedom by freezing the averaged curvature to the `Friedmannian' value (compare an attempt to work with the Friedmannian curvature parameter that, of course, needs an extra scalar field as a source for Dark Energy \\cite{franca}). From these remarks we understand the importance of a strongly evolving averaged 3--Ricci curvature, if the attempt to explain Dark Energy through backreaction should be successful. In this line there is clear support for the importance of a strongly evolving 3--Ricci curvature from the Lema\\^\\i tre--Tolman--Bondi solution, see \\cite{rasanen:LTB}, and in particular the recent work \\cite{curvatureLTB} and references to other papers on this solution therein. Recently, R\\\"as\\\"anen \\cite{rasanen:model} provided an illustrative example and a comprehensive discussion of the physics of backreaction--driven accelerated expansion.  Weakly perturbed Friedmann models have a special status if placed into the full solution space of the averaged inhomogeneous models: only if  the averaged fluctuations {\\it decouple} from the averaged curvature, i.e. if they evolve independently, then the averaged curvature evolves as a constant `Friedmannian' curvature, and fluctuations decay in proportion to the square of the inverse volume. Hints that point to the likely existence of a curvature--fluctuation--coupling come again from averaging the LTB solution \\cite{nambu} (see also \\cite{chuang}) and other global solutions \\cite{buchert:static} that all show an extra term in the averaged scalar curvature. In these examples the extra term  evolves in proportion to the inverse volume, hence deviates from a `Friedmannian' curvature and implies maintainance of large kinematical fluctuations that decay  only in proportion to the inverse volume.  While this particular behavior of the backreaction terms still requires large backreaction to start with \\cite{buchert:darkenergy}, e.g. a feature of globally stationary solutions \\cite{buchert:static}, a stronger evolution of averaged curvature, i.e. injecting more curvature energy into kinematical fluctuations, would allow to start with `near--Friedmannian' initial conditions and still explain current observations. An example for the latter possibility has been analyzed in this paper. This possibility can be interpreted such that there is some hope to find enough backreaction by starting with almost FRW initial data, as suggested by \\cite{kolbetal}, although here the conversion of curvature energy into backreaction must be very efficient; a more moderate evolution into backreaction--dominated phases would result, if we already start with `out--of--equilibrium' initial data.  In this line we have also pointed out, and this is worth stressing again, that the `Friedmannian' curvature parameter $k_{\\initial\\CD}$ is in general unrelated to the physical averaged scalar curvature, it is an integration constant and actually an integral of motion \\cite{buchert:static}: the dynamics of the averaged physical curvature is not represented by and, in most examples, deviates substantially from it. The restriction to a `Friedmannian' evolution of the averaged curvature singles out the special case where the curvature--fluctuation--coupling is {\\it strictly absent}.  Summarizing these thoughts we can say that there are two conceivable scenarios that remove the need for Dark Energy: first, a `soft scenario' that was discussed in detail in \\cite{buchert:static}. Here we already have an initial global state with strong expansion fluctuations, so that regionally a moderate evolution of backreaction and averaged curvature suffices to explain the observed acceleration. It was, however, pointed out that such models imply paradigmatic changes, and observational data have to be reinterpreted in order to put firm constraints on ${\\cal Q}_\\CD$ at the CMB epoch. Second, there is a `hard scenario' (implied by the suggestion of Kolb et al. \\cite{kolbetal}) that literally creates enough backreaction out of `nothing', {\\it cf.} Eq.~(\\ref{earlydarkenergy}). In our example of a particular scaling solution a phantom quintessence scenario arises that complies with the strong energy condition and also with  constraints in accord with those already put on the standard model. Consequently, this scenario needs strong evolution of backreaction and averaged curvature as seen best for the dimensionless cosmological parameters in Fig.~\\ref{fig:omegasSNLS}.  The present work has shown that the `hard' version can be made consistent with the framework of the averaged Einstein equations.  \\subsection{Outlook}  There are obviously a number of possible routes for generalizing the scalar field correspondence. Let us briefly discuss some of them.  Spatial averages are scale--dependent, since we integrate over a given compact volume of the space sections. In this paper we left the domain--dependence untouched, all of our considerations were focussing on a fixed scale. However, the scalar field correspondence holds on every scale, which is not only reminiscent of, but also physically a manifestation of renormalizable quantities. In this respect the morphon analogues of quintessence are different from standard quintessence models. In this context we could `Ricci flow' the averages to averages on a constant curvature geometry \\cite{klingon}, or we could study other renormalization group methods to control the scale--dependence, e.g. \\cite{calzetta}, \\cite{gaite:RG}, \\cite{carfora:focker}.  A possibly fruitful investigation would consider string--motivated gravitational theories and, employing the proposed correspondence, would aim at determining the scalar field theory from the higher--dimensional geometry of extrinsic and intrinsic curvature. For this end one would have to derive the averaged equations for the extended spacetime. Brane world cosmologies could be analyzed in the spirit of this correspondence. Not only in this context it will be interesting to understand, in which cases a non--minimally coupled morphon field would arise.  Furthermore, there are other interesting strategies of a more technical nature that, however, widen the physical context of applications. For example,  the introduction of a complex scalar field, which is necessary to access solutions with $\\Omega^\\CD_m >1$ with a positive averaged 3--Ricci curvature; the discussion of globally stationary cosmologies in \\cite{buchert:static} shows that a stationary cosmos obeys these conditions at early times, and later the averaged scalar curvature decreases fast until it becomes negative, while maintaining large kinematical backreaction.  Another example is the averaging on a different foliation of spacetime, i.e. introducing coordinate degrees of freedom. The latter is actually necessary if matter sources other than `irrotational dust' are considered; it results in an extra coupling of the scalar field to the matter sources stemming from a backreaction term due to the pressure gradient that is absent in the present paper \\cite{buchert:grgfluid}. This is the subject of a follow--up work that will also allow access to realistic models of backreaction--driven inflation \\cite{khlopov}. Here, the present investigation in principle allows to model inflationary scenarios too by translating typical characteristics of inflaton fields into corresponding morphon fields, however, within the restricted cases of `dust matter' or `vacuum'.  Since $\\Phi_\\CD$ models the inhomogeneous `vacuum part' of the sources,  there might be an interesting connection to the energy budget of gravitational waves that could be fruitfully exploited \\cite{dautcourt}, and, since $\\Phi_\\CD$ represents a `mean field' of fluctuations, also a connection to statistical thermodynamics is implicit. In order to establish this latter link, however, we have to note that ${\\cal Q}_\\CD$ models the averaged spatial variance of extrinsic curvature and so far not `fluctuations' in the thermodynamical sense. Another intimate, but on the level of a `dust matter model' formal, relation to effective thermodynamic models is suggested in the case of imperfect fluid models. They imply non--equilibrium effects that in turn can be associated with a `friction coefficient' proportional to the `cosmic equation of state' $p_{\\rm eff}/\\varrho_{\\rm eff}$. Thermodynamic arguments that were used in \\cite{zimdahl,schwarz} within the imperfect fluid picture lead to further possibilities of interpreting the different scaling regimes, e.g. by employing the second law of thermodynamics the authors of \\cite{zimdahl,schwarz} would conclude $p_{\\rm eff} < 0$, and thus ${\\cal Q}_\\CD > 1/3 \\langle \\CR\\rangle_{\\CD}$, i.e. $r > 1/3$ for positive averaged scalar curvature and $r < 1/3$ for negative one. Since the first case does not give rise to acceleration one would be led to conclude that the scalar curvature has to be negative on average, a conclusion shared by other work (in Appendix A we also reach this conclusion and provide references).  A final comment related to scalar field perturbations, often investigated in a post--Newtonian setting, is in order. In those perturbative approaches, the scalar degree of freedom that emerges is a combination of matter and metric inhomogeneities. Within the well--developed standard perturbation theory \\cite{mukhanov}, it is natural to first restrict the analysis to a tight range around a FRW background in order to calculate backreaction effects \\cite{futamase1,futamase2}, \\cite{abramo1,abramo2}. The amplitude of scalar perturbations, but also their  derivatives and, hence,  all curvature terms have to be small \\cite{wald}. Besides considerations of long--wavelength perturbations (e.g., \\cite{branden1}), the analysis of sub--horizon perturbations  with the aim to explain Dark Energy in a perturbative setting is the focus of many recent research papers and, here, we do not enter into the details of these works. We wish to point out two aspects that emerged from the present investigation, and which may be relevant for future strategies in a perturbation analysis.  First, we address the averaging issue: a post--Newtonian approximation may (and in most cases of cosmological relevance will) be adequate {\\it locally} or piecewise on a small range of scales. However,  as soon as we are looking at integral properties, i.e. integrating out a wider range of scales, its applicability should be verified.  In standard perturbation theory spatial averages are taken with respect to a `background observer'. Since a major player in the mechanism that can produce large kinematical fluctuations is the averaged 3--Ricci curvature, we have to be careful in relating Euclidean `background averages' to the Riemannian volume averages that govern the dynamics of the averaged cosmology.  Second, in the present approach we do not specify a metric of the space sections; the formulation and the correspondence holds for arbitrary 3--metrics. Also, when we spoke about a `scale--dependent Friedmannian model', we referred to the kinematics of the volume scale factor, we did not refer to the FRW metric. We can, however, specify a spatial metric to  establish a dynamical model.  Here, we think that the metric setting must allow for large deviations of the 3--Ricci curvature from the constant Friedmannian curvature. As a `rule of thumb' (another was recently given in \\cite{parry} for super--Hubble perturbations), we may say: {\\it if the averaged scalar curvature evolves at or near the constant curvature model, then there is no hope to model cases that lead to enough late--time acceleration.}  A future strategy related to perturbation theory could be motivated by Newtonian cosmological models for structure formation: in the Newtonian framework, an Eulerian perturbation theory does not provide access to the highly non--linear regime. Instead, the Lagrangian point of view offers a way to move with the largely perturbed fluid. A relativistic Lagrangian  perturbative approach is currently worked out aiming at generalizing the Newtonian work \\cite{bks} that has employed the exact averaged equations to construct a non--perturbative model for inhomogeneities out of perturbatively calculated fluctuations. A Lagrangian perturbation scheme itself does not include non--perturbative features, that may be needed in this context; non--perturbative effects have been recently discussed in the Newtonian framework \\cite{buchert:nonperturbative}.  \\vspace{6pt}  All these efforts aim at constructing a generic evolution model. The morphon field is an effective description without any perturbative assumption. But, one would like to establish the underlying inhomogeneous dynamics in order to understand, to what the realistic case studies that `explain away' the Dark Energy problem actually correspond.  \\newpage  \\subsection*{\\it Acknowledgements:}  {\\footnotesize TB and JL  acknowledge support by the Sonderforschungsbereich SFB 375 `Astroparticle physics' by the German science foundation DFG, JL during a visit to ASC Munich. TB would like to thank the Observatory of Meudon, Paris, and the University of Bielefeld, for hospitality and support. In particular, he would like to thank Dominik Schwarz for his invitation to temporarily hold a chair in Bielefeld university, where an excellent working environment and stimulating discussions with collegues in the physics department made this stay most enjoyable. TB and JL would like to thank Matthew Parry and Herbert Wagner for fruitful discussions during the preparation stage of this work; TB would like to also thank Misao Sasaki for encouraging conversations during a Dark Energy workshop held in Munich, and Mauro Carfora, Martin Kerscher, Slava Mukhanov, Syksy R\\\"as\\\"anen, Varun Sahni and Dominik Schwarz for valuable remarks and discussions.}  \\renewcommand{\\theequation}{A.\\arabic{equation}} \\setcounter{equation}{1}"
    },
    "0606/astro-ph0606067_arXiv.txt": {
        "abstract": "We use deep multi-band optical and near-infrared data for four general fields, GOODS-South, HDF North/South, and IRAC UDF in GOODS-North to investigate the evolution of the observed rest-frame $U-V$ color of field galaxies as a function of the stellar mass evaluated by fitting the galaxy spectral models to the observed broad-band SEDs. In these four fields, the $U-V$ color distributions of the galaxies at each mass and redshift interval are very similar. At $0.3 < z < 2.7$, we found that more massive galaxies always tend to have a redder $U-V$ color. High- and low-mass galaxies exhibit quite different color evolutions. As seen in our previous study in HDF-N, the color distribution of low-mass (M$_{stellar} \\lesssim 3 \\times 10^{9}$M$_{\\odot}$) galaxies becomes significantly bluer with an increase in the redshift. This evolution of the average color can be explained by a constant star formation rate model with $z_{\\rm form} \\sim 4$. On the other hand, the average color of high-mass galaxies (M$_{stellar} \\gtrsim 3 \\times 10^{10}$ M$_{\\odot}$) evolves more strongly at a high redshift. Such mass-dependent color distribution and its evolution indicate that galaxies with a larger stellar mass appear to have shorter star-formation timescales, and on an average, they form the larger fraction of their stars in the earlier epoch. ",
        "introduction": "Stellar mass is one of the most important quantities of galaxies that facilitate the understanding of their formation and evolution. The evolution of the number density of galaxies as a function of the stellar mass conveys the overall history of star formation and the mass assembly process in the universe. Many authors have investigated such evolution of the stellar mass distribution function back to $z\\gtrsim1$ (e.g., \\citealp{dic03}, \\citealp{bel03}, \\citealp{fon04}, \\citealp{dro04}, \\citealp{dro05}).    On the other hand, the study of properties such as the spectral energy distribution (SED), colors, and line fluxes of galaxies as a function of the stellar mass at various epochs helps us to understand star formation histories in additional detailed, such as, in what manner the star formation occurs in low- and high-mass galaxies and how it contributes to the growth of their stellar mass. In the case of local galaxies detected in the Sloan Digital Sky Survey, \\cite{hea04} and \\cite{jim05} argued that most of the stars in the present-day massive galaxies have been formed at a high redshift, while star formation histories of low-mass galaxies have relatively longer timescales. In more direct observations of high-redshift galaxies, \\cite{cow96} originally found that the mass of the galaxies undergoing rapid star formation has been declining smoothly with a decrease in the redshift. Recently, \\cite{jun05} also investigated the stellar mass dependence of star formation rates of galaxies at $0.8 < z < 2.0$ using the Gemini Deep Deep Survey data. They found the trend to be similar to that predicted in the studies of local galaxies \\citep{hea04}. Similar results are obtained in the studies on the evolution of specific star formation rates back to $z \\sim $1-1.5 (\\citealp{bri00}, \\citealp{bau05}, \\citealp{feu05a}). Very recently, \\cite{feu05b} extended the analysis to $z \\sim 5$, and found that a similar trend continues even at a higher redshift.  The samples in these previous studies, however, are too shallow to reach a sufficiently low mass range below $\\sim 5 \\times 10^{9}$ M$_{\\odot}$, a characteristic ``transition'' mass (Kajisawa and Yamada 2005, and Section 4 below), or are biased toward objects with typically low M/L ratios, such as more actively star-forming bluer populations due to the lack of very deep near-infrared (NIR) data. The evolution of low-mass galaxies is, however, important to understand the overall picture of galaxy evolution not only because they may be the building blocks of more massive galaxies but also because they are more sensitive probes to understand the thermal history of the universe.  In Kajisawa and Yamada (2005, hereafter referred to as KY05), we investigated the evolution of the rest-frame color and morphology distribution of the galaxies in HDF-N as a function of the stellar mass back to $z \\sim 3$ over a wide range of stellar masses down to $10^{9}$ M$_{\\odot}$. We found that the behavior of the rest-frame $U-V$ color distribution and its dependence on the stellar mass change at around $\\sim 5 \\times 10^{9}$ M$_{\\odot}$. In addition, we discovered that the sequence of the color distribution of low-mass galaxies systematically becomes bluer with the redshift, while the colors of high-mass galaxies do not show a significant redshift evolution. However, the area of the HDF-N data used in our analysis was only $\\sim$ 3.5 arcmin$^{2}$, and it was inferred that the effect of field-to-field variation was large.  \\begin{figure*} \\vspace{10mm} \\hspace{15mm} \\includegraphics[angle=90, scale=1.1]{Figure1r.eps} \\caption{Rest-frame $U-V$ color distribution as a function of stellar mass of $K$-selected galaxies. Each row depicts different redshift bin, and each column depicts different fields. Blue symbols indicate objects with spectroscopic redshift, while red ones indicate the phot-z sample. Error bars show the range of 90\\% confidence level (the redshift uncertainty is considered for the phot-z sample). Solid lines (and dashed lines for the most shallow ISAAC field in GOODS-S) show the detection limit that corresponds to the 90\\% completeness limit for point sources (see KY05 for the method of estimating the lines). The dotted-dash lines show the weighted average for the GOODS-S sample at $0.3<z<0.7$ (top-left panel). \\label{MsUVall}} \\end{figure*}  Therefore, in this paper, we have extended the study area of KY05 to include some other fields, GOODS-S \\citep{gia04}, a part of the IRAC Ultra Deep Field in GOODS-N \\citep{dic03b}, and FIRES (HDF-South, \\citealp{lab03}). Using the deep optical and NIR data for these fields, we investigate the rest-frame $U-V$ color evolution of galaxies at $0.3 < z < 2.7$ as a function of the stellar mass. These data allow us to extend our previous analysis to a larger sample in order to study the evolution of galaxies with a high statistical accuracy. We used the Vega magnitude system throughout the paper.  \\begin{figure*}[t] \\epsscale{1.1} \\hspace{-20mm} \\plotone{Figure2r.eps} \\caption{Rest-frame $U-V$ color distribution as a function of the stellar mass for the sum of all the four fields in each redshift bin. Data from different fields are plotted by different symbols. Blue and red symbols represent the spec-z and phot-z samples, respectively. Solid lines show the fitted ``tanh'' function in each redshift bin. The fitting parameters, UV$_{MIN}$ and M$_{tr}$ (see text) are also shown. \\label{MsUVfit}} \\end{figure*} ",
        "conclusions": "In Figure \\ref{MsUVall}, for each of the four fields, we plotted the rest-frame $U-V$ color of the K-selected galaxies in each redshift bin as a function of the stellar mass. First, we can see from this figure that the color distributions of the galaxies in the four fields are very similar. For reference, we estimated the weighted mean colors with a bin of $\\pm$0.15 dex of the GOODS-S sample at $0.3 < z < 0.7$ (in the top-left panel), and plotted them in all the panels by dotted-dash lines. The $U-V$ color is always very blue for M$_{stellar} \\lesssim 10^{9}$ M$_{\\odot}$ and becomes redder as the stellar mass increases toward $\\sim 10^{11}$ M$_{\\odot}$. Most of the galaxies with M$_{stellar} \\gtrsim 10^{11}$ M$_{\\odot}$ have the reddest color.  In fact, it is evident from Figure \\ref{MsUVall} that over the entire redshift range at $0.3 < z < 2.7$, the more massive galaxies always have the redder $U-V$ color. For example, while most of the galaxies with M$_{stellar} \\lesssim 3 \\times 10^{9}$ M$_{\\odot}$ have $U-V < 0.3$ at any redshift, many objects with M$_{stellar} \\gtrsim 3 \\times 10^{10}$ M$_{\\odot}$ have $U-V \\sim 1$. This trend is consistent with the results found in \\cite{feu05b} and other previous studies. If we observe the redshift evolution in Figure \\ref{MsUVall} more closely, we notice the gradual bluing of the $U-V$ color along the redshift within the mass range of $10^{9}$ M$_{\\odot} \\lesssim$ M$_{stellar} \\lesssim 10^{10}$ M$_{\\odot}$. While this systematic evolution of the rest-frame color distribution of low-mass galaxies has already been reported in KY05 for HDF-N, it is now confirmed by the larger sample that includes three other different fields.  In Figure \\ref{MsUVfit}, we try to quantify the observed mass-dependent color evolution by using an analytical fitting function. The data for all the four fields are combined. We excluded the lowest mass galaxies in each field from the sample so that the color bias near the detection limit (solid lines in Figure \\ref{MsUVall}) does not affect the fitting significantly; the limiting masses are set so that the $K$-band source detection completeness at $U-V \\sim 0.8$ is 50\\%. These limiting stellar masses are tabulated in Table \\ref{cutMs}. Considering the shape of the mass-weighted average color profile in GOODS-S (the top-left panel in Figure 1), we chose a hyperbolic tangent as the fitting function. This was in keeping with \\cite{bal04} who used a similar form to quantify the bimodality in the color-magnitude distribution of the SDSS sample. The fitting function is written as follows: \\begin{table} \\caption{Lower limit of stellar mass in each field (M$_{\\odot}$) \\label{cutMs}} \\begin{tabular}{ccccc} \\tableline\\tableline Redshift & GOODS-South & IRAC UDF & HDF-N & FIRES\\\\ \\tableline 0.3-0.7 &1.4-3.3$\\times 10^{8}$ &1.8$\\times 10^{8}$ &1.1$\\times 10^{8}$ &6.6$\\times 10^{7}$\\\\ 0.7-1.1 &4.2-10$\\times 10^{8}$ &5.5$\\times 10^{8}$ &3.5$\\times 10^{8}$ &2.0$\\times 10^{8}$\\\\ 1.1-1.9 &1.2-2.9$\\times 10^{9}$ &1.6$\\times 10^{9}$ &1.0$\\times 10^{9}$ &5.8$\\times 10^{8}$\\\\ 1.9-2.7 &3.1-7.6$\\times 10^{9}$ &4.1$\\times 10^{9}$ &2.6$\\times 10^{9}$ &1.5$\\times 10^{9}$ \\\\ \\tableline \\end{tabular} \\end{table} \\\\ \\vspace{1mm} $U-V$$({\\rm M}_{stellar})$=$ \\left(\\frac{UV_{MIN}+UV_{MAX}}{2}\\right)+\\\\\\left(\\frac{UV_{MAX}-UV_{MIN}}{2 }\\right)$$\\times\\tanh{(R({\\rm M}_{stellar}-{\\rm M}_{tr}))}$ \\vspace{1mm} \\\\ where $UV_{MIN}$ and $UV_{MAX}$ are the asymptotic minimum and maximum values of the rest-frame $U-V$ color at the low- and high-mass ends. M$_{tr}$ and $R$ represent the mass and mass range at which the color transition occurs. For simplicity, we fixed $R$ and $UV_{MAX}$, which seem to change little with the redshift in our sample, to the values fitted at $0.3 < z < 0.7$ ($R = 1.0$, $UV_{MAX} = 1.3$). Next, we fitted the color distribution in each redshift bin by changing the values of M$_{tr}$ and $UV_{MIN}$. The result is shown by the solid lines in Figure \\ref{MsUVfit}. The best-fitted values of the parameters are M$_{tr} = 1.57^{+0.39}_{-0.23} \\times 10^{10}$ M$_{\\odot}$ and $UV_{MIN} = -0.036 \\pm 0.055$ at $0.3 < z < 0.7$, M$_{tr} = 1.71^{+0.28}_{-0.24} \\times 10^{10}$ M$_{\\odot}$ and $UV_{MIN} = -0.277 \\pm 0.056$ at $0.7 < z < 1.1$, M$_{tr} = 3.56^{+1.09}_{-0.84} \\times 10^{10}$ M$_{\\odot}$ and $UV_{MIN} = -0.317 \\pm 0.128$ at $1.1 < z < 1.9$, M$_{tr} = 5.77^{+1.79}_{-1.37} \\times 10^{10}$ M$_{\\odot}$ and $UV_{MIN} = -0.344 \\pm 0.084$ at $1.9 < z < 2.7$. These results again show that the $U-V$ color at the low-mass end gradually evolves blueward with the redshift. In addition, the transition mass also seems to increase along the redshift, although the uncertainty is relatively large. It is noteworthy that \\cite{bun06} also recently discovered similar transition mass increments with the redshift at $0.4 < z < 1.4$ in the study using the DEEP2 data.  \\begin{figure*} \\epsscale{0.9} \\hspace{-30mm} \\plotone{Figure3r.eps} \\caption{rest-frame $U-V$ color vs redshift for galaxies within each stellar mass range. The plotted sample and the meaning of the symbols are similar to those in Figure \\ref{MsUVfit}. Dotted-dash lines and solid lines show the weighted average and the best-fit GALAXEV models with exponentially decaying SFR. Dashed lines show the upper and lower envelope of the color distribution. \\label{UVz}} \\end{figure*}  Figure \\ref{UVz} shows the evolution of the rest-frame $U-V$ color at each stellar mass range for the combined data. We plotted the weighted average colors binned over the redshift range of $\\pm$0.15 (dotted-dash lines). In the top panel, it is more clearly evident that the color distribution of the galaxies with M$_{stellar} < 3 \\times 10^{9}$ M$_{\\odot}$ becomes gradually bluer with an increase in the redshift. In KY05, we discussed that the star formation activity in these low-mass galaxies occurs rather continuously since they exhibit a very blue color at {\\it any} redshift. A bluer color even at a higher redshift also suggests younger average ages of the galaxies at $z =$2-3. In the lower panels, we see that the colors of galaxies with a larger stellar mass also become bluer at a high redshift. There seems to be a trend where the more massive populations show a stronger color evolution at a high redshift, which in fact corresponds to the evolution of the transition mass found in Figure \\ref{MsUVfit}.  We try to fit such color evolutions using GALAXEV population synthesis models with simple star formation histories. We simply consider an exponentially decaying SFR and assume solar metallicity and no dust extinction. We estimate the $U-V$ color for the model grid with the star formation timescale $\\tau=$ 0.01, 0.05, 0.1, 0.5, 1, 1.5, 2, 3, 5, 15, and $\\infty$ (constant SFR) Gyr and the formation redshift $z_{f} =$ 3, 4, 5, 6, 7, and 8. Eventually, we found that $\\tau$ mainly affects the overall slope of the color evolution, and $z_{f}$ strongly affects the color at a high redshift (e.g., $z > 2$). The resultant best-fit models and their $\\tau$ and $z_{f}$ values are shown in Figure \\ref{UVz}. It is found from the results that the more massive populations appear to have a shorter star formation timescale. In other words, they form the larger fraction of their stellar contents at a high redshift. Such a trend is very similar to the concept of mass-dependent star formation histories argued by \\cite{hea04}. Further, the trend is also consistent with the studies on mass-dependent evolutions of specific star formation rates at a high redshift (\\citealp{feu05b}, \\citealp{jun05}). However, in this study, the trend is depicted by more directly observable quantities.  In these analyses, it is important to check the effect of the color bias near the detection limit. The bias against a red high-M/L population is expected to be severe, especially for low-mass galaxies at a high redshift since the $K$-band filter samples relatively short rest-frame wavelengths. We finally discovered, however, that our deep $K$-selected samples do not suffer considerably from the bias. It is true that our 50\\% completeness limit at $U-V = 0.8$ could neglect some fractions of relatively red low-mass galaxies in the highest redshift bin. However, the color distribution tends to concentrate at considerably bluer values ($U-V \\lesssim 0$), and there is few galaxies that are redder than $U-V = 0.3$, where the completeness is greater than 90\\%. This can be demonstrated by the lowest mass objects near the detection limit ($K \\sim 23.5$) in the highest redshift bin of the deepest FIRES data. For example, out of 20 galaxies with the stellar mass of $1.5 \\times 10^{9}$ M$_{\\odot} <$ M$_{stellar} < 3.0 \\times 10^{9}$ M$_{\\odot}$ in the $1.9 < z < 2.7$ bin of FIRES, 19 galaxies have $U-V < 0$, and one galaxy has $0 < U-V < 0.1$; further, no galaxy has $0.1 < U-V < 0.3$.  Indeed, a sample selection with very deep NIR data is essential not only for evaluating the stellar mass of galaxies but also for generating relatively unbiased samples to study the color distribution of these low-mass galaxies at a high redshift. For example, $I$-band selections such as those used for the FORS Deep Field data in \\cite{feu05b} could introduce a severe color bias. The detection limit of $I(AB) = 26.4$ in \\cite{feu05b} can only sample galaxies with $U-V < -0.1$ for M$_{stellar} = 2 \\times 10^{9}$  M$_{\\odot}$ at $z > 2$. Such a selection does not reveal whether red galaxies exist.  In Figure \\ref{UVz}, we found that the stronger color evolution of massive galaxies at a high redshift was explained by the exponentially decaying SFR models with $z_{f} \\sim $4-5, for which the star formation timescale decreases with mass. On the other hand, the evolution of the average color of low-mass ($\\lesssim 3 \\times 10^{9}$ M$_{\\odot}$) galaxies can actually be explained by the constant SFR model. However, the best-fit formation redshift occurs at $z_{f} \\sim 4$, which is similar to those for high-mass galaxies. If this is true, the differences in the star formation histories among galaxies with different stellar masses are not mainly in the epochs when the star formation began, but rather, in the star formation efficiencies and/or mechanisms by which star formation activities are suppressed. The long star-formation timescale for low-mass galaxies can be caused by the self-regulation of star formation, for example, by a supernova feedback \\citep{dek05}. The mass dependence of the star formation efficiencies and suppression mechanisms can be related with processes such as the efficient exhaustion of gas reservoirs (e.g., \\citealp{men05}) or shock heating of cold gas during a major merger \\citep{cox04}.  It is noteworthy that some of the assumptions in our analysis may be oversimplifications. If we consider dust extinction, the estimated formation redshift shifts to smaller values. For example, if we assume the extinction of E($B - V$) $= 0.13$ (median-fitted value of our sample) for low-mass galaxies, the best-fit formation redshift of these galaxies shifts from $z_{f} = 4$ to $z_{f} = 3$. We also verified which variations of star formation histories are permissible if the scatter of the colors around the average values is considered. To illustrate this, we plotted the GALAXEV models that trace the upper and lower envelopes of the color distribution by dashed lines in each panel in Figure \\ref{UVz}.  Finally, we note that several other surveys have investigated the evolution of the stellar mass or luminosity function of field galaxies and found that the normalization decreases at $z \\gtrsim 1$ (e.g., \\citealp{fon04}, \\citealp{yam05}, \\citealp{cap06}, \\citealp{sar06}). Several studies have also suggested mass-dependent number density evolution (e.g., \\citealp{gla04}, \\citealp{dro04}, \\citealp{dro05}). The next step is to consistently explain both the color (SED) evolution and the number density evolution as functions of the stellar mass."
    },
    "0606/astro-ph0606584_arXiv.txt": {
        "abstract": "We demonstrate that if k-essence can solve the coincidence problem and play the role of dark energy in the universe, the fluctuations of the field have to propagate superluminally at some stage. We argue that this implies that successful k-essence models violate causality. It is not possible to define a time ordered succession of events in a Lorentz invariant way. Therefore, k-essence cannot arise as low energy effective field theory of a causal, consistent high energy theory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606251_arXiv.txt": {
        "abstract": "$\\eta$~Carinae (\\etc) is a stellar binary system with a period of 5.54 years.  It harbors one of the brightest and most massive stars in our galaxy.  This paper presents spectroscopic evidence for a fast (up to 2,000 \\kms) X-ray outflow of ionized gas launched from \\etc\\ just before what is believed to be the binary periastron (point of smallest binary separation). The appearance of this high-velocity component, just as the irregular flares in the \\x\\ light curve, can not be explained by the simple continuous binary wind interaction, adding to the intrigue of the \\etc\\ system. ",
        "introduction": "\\label{sec:intro}  Despite being studied for many years and in all wavebands, the $\\eta$ Carinae system (\\etc), best known for its spectacular nebula \\citep{morse98} that can be tracked back to the Great Eruption of 1840 \\citep{davidson97} is still much of a mystery. It is now believed to be a stellar binary system with a period of 5.54 years \\citep{damineli96}. The primary star is one of the brightest and most massive ($\\sim$ 120~\\msol) stars in our galaxy. The system is currently in a phase of immense mass ejection. The \\x s detected from \\etc\\ can be explained by the collision of the two stellar winds \\citep{corcoran01, pittard02}, not so the 70 day X-ray minimum observed periodically around periastron \\citep{corcoran05}.  Currently, the mass loss rate of the primary exceeds $10^{-4}$\\msolyr\\ \\citep{smith03} and that of the secondary is an order of magnitude lower \\citep{pittard02}, but still unusually high. The wind of the secondary is much faster \\citep[few 1,000 \\kms,][]{pittard02} than that of the primary \\citep[$\\sim$~600 \\kms\\ and depending on latitude,][]{smith03}. The \\x\\ flux between 2 and 10~keV during the past 12 years is normally at a level of 5$\\times 10^{-11}$ erg~s$^{-1}$~cm$^{-2}$, as expected from the collision of the two stellar winds \\citep{usov92, corcoran01, pittard02}. The \\x\\ light curve is roughly constant throughout most of the 5.54 year orbital period, but rises gradually by a factor of 3--4 upon approach to periastron at which time short flares appear \\citep{corcoran97}, before it drops sharply and stays in an \\x\\ low state ($<$~20\\% of normal brightness) for approximately 70~days. The full X-ray light curve can be found in \\citet{corcoran05}. Although absorption by the dense primary wind may play a certain role in attenuating the X-ray emission from \\etc, both spectral and temporal arguments have been put forward to reject absorption or an eclipse as the prime reason for this persisting low state \\citep{hamaguchi04, soker05, akashi06}, which is unique to \\etc. Recently, it has been proposed that if the binary separation during periastron is small enough, the massive primary wind could smother the secondary wind and thus shut down the X-ray emission \\citep{soker05, akashi06}. However, so far there has been no direct evidence for accretion.  In this work, we wish to add to the temporal aspects of the \\etc\\ tale a high-resolution \\x\\ spectroscopic dimension. This is achieved by using five deep exposures of \\etc\\ with the High Energy Grating Spectrometer (\\hetgs) on board the \\chandra\\ \\x\\ Observatory (CXO), only the first of which has been published \\citep{corcoran01,corcoran01b}. Theoretical variations of spectral line profiles in binary colliding wind systems have been modeled by \\citet{henley03}. The varying line profiles presented here for \\etc, however, do not follow these models, at least not in a straightforward way, as we explain in \\S 3. ",
        "conclusions": "The colliding wind binary model, which is generally accepted for explaining the \\x\\ emission of \\etc, has several clear predictions regarding the spectral line profiles expected at different orbital phases.  The Doppler shifts are the result of shocked gas flowing along the contact discontinuity (CD) surface and away from the stagnation point (SP). If the shock opening angle is between $45^\\circ\\ - 90^\\circ$ \\citep{pittard02}, the broadest line profiles are expected when the system is observed perpendicular to the line connecting the two stars, as the hottest gas formed around the SP flows directly towards and away from the observer. In this case, the lines are centered at zero velocity and are intrinsically symmetrical, but could be skewed by absorption of the far (red) emission.  On the other hand, when the system is observed from behind one of the stars, shifted centroids are expected. If the system is observed from behind the weaker (stronger) wind, downstream shocked gas is flowing along the CD surface in the general direction of (opposite) the observer, producing blue- (red-) shifted centroids. The maximum velocities in this case would be less than those in the first (orthogonal) orientation. For a complete and detailed calculation of line profiles from colliding winds see Henley et al. (2003). In the following, however, we argue that the current line profiles observed at the different phases of the \\etc\\ orbit are in contradiction, even qualitatively, with the simple colliding wind scenario.  If the major axis of the \\etc\\ binary is perpendicular to our line of sight \\citep{smith04}, the broadest profiles are expected at apastron (see above), but that is when the observed profiles are narrowest ($\\phi =-0.470$).  Also, with this geometry, the secondary would have to pass in front of the primary before periastron (in contrary to the picture of \\citet{smith04}) to explain the blue-shifted centroids at $\\phi = -0.028, -0.006$, but that would produce symmetrical (blue-shifted) profiles as absorption through the secondary wind is weak, while strongly asymmetrical profiles are observed.  In any case, the lower velocities observed near apastron appear to rule out this orientation if the colliding wind geometry is to produce the observed \\x\\ line profiles.  If alternatively, the projection of the line of sight on the orbital plane is more or less along the major axis and the viewing angle at apastron is from behind the secondary \\citep{corcoran01}, the observed blue-shifts at $\\phi = -0.028, -0.006$ could be ascribed to gas flowing from the SP towards the observer and the unobserved red wing of the line to absorption by the primary wind. However, according to this scenario, at apastron one would expect blue-shifted centroids of at least a few 100~\\kms\\ due to downstream gas (see above), which is not observed ($\\phi = -0.470$). For an observer above the binary plane ($i < 90^\\circ$), as seems to be the case in \\etc\\ \\citep{corcoran01}, the problem becomes more severe (higher blue-shift is expected, but not observed).  In short, the observed behavior of the \\x\\ line profiles of \\etc\\ appear qualitatively inconsistent with a naive wind-collision scenario, since the interpretation of the high blue-shifts from a post-shock stream can not be made consistent with all phases. \\citet{henleyphd}, comparing with line profiles calculated from hydrodynamical simulations, also reached the conclusion that the observed profiles are not as expected from simple models of the geometry of the wind-wind collision. He suggested that including the Coriolis effect may bring the models to better agreement with observations. \\citet{kashi07} suggest the effect of orbital motion, which can account for small centroid shifts, but can not explain the currently observed \\x\\ profiles. Here, we wish to propose an alternative explanation. In particular, we note that the asymmetric line profiles are observed during strong peaks in the X-ray light curve. These peaks too can not be accounted for by the simple colliding wind model \\citep{pittard02, henley03}, or by the Coriolis force.  We therefore prefer to interpret the observed high-velocity gas up to --2000~\\kms\\ in terms of a transient, collimated fast wind, or a jet, ejected from the immediate vicinity of the binary system through the interaction of the two stars. The peaks in the light curve hint that the outflow may be in the form of blobs. The unresolved \\chandra\\ images and the rapid variations both indicate that the outflow is restricted to within approximately 2$\\times$10$^{16}$~cm from the center (assuming 0.5\\arcsec\\ at \\etc 's distance of 2.3~kpc). The appearance of the same charge states in the spectra attained during all phases implies that the temperatures of the X-ray gas and line-emitting outflow remain in the range of $kT$~= 2~-- 5~keV. Consequently, we speculate that the fast component of the outflow consists of gas shocked by the collision of the winds (as observed throughout the orbit) that is launched and collimated near periastron. The widths of the major peaks seen in Figs.~2 and 3 suggest that the outflow is only moderately collimated to within $\\sim 30^\\circ$. The condensation of shocked gas provides a natural explanation for the rise in X-ray intensity during the short intermittent flares observed in the X-ray light curve. A collimated outflow from the secondary has also been suggested to explain the enhanced He~II $\\lambda 4686$ emission before periastron \\citep{soker06}. Interestingly, both proper motion and Doppler shifts have been measured for high-velocity visible-light knots much further out from the center \\citep{walborn78, meaburn93}. The optical knots were ejected more than a hundred years ago along the minor axis of the nebula \\citep{meaburn93}. We would expect the present outflow to be launched along the major axis of the nebula although we have no pertinent information on its present direction with respect to the system's geometry. If indeed that is the case and if the inclination angle of the binary plane to our line of sight is $i= 42^\\circ$ \\citep{smith02}, the actual outflow velocity would be as high as $\\sim$2,700~\\kms.   As the massive \\etc\\ primary is a short-lived star shedding considerable mass and as it is now understood that core collapse supernovae can not expel more than a few solar masses, implying that higher mass stars must shed most of their mass prior to the explosion, the winds and fast outflow of \\etc\\ may be a supernova in the making. At the very least, it highlights the next periastron passage on 2009, January 12 as a faithfully scheduled experiment for the astrophysically common, but poorly understood phenomenon of jet launching."
    },
    "0606/astro-ph0606298_arXiv.txt": {
        "abstract": "{ We present broad-band imaging with the Subaru Telescope of a $25' \\times 25'$ field surrounding the radio galaxy TN~J1338--1942 at redshift $z = 4.1$. The field contains excesses of Lyman-$\\alpha$ emitters (LAEs) and Lyman break galaxies (LBGs) identified with a protocluster surrounding the radio galaxy. Our new wide-field images provide information about the boundary of the protocluster and its surroundings. There are 874~candidate LBGs within our field, having redshifts in the range $z = 3.5 - 4.5$. An examination of the brightest of these (with $i' < 25.0$) shows that the most prominent concentration coincides with the previously discovered protocluster. The diameter of this galaxy overdensity corresponds to $\\sim 2$~\\mpc{} at $z = 4$, consistent with the previous estimation using LAEs. Several other concentrations of LBGs are observed in the field, some of which may well be physically connected with the $z = 4.1$ protocluster. The observed structure in the smoothed LBG distribution can be explained as the projection of large-scale structure, within the redshift range $z = 3.5 - 4.5$, comprising compact overdensities and prominent larger voids. If the $5 - 8$ observed compact overdensities are associated with protoclusters, the observed protocluster volume density is $\\sim 5 \\times 10^{-6}$~\\mpc$^{-3}$, similar to the volume density of rich clusters in the local Universe. ",
        "introduction": "\\label{sec:int}  There is considerable evidence for galaxy overdensities at high redshifts ($z > 2$; e.g. Steidel et al. \\cite{bib:steidel1998}; M\\o{}ller \\& Fynbo \\cite{bib:moller2001}; Shimasaku et al. \\cite{bib:shimasaku2003}; Palunas et al. \\cite{bib:palunas2004}; Ouchi et al. \\cite{bib:ouchi2005}). In many cases these overdensities have been presumed to be associated with the ancestors of rich local clusters. At $z > 2$ the Universe is $\\lesssim 3$~Gyr old, too short for these structures to have virialized (e.g. Venemans \\cite{bib:venemans2005}). Hence these structures are often called \\emph{protoclusters}. Most searches for protoclusters have been limited by relatively small fields (typically smaller than $10' \\times 10'$).  An efficient way of finding protoclusters is to use high redshift radio galaxies (HzRGs; e.g. R\\\"ottgering et al. \\cite{bib:rottgering1994}) as tracers (Venemans et al. \\cite{bib:venemans2002}; Kurk et al. \\cite{bib:kurk2004}). HzRGs are large massive objects with many of the properties expected of forming dominant cluster (cD) galaxies (West \\cite{bib:west1994}). Although most protoclusters are not radio-loud, radio-selected protoclusters may be typical. Because radio-sources are relatively short-lived ($\\sim 10^7 $ years; Blundell \\& Rawlings \\cite{bib:blundell1999}), the statistics are consistent with the progenitor of every rich local cluster having harboured a luminous radio galaxy at some stage in its existence.  Using the VLT, Venemans et al. (\\cite{bib:venemans2002}) spectroscopically confirmed 20 Lyman-$\\alpha$ emitters (LAEs) in a $7' \\times 7'$ field around HzRG TN~J1338--1942 at a redshift of $z = 4.1$ and identified these LAEs with a $z = 4.1$ protocluster. Further evidence that this LAE overdensity was indeed associated with a protocluster was provided by observations with the HST/ACS that revealed an excess and non-uniform distribution of candidate Lyman break galaxies (LBGs) around the radio galaxy (Miley et al. \\cite{bib:miley2004}; Overzier et al. \\cite{bib:overzier2005}).  Although the LAE search around TN~J1338--1942 was extended with a second $7' \\times 7'$ field (Venemans \\cite{bib:venemans2005}), this was insufficient to determine the boundary of the protocluster. Here we present the results of a multi-color study of candidate LBGs from a $25' \\times 25'$ region surrounding TN~J1338--1942. The large field-of-view (FOV) facilitates searches for LBGs out to the boundary of the protocluster structure and beyond. The data also provide new information about large-scale structure and voids at $z \\sim 4$.  Throughout this paper, we use AB-magnitudes, $1 \\sigma$ errors and adopt a flat, $\\Lambda$-dominated cosmology with $\\Omega_\\mathrm{M} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $H_0 = 100~h$~km~s$^{-1}$~\\mpc$^{-1}$ with $h = 0.7$. ",
        "conclusions": "\\label{sec:con}   We draw several conclusions from the present observations: \\\\ (i) TN~J1338--1942 is located in an overdensity of LBGs, both in projection and in redshift space. The new wide-field results are consistent with previous observations which revealed the presence of a protocluster through the overdensities of LAEs (Venemans et al. \\cite{bib:venemans2002}; Venemans \\cite{bib:venemans2005}) and LBGs (Miley et al. \\cite{bib:miley2004}; Overzier et al. \\cite{bib:overzier2005}). This further supports the hypothesis that HzRGs are located in dense environments. The apparent size of the overdensity and the relative position of TN~J1338--1942 within the overdensity are similar to that found by Venemans et al. for LAEs. \\\\ (ii) There are $4 - 7$ additional overdensities in the projected LBG distibution, similar to the one harbouring TN~J1338--1942. These may well be due to protoclusters at $z = 3.5 - 4.5$ and one or more of these overdensities could well be physically related to the TN~J1338--1942 protocluster. \\\\ (iii) The spatial distribution of our complete LBG sample is consistent with a Universe at $z \\sim 4$ that comprises a web of compact galaxy overdensities (protoclusters) embedded in larger regions of galaxy underdensities (voids). The statistics of the overdensities are consistent with the local volume density of rich clusters.  Spectroscopic measurements are needed to investigate whether there is a physical connection between some of the outlying observed galaxy overdensities and the overdensity corresponding to the TN~J1338--1942 protocluster. Such observations could enable the cosmic web at $z = 4.1$ to be traced over distances of tens of megaparsec. Furthermore, similar measurements on other $z > 2$ protoclusters would be useful for constraining the development of large-scale structure in the early Universe."
    },
    "0606/astro-ph0606121_arXiv.txt": {
        "abstract": "We discuss our high precision radial velocity results of a sample of $90$ M~dwarfs observed with the Hobby-Eberly Telescope and the Harlan J. Smith 2.7~m Telescope at McDonald Observatory, as well as the ESO VLT and the Keck I telescopes, within the context of the overall frequency of Jupiter-mass planetary companions to main sequence stars. None of the stars in our sample show variability indicative of a giant planet in a short period orbit, with $a\\leq1$~AU. We estimate an upper limit of the frequency $f$ of close-in Jovian planets around M dwarfs as $< 1.27\\%$ (at the $1 \\sigma$ confidence level). Furthermore, we determine the efficiency of our survey to have noticed planets in circular orbits as $98\\%$ for companions with $m \\sin i > 3.8~{\\rm M}_{\\rm Jup}$ and $a\\leq0.7$~AU. For eccentric orbits ($e=0.6$) the survey completeness is $95\\%$ for all planets with $m \\sin i > 3.5~{\\rm M}_{\\rm Jup}$ and $a\\leq0.7$~AU. Our results point toward a generally lower frequency of close-in Jovian planets for M dwarfs as compared to FGK-type stars. This is an important piece of information for our understanding of the process of planet formation as a function of stellar mass. ",
        "introduction": "Despite the stunning success of the radial velocity (RV) technique in finding extrasolar planets (e.g. Mayor \\& Queloz 1995), which resulted in the discovery of more than $160$ planetary companions over the past decade, our understanding of planet formation is far from complete. We are especially lacking a general overview of the frequency of planetary companions to stars throughout the entire HR-diagram. Doppler surveys have traditionally targeted bright solar type main sequence stars and it is no big surprise that most planets were found around G-type stars. But, is this entirely a result of an observational bias, or is it a true effect which could lead to a better understanding of the underlying physics of planet formation?  The majority of the stars in the solar neighborhood are M~dwarf stars with masses of $0.5~{\\rm M}_{\\odot}$ or less (Henry~1998). In order to determine the overall galactic population of planets it is important not to ``overlook'' the faint and low mass regime of the HR-diagram and to control observational biases in order to arrive at statistically meaningful results. Because of the intrinsic faintness of M dwarfs it is generally more difficult and time consuming to achieve the high quality RV data for the detection of planetary companions. This led to the situation that M dwarfs constitute only small subsets in the target samples of most precision Doppler surveys.  So far, we know of only one M dwarf, GJ~876 (M4~V, M=$0.3~{\\rm M}_{\\odot}$), to harbor a planetary system with Jupiter-mass companions (Delfosse et al.~1998; Marcy et al.~1998; Marcy et al.~2001; Benedict et al.~2002). There is the possibility that the giant planet detected by microlensing (Bond et al.~2004) also orbits an M~dwarf, but the exact spectral type of the primary lens has not been determined yet. Butler et al.~(2004) announced the discovery of a short periodic Neptune mass planet around the M dwarf GJ~436 and Rivera et al.~(2005) presented evidence for an extremely low mass ($M\\approx7.5~{\\rm M}_{\\rm Earth}$) third planetary companion to GJ~876. Bonfils et al.~(2005b) reported the detection of a Neptune mass companion to the southern M3~V star GJ~581.  Endl et al.~(2003) described the dedicated M dwarf survey at the Hobby-Eberly Telescope (HET) which targets exclusively M dwarfs and presented the data of the first year of the survey. This paper now contains three years of RV results from the on-going HET program with the addition of five years of the M~dwarf RV results from our planet search program at the ESO Very large Telescope (VLT), as well as data from the McDonald 2.7~m telescope program and from the Keck Hyades survey. We discuss these results and their implications on the total frequency of detectable giant planets along the main sequence. ",
        "conclusions": "Marcy et al. (2005) find a frequency of $1.2\\pm0.2\\%$ of ``hot Jupiters'' with $a<0.1$~AU around FGK-type stars and according to their Fig.2 a frequency of $2.5\\pm0.4\\%$ of planets with $a<1$~AU. While the frequency of ''hot Jupiters'' is still consistent with our results, it appears that there is a difference emerging between the detection rate of Jovian planets with $a<1$~AU around FGK-type stars and M dwarfs. Also Butler et al.~(2004) noted that the combined results from all radial velocity planet search programs, which include M dwarfs in their target samples, point toward an upper limit for $f$ of $<0.5\\%$. Lineweaver \\& Grether (2003) presented an analysis and extrapolation of the planet frequency for FGK-type stars currently monitored by Doppler surveys. They find that the fraction of detected planets increases from $\\approx3.9\\pm1\\%$ for a survey duration of 2 years to $5.5\\pm1.5\\%$ for a 4 year survey. Again, both values are higher than the upper limit we find. However, our upper limit of $1.27\\%$ for M dwarf planets is only valid in the planet mass-separation range, to which our program is most sensitive (with a $>98-99\\%$ completeness, see Fig.~\\ref{sens1} and Fig.~\\ref{sens2}). But, also FGK-star surveys are not ideal surveys and they miss a small fraction of planets.  Gaudi et al.~(2002) also finds a low frequency for Jupiter-mass companions around M dwarfs in the galactic bulge based on the results of the PLANET microlensing survey. However, one has to bear in mind that the PLANET survey samples a different region of our galaxy and a comparison with the solar neighborhood could be inadequate.  A general observational bias as explanation for a lower M dwarf Jovian planet frequency becomes increasingly unlikely but is not completely ruled out. Doppler surveys of FGK-type stars usually have a higher RV precision than M dwarf programs. The bias introduced by a somewhat lower RV precision for M dwarfs is partly compensated by the fact that the RV amplitudes induced by giant planets around M dwarfs are larger because of the lower mass of the host star and thus easier to detect. Using large aperture telescopes allows us to overcome their intrinsic faintness and to obtain high quality RV data also for these late spectral types. Continuation, improvement and expansion of current M dwarf Doppler surveys is necessary to allow a better determination of the planet frequency around these stars and to perform a more detailed comparison with the results from the FGK-star surveys.  Another issue which could still lead to an unintentional observational bias in our M dwarf sample is stellar metallicity. In recent years, it has became more and more obvious that the frequency of detectable planets is a function of the metal content of the stars included in Doppler surveys (Fischer \\& Valenti~2005 ; Santos, Israelian, \\& Mayor~2004). Metal poor stars appear to harbor fewer planets detectable by the RV technique (close in massive planets) than metal rich stars. This could mean that the formation mechanism is somehow linked to the metal content in the protoplanetary disk (at least for the stars included in these surveys). Is it possible that we mostly targeted metal poor M dwarfs and that this is the reason for the observed low planet frequency? Our HET and VLT samples are biased toward inactive and thus presumably older M dwarfs. Thus, it is conceivable that the majority of these M dwarfs are metal poor. A survey including also more active M dwarfs as well as the determination of the metallicities of the targets can solve this issue. Unfortunately, there are no large scale spectroscopic surveys to determine precise metallicities of M dwarfs, although there are efforts under way to find suitable techniques (e.g. Valenti et al.~1998). Woolf \\& Wallerstein~(2005) presented [Fe/H] measurements for 35 M and K dwarfs. Three M dwarfs from our sample, GJ~411, GJ~412~A, and GJ~687 are included in their study. They find that GJ~411 and GJ~412~A are metal poor (${\\rm [Fe/H]}\\approx-0.4$), while GJ~687 is metal rich (${\\rm [Fe/H]}\\approx0.15$). Iron abundances for three more stars (GJ~15~A, GJ~109, GJ~849) of our sample are estimated by Bonfils et al.~(2005a). These authors used visual binaries with M dwarf secondaries to calibrate a photometric method to derive M dwarf metallicities. GJ~15~A (${\\rm [Fe/H]}\\approx-0.45$) and GJ~109 (${\\rm [Fe/H]}\\approx-0.2$) turn out to be metal poor stars, while GJ~849 (${\\rm [Fe/H]}\\approx0.14$) has a higher than solar iron abundance. So far, four out of six target stars with [Fe/H] measurements thus turn out to be metal poor stars. A detailed determination of the metallicity of our sample stars using the method of Bean et al.~(2005) is currently in progress. The Hyades M dwarfs are of course young and active stars. They should have a similar metal content as the cluster mean of [Fe/H] =$0.13\\pm0.01$ (Paulson, Sneden, Cochran~2003). But the small number of M dwarfs included in the overall Hyades sample prevents a meaningful comparison with the planet frequency of earlier type stars in that metallicity range.  How do our observations agree with current models for planet formation? Laughlin, Bodenheimer \\& Adams~(2004) explored the formation of gas giants around M dwarfs within the framework of the core accretion model (e.g. Pollack et al.~1996). These authors conclude that this model has severe problems in forming Jupiter-class planets in the less massive protoplanetary disks of M dwarfs. Also, Ida \\& Lin~(2005) show that close-in Jovian planets should be relatively rare around M dwarfs. Our data confirm so far their theoretical prediction that M dwarfs should harbor fewer gas giant planets. On the other hand, Boss~(2006) shows that the gravitational instability model has a higher efficiency in forming giant planets in less massive M dwarf disks. GJ~876 remains the only unambiguous case of an M dwarf orbited by Jupiter-mass planetary companions despite the accumulation of high quality RV data for more M dwarfs over the past years. So what makes GJ~876 so special? At the moment we can only speculate that GJ~876 had either a more massive disk than other M dwarfs, albeit for unknown reasons, or these planets indeed formed by gravitational instability. As suggested by Boss~(2006), M dwarfs could be a used as a test ground for competing giant planet formation models.  In the case that orbital migration (e.g. Lin et al.~1996) results in the observed small semi-major axes for many of the planets detected by Doppler surveys, we can pose the question of whether this mechanism might be less efficient for M dwarfs. This would explain the current null detections (again with the exception of GJ~876) because high quality RV data for most M dwarfs do not have sufficient time coverage to find (or exclude) Jupiter-type companions in long periodic orbits. In our sample we do observe RV trends indicating more distant companions, but most of them are quite large and most likely caused by stellar companions. Highly precise astrometric studies using data from the upcoming Space Interferometry Mission (SIM) should allow the detection of planets at large orbital separations around nearby M dwarfs. Firm upper limits from astrometry have been placed by Benedict et al.~(1999) on the masses of planetary companions in the period range from $60$ to $600$ days, orbiting Barnard's star and Proxima Centauri, using HST Fine Guidance Sensor data. Benedict et al.~(2002) succeeded in detecting the astrometric perturbation caused by the outermost planet in the GJ~876 system, and derived a mass for the companion of $1.89\\pm0.34~{\\rm M}_{\\rm Jup}$.  Despite their lack of close-in Jovian planets, M dwarfs remain attractive targets for current Doppler surveys. Because of their low primary masses the RV amplitude induced by an orbiting companion is higher than for F,G, or K type stars. The discovery of planets with extremely low masses (a few ${\\rm M}_{\\rm Earth}$) using precise radial velocity measurements is feasible in the M dwarf regime (e.g. K\\\"urster et al. 2003). Model calculations (e.g. Wetherill~1996 ; Laughlin et al.~2004) also show no difficulties in forming low mass planets via planetesimal accretion in disks around M dwarfs. Ida \\& Lin (2005) even predict a higher frequency of icy giant planets with masses comparable to Neptune in short periodic orbits for M dwarfs than for G type stars. The lowest mass extrasolar planet our group has found so far is the fourth companion in the $\\rho^{1}$~Cnc (G8~V) system (McArthur et al.~2004). The RV semi-amplitude induced by this planet is only $6.7~{\\rm m\\,s}^{-1}$ and hence a difficult signal to detect. If the same planet (${\\rm P}=2.81~{\\rm days}, m \\sin i = 14.2~{\\rm M}_{\\rm Earth}$) were to orbit a high mass M dwarf ($0.5~{\\rm M}_{\\odot}$) the RV semi-amplitude would increase to $10.2~{\\rm m\\,s}^{-1}$ and for Proxima Centauri (M5~V, $0.12~{\\rm M}_{\\odot}$) the signal would be $26.3~{\\rm m\\,s}^{-1}$. The discoveries of the short-period Neptunes by Butler et al.~(2004) and Bonfils et al.~(2005b) and of an additional planet with an extremely low mass of $m \\approx 7.5~{\\rm M}_{\\rm Earth}$ in the GJ~876 system (Rivera et al.~2005) further support the notion that M dwarf stars remain a fruitful and interesting hunting ground for high precision radial velocity surveys. In the future high resolution spectrometers working in the near infrared will be the ideal tools for a more thorough exploration of the red (and low mass) part of the main sequence.  The NASA Kepler mission (Borucki et al.~2003), planned for launch in late 2008, should give much better insight into the true frequency of short-period planets ($P < 1$ year) around M~dwarfs, provided a sufficiently large sample of M dwarfs is included in the Kepler target list.  M~dwarfs are particularly good targets for large photometric transit surveys such as Kepler because their significantly smaller radius than solar-type stars results in a much larger photometric signal for a given planet size."
    },
    "0606/astro-ph0606317_arXiv.txt": {
        "abstract": "% {}% {Our aim is to measure the time delay between the two gravitationally lensed images of the \\zqso\\ = 1.547 quasar \\obj, in order to estimate the Hubble constant \\ho.} {Our  measurement  is  based on $R$-band light curves with 57 epochs obtained at Maidanak Observatory,  in  Uzbekistan, from May 2004 to September  2005.   The  photometry is    performed using   simultaneous deconvolution  of the data,  which provides the individual light curves of the  otherwise blended quasar images. The  time delay is determined from the light curves using two very different numerical techniques, i.e., polynomial fitting and direct cross-correlation. The time delay is converted into \\ho\\ following analytical modeling of the potential well.} {Our best  estimate  of the time  delay  is $\\Delta t  = 49.5 \\pm 1.9$~days, i.e., we reach a 3.8\\,\\% accuracy. The $R$-band flux ratio between the quasar images, corrected for the time delay and for slow microlensing, is $F_{\\rm A}/F_{\\rm  B} = 6.2 \\pm 5$\\,\\%.} {The  accuracy reached on the  time delay allows us to discriminate well between  families of  lens models.  As  for most other multiply imaged quasars, only models  of the   lensing galaxy  that  have a de Vaucouleurs mass profile plus external shear give a Hubble constant compatible with the current most popular value (\\ho\\  = 72 $\\pm$ 8 \\kmsmpc). A more realistic singular isothermal sphere model plus external shear gives \\ho\\ = 51.7$^{+4.0}_{-3.0}$ \\kmsmpc.} ",
        "introduction": "\\begin{figure*}[ht] \\begin{center} \\includegraphics[width=17.8cm]{5823fig1.ps} \\caption{A $R$-band image of \\obj\\ obtained at Maidanak Observatory. This image  is a combination of 376  frames, totalising  31 hours of exposure.  The mean seeing  is 1.0\\arcsec, and the  field of view is $3.5\\arcmin \\times   8.9\\arcmin$.  The two  stars  labeled PSF1 and PSF2 are used  to model the Point Spread  Function required for  the MCS  deconvolution method.  The 4  reference stars used for the photometric calibration are star \\#3, \\#4, PSF1 and PSF2.  Star \\#5 is used as a cross-check of the deconvolution photometry.} \\label{field} \\end{center} \\end{figure*}  The so-called  time-delay method in  gravitationally  lensed quasars (Refsdal \\cite{Refsdal64})  is one  of the  rare techniques   that can yield a measurement of the Hubble constant  \\ho\\ at truly cosmological distances, independently of any local calibration or standard  candle. Until now, however, no  concerted and long-term project has  succeeded in applying it in a systematic way, at a level really competitive with other techniques, such as Cepheids, supernovae  or the fluctuations of the cosmic   microwave background radiation.  The   COSMOGRAIL project aims  at measuring time  delays  for most known gravitationally lensed quasars,  as well  as  improving  their lens  models, using  deep high resolution imaging and spectroscopy.  The contribution of the error on the time delay measurement to the total error budget  on \\ho\\ is about 50\\,\\%,  the  other half  being   due  to degeneracies   in  the  lens models. The  immediate goal of COSMOGRAIL  is to make  the former close to negligible in front  of the latter.  While  this is very hard for lenses with short time delays (Kochanek et al.~\\cite{Kochanek06a}, Morgan et al.~\\cite{Morgan06}),  longer  time delays can  be  measured accurately (Eigenbrod et al.~\\cite{Eigenbrod05}).  In this paper, we present the first result from our monitoring program: the  time delay   measurement in  the   gravitationally lensed  quasar \\obj\\ (16$^\\textrm{\\scriptsize{h}}$ 50$^\\textrm{\\scriptsize{m}}$ 43\\fs5, +42\\degr\\ 51\\arcmin\\ 45\\farcs 0; J2000.0), discovered by Morgan et  al.  (\\cite{Morgan03})  as a doubly imaged \\zqso\\  = 1.547 quasar,   with an angular  separation of 1.2\\arcsec.  At the time of the discovery,  the $B$-band magnitudes of components  A and B were $17.8$  and $20.0$ respectively.  The lensing galaxy  is  detected  by  Morgan  et   al.   (\\cite{Morgan03})  in the $I$ band. Absorption  lines seen in the spectrum  of the quasar images suggest a lens redshift of \\zl\\ = 0.577. In a cosmology with \\ho\\ = 75 km s$^{-1}$ Mpc$^{-1}$,  Morgan et al.  (\\cite{Morgan03}) predict that the time delay between the  quasar images is  of the order of a month, assuming a Singular Isothermal  Sphere (SIS) potential for the lensing galaxy,  plus an external shear.   Following the non-parametric models of Saha  et al.  (\\cite{Saha06})  and the same cosmology, the expected time delay  between  the two quasar   images is $\\Delta  t \\sim 30-60$ days.   The  first two  seasons of the  photometric  monitoring of  \\obj\\ were carried  out with the  1.5-m   telescope at Maidanak  Observatory,  in Uzbekistan.  Given  the  high declination of  \\obj\\   and the pointing limits of the telescope, \\obj\\ can be followed for about 8 months per year under  good airmass and  seeing conditions.  This combination of time delay and visibility window makes \\obj\\ an excellent target for an  accurate time delay measurement.  From numerical simulations  using artificial light  curves, Eigenbrod et al. (\\cite{Eigenbrod05}) predict that the  accuracy on the time  delay can be  as good as  1\\,\\%, assuming that  the peak to  peak amplitude of the quasar light variation  is  0.2~mag over  the two  years  of observation and using  a  temporal  sampling of  at   least  one observing  point  per week. Although  more than two  observing seasons will  be necessary to achieve this final  goal, it is  already possible to measure the  time delay with 3.8\\,\\% accuracy, which is sufficient to show that lens models with constant mass-to-light ratio and models with dark matter profiles that do not trace light give discrepant values of \\ho.   The  photometric monitoring, the data reduction   and the light curves are presented in Sects. 2 and 3 respectively. The determination of the time delay between the two lensed images is described in Sect. 4.   In Sect.  5, we discuss the mass models for the lensing galaxy and the implications for   the  Hubble constant.   Finally, Sect.   6  summarises  the main results. ",
        "conclusions": "The  main result   of  the  present  work   is the  first  time  delay measurement of the COSMOGRAIL project, for the doubly imaged quasar \\obj. Our best estimate of the time delay is $\\Delta t  = 49.5 \\pm 1.9$~days (1$\\sigma$). This corresponds to a relative accuracy of 3.8\\,\\%. The $R$-band flux ratio of the two quasar images, corrected for microlensing  and for the time delay, is $F_{\\rm A}/F_{\\rm B} = 6.2 \\pm 5\\%$.  The amplitude of the external shear in  the circular lens models is larger than $\\gamma=0.15$, and suggests that both external shear and ellipticity in the  main  lensing galaxy are  necessary to  model the total potential well. However,  the  present observational constraints  available  for \\obj\\, prevent   us from introducing  a lens ellipticity and position angle  in the models.   Deep, high resolution  images of \\obj\\ will be necessary to estimate the latter two parameters.  With  the present   3.8\\,\\%   accuracy  on   the time  delay,  we  can efficiently   discriminate   between  families  of   lens models   or, conversely, estimate \\ho\\ assuming a  lens model. Our results suggest that SIS models are not acceptable in order to match the current favored value of the Hubble constant (\\ho\\  = 72 $\\pm$ 8 \\kmsmpc) and that only models of the lensing galaxy that have a de Vaucouleurs mass profile can reproduce it.   Slow microlensing  is  negligible   in   the first   two  seasons   of monitoring, with a  global variation  of  less than 0.02 mag  over 500 days.  Finally, note that the redshift of the lensing galaxy is based on MgII and  FeII  absorption  lines  in     the   spectrum of  the     quasar images. Although the impact parameter necessary to form these lines is small,   we cannot  exclude   that the true   lens  redshift does not correspond to the absorption lines, as in other systems (HE~1104-1805; Lidman et  al.~\\cite{lidman2000}). A  genuine spectrum of  the lensing galaxy must be obtained to confirm its redshift."
    },
    "0606/astro-ph0606385_arXiv.txt": {
        "abstract": "{Shells in Elliptical Galaxies are faint, sharp-edged features, believed to provide evidence of a recent ($\\sim 0.5 - 2 \\times 10^9$ years ago) merger event. We analyse the Globular Cluster (GC) systems of six shell elliptical galaxies, to examine the effects of mergers upon the GC formation history. } {We examine the colour distributions, and investigate differences between red and blue globular cluster populations. We present luminosity functions, spatial distributions and specific frequencies ($S_N$) at 50 kpc radius for our sample. } {We present V and I magnitudes for cluster candidates measured with the HST Advanced Camera for Surveys (ACS). Galaxy background light is modelled and removed, and magnitudes are measured in 8 pixel (0.4 arcsec) diameter apertures. Background contamination is removed using counts from Hubble Deep Field South. } { We find that the colour distributions for NGC 3923 and NGC 5982 have a bimodal form typical of bright ellipticals, with peaks near $V-I=0.92 \\pm 0.04$ and $V-I=1.18 \\pm 0.06$. In NGC 7626, we find in addition a population of abnormally luminous clusters at $M_I=-12.5$. In NGC 2865 we find an unusually blue population, which may also be young. In NGC1344 and NGC474 the red cluster population is marginally detected. The radial surface density profiles are more flattened than the galaxy light in the cores. As already known, in NGC3923, which has a high $S_N$ of 5.6, the radial density distribution is more shallower than the diffuse galaxy light. } {The clusters in NGC 2865 and NGC 7626 provide evidence for formation of a population associated with a recent merger. In the other galaxies, the properties of the clusters are similar to those observed in other, non-shell, elliptical galaxies. } ",
        "introduction": " ",
        "conclusions": "The properties of the globular cluster systems of six shell galaxies were analysed. For NGC 2865 and NGC 7626 we observe anomalous features in the (V-I) histograms, in addition to the bimodal structure which is normal for ellipticals. The features represent excesses at intermediate colour in NGC 7626, and at very blue colours in NGC 2865, consistent with a small population of GCs, of age 2-5 Gyr and 0.5-1 Gyr respectively, possibly formed in the merger event which created the shells. In each case the young population is dominated by the much larger, old, bimodal distribution.  The data for two galaxies (NGC 1344 and NGC 3923) allow the determination of their globular cluster luminosity function distances. Fitting Gaussians to their GCLF give distances moduli of 31.72 and 31.50 for NGC 1344 and NGC 3923 respectively.  The properties of NGC 3923 and NGC 5982 are very similar. Their bimodal V-I distributions and radial density profiles of blue and red clusters are typical for old GC systems in ellipticals. NGC 474 and NGC 1344  show one single blue peak and very shallow red peaks in their V-I histograms. These properties are unusual for bright ellipticals (\\cite{kundu}, \\cite{peng06}) and may indicate less early mergers in their formation history.  NGC 3923 and NGC 7626 have higher specific frequencies (respectively 5.6 and 3.9 within 50kpc) than normal for galaxies located in a low density environment. The luminosity of the X-ray halo detected in the former galaxy is probably not sufficient to explain its high $S_N$ as proposed by \\cite{mcl}. The $S_N$ of the other galaxies have values of around 2 and are typical for galaxies located in a low density environment.  Although for some of these six galaxies new GCs may have been formed recently, the general properties (like V-I distributions and flattening of GC density profile) of the globular cluster systems of these shell galaxies do not deviate systematically from 'normal' elliptical galaxies.  In a future paper (Sikkema et al., in preparation) we will investigate the morphology and stellar populations of the diffuse galaxy, in particular in the shells."
    },
    "0606/astro-ph0606450_arXiv.txt": {
        "abstract": "The conditions for the leading r-process site candidate, neutrino-driven winds, can not be reproduced self-consistently in current supernova models. For that reason, we investigate an alternate model involving the mass ejected by fallback in a supernova explosion, through hydrodynamic and nucleosynthesis calculations. The nucleosynthetic products of this ejected material produces r-process elements, including those in the vicinity of the elusive 3rd peak at mass number 195. Trans-iron element production beyond the second peak is made possible by a rapid ($<1ms$) freezeout of $\\alpha$ particles which leaves behind a large nucleon (including protons!) to r-process seed ratio. This rapid phase is followed by a relatively long ($\\gtrsim 15\\mathrm{ms}$) simmering phase at $\\sim$ 2$\\times$10$^9$ K, which is the thermodynamic consequence of the hydrodynamic trajectory of the turbulent flows in the fallback outburst. During the slow phase high mass elements beyond the second peak are first made through rapid capture of both protons and neutrons. The flow stays close to valley of stability during this phase. After freeze-out of protons the remaining neutrons cause a shift out to short-lived isotopes as is typical for the r-process.  A low electron fraction isn't required in this model, however, the detailed final distribution is sensitive to the electron fraction. Our simulations suggest that supernova fallback is a viable alternative scenario for the r-process. ",
        "introduction": "Production mechanisms for elements heavier than iron, by fast (r-process) and slow (s-process) captures of neutrons, have been known for a long time (Cameron 1957, Burbidge et al. 1957), yet finding conditions in Nature and understanding the physics that allow these mechanisms to robustly operate has been more elusive.  The search for an r-process production site has proven especially difficult.  The supernova wind model (see Qian \\& Woosley 1996 for a detailed description), invoking the neutrino-driven wind produced by the cooling proto-neutron star formed in a core-collapse supernovae, is the best-studied r-process mechanism.  However, the wind scenario generally seem to require uncommon conditions (e.g. $> 2$\\,M$_\\odot$ neutron star; Argast et al. 2004; Suzuki \\& Nagataki 2005) to achieve a reasonable r-process signature.  Difficulties with the wind mechanism have led to investigations about other possible sites for the r-process; the best-studied being the coalescence of two neutron stars (Freiburghaus 1999).  Unfortunately, the low event rate of merging neutron stars appears to rule out binary coalescence as a primary production site (Argast et al. 2004).  Most other mechanisms have been based on the suggestions by Qian \\& Woosley (1996) from changes in the neutrino/nuclear physics (e.g.  neutrino oscillations - Qian et al. 1993) to magnetic fields (e.g. Suzuki \\& Nagataki 2005).  In this paper we explore supernova fallback, yet another of Qian \\& Woosley (1996).  After the launch of a supernova explosion, some of the material initially ejected in the blast can decelerate and fall back onto the proto-neutron star.  This ``fallback'' material was initially proposed by Colgate (1971) who argued that a rarefaction wave would catch the shock and the lack of pressure support would cause the shock to fall back onto the neutron star.  Essentially, the material with ejecta velocities below the escape velocity will fall back onto the neutron star.  Shigeyama, Nomoto, \\& Hashimoto (1988) and Woosley (1989) proposed a scenario where fallback occurs when the shock decelerates as it moves through the star.  This deceleration sends a reverse shock through the star, decelerating the innermost material and causing it to accrete.  This mechanism is strongest when the shock hits the hydrogen envelop, attaining its most dramatic deceleration. Simulations show that fallback occurs early (Fryer et al. 1999; MacFadyen et al. 2001), establishing that sub-escape velocity ejecta dominates fallback material (Fryer \\& Kalogera 2001).  Indeed, Fryer \\& Heger (2000) found fallback in a disk in the first second after the launch of the explosion.  Although simulations of supernova explosions suggest that fallback occurs in all simulations, this is not yet the prevailing view among core-collapse theorists.  Fallback is strongest when the explosion is weak and it may be that only for weak supernovae, M $\\ge$ 20\\,M$_\\odot$, that the fallback mechanism can work (Fryer 1999). This does not preclude fallback as an r-process source!  In fact, Argast et al. (2004) found that narrow ranges of massive stars can explain the entire r-process abundance pattern.  In this paper, we present the first calculations of the r-process based on hydrodynamic simulations of fallback, the ejection of fallback material (\\S 2), and detailed nucleosynthesis calculations of the ejecta (\\S 3).  Without tuning our initial conditions, we find that fallback leads to ejecta that carry abundance signatures characteristic for r-process, thereby demonstrating that ejecta from the fallback of material onto neutron stars is a viable r-process site. ",
        "conclusions": ""
    },
    "0606/astro-ph0606499_arXiv.txt": {
        "abstract": "We report the independent discovery and follow-up observations of supernova 2005gj by the Nearby Supernova Factory. This is the second confirmed case of a ``hybrid'' Type Ia/IIn supernova, which like the prototype \\ic, we interpret as the explosion of a white dwarf interacting with a circumstellar medium. Our early-phase photometry of \\gj\\ shows that the strength of the interaction between the supernova ejecta and circumstellar material is much stronger than for \\ic. Our first spectrum shows a hot continuum with broad and narrow \\Halpha\\ emission. Later spectra, spanning over 4 months from outburst, show clear Type~Ia features combined with broad and narrow \\Hgamma, \\Hbeta, \\Halpha\\ and \\ion{He}{1}~\\wl\\wl5876,7065 in emission.  At higher resolution, P~Cygni profiles are apparent. Surprisingly, we also observe an inverted P~Cygni profile for [\\ion{O}{3}]~\\wl5007. We find that the lightcurve and measured velocity of the unshocked circumstellar material imply mass loss as recently as 8 years ago. This is in contrast to \\ic, for which an inner cavity in the circumstellar material was inferred.  Within the context of the thin-shell approximation, the early lightcurve is well-described by a flat radial density profile for the circumstellar material. However, our decomposition of the spectra into Type~Ia and shock emission components allows for little obscuration of the supernova, suggesting an aspherical or clumpy distribution for the circumstellar material. We suggest that the emission line velocity profiles arise from electron scattering rather than the kinematics of the shock. This is supported by the inferred high densities, and the lack of evidence for evolution in the line widths. Ground- and space-based photometry, and Keck spectroscopy, of the host galaxy are used to ascertain that the host galaxy has low metallicity ($Z/Z_{\\odot}<0.3$; 95\\% confidence) and that this galaxy is undergoing a significant star formation event that began roughly $200\\pm70$~Myr ago. We discuss the implications of these observations for progenitor models and cosmology using Type~Ia supernovae. ",
        "introduction": "Type~Ia supernovae (\\sneia) are very important in astrophysics. As standardized candles they played the leading role in the determination that the mass-density of the Universe is sub-critical and that the expansion of the Universe is accelerating \\citep{perl98,garnavich98,riess98,perl99}.  They continue to play a major role in on-going efforts to determine the cause of this acceleration \\citep{knop03,tonry03,riess04,astier06} and in refining the determination of the Hubble constant \\citep{riess05}. \\sneia\\ are a major source of chemical enrichment in galaxies, and may play a major role in the development of low-mass galaxies \\citep[e.g.][]{ferrara00}.  For these many reasons, a long-standing goal has been to precisely identify the progenitor systems of \\sneia. Of special interest for cosmology --- where high-redshift \\sneia\\ are compared to low-redshift \\sneia\\ to deduce cosmological parameters --- is the question of whether or not \\sneia\\ may arise from multiple channels \\citep[e.g.,][]{howell01,scannapieco05}, acting on significantly different timescales, and in particular whether the relative contributions of multiple channels could act in such a way as to lead to biased cosmology results \\citep{drell00,leib01}.  Historically speaking, the two most strongly favored \\snia\\ progenitor scenarios both involve the thermonuclear disruption of white dwarf stars \\citep[for an overview, see][]{bra95}.  In the \\emph{single degenerate} model, a carbon-oxygen white dwarf accretes matter shed by a less-evolved companion star until it reaches the Chandrasekhar mass and explodes.  The \\emph{double degenerate} scenario consists of the coalescence of a binary white dwarf system.  The role of mass loss in the former scenario suggests an observational test to discriminate between the two models.  The unambiguous detection of circumstellar hydrogen in the spectrum of a \\snia\\ would be interpreted as support for the single degenerate scenario.  The search for this signature has resulted in mostly nondetections and upper limits on mass loss from the progenitor system, until recently.  The spectra of the peculiar \\snia~2002ic \\citep{woo02,ham03} were the first to clearly reveal the interaction of the ejecta of one of these events with its circumstellar medium (CSM).  The spectra of \\ic\\ resembled those of the spectroscopically peculiar, overluminous \\snia~1991T, with features characteristic of a \\snia\\ though ``diluted'' \\citep{ham03}.  Most remarkably, these spectra exhibit narrow \\Halpha\\ emission consistent with the redshift of the SN \\citep[$z = 0.0667$,][]{kotak04} strongly reminiscent of \\sneiin\\ \\citep[``n'' for narrow hydrogen emission,][]{sch90}.  More circumstantial evidence obtained by comparing late time spectra of \\sneiin\\ 1997cy and 1999E with \\ic\\ \\citep{ham03,deng04} suggests that some fraction of \\sneiin, ostensibly core-collapse SNe, may in fact be \\sneia\\ veiled by a substantial CSM.  Various interpretations of the observations of \\ic\\ have appeared in the literature.  Motivated by the large mass inferred for the CSM, \\citet{ham03} suggested that the progenitor system contained a massive asymptotic giant branch (AGB) star which shed the material prior to explosion.  The AGB star could be a binary companion star, or the SN progenitor itself \\citep[a Type 1.5 SN,][]{ibe83}.  Most analyses agree that the density of the CSM was high, that a few tenths to several solar masses of material were shed, and that the mass loss rate was high or the CSM was clumpy \\citep{kotak04,chugai04c,wang04}.  Spectropolarimetry suggested that the geometry of the CSM was aspherical, and perhaps equatorially condensed \\citep{wang04,uenishi04}.  A light curve analysis \\citep{woo04} implied a delay in the onset of the circumstellar interaction, suggesting that the mass loss phase had ended or trailed off in the years preceding the SN or that the inner CSM had been cleared by a fast wind, both consistent with the formation of a proto-planetary nebula (PPN; \\citet{kwok93}) by an AGB star.  Instead of an AGB scenario, \\citet{livio03} suggested that \\ic\\ was in fact the product of a double degenerate merger during the common envelope (CE) phase of a binary; the white dwarf spirals in to coalesce with its companion's core.  However, the disk-like geometry and large radial extent of the CSM inferred from spectropolarimetry \\citep{wang04} and coalescence timescales that are incompatible with the timescale inferred for the mass ejection \\citep{chu04b} seem to rule out the CE coalescence model for \\ic.  Motivated in part by the inferred presence of considerable CSM surrounding \\ic, \\citet{wang05a} advanced a hypothesis to explain the low values $R_B \\equiv A_B /$\\ebmv\\ often derived from observations of \\sneia. Instead of the value $R_B\\sim4.1$ expected for dust extinction in the Galaxy and LMC, values $R_B\\sim2$--3 are often found \\citep{tripp99,phillips99,wang03,knop03,wang05b}.  In this hypothesis, dust in the immediate vicinity of the SN or at the highest velocities of its ejecta scatter SN light into the observer line of sight, thus adding back light and effectively reducing derived values of $R_\\lambda$ in the optical.  The source of the dust may be a PN or PPN formed during the evolution of the progenitor system.  This explanation suggests a continuum of models for ejecta-CSM interaction in \\sneia.  \\sneia\\ with more normal values of $R_B$ could represent cases where the PN has had enough time to disperse into the host interstellar medium.  Cases where $R_B$ is low but there is no CSM interaction signature could occur in diffuse PNe.  Still others (such as \\ic) with a strong CSM interaction signature might originate in more compact PNe or PPNe.  Here we report on our observations of the recent \\gj, which has much in common with \\ic.  \\citet{pri05} observed the development of spectral features characteristic of a \\snia, during their follow-up of the discovery of \\gj\\ \\citep{cbet247,fri05}, in what had originally appeared to be a \\sniin.  The Nearby Supernova Factory \\citep[SNfactory]{ald02} had independently discovered \\gj\\ in images taken 2005 September~29, and acquired optical spectroscopy on 2005 October~3 which showed the strong, comparatively narrow Balmer \\Halpha\\ line that is the defining characteristic of the IIn class, and no visible \\snia\\ features.  Both \\gj\\ and \\ic\\ were discovered by ``blind'' SN surveys (conducted without targeting specific galaxies) --- \\gj\\ by the SDSS and the SNfactory, and \\ic\\ by the SNfactory prototype search \\citep{woo02}.  The two SNe are hosted by low-luminosity galaxies at about the same redshift, and their spectra at about 70 days after explosion are remarkably similar.  There are, however, important differences between the two events.  Our earliest spectrum of \\gj\\ indicates a more intense, earlier interaction with the CSM than in the case of \\ic.  Additionally, the light curve of \\gj\\ is brighter than that of \\ic, supporting the conclusion that the \\gj\\ ejecta-CSM interaction is stronger.  Still, X-ray \\citep{imm05} and radio \\citep{sod05} observations resulted in non-detections, as was the case for radio observations of \\ic\\ \\citep{berger03,stockdale03}.  The organization of this paper is as follows.  In \\S2 we present new photometry and spectroscopy of \\gj\\ covering numerous epochs extending more than four months past the explosion date. We also describe spectroscopic observations of the host galaxy. In \\S3 we analyze the \\gj\\ observations in the context of a two-component SN plus CSM interaction model in order to place constraints on the amount of CSM and the relative contributions of SN and CSM light to the observations.  \\S4 presents an analysis of the host galaxy spectroscopy and photometry. These are brought together in \\S5 in a discussion of our results and their possible implications for SN progenitor systems and cosmology.  \\S6 summarizes our findings. We take $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ throughout this paper. ",
        "conclusions": "First classified as Type IIn and then displaying clear Type Ia features, \\gj\\ is the second supernova spectroscopically confirmed to possess Type Ia features combined with a strong CSM interaction. Of these two confirmed events and two suspected such events, \\gj\\ presented the largest observed CSM interaction, and its early behavior raises the question of how many of the early detected \\sneiin\\ might be similar events.  The lightcurve and spectroscopy presented here offer a wealth of information.  Much more detailed modeling will help in understanding the nature of progenitor system, and could be a key to understanding important aspects of the formation of Type~Ia supernovae. Better statistics from larger unbiased nearby supernova searches will help greatly in understanding the possibly multiple channels that produce SNe~Ia."
    },
    "0606/astro-ph0606720_arXiv.txt": {
        "abstract": "{} { This study {{seeks to provide empirical constraints on the}} different physical components that can {{act to yield}} temporal variability in predominantly or partially wind-formed optical lines of luminous OB stars, and thus potentially affect the reliable determination of fundamental parameters, {{ including mass-loss rates via clumped winds.}} } {Using time-series spectroscopy from epochs spread over $\\sim$ 4 years, we present a case study of the O9.5 supergiant $\\alpha$ Cam. We demonstrate that the \\he (2$^3$P$^0$--3$^3$D) line is an important diagnostic for photospheric and wind variability in this star. The actions of large radial velocity shifts (up to $\\sim$ 30 km s$^{-1}$) in the photospheric absorption lines can also affect the morphology of the \\Ha line profile, which is commonly used for measuring mass-loss rates in massive stars.} {We identify a 0.36-day period in subtle absorption profile changes in \\he, which likely betrays {{photospheric structure, perhaps due to low-order non-radial pulsations.}} This signal persist over $\\sim$ 2 months, but it is not present 2 years later (November 2004); it is also not seen in the stellar wind components of the line profiles. Using a pure \\Ha line-synthesis code we interpret maximum changes in the red-ward and peak emission of $\\alpha$ Cam in terms of mass-loss rate differences in the range $\\sim$ 5.1 $\\times$ 10$^{-6}$ to 6.5 $\\times$ 10$^{-6}$ M$_\\odot$ yr$^{-1}$. However, the models generally fail to reproduce the morphology of blueward (possibly absorptive) regions of the profiles.} {The optical line profiles of $\\alpha$ Cam are affected by (i) {{deep-seated fluctuations close to, or at, the photosphere}}, (ii) atmospheric velocity gradients, and (iii) large-scale stellar wind structure. This study provides new empirical perspectives on accurate line-synthesis modelling of stellar wind signatures in massive luminous stars.} ",
        "introduction": "Through their powerful stellar winds, high ionizing fluxes, and ultimate demise as supernovae, massive stars have a profound effect on the dynamical and chemical evolution of our own and other galaxies. The hot star winds are driven by line scattering of radiation, and the extent of the mass-loss is a major factor in the evolution of these stars; it can for instance influence whether the end-state is a neutron star or a black hole \\citep[e.g.][]{Heger03}.  Therefore it is important that the nature of these outflows is well understood, and in particular the mass-loss rates are reliably determined.  One of the most serious obstacles in this path arises from detailed optical and ultraviolet observations that reveal widespread line profile variability, implying clumped structure in the winds \\citep[see][plus references within]{Prinja02, Markova04, Morel04}. The developing picture of massive star winds as a clumped medium is further supported by recent model atmosphere calculations; e.g. \\citet{Bouret05} need to include substantial clumping factors in their models to match key wind features in the spectra of O-type stars, implying significant downward revisions in the canonical mass-loss rates. Clumping may be a consequence of the inherent instability of the radiation driving mechanism of the wind, or of the creation of coherent structures arising from photospheric disturbances, or both. In either case, wind clumping is inextricably linked to the fundamental physical problem of radiation driven winds.   Our objective in the study presented here is to exploit detailed time-series optical observations of the late-O supergiant $\\alpha$ Cam (HD~30614) to provide empirical constraints for {\\it ab initio} model atmosphere codes that are commonly used to synthesise spectra in order to derive fundamental stellar and wind parameters (e.g \\citealt[CMFGEN] {Hillier1998},\\citealt[WMbasic] {Pauldrach01}, \\citealt[FASTWIND]{Puls05}). Our approach is to constrain the components that contribute to changes evident in diagnostic wind and photospheric lines. We seek to `decompose' the variability patterns and assess in particular the degree to which fluctuations in the wind-formed features are due to `contamination' by photospheric and deep-seated atmospheric activity. This is particularly critical in late O and early B stars, where the \\Ha \\,emission may be weak; see e.g. the imprints of photospheric signals seen in the Balmer lines of $\\epsilon$ Ori \\citep[B0 Ia,]{Prinja04}. Our study of $\\alpha$ Cam is also designed to probe causal links between photospheric disturbances and spatial wind structure.  The target $\\alpha$ Cam is previously known to exhibit line-profile variability in \\Ha \\, (e.g. \\citet{Ebbets82, Fullerton1990, Kaper1997}) as well as in some other absorption lines in the optical spectrum \\citep{Fullerton96, Markova02}. Whist its UV resonance lines are strong and saturated, \\citet{Lamers1988} reported on weak low-velocity fluctuations and changes in the soft blue wings of saturated (`black') absorption troughs. Some fundamental parameters of $\\alpha$ Cam are listed in Table~\\ref{para}.  \\begin{table} \\centering \\caption{Fundamental parameters for $\\alpha$ Cam.} \\label{para} \\begin{tabular}{l@{\\hspace{2mm}}l@{\\hspace{2mm}}l@{\\hspace{2mm}}} Parameter & Value & Reference\\\\ \\hline \\hline \\\\ Spectral Type & O9.5 Ia & \\citet{Walborn1973} \\\\ $T_{\\rm eff}$ & 29000 & \\citet{Repolust04} \\\\ $R_{\\star}/R_\\odot$ & 32.5 & \\citet{Repolust04} \\\\ \\vesini & 115 km s$^{-1}$ & \\citet{Penny1996} \\\\ $M_{\\star}/M_\\odot$ & 20 & \\citet{Schaller1992} \\\\ $V_\\infty$ & 1550 km s$^{-1}$ & \\citet{Haser1995} \\\\ $P_{\\rm rot}(max)$ & $\\sim$ 14.3 days & \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure*} \\centering \\includegraphics[scale=0.67]{fig1r.ps} \\caption{Examples of variability in the \\he and \\Ha line profiles of $\\alpha$ Cam are shown for two characteristic time-scales; (i) night-to-night mean profiles (upper panels) and (ii) $\\sim$ hourly changes within a single night (lower panels).} \\label{fig1}% \\end{figure*} ",
        "conclusions": "{We have demonstrated in this study that the primarily, or partially, wind-formed spectral lines of the late O supergiant $\\alpha$ Cam can exhibit variability that arises in different physical regimes of the star. This `contamination' of wind signatures by variable non-wind components is directly relevant to line-synthesis modelling, particularly in the context of revising mass-loss rates due to the presence of small- and large-scale structure in the outflows.} In particular, we have found that the temporal behaviour of $\\alpha$ Cam's optical spectrum is characterised by at least three different variability patters: \\begin{enumerate} \\item[i)] {systematic changes in the photospheric lines, with some evidence that the behaviour may be linked to surface velocity fields due to non-radial pulsations}, \\item[ii)] short time-scale ($\\sim$ hourly) perturbations that differentially shift the radial velocity of He~I % lines, and potentially give rise to localised changes in \\Ha emission, \\item[iii)] a variable stellar wind, which may be parameterised typically in terms of a $\\sim$ 5{\\%} change in mass-loss rate over several days, but very occasionally this difference can rise to $\\sim$ 30{\\%}. The red-ward emission changes are accompanied by more enigmatic blue-ward changes that may result from absorption effects linked to localised density structures in the outflow. \\end{enumerate}  \\begin{figure} \\centering \\includegraphics[scale=0.67]{fig7r.ps} \\caption{`Representative' model \\Ha profiles, demonstrating the effect of shifting the TLUSTY photospheric models by $-$30 km s$^{-1}$ (dashed line; see Sect. 3.3).} \\label{fig7}% \\end{figure}    \\begin{figure} \\centering \\includegraphics[scale=0.67]{fig8r.ps} \\caption{\\Ha model fits (solid and dashed lines) to extreme changes observed in 2004 and 1999 (filled and open circles).} \\label{fig8}% \\end{figure}    Hot massive stars are expected to be subject to photospheric instabilities and/or pulsations.  During the last decade the  search for non-radial pulsations in OB stars has been largely motivated by the need to test the hypothesis of the ``photospheric connection\" which aims at linking  the cyclic variability observed in the winds of these  stars to processes in the stellar photosphere. The task however turned out to be  observationally very demanding  and as a result  only a limited number of OB stars are currently flagged as confirmed or suspected pulsators \\citep[e.g.][]{Fullerton96, Reid1996, deJong01, Kaufer06}.  The results derived in this study {raise the possibility} that the supergiant $\\alpha$~Cam might also be a non-radial pulsator.  Radial velocity variations  in  optical absorption lines of $\\alpha$~Cam have been reported previously by \\citet{Zeinalov86} and \\citet{ Musaev88}  who argued that these variations are likely  periodic  with a possible cycle length of 1.35, 1.54 or 2.93 days (with no preference to any of them).  However, our analysis of the 2002 and 2004 time-series  did not support any of  these periods but instead revealed a new one of 0.36 days, which may be interpreted in terms of prograde low-order non-radial pulsational behaviour. {The $\\sim$ 8--9 hour modulation may persist over two months in 2002, albeit with different phase shifts.} This signal is not evident in data collected more than 2 years later (November 2004); it is also not seen in the stellar wind components of the line profiles. {The interpretation of the 0.36-day  variation in terms of  photospheric velocity fields is also  supported  by weak fluctuations in the shallow photospheric lines of C~IV $\\lambda\\lambda$~5801, 5812.}   Apart from the 0.36-day periodic variation detected between  Jan- March 2002, our observations also documented occasional  large-amplitude (of about 15 to 30 \\kms) differential radial velocity shifts in \\he and \\hee absorption lines, accompanied by an increase in the \\Ha emission flux centred at \\vsys and localised within $\\pm$ \\vesini.  Large-amplitude differential  velocity shifts have been discovered by \\citet{Prinja04} in the optical absorption lines of $\\epsilon$ Ori (B0 Ia) and they are also present  in the photospheric lines of $\\beta$ Cepheid stars. However in both these cases the velocity shifts are regular and periodic while those seen in  $\\alpha$~Cam occurred over at least several  hours, during a specific night in 2004,  with no recurrence during  the other nights of the run. This phenomenon can be interpreted either in terms of large-amplitude changes in photospheric velocity fields or as the result of deep-seated, large-scale wind structure that occasionally modulates the density distribution in the innermost part of the wind.    Discrete Absorption Components moving from red to blue across the unsaturated  UV resonance lines are the characteristic signature of large scale structure in O star winds. Unfortunately the UV resonance lines of $\\alpha$ Cam  are saturated and do not provide any detailed information on variable wind features \\citep{Kaper1997}. We have re-examined the $IUE$ time-series of $\\alpha$ Cam from February 1991 and December 1994. Whilst there is subtle evidence for low velocity fluctuations in Si~IV $\\lambda\\lambda$1400 (and Si~III $\\lambda$1207), the clearest variability is in the soft blue wings of the saturated UV lines. For example, during the 1991 $IUE$ observing run of $\\sim$ 5 days, additional absorption is evident in the blue wing of Si~IV beyond $v_\\infty$ (= 1550 km s$^{-1}$) out to $\\sim$ $-$1700 km s$^{-1}$. This enhancement persists for $\\sim$ 24 hours and does not recur during the following 4 days. This finding is at least consistent with the notion based earlier on optical observations, that large scale structure may develop in the wind of $\\alpha$ Cam; the present study suggests however that these structures are somewhat unstable and thus do not give rise to longer-term periodic phenomena in the wind.   The possibility that the wind of $\\alpha$ Cam is not smooth but structured has got additional support by the recent results derived by \\citet{Fullerton06} and \\citet{Puls06}. In particular,  \\citet{Fullerton06} have modelled the P~V lines of $\\alpha$ Cam using the SEI method to derive products of mass-loss rate and ionization fraction. Adopting mass-loss rates derived from radio observations, they find resulting ionization fractions, q(P$^{4+}$), = 7$\\times$10$^{-3}$ for $\\alpha$ Cam. In the case of $\\alpha$ Cam the low ion fraction is highly discrepant with results from NLTE stellar atmosphere models which predict that the P~V ion fraction should be dominant and thus close to unity at these effective temperatures. The discrepancy is thought to indicate that the wind is highly clumped and has very small volume filling factors.  Further evidence in support of the clumped structure of $\\alpha$~Cam wind has been derived by \\citet{Puls06}. By means of a  simultaneous model analysis of \\Ha, IR and radio data these authors  showed  that the wind of $\\alpha$ Cam likely has a clumping factor in the innermost region ( between 1.05  \\footnote{The wind below 1.05\\Rstar was assumed to be unclumped in agreement with results from hydrodynamical simulations which showed that the line-driven  instability needs some time to grow before significant structure is formed.} and  2\\Rstar),  being a factor of 2.6 larger than in the outermost wind (beyond 50\\Rstar). Using the \\Ha emission wings as a tracer of wind clumping in the intermediate wind (between 2 and  15\\Rstar), Puls et al. furthermore found that the clumping factor in this region is similar or fractionally larger than in the inner wind.  Interestingly, the properties of the TVS of \\Ha (derived in the present study  and by \\citet{Markova02}) also seem to be consistent with results obtained via line-profile simulations including wind clumps \\citep{Markova05}. In particular, the velocity limits as well as the established asymmetry of the TVS are both consistent with a model  including 'broken shells\" (i.e.~ clumps), uniformly  distributed over the wind volume, with density contrast  $\\delta \\rho/\\rho_0$ of $\\pm$0.35 (model series BS21 in \\citep{Markova05}) where $\\rho_0$ is the stationary density. Thus, it may be  that \\Ha night-to-night variability observed in $\\alpha$ Cam is at least partially connected with presence of small-scale structure (clumps, blobs) in the innermost part of the wind.  Overall, there is now substantial observational evidence which suggests that the wind of $\\alpha$ Cam is not smooth but structured over, both, large and small spatial scales.  In view of this possibility it is therefore not surprising that we failed to reproduce the observed set of \\Ha profiles (Sec. 4) by means of model calculations. Indeed, the approximate method of \\citet{Puls1996} employed here, as well as most of the state-of-art model atmosphere codes,  rely on the assumption of a globally stationary wind with a smooth density /velocity stratification, {coupled to a `stable' photospheric line contribution}, and are therefore incapable in principle of describing the observed {\\it blue-ward (absorptive)} behaviour of \\Ha in $\\alpha$ Cam and of any other star with a substantially structured wind.    \\acknowledgement{ This study was in part supported by funds from the National Scientific Foundation to the Bulgarian Ministry of Education and Science (F-1407/2004). We are also grateful to Bulgarian Academy of Sciences and the Royal Society (UK) for a collaborative research grant.  NM acknowledges the financial support of the Bulgarian Academy of Sciences and PPARC. NM is also grateful to her colleague Sabotinov (BAS) for his understanding and moral support during this study.}"
    },
    "0606/astro-ph0606516_arXiv.txt": {
        "abstract": "It is shown that the two-part Minimum Description Length Principle can be used to discriminate among different models that can explain a given observed dataset. The description length is chosen to be the sum of the lengths of the message needed to encode the model plus the message needed to encode the data when the model is applied to the dataset. It is verified that the proposed principle can efficiently distinguish the model that correctly fits the observations while avoiding over-fitting. The capabilities of this criterion are shown in two simple problems for the analysis of observed spectropolarimetric signals. The first is the de-noising of observations with the aid of the PCA technique. The second is the selection of the optimal number of parameters in LTE inversions. We propose this criterion as a quantitative approach for distinguising the most plausible model among a set of proposed models. This quantity is very easy to implement as an additional output on the existing inversion codes. ",
        "introduction": "When a scientist tries to analyze a given observed data set, it is customary to begin by examining the data in various ways, such as plotting the data and looking for patterns. In a second step, the scientist proposes a number of physically plausible models that can reproduce the observed data set. A model fitting procedure is then applied to the data set so that the parameters that characterize each model are inferred. After all the proposed models are fitted to the data, the next step is to compare them and infer which of the fitted models is the most suitable. Methods for such a task have been developed. For instance, for models that represent a hierarchical structure, we can use Akaike's Information Criterion \\citep{akaike74} or Mallows' $C_p$ \\citep{mallows73}. These methods can be applied even in the case that the set of proposed models does not contain the perfect model. In this case, the aim is to select the most optimal one. If the models are of completely different type and do not belong to a hierarchical structure of models, cross-validation type methods can be applied. However, they can be computationally demanding.  The standard procedure to select the optimum model is to make use of the Occam's Razor Principle (also known as the Principle of Parsimony). This principle is commonly applied in science and is usually thought to be an heuristic approach for eliminating unnecessary complex hypothesis. The principle states that the selected model has to provide an equilibrium between the model complexity and its fidelity to the data. Nevertheless, determining the simplest model is often very complicated. It is usually argued that the number of parameters that parameterizes the model should be smaller than the number of degrees of freedom of the data. However, it is a difficult matter to estimate the number of degrees of freedom of the data.  An alternative and successful procedure is to consider the problem in terms of a communication process. Assume that a sender S is interested in sending a given observation to a receiver R. Several techniques are available for such a task. The most trivial one is to send the whole amount of data through a given channel. If the data set is very large, the length of the message will be consequently very large. If S is able to model the observation using a given model, it is wiser and shorter to transmit the model followed by the points in the observation that are not correctly reproduced by the model. If the model is good and simple enough, the length of the message between S and R will be much shorter than sending the complete observed data set. This compression will be degraded when the model needed for explaining the observations is made unnecessarily complex. \\cite{rissanen78} was the first in suggesting that the code length could be used for model comparison. This principle is nowadays known as the Minimum Description Length (MDL) Principle, which states that we should \\emph{choose the model that gives the shortest description of the data}. In this framework, there is an interplay between models that are concise and easy to describe and models that produce a good fit to the data set and captures the important features evident in the data. Of course, this is neither the only nor the best strategy for such purposes. However, its application is desirable in comparison with ad-hoc or trial-and-error ways of performing model selection. The reason is that MDL principle has strong theoretical roots that lie on the Kolmogorov complexity theory \\citep{vitanyi00,gao00}. As we will show, the practical application of the MDL principle is very easy to implement and it constitutes an ideal approach for model selection.  This paper will be focused on one of the version of the MDL principle, the so-called two-part MDL. This strategy was developed by Rissanen in a series of papers published during the seventies and eighties \\citep[e.g.,][]{rissanen78,rissanen83,rissanen86} and summarized in \\cite{rissanen89}. We give a brief summary of the main results in section \\ref{sec:theory}. As shown below, the main idea is to write the description length of a given model applied to a data set as the sum of the length of the code for describing the model and the length of the code for describing the data set fitted by the model.  The interpretation of spectropolarimetric data in solar and stellar observations allows to infer information about the properties of the magnetic field. Almost always, the recovery of the magnetic field vector is based on the assumption of a model. These models depend on a given amount of variables that can be estimated by fitting the observed Stokes profiles. Among these models, we can find the Milne-Eddington approximation, the LTE approximation, the MISMA hypothesis \\citep{sanchez_almeida_misma96}, etc. Sometimes, it is clear that a model is the most appropriate when observational clues are available. However, it is customary that these clues are not present and one relies on one model based on completely or partially subjective reasons.  Some of these models are based on an extraordinarily large number of parameters. Until now, there is not any critical investigation toward analyzing whether enough information is available in the observations to constraint such a large amount of parameters. This work is a first step in this direction. We propose the application of the MDL criterion to quantitatively differentiate between possible models taking into account the information available in the observations. ",
        "conclusions": "We have presented the Minimum Description Length Principle developed by \\cite{rissanen78} to discriminate between a set of available models that can approximate a given data set. We have briefly presented its relation with the Bayesian approach of model selection through the application of the Shannon's theorem. In our opinion, the MDL principle presents a user friendly procedure for model selection. We have presented simple ways of estimating the message length that can be applied for communicating integer and real numbers. A more computer-oriented and easy to implement procedure has been also shown.  For the sake of clarity, we have applied the MDL principle to simplified problems. The selection of the optimal number of PCA components when de-noising observed Stokes profiles and the selection of the optimal number of nodes in the temperature and magnetic field depth profile obtained from LTE inversion of Stokes profiles. The results show the potential of this technique for model selection, with the advantage of being very simple to calculate.  As the main conclusion of this paper, we propose using the MDL principle as a way to quantitatively select among different competing models. We propose to include the description length as one of the final outputs of any inversion code. This makes it very easy to select the optimal model from a proposed set of models based on the framework of the MDL principle. It is of interest to stress that this principle can be used to select among models that are based on different scenarios. Usually, most complex scenarios translate into more degrees of freedom that may not be constrained by the observations. Using the MDL principle, the selected model among all the possibilities might not give the best fit to the observations but it represents the model that produces a good fit with a conservative number of parameters.  As a final remark, a possible future way of research may be how the MDL principle can be implemented as a regularization of the merit function in the existing inversion techniques. As a consequence, one would end up with an inversion code that automatically selects the optimal model."
    },
    "0606/astro-ph0606270_arXiv.txt": {
        "abstract": "{} {To test if the irregularly variable optical emission of hot stars in the Small Magellanic Cloud is due to variable amounts of Brehmstrahlung by ionized material in the circumstellar environment. The light curve variability of this sample of hot stars ($\\sim 200$ objects) is described by Mennickent et al. (2002).} {(1) Probing the relation between flux excess in the optical and in the near-infrared $J$-band using EROS light curves (time span of 5 years) and 2MASS photometry. (2) Analyzing the relation between colour and magnitude of the variability by means of an analytical, time-dependent model of emission from ionized material in an outflowing circumstellar disk.} {(1) A correlation between optical and near-IR flux excess is found. (2) A bi-valued relation between excess colour and magnitude for $\\sim 100$ objects is discovered. It results in a loop structure in a Colour-Magnitude diagram. Significantly, this loop is traversed in a clockwise sense for $\\sim90\\%$ of the stars, and anti-clockwise for the remainder. (3) A simple analytical model of an outflowing disk and a variable mass loss rate is capable of reproducing observed light curves and the bi-valued colour-magnitude relation.} {The optical variability of a large fraction of the hot stars is due to variations in the amount of Brehmstrahlung. The circumstellar envelope can contribute up to 50\\% of the total optical flux. This variability can be interpreted as due to an outflowing ionized CS disk, that evolves into a ring structure. The observed bi-valued colour-magnitude relation is an optical depth effect at the various wavelengths of interest. The loop is traversed clockwise by outflowing matter, but anti-clockwise by infalling material. As the material is generally outflowing, the CS environment of the star is progressively cleared once the star stops losing mass. This study makes use of public OGLE, MACHO and 2MASS data.} ",
        "introduction": "\\label{parintro} An important contribution to observational astronomy and in particular to the studies of variable stars was made in the nineties by the microlensing experiments: EROS (e.g. Aubourg et al. 1995)\\nocite{1995A&A...301....1A}, MACHO (e.g. Alcock et al. 1993)\\nocite{1993Natur.365..621A}, OGLE (e.g. Udalski et al. 1997)\\nocite{1997AcA....47..319U}. The projects produced an enormous quantity of photometric data for millions of stars in the Galactic Bulge and the Magellanic Clouds covering nearly a decade of continuous observations. Among the interesting discoveries reported are the Magellanic Clouds' blue variable stars. Two papers in 2002 by Keller et al. (MACHO)\\nocite{2002AJ....124.2039K} and Mennickent et al. (OGLE)\\nocite{2002A&A...393..887M} reported the huge variety of photometric variability displayed by hundreds of blue stars in the Large and Small Magellanic Cloud, respectively. Both papers type the various light curves according to their morphology. A statistical comparison of various characteristic quantities of the variability between LMC and SMC blue variable stars resulted in pronounced differences for phenomenologically similar light curves (Sabogal et al. 2005). Keller et al. (2002; hereafter K02) examine the presence of Balmer line emission in a subset of their sample and find a detection rate of 91\\%. They conclude that the Large Magellanic Cloud population of blue variable stars is consistent with the Be star phenomenon (see also Lamers et al. 1999\\nocite{1999A&A...341..827L} and de Wit et al. 2003\\nocite{2003A&A...410..199D}).  In the present work, we focus on the phenomenal optical brightness and colour variability of blue variable stars in the Small Magellanic Cloud. The stars and their light curves were first described by Mennickent et al. (2002; hereafter M02).  About 7\\% of the SMC Be stars was actually found by M02 to have well-defined {\\it periodically} varying light curves (see Mennickent et al. 2003\\nocite{2003ipc..conf...89M}). These stars are not considered here. Instead we focus on the {\\it irregular} photometric variability displayed by the majority of the Be stars on time scales of 1 to $~1000$ days and amplitudes up to $1^{m}$. Given that the observed Balmer line emission variability correlates with the photometric variability for similar LMC blue variables (K02) and the fact that generally the stars have redder colours at brighter phases, bound-free and free-free (bf-ff) emission originating in an ionized circumstellar envelope (CSE) is a strong candidate for the observed optical variability. This is generally the case for Galactic stars showing the Be phenomenon. A correlation between near-infrared (near-IR) flux excess and optical flux excess should therefore exist for the SMC Be stars.  Such a correlation was already shown to exist for Galactic Be stars by Dachs \\& Wamsteker (1982). We proceed therefore by cross-correlating the SMC blue variable stars discovered and catalogued by M02 with the near-IR 2MASS catalogues.  The paper is organized as follows. In Sect.\\,\\ref{defi} we specify the way the sample of blue variable SMC stars used in this contribution is defined. We note that in the following we refer to the stars simply by Be stars. In Sect.\\,\\ref{feon} the presentation of the cross-correlation between 2MASS and the sample stars follows. Then in Sect.\\,\\ref{colour} we describe the remarkable change in optical colour as a function of the brightness. This behaviour is modeled in Sect.\\,\\ref{model}. We discuss the results and conclude in Sects.\\,\\ref{disc} and \\ref{concl}. ",
        "conclusions": "\\label{concl} Irregular optical variability observed in a large ensemble of blue stars in the Small Magellanic Cloud is due to variations in the amount of Brehmstrahlung, most likely located in a circumstellar envelope or disk in analogy with classical Be stars. The CS disk can thus contribute up to 50\\% of the total optical luminosity of the stars, as derived from the amplitude of optical variability. Using a simple model we have shown that these flux changes can be interpreted as due to a time variable mass loss of the central star. The disk material is always outflowing. The outflowing material is observed to produce a bi-valued colour-magnitude relation (a loop in a colour-magnitude diagram), that can be ascribed to optical depth effects. The loop is traversed in a clock-wise sense for outflowing material, but in a anti clock-wise sense for accreted material. The data is compatible with disks that have a (small) radial outflow velocity that can mimic a stationary situation. In a small number of instances the material may actually be falling back on the star."
    },
    "0606/astro-ph0606046_arXiv.txt": {
        "abstract": "% {A specific example of Bianchi Type VII$_h$ models, i.e. those including universal rotation (vorticity) and differential expansion (shear), has been shown in Jaffe et al. (2005) to correlate unexpectedly with the \\emph{WMAP} first-year data.} {We re-assess the signature of this model in the \\emph{WMAP} 3-year data.} {The cross-correlation methods are described in Jaffe et al. (2006a). We use the \\emph{WMAP} 3-year data release, including maps for individual years, and perform additional comparisons to assess the influence of both noise and residual foregrounds and eliminate potential non-cosmological sources for the correlation.} {We confirm that the signal is detected in both the combined 3-year data and the individual yearly sky maps at a level consistent with our original analysis. The significance of the correlation is not affected by either noise or foreground residuals.  } {The results of our previous study are unchanged.} ",
        "introduction": "\\label{sec:introduction}  In \\citet{jaffe:2005,jaffe:2006a}, we reported on the unexpected detection of a correlation between the cosmic microwave background (CMB) sky measured by the \\emph{WMAP} first-year data release and a Bianchi Type VII$_h$ template. This result is particularly provocative % in that such a signal may provide an explanation for several independently discovered anomalies: the low quadrupole amplitude \\citep{doc:2004} relative to the best-fit cosmological model, the curious alignment of several large angular scale multipoles along the so-called ``axis of evil'' \\citep{land:2005a}, a significant power asymmetry between two hemispheres on the sky \\citep{hansen:2004a,eriksen:2004a} and a large temperature decrement towards a specific part of the southern sky as revealed by a wavelet analysis \\citep{cruz:2005,vielva:2004}.  The subtraction of the Bianchi component leaves a statistically isotropic sky.  In this note, we examine the newly released three-year data described in \\cite{hinshaw:2006} and show that the signal is detected both in the combined 3-year data and in each individual yearly map, despite differences in the data processing pipelines, calibration and corrections for foreground contamination. ",
        "conclusions": ""
    },
    "0606/astro-ph0606114_arXiv.txt": {
        "abstract": "{}{We report the \\integ/\\ibis observations of the \\coma in the hard \\xray/soft \\gammaray domain.}{Since the \\coma appears as an extended source, its global intensity and significance cannot be directly extracted with standard coded mask analysis. We used the method of imaging the extended sources with a coded mask telescope developed by Renaud \\etal (2006).}{The imaging capabilities and the sensitivity of the \\ibis/\\isgri coded mask instrument allows us to identify for the first time the site of the emission above $\\sim$ 15\\un{keV}. We have studied the \\coma morphology in the 18-30\\un{keV} band and found that it follows the prediction based on \\xray observations. We also bring constraints on the non-thermal mechanism contribution at higher energies.} {} ",
        "introduction": "\\label{s:intro}  Clusters of galaxies are the largest bound structures in the visible universe, and amongst the most luminous ones. This makes them important cosmological probes. The intercluster medium (\\icm) is heated to high temperatures by the initial collapse. In addition, it is believed that as groups and smaller clusters of galaxies merge to form larger ones, violent shocks compress and heat the intercluster medium to \\xray emitting temperatures ($10^7$--$10^8$\\un{K}). These shocks could also drive the acceleration of ions and electrons to relativistic energies, thereby producing a spectrum extending from the radio to the \\gammaray domains through processes such as synchrotron emission and inverse Compton (hereafter, IC) scattering.  The observations of diffuse radio halos from cluster cores \\cite{c:giovannini93}, and the more recent reported detection of extreme ultraviolet (EUV) and hard \\xray (\\hxr) emission from clusters \\citep{c:bowyer99,c:fusco99,c:rephaeli99} seem to point to the possible presence of non-thermal and/or supra-thermal particles. Although the origin of the radio halos is most probably synchrotron emission from high energy electrons, that of the EUV photons and hard \\xrays has been the source of much debate and is still uncertain. Sarazin (1999) suggested that the extreme UV and hard \\xrays could arise from IC scattering of the energetic, synchrotron producing electrons off the cosmic microwave background radiation (see also Bykov \\etal 2000 for a detailed discussion of nonthermal high energy emission in galaxy clusters). Alternatively, this emission could also be due to bremsstrahlung radiation of supra-thermal electrons accelerated by turbulent gas dynamics in the \\icm \\cite{c:ensslin99,c:sarazin00}. If the \\hxr cluster component is confirmed, then the question remains: is it thermal or non-thermal in origin?  Thus far, \\xray imaging observations of clusters such as Coma, have shown that the emission in the \\xray range is predominantly thermal, originating from bremsstrahlung in the hot \\icm. \\sax and \\rxte detected a hard \\xray excess that extends beyond 10\\un{keV} and apparently deviates from the steep thermal spectrum expected from pure bremsstrahlung emission \\cite{c:fusco99,c:fusco04,c:rephaeli99,c:rephaeli02}. Given that these observations were carried out with non-imaging instruments, contamination from point sources contained in their field of view is not unlikely. Moreover, a recent second look at the \\sax data on Coma seems to show that there is no evidence of a \\hxr excess \\cite{c:rossetti04}. This debate can only be convincingly resolved and put to rest through sensitive imaging observations at high energies beyond the thermally dominated \\xray emission.  In this letter, we report on hard \\xray/soft \\gammaray observations of the \\coma with the \\ibis/\\isgri \\cite{c:ubertini03,c:lebrun03} coded mask instrument onboard \\integ \\cite{c:winkler03}. A large field of view (29$^{\\circ}$\\,$\\times$\\,29$^{\\circ}$, 8$^{\\circ}$\\,$\\times$\\,8$^{\\circ}$ fully coded), fine imaging capabilities (PSF of 12\\am FWHM, Gros \\etal 2003), and an unprecedented 3$\\sigma$ broadband sensitivity of $\\sim$\\,1\\un{mcrab} at 20\\un{keV} ($10^5$\\un{s}, $\\Delta$E\\,=\\,E/2), make \\ibis/\\isgri ideal for probing the nature of the \\hxr excess in the \\coma: {\\bf the first extended source detected with  this instrument}. ",
        "conclusions": "\\label{s:concl}  The unequalled imaging capabilities and sensitivity of the \\ibis/\\isgri telescope on \\integ give us the opportunity of imaging extended sources in the hard \\xray and soft \\gammaray domains for  the first time. We have shown that the morphology of the emission from the \\coma in the 18--30\\un{keV} energy range is akin to that in \\xrays when the temperature variations measured by \\xmm across the cluster's emission are taken into account. There is no evidence for the presence of non-thermal emission arising from a region with the same morphology as that associated with the thermal emission. We find that the integrated intensity is in good agreement with previous \\rxte and \\sax observations of Coma.  Future \\integ observations will surely reveal more details about the morphology and nature of the emission from Coma. Our analysis was based on the assumption that the morphology of the emission is the same in the 30--50\\un{keV} range as it is between 18 and 30\\un{keV}. Given that the upper limit we quote for the 30--50\\un{keV} flux is only a factor of 1.5 above the mean \\rxte spectrum, and that non-thermal mechanisms such as IC scattering and non-thermal bremsstrahlung are expected to  contribute half the flux in this range, deeper \\integ observations ($\\sim$1.5\\un{Ms}) will permit a  fine morphological study of the emissive region. This will provide the opportunity to definitively determine whether there is or not a \\hxr tail, and if so to identify the site of particle  acceleration in the \\coma for the first time."
    },
    "0606/astro-ph0606608_arXiv.txt": {
        "abstract": "We present a study of the nature of the blue early-type galaxies (BEGs) in the GOODS north and south fields using the GOODS HST/ACS archival data. Using visual inspection, we have selected 58 BEGs and 113 normal red early-type galaxies (REGs) in the sample of 1,949 galaxies with spectroscopic redshifts. We find that the BEGs are generally bluer, fainter, and less-massive than the REGs, although a few BEGs are exceptionally bright and massive. The number fraction of the BEGs to total early-type galaxies is almost constant ($\\sim0.3$) at $z \\le 1.1$. In addition, we find that the size of the BEGs in given redshift bin decrease as redshift decreases. The BEGs look similar to the REGs in the images and surface brightness profiles. However, at least 27 BEGs show traces of tidal disturbances in their fine structures: elongated cores, off-centered cores, asymmetric internal color distributions, tidally distorted outer structures, collisional rings, or very nearby companions. Twenty-one BEGs are detected in the X-ray bands and eleven of them are as luminous as $L_{0.5-10 {\\rm keV}} \\ge 10^{43.5} ~{\\rm erg~s}^{-1}$, indicating the existence of AGNs in their centers. These results show that at least a half of the BEGs may be descendants of mergers/interacting-galaxies and that at least a quarter of the BEGs may be AGN-host galaxies. The BEGs may evolve into REGs, and the size evolution of the BEGs is consistent with the galactic \\emph{downsizing} scenario. ",
        "introduction": "Early-type galaxies are one of the key objects in the modern observational cosmology. One of the fundamental questions about early-type galaxies is their formation history, which is closely related to the history of cosmic mass assembly and cosmic star-formation \\citep{lee03}. There are two competing models for the formation of massive elliptical galaxies. One is a monolithic collapse model, according to which massive elliptical galaxies were formed via a single collapse of proto-galactic cloud followed by a rapid formation of stars \\citep{par67,tin72,lar74}. The other is a hierarchical merging model where massive galaxies were formed through continuous merging and accretion of smaller galaxies \\citep{too77,sea78}.  The properties of early-type galaxies at low redshift are relatively well known in many aspects thanks to a tremendous number of studies for several decades: the surface brightness profiles of early-type galaxies follow the $R^{1/4}$ law \\citep{dev48}; there is a strong correlation between luminosity and central velocity dispersion of early-type galaxies \\citep{fab76b}; early-type galaxies show very slow rotation \\citep{ber75,ill77,dav83}; there are poor cold-gas contents in early-type galaxies, implying rare star-formation events \\citep{fab76a}; the color and magnitude of early-type galaxies are strongly correlated \\citep{vis77}; there is a tight correlation among size, surface brightness, and velocity dispersion of early-type galaxies \\citep[fundamental plane; ][]{luc91,jor93}; and the center of an early-type galaxy is redder than its outer part \\citep{vig88,fra89}. However, it is still unclear whether all early-type galaxies have such `typical' properties. Because understanding the property variation of early-type galaxies along redshift makes it possible to constrain the model for the evolution of early-type galaxies more strongly, it is needed to observe and investigate early-type galaxies in the deep universe.  Recent studies based on the HST images found a new class of early-type galaxies in the distant universe that have different properties from nearby early-type galaxies. \\citet{abr99} showed from the analysis of internal color dispersion that 40\\% of early-type galaxies (four out of eleven) at $0.4<z<1$ in the Hubble Deep Field \\citep[\\hereafter HDF; ][]{wil96} have evidence of current star-formation. \\citet{men99} also identified three star-forming E/S0 galaxies among the HDF galaxies. Subsequently, \\citet{men01a} estimated the surface color dispersion of the HDF objects and reported that about 30$\\%$ of 79 spheroidals show strong internal-color variation, indicating the existence of not only old stellar populations but also young stellar populations. These objects generally show a bluer center than its outer part, unlike typical nearby early-type galaxies. \\citet{men01b} tried to explain the origin of the blue centers with a multizone single-collapse model, in which late infalls of material into the high-potential core could cause prolonged star-formation in the central region of a galaxy. Blue E/S0 galaxies were also found in other HST fields. \\citet{im01} found ten blue spheroids at $0.2<z<1$ from the HST ``Groth Strip'' \\citep{gro94} data. Based on the small velocity dispersions of those spheroids, \\citet{im01} concluded that they are low-mass systems suffering late star-formation events. Later, \\citet{men04} identified early-type systems with inhomologous internal colors (30\\%--40\\% of 116 spheroids) in the HST/ACS UGC 10214 field using photometric redshifts. They suggested that the internal color distribution of these objects can be explained by recent star-formations, and that the fraction of `active systems' does not vary severely at $z < 1$. In addition, \\citet{men05} pointed out in a study of early-type galaxies in Abell 1689 that the existence of an AGN in a galaxy is another possible origin of the blue center in these galaxies.  Blue E/S0 galaxies were also found in two recent great missions: UDF \\citep[Hubble Ultra Deep Field; ][]{bec03} and GOODS \\citep[Great Observatories Origins Deep Survey; ][]{dic03}. \\citet{elm05a} found 30 early-type galaxies with blue central clumps in the UDF, and estimated the sizes, masses, and ages of six galaxies with spectroscopic redshifts. They found that the six galaxies are 1 -- 5 Gyr old (and their clumps are younger than 100 Myr) and a half of them are more massive than $10^{10}$ M$_{\\odot}$, although there could be some degeneracy and model-dependence in those estimations. \\citet{fer05} analyzed the statistical properties of 249 early-type galaxies in the southern GOODS field, and reported that one third of the early-type galaxies at intermediate redshifts have blue colors. Based on the analysis of the color gradients, color scatters, and redshift-variations, they concluded that there exist a clear cut between blue E/S0 galaxies and red E/S0 galaxies in their color, and that the `two-burst formation scenario' can reproduce the blue E/S0 galaxies. However, fine structures of these blue galaxies were rarely investigated in previous studies.  In this paper, we present a study of the nature of the blue early-type galaxies in the GOODS north and south fields, based on the GOODS HST/ACS archival data. We use only galaxies with spectroscopic redshift and focus on studying the fine structures of the target galaxies. The outline of this paper is as follows. Section 2 describes the data set we used, and \\S3 explains  our selection process for blue early-type galaxies. We present the results of our photometric, structural, and statistical analyses of the blue early-type galaxies in \\S4. Primary results are discussed in \\S5, and a summary and conclusions are given in \\S6. Throughout this paper, we adopted the cosmological parameters: $h=0.7$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{M}=0.3$. ",
        "conclusions": "In this paper, we present a study of the properties of BEGs in the GOODS HST/ACS fields. We selected 171 early-type galaxies visually in the GOODS HST/ACS archival data with spectroscopic redshift. We divided the early-type galaxies into 58 BEGs and 112 REGs using their $(i-z)_{AB}$ color distribution at given redshift. The BEGs have well-defined elliptical shapes and $R^{1/4}$ surface profiles similar to the REGs, with just a few exceptions. However, the analysis of internal color distribution, X-ray luminosity, and spectral line features show that almost a half of the BEGs have evidence of tidal events and at least a quarter of the BEGs probably contain an AGN in their centers. From the analysis of the size and magnitude of the BEGs and the REGs, we have found the evidence that the sizes of BEGs are decreasing as redshift decreases, which is consistent with the \\emph{downsizing} scenario with hierarchical merging.  We conclude that the BEGs may be primarily descendants of past merger/interacting-galaxies and secondarily the AGN-host early-type galaxies, and that BEGs may evolve into normal or dwarf elliptical galaxies. However, there are still several open questions to answer. First, we could not find any BEGs at $z>1.2$ as large as REGs. If larger REGs formed at the higher redshift and if all REGs were BEGs at one time, many BEGs as large as giant REGs are expected to exist at $z>1.2$. Therefore, the question is whether all REGs evolved from BEGs or not. If not, what is the difference between those two (or more) building processes of early-type galaxies? The mechanism that makes the central deep blue clumps in almost a half of the BEGs is another question. It is hard to tell whether all \\emph{blue clumpy} BEGs contain AGNs or not, since the X-ray detection or matching with optical catalog is limited. Deep observation in other bands (e.g., mid-infrared spectroscopy using JWST \\citep{rie05}, which may enables to identify the PAH emission of faint objects) will be helpful to resolve which is dominant in the blue clumps between AGN activity and star-forming activity. In addition, the origin of the red inner rings in some BEGs is also a mystery to be solved with further studies."
    },
    "0606/astro-ph0606322_arXiv.txt": {
        "abstract": "The electrostatic potential that keeps approximate charge neutrality in neutron star matter is self-consistently introduced into the formalism for \\emph{rotochemical heating} presented in a previous paper by Fern\\'andez and Reisenegger. Although the new formalism is more rigorous, we show that its observable consequences are indistinguishable from those of the previous one, leaving the conclusions of the previous paper unchanged. ",
        "introduction": "\\label{sec:intoduction}  Rotochemical heating is one of several mechanisms that can keep neutron stars (NSs) hot beyond their standard cooling time of $\\sim 10^7~\\mathrm{yr}$ (see \\citealt{fernandez05} [hereafter FR05] for further references, and \\citealt{yakovlev04} for a recent review on neutron star cooling). It originates in deviations from beta equilibrium produced as the star gets compressed due to its decreasing rotation rate, implying a conversion of rotational energy into thermal energy. This effect is capable of keeping old millisecond pulsars (MSPs) at surface temperatures $\\sim 10^5$~K, consistent with likely thermal ultraviolet emission from PSR J0437-4720 (\\citealt*{kargaltsev04}; FR05).  The heating mechanism works as follows. As a NS spins down, the centrifugal force acting on each fluid element diminishes, raising the local pressure. This compression results in a change of the equilibrium concentration of each particle species. Reactions that modify the chemical composition drive the system towards the new equilibrium configuration. But if the rate at which reactions do this task is slower than the change of the equilibrium concentrations due to spin-down compression, the matter remains out of chemical equilibrium. The excess energy stored in the chemical imbalance is dissipated, partly by enhanced neutrino emission and partly by heating the matter in the star. Thus, the star is kept warm and can radiate for long after the standard cooling time.  After its introduction \\citep{reisenegger95}, this heating mechanism was studied by several authors: \\citet{cheng96} applied it to quark stars, \\citet{reisenegger97} estimated the effects of superfluidity, and \\citet{iidasato97} studied the heating due to compositional transitions in the crust due to spin-down compression.  The most striking prediction associated with rotochemical heating is that, if the spin-down timescale is substantially longer than any other timescale involved (with the likely exception of magnetic field decay), the star arrives at a \\emph{quasi-steady state}, in which the temperature depends only on the current, slowly changing value of $\\Omega \\dot{\\Omega}$ (the product of the angular velocity and its time derivative), and not on the star's previous history \\citep{reisenegger95}. This provides a simple way to constrain the physics involved in theoretical models, once the spin parameters and observed surface temperature of a given, sufficiently old NS are known.  In spite of the previous work, FR05 were the first to go beyond order-of-magnitude estimates and calculate the thermal evolution of NSs with rotochemical heating, taking the structure of the star fully into account in the frame of general relativity, and using realistic equations of state of dense matter. They developed a general formalism to treat the evolution of the temperature and departures of chemical equilibrium, implementing it for the simplest possible core composition, namely neutrons, protons, electrons, and muons ($npe\\mu$ matter). Using state-of-the-art equations of state, and modified Urca reactions as the dissipative process, they obtained results consistent with the observed ultraviolet emission from the millisecond pulsar PSR J0437-4720 \\citep{kargaltsev04}.  Here, we show that their formalism is not completely rigorous, in the sense of ignoring the perturbations to the electrostatic potential that must be present inside the NS in order to keep nearly exact charge neutrality among the different charged particles present inside the star.  We start (in \\S 2) by discussing the physical origin, importance, and strength of the electrostatic potential expected inside a NS, as well as its relation to the (nearly exact) condition of local charge neutrality. Then (in \\S 3), we review the main components of the rotochemical heating formalism developed by FR05, which are rigorously correct and not directly affected by the presence of electrostatic potential perturbations, showing that their main result, namely the existence and properties of the quasi-steady state, are unchanged. In \\S 4, we introduce the electrostatic potential into the equations of FR05 that relate the different chemical imbalance variables (chemical potential imbalances and deviations from equilibrium particle abundances) to each other. The quantitative consequences of these modifications are explored in \\S 5, whereas \\S 6 contains our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that the electrostatic potential inside a NS plays an important role in determining the spatial distribution of charged particles and their nearly exact local charge neutrality everywhere in the star. We take it into account explicitly in an extension of the formalism for rotochemical heating of FR05, in which the relations between the two sets of variables quantifying chemical imbalances (chemical potential differences and departures from equilibrium abundances) are affected, without changing the rest of the formalism. The quasi-steady state reached by old NSs is strictly unaffected by these corrections, and the time-dependent evolution at earlier times shows quantitatively negligible changes."
    },
    "0606/astro-ph0606578_arXiv.txt": {
        "abstract": "{} {The clustering properties of a large sample of $U$-dropouts are investigated and compared to very precise results for $B$-dropouts from other studies to identify a possible evolution from $z=4$ to $z=3$.} {A population of $\\sim8800$ candidates for star-forming galaxies at $z=3$ is selected via the well-known Lyman-break technique from a large optical multicolour survey (the ESO Deep Public Survey). The selection efficiency, contamination rate, and redshift distribution of this population are investigated by means of extensive simulations. Photometric redshifts are estimated for every Lyman-break galaxy (LBG) candidate from its $UBVRI$ photometry yielding an empirical redshift distribution.  The measured angular correlation function is deprojected and the resulting spatial correlation lengths and slopes of the correlation function of different subsamples are compared to previous studies.} {By fitting a simple power law to the correlation function we do not see an evolution in the correlation length and the slope from other studies at $z=4$ to our study at $z=3$. In particular, the dependence of the slope on UV-luminosity similar to that recently detected for a sample of $B$-dropouts is confirmed also for our $U$-dropouts. For the first time number statistics for $U$-dropouts are sufficient to clearly detect a departure from a pure power law on small scales down to $\\sim2\\arcsec$ reported by other groups for $B$-dropouts.} {} ",
        "introduction": "\\label{sec:intro}  For more than a decade now, the high-redshift universe has become reachable by observations mainly due to the development of efficient colour selection techniques.  The group around C.\\,C.  Steidel \\citep{1993AJ....105.2017S,1996ApJ...462L..17S,1999ApJ...519....1S} has introduced the Lyman-break technique selecting high-redshift, star-forming galaxies from optical multicolour data by their pronounced Lyman-break. Many groups have used this technique and analysed various properties of these galaxy populations from redshifts $z\\approx3$ up to $z\\approx7$.  A particular emphasis in these studies was given to the clustering properties of the Lyman-break galaxy (LBG) samples \\citep{1998ApJ...492..428S,1998ApJ...503..543G,1998ApJ...505...18A,2005ApJ...619..697A,2001ApJ...550..177G,2001ApJ...558L..83O,2004ApJ...611..685O,2005ApJ...635L.117O,2002ApJ...565...24P,2004ApJ...609..513B,2003A&A...409..835F,2005MNRAS.360.1244A,2006ApJ...642...63L,2006ApJ...637..631K,2006ApJS..162....1G}. Going back in cosmic time the correlation strength of these high-$z$ galaxies can be directly compared to $N$-body simulations or semi-analytical predictions yielding estimates of, e.g., the galaxy bias for early epochs. These results can then be used as an input for models of galaxy formation, and help to constrain the large number of free parameters to adjust. The more precise our knowledge of the clustering evolution becomes the more accurate our understanding of galaxy evolution will be.  In this context it is of crucial importance to reach a similar precision for the same galaxy populations at different redshifts. While in the beginning most LBG studies concentrated on relatively bright $z=3$ $U$-dropouts, in the last years more and more groups have focused on the investigation of LBGs at redshifts around $z=4$ selected as $B$-dropouts. This is obviously due to the fact that deep and wide $U$-band images are still a very telescope-time consuming task and some recent wide-field cameras like \\emph{Suprimecam} or space-based surveys like GOODS even lack a $U$-filter entirely.  By estimating the angular correlation function of nearly 17\\,000 $B$-dropouts selected from the Subaru/XMM-Newton Deep Field, \\cite{2005ApJ...635L.117O} find evidence for a departure from a pure power law on small scales. The same trend is reported by \\cite{2006ApJ...642...63L} for $B$- and $V$-dropouts from the very deep GOODS ACS data. Neither the number statistics in the former study mentioned nor the depth and angular resolution of the GOODS data have a comparable counterpart at slightly lower redshifts. The most precise estimates of $U$-dropout clustering to date come from \\cite{2005ApJ...619..697A} estimating the angular correlation function but not reporting any obvious excess on small scales. The question whether this is an evolutionary effect or whether this feature is only visible in the $z=4$ data because of their superior quality can only be answered by a $U$-dropout survey comparable in size to the $B$-dropout surveys mentioned above.  In this paper we describe our investigations of $z=3$ LBGs in the ESO Deep Public Survey (DPS). The methods presented here are based on our investigations in the Chandra Deep Field South (the DPS field Deep2c) presented in \\cite{2005A&A....Hildebrandt}. In Sect.~\\ref{sec:data} the data of the DPS and the selection of the LBGs are described. Simulations to assess the performance of our LBG selection are presented in Sect.~\\ref{sec:simulations}. The clustering ana\\-ly\\-sis is covered in Sect.~\\ref{sec:clustering}.  A summary and conclusions are given in Sect.~\\ref{sec:conclusions}.  Throughout this paper we adopt a standard $\\Lambda$CDM cosmology with $\\left[ H_0, \\Omega_{\\mathrm{m}}, \\Omega_{\\Lambda}, \\sigma_8 \\right]=\\left[70\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}},0.3,0.7,0.9\\right]$. We use Vega magnitudes if not stated otherwise. ",
        "conclusions": "\\label{sec:conclusions} We measure the clustering properties of a large sample of $U$-dropouts from the ESO Deep Public Survey with unprecedented statistical accuracy at this redshift.  Candidates are selected via the well-known Lyman-break technique and the selection efficiency is investigated and optimised by means of simulated colour catalogues. The angular correlation function of LBGs is estimated over an area of two square degrees, depending on depth, and a deprojection with the help of the photometric redshift distribution yields estimates for the correlation lengths of different subsamples.  We find clustering segregation with restframe UV-luminosity indicated by a decreasing correlation length and a decreasing slope of the correlation function with increasing limiting magnitude. The latter result was reported at redshift $z=4$ and is now confirmed at redshift $z=3$ for the first time. Furthermore, the unprecedented statistical accuracy of our survey at $z=3$ allows us to study the small-scale clustering signal in detail. We find an excess of the angular correlation function on small angular scales similar to that found previously at $z=4$.  Applying a halo model we find average masses for LBG-hosting halos at $z=3$ which are slightly larger than literature values for $z=4$ implying decreasing star-formation efficiency with decreasing redshift."
    },
    "0606/astro-ph0606052_arXiv.txt": {
        "abstract": "Using the 3-telescope IOTA interferometer on Mt. Hopkins, we report results from the first near-infrared ($\\lambda=1.65\\mu$m) closure-phase survey of Young Stellar Objects (YSOs).  These closure phases allow us to unambiguously detect departures from centrosymmetry (i.e., skew) in the emission pattern from YSO disks on the scale of $\\sim$4~milliarcseconds, expected from generic ``flared disk'' models.  Six of fourteen targets showed small, yet statistically-significant, non-zero closure phases, with largest values from the young binary system MWC~361-A and the (pre-main sequence?) Be star HD~45677.  Our observations are quite sensitive to the {\\em vertical} structure of the inner disk and we confront the predictions of the ``puffed-up inner wall'' models of Dullemond, Dominik, and Natta (DDN).  Our data support disks models with {\\em curved} inner rims because the expected emission appear symmetrically-distributed around the star over a wide range of inclination angles. In contrast, our results are {\\em incompatible} with the models possessing {\\em vertical} inner walls because they predict extreme skewness (i.e., large closure phases) from the near-IR disk emission that is not seen in our data.  In addition, we also present the discovery of mysterious H-band ``halos'' ($\\sim$5-10\\% of light on scales 0.01-0.50\\arcsec) around a few objects, a preliminary ``parametric imaging'' study for HD~45677, and the first astrometric orbit for the young binary MWC~361-A. ",
        "introduction": "The study of accretion disk evolution in Young Stellar Objects (YSOs) promises to directly link the fields of star formation with planet formation.  Spatially-resolved, multi-wavelength observations can directly probe the dust and gas distribution on solar system scales, allowing precise testing of disk models while simultaneously providing the crucial physical parameters for planet formation initial conditions.  Young Jupiter-mass planets embedded in natal disks should even be detectable in the future.  Already observations from long-baseline interferometers (IOTA, PTI, Keck) have empirically determined the near-infrared (NIR) emission sizes of YSOs, \\citep[e.g.,][]{rmg1999a,rmg2001,akeson2000,eisner2003}, finding the sizes to be strongly linked to the luminosity of the central stars \\citep{mm2002,monnier2005}.  These results are qualitatively explained by a model where the NIR emission arises entirely from the hottest dust at the inner edge of an accretion disk, and the measured sizes correspond to the dust-destruction radius around the young stars \\citep{lkha2001,natta2001}.  Recently, \\citet{eisner2004} reported convincing signs that NIR dust emission from many intermediate-mass Herbig Ae/Be stars are elongated, implicating disk-like structure as expected from star and planet formation theory.  The simple detection of ``elongation'' by an interferometer, however, does not well-constrain the geometry of the emission region.  For instance, current data can not distinguish between a) a centrosymmetric emission pattern as expected from a perfectly flat, inclined disk, b) elongated emission arising preferentially from one side of the disk \\citep[as expected from optical depth effects associated with flaring;][]{malbet2001}, or c) hybrid disk+halo models with multiple emission components \\citep{vinkovic2006}.  All current results are based on measuring interferometer fringe {\\em amplitudes}, while interferometric fringe {\\em phase information} is required to unambiguously detect deviations from simple symmetries. Ultimately, the phases and amplitudes are all needed to permit model-independent imaging using aperture synthesis techniques \\citep{tms2001}.  While atmospheric turbulence corrupts the direct measurement of fringe phase, interferometrists using three or more telescopes can measure the ``closure phase'' (CP), a phase quantity that is unspoiled by telescope-specific phase errors \\citep{jennison58,monnier_mss}. Recently, the Infrared Optical Telescope Array (IOTA) became the third facility to successfully achieve closure phase operation \\citep{traub2003,monnier2004b} and is the only such facility that has demonstrated the capability to study Young Stellar Objects \\citep{rmg2006}.   Here, we expand the original Herbig Ae/Be survey work of \\citet{rmg2001} using three IOTA telescopes simultaneously, tripling the rate at which visibility measurements are collected and allowing  closure phases to be measured for the first time for a sample of YSOs \\citep[see][for the first results for the prototype Herbig Ae star AB~Aur]{rmg2006}.  This work has been made possible by the advanced beam combiner IONIC3 \\citep{berger2003} which exploits integrated optics technology developed for the telecommunication industry.  This article presents a large volume of data and analysis and is organized in the following sections.  First, we describe the interferometer observations and present the data results in a summary form.  Second, we analyze the spectral energy distribution and visibility data together for each target in order to estimate the NIR emission size; this step is crucial for interpreting the closure phase results.  Sections 4 \\& 5 contain the most novel aspects of this paper. Here we outline our method for using closure phases to probe the amount of skew in the YSO disk emission.  Then, we apply our method to quantify the {\\em skewed} disk emission and discuss the implications on current disk models. Lastly, we treat the special case of HD~45677 and report the first astrometric orbit for the YSO binary MWC~361-A. ",
        "conclusions": "\\label{sectionc:conclusions}  Here we have presented the results of the first closure phase survey of Young Stellar Objects using an infrared interferometer.  We find most YSO targets show no signs of skewness in their emission on the scale of $\\sim$4 milliarcseconds.  Our data is incompatible with vertical wall models of the puffed-up inner rim \\citep[e.g.,][]{dullemond2001}. A ``skewed ring'' model was developed in order to quantify our results and the observed small closure phases favor inner dust rims which are smoothly curved \\citep[e.g.,][]{isella2005}.  At least a crude understanding of the inner AU of YSO disks is crucial for interpreting a host of other YSO observations, including spectrally-resolved CO data, spectro-astrometry, mid-IR sizes and spatially-resolved silicate feature minerology.  Baselines significantly longer than the maximum IOTA length of 38\\,m (i.e., from CHARA and VLTI) will be required to investigate this further for specific sources.  An independent and complete reanalysis of new H band \\viss data found fresh evidence for extended ``halos'' of emission with about 5-10\\% of the emission for a few targets \\citep[see also][]{leinert2001}.  Halos appear to be quite common in many young disk systems (T~Tauri, Herbig, FU~Ori), although the physical origin of this large-scale emission appears, as of yet, not understood.  These new data may re-open the debate concerning the relative importance of disks, halos, envelopes, localized thermal emission, and the applicability of hybrid models.  We also took the first steps towards imaging young stellar objects by studying the unusual system HD~45677, finding this Be star to be surrounded by an elongated and highly-skewed dust ring.  We agree with earlier suggestions that HD~45677 may not be a young object, but rather a member of the enigmatic class of main sequence hot stars with large IR excess.  Lastly, we calculated the first astrometric orbit for the short-period binary MWC~361-A, finding good agreement with recent spectroscopic work. We found mass estimates in line with expectations and discuss the exciting science potential from future studies of this unique system."
    },
    "0606/astro-ph0606264_arXiv.txt": {
        "abstract": "{\\emph{Context}. In the framework of the interstellar PAH hypothesis, far\\textendash IR skeletal bands are expected to be a fingerprint of single species in this class.\\\\ \\emph{Aims}. A detailed model of the photo\\textendash physics of interstellar PAHs is required for such single\\textendash molecule identification of their far-IR features in the presently available Infrared Space Observatory data and in those of the forthcoming Herschel Space Observatory mission.\\\\ \\emph{Methods}. We modelled the detailed photophysics of a vast sample of species in different radiation fields, using a compendium of Monte\\textendash Carlo techniques and quantum\\textendash chemical calculations. This enabled us to validate the use of purely theoretical data and assess the expected accuracy and reliability of the resulting synthetic far\\textendash IR emission spectra.\\\\ \\emph{Results}. We produce positions and intensities of the expected far\\textendash IR features which ought to be emitted by each species in the sample in the considered radiation fields. A composite emission spectrum for our sample is computed for one of the most favourable sources for detection, namely the Red Rectangle nebula. The resulting spectrum is compared with the estimated dust emission in the same source, to assess the dependence of detectability on key molecular parameters.\\\\ \\emph{Conclusions}. Identifying specific PAHs from their far\\textendash IR features is going to be a difficult feat in general, still it may well be possible under favourable conditions.\\\\ ",
        "introduction": "The hypothesis of the ubiquitous presence of free gas\\textendash phase polycyclic aromatic hydrocarbons (PAHs) in the interstellar medium (ISM) originated about 20 years ago \\citep{leg84,all85}. Due to their spectral properties and their high photostability, these molecules were suggested as the most natural interpretation for the so\\textendash called ``Aromatic Infrared Bands'' (AIBs), a set of emission bands observed near 3.3, 6.2, 7.7, 8.6, 11.3 and 12.7~$\\mu$m, in many dusty environments excited by UV photons \\citep{leg89,all89}. The AIBs are the spectral fingerprint of the excitation of vibrations in aromatic C\\textendash C and C\\textendash H bonds \\citep{dul81}. Furthermore, PAHs and their cations were also supposed to account for a subset of ``Diffuse Interstellar Bands'' (DIBs) \\citep{leg85,van85,cra85}, more than $300$ absorption features observed in the near-UV, visible and near-IR in the spectra of reddened stars \\citep{her95,ehr00,dra03} in the Milky Way as well as in external galaxies \\citep{ehr02,sol05}.  PAHs are believed to play an important role in the physics and chemistry of the ISM, showing intermediate properties between gas and dust phases, i.e. behaving at the same time both as very small grains and large molecules. This has motivated many observational, experimental and theoretical efforts in the last two decades, which confirmed PAH\\textendash related species to be promising candidates to explain AIBs and some of the DIBs, accounting for a substantial fraction of the total interstellar carbon budget \\citep{sal99b}. However, despite the large number of studies and the tentative identification of neutral anthracene and pyrene in the Red Rectangle nebula \\citep{vij04,vij05,mul06}, no \\emph{definitive} spectral identification of any \\emph{specific} individual member in this class exists to date.  Concerning optical absorption spectra, in fact, while the large number of low\\textendash temperature matrix isolation spectroscopy studies of PAHs are of fundamental importance to assess the link between PAHs and DIBs \\citep{job95,sal96,sal99a,sal99b,rui02,rui05}, at the same time they do not permit an unambiguous identification of any single PAH, due to the unpredictable matrix\\textendash induced broadening and shift of the absorption bands. Measurements of the electronic absorption spectrum of cold PAH ions in gas\\textendash phase \\citep{rom99,pin99,bre99,bie03,bie04,suk04,tan05b,tan05a} are ``the right way'' to identify specific PAHs based on their optical absorption spectrum, but each of them still represents a demanding experimental task, which can hardly be generalised in a systematic way to a representative sample of this vast class of molecules.  As to IR spectroscopy, ``classical'' AIBs do not permit an unambiguous identification of any single species, because vibrational transitions in the near and medium IR, are a common feature of the whole class of PAHs \\citep{lan96,sal99b}. These vibrational modes just probe specific chemical bonds and not the overall structure of the whole molecule. Indeed, the emission in these bands is usually explained assuming a whole population of different PAHs to contribute to them \\citep{shu93,coo98,all99,bak01b,bak01a,pec02}.  On the other hand, every single such molecule ought to show a unique spectral fingerprint in the far\\textendash IR spectral region, which contains the low\\textendash frequency vibrational modes associated with collective oscillations of the whole skeletal structure \\citep{zha96,job02,mul03}. Moreover, far\\textendash IR bands are expected to be emitted preferentially when the excitation energy of the molecule is relatively low \\citep{job02}; this results in slower internal vibrational redistribution of energy and thus smaller lifetime broadening, implying that their rotational envelopes might be still discernible, as hinted by the gas phase laboratory measurements of \\citet{zha96}. It might then be possible to resolve the rotational structure with a high resolution spectrometer, such as the HIFI heterodyne spectrometer on board of the forthcoming Herschel Space Observatory (HSO) mission\\footnote{\\texttt{www.sron.nl/divisions/lea/hifi/directory.html}}, providing one more crucial identification element for interstellar PAHs \\citep{job02}.  The far\\textendash IR spectrum emitted by a given interstellar PAH will depend essentially just on its molecular properties and on the exciting radiation field, and can be modelled in detail if all these ``ingredients'' are known \\citep{job02,mal03c,mul03,mul06}. Here we make use of a database of theoretical spectral properties of a sample of PAHs \\citep{mal06} to derive their expected far\\textendash IR emission in a grid of radiation fields (RFs) covering some of the environments in which the AIBs, commonly attributed to PAHs, have been observed.  In this paper we describe the modelling procedure, we detail the approximations used and assess their impact on the resulting calculated spectra, both in terms of absolute fluxes and band ratios. Then we present the calculated spectra for a large sample of PAHs and their cations. Our approach can be used also to model molecular rotation, and thus the expected rotational envelopes of far\\textendash IR emission bands \\citep{job02} This latter calculation is very demanding from a computational point of view % and is the subject of more forthcoming work \\citep{mul06}.  In Sect.~\\ref{modelling} we summarise the modelling procedure we used. We then proceed in Sect.~\\ref{validation} to validate our modelling approach with respect to the use of theoretical UV\\textendash visible photo\\textendash absorption spectra (Sect.~\\ref{validation1}), the impact of some necessary simplifying assumptions for the molecular relaxation processes following electronic excitation (Sect.~\\ref{relaxation}) and its comparison with an independently developed model (Sect.~\\ref{dominique}). Results are presented in Sect.~\\ref{results}, while Sect.~\\ref{discussion} discusses the diagnostic potential and limitations of the calculated spectra. ",
        "conclusions": "\\label{discussion}  The results in Appendix~\\ref{tableappendix} show the diagnostic potential of far\\textendash IR observations for the identification of \\emph{specific} PAHs. All the molecules in our sample, in the RFs considered, exhibit a few relatively strong IR\\textendash active bands which dominate its far\\textendash IR spectrum. These bands, in most cases flopping modes of the whole molecule, strongly depend on its size and shape and on the rigidity of its bonds. This latter dependence usually makes the far\\textendash IR spectrum of each cation distinguishable from that of its parent neutral, even if similar to it.  For each species the IR emission spectrum expected for the RF of the Red Rectangle halo is considerably ``colder'' than the one expected for the ISM, which in turn is ``colder'' than that in the PDR of IRAS~21282+5050, i.e. a larger and larger flux fraction is emitted in higher energy bands going from the first to the last. This is just the effect of the spectral energy distribution of the exciting RF, which in the Monte\\textendash Carlo simulations translates to a higher average excitation energy $\\mathcal{E}$ of the molecules (see Table~\\ref{all_rates}).  Another similar trend which can be seen is that the IR emission spectrum of larger molecules is ``colder'', in the same sense as above, than that of smaller ones in the same RF. Again, this is expected due to the larger number of vibrational degrees of freedom for increasing number of atoms.  An interesting result of our simulations is the prediction of non\\textendash negligible emission in far\\textendash IR electric\\textendash dipole forbidden vibrational transitions. Such emission is essentially the excitation energy remaining ``trapped'' in IR\\textendash inactive modes at the end of the vibrational cascade following UV\\textendash visible absorption. In perfect isolation conditions, this energy has no other way out but to be slowly emitted via forbidden or very weakly permitted vibrational transitions. The absorption of an additional UV\\textendash visible photon by the molecule before it can emit the energy trapped in its IR\\textendash inactive modes effectively quenches the emission of far\\textendash IR photons. Hence emission in an IR\\textendash inactive mode only occurs if its time scale is shorter than or at least comparable to the rate of UV\\textendash visible photon absorption. Such conditions are by and large met in the ISRF, in which the calculated fluxes in such bands can be comparable to the lowest energy permitted ones. In the Red Rectangle halo, on the other hand, the radiation density is high enough to make the absorption of another UV\\textendash visible photon happen on a timescale comparable to that of the emission of photons in IR\\textendash inactive modes, which is reduced. In the PDR of IRAS~21282+5050 the absorption of UV\\textendash visible photons completely dominates over the emission in all weak IR bands, essentially suppressing them. An increase in molecular size produces the same effect, described above, of an increase in RF intensity: this happens because the photo\\textendash absorption cross\\textendash sections and, consequently, the photon absorption rates $\\mathcal{R}$, roughly scale with the number of carbon atoms in the molecule \\citep{job92a,job92b}, while the total IR emissivity at the low\\textendash energy end does not scale in the same way for this sample of molecules.  Quite generally, emission in far\\textendash IR bands can be divided in three different regimes, depending only on the photon absorption rate: \\begin{enumerate} \\item for very dilute RFs, e.~g. in the diffuse ISM, a molecule always has time to completely relax after each photon absorption, and consequently the fluxes in \\emph{all} far\\textendash IR bands scale linearly with the photon absorption rate; \\item for intermediate RFs, e.~g. in the Red Rectangle halo, a molecule always has time to emit all quanta in IR\\textendash active modes, but absorbs another photon before being able to completely relax IR\\textendash inactive modes; hence, fluxes in IR\\textendash active bands still scale linearly with the photon absorption rate, while some emission in electric\\textendash dipole forbidden bands is suppressed; \\item for very strong RFs, e.~g. in the PDR of IRAS~21282+5050, a molecule is always vibrationally ``warm'', since photon absorption events are very frequent, and it almost never reaches the bottom of its IR emission cascade; as a result, \\emph{all} weak IR bands tend to be progressively suppressed in favour of the stronger ones (i.e. classical AIBs). \\end{enumerate} The boundaries between these regimes will depend on the molecule, since different molecules will have different photon absorption rates in the same RF.  Within the same regime, far\\textendash IR fluxes scale linearly with the dilution of the RF adopted, which means that the data we present here cover a relatively wide range of environments. For example, the ISRF estimated by \\citet{dra78} is very similar to the interstellar RF we used, with just a different dilution factor; the former has an intensity, in Habing units, of $G_1\\simeq1.68~G_0$ \\citep{wein01} while the latter corresponds to $G_2\\simeq3.49~G_0$: therefore, the expected far\\textendash IR spectrum for the Draine ISRF can be obtained simply by multiplying the fluxes we calculated by the ratio $G_1/G_2\\simeq0.48$. As another example, the RF in the Orion Bar is very similar to the one of IRAS~21282+5050, with a $\\sim$10\\textendash fold relative dilution, hence the expected far\\textendash IR spectrum for the Orion Bar can be obtained, as a first order approximation, dividing the fluxes we calculated for IRAS~21282+5050 by a factor of 10.  PAH emission in IR\\textendash inactive modes, if detectable, will provide a powerful probe both of the absolute intensity of the RF the emitting molecule is embedded in and of some poorly known molecular parameters \\citep{job02,mal03c,mul03,mul06}. However, as a cautionary remark, we recall that in the absence of any information on the intensity of electric\\textendash dipole forbidden vibrational bands, in the model we made the simplification of considering them as electric\\textendash dipole permitted bands with an oscillator strength $10^4$ times smaller than the average IR\\textendash active bands. Therefore, while both the asymptotic results calculated for the ISRF (perfect isolation) and for the PDR of IRAS~21282+5050 (complete suppression) are expected to be correct, as well as the general trends predicted, accurate quantitative results for intermediate cases will require either the measurement or the calculation of IR\\textendash inactive band intensities. On the plus side, the impact of this uncertainty on the accuracy of fluxes in IR\\textendash active bands is by and large negligible in all cases.   For comparison with observational data, in the optically thin case, the spectra we presented here must be multiplied by the assumed column density and the result integrated over the instrumental aperture on the sky for extended sources. A detailed radiative transfer analysis is needed for the optically thick case. We previously performed such a comparison for three small, neutral PAHs in the Red Rectangle \\citep{mul06}, and in this case found the expected band intensities to be well within the sensitivity of available Infrared Space Observatory (ISO) data. This example clearly demonstrated the discriminating capabilities of our synthetic spectra with presently available ISO observations and even more so for the forthcoming HSO mission, which is expected to go about two orders of magnitude deeper than ISO in the far\\textendash IR.  \\subsection{Detectability}\\label{detectability}  To assess the possibility to identify specific interstellar PAHs in a large population of species, we computed the spectrum of a weighted sum of all the species in our sample, using the inverse of the number of carbon atoms in each molecule as its weight, i.e., \\begin{equation} \\label{sum} \\mathrm{I}_\\lambda^\\mathrm{tot} = k \\sum_\\mathrm{PAH} \\frac{1}{\\mathrm{N}_\\mathrm{C}(\\mathrm{PAH})}\\mathrm{I}_\\lambda^\\mathrm{PAH}, \\end{equation} the summation being over all molecules and charge states in our sample, and the single spectra calculated for the Red Rectangle halo; k is a normalisation constant which is defined later.  To assess the problem of detectability, some assumption on band shapes had to be made. One major cause of band broadening is the overlap of many nearly degenerate vibrational modes due to different occurences of the same  chemical bond in the molecule, such as e.g. out\\textendash of\\textendash plane C\\textendash H bends. This is automatically taken into account in our model, which calculates separate expected fluxes for all bands, regardless of their near degeneracy.  Another cause of broadening is anharmonicity: almost all emitted photons are not due to a transition from a state in which one single vibrational mode was singly excited to the ground state, but from highly excited states. As an example, Figs.~\\ref{histo_bisanthene_n} and \\ref{histo_bisanthene_c} show the distribution of the number of photons emitted in three vibrational modes as a function of the excitation energy of neutral and cationic bisanthene (C$_{28}$H$_{14}$), a medium\\textendash sized molecule of our sample, in the three modelled RFs. \\begin{figure}[b!] \\includegraphics[width=\\hsize]{enerhists_bisanthene_n.eps} \\caption{Distribution of the excitation energies at which photons are emitted by neutral bisanthene in three RFs and three vibrational bands.} \\label{histo_bisanthene_n} \\vspace{-0.4cm} \\end{figure}  \\begin{figure}[b!] \\includegraphics[width=\\hsize]{enerhists_bisanthene_c.eps} \\caption{As Fig.~\\ref{histo_bisanthene_n} for the bisanthene cation.} \\label{histo_bisanthene_c} \\vspace{-0.4cm} \\end{figure} Figures~\\ref{quantahisto_bisanthene_n} and \\ref{quantahisto_bisanthene_c} show the distribution of the number of photons emitted in the same three specific vibrational modes as a function of the excitation energy in that specific vibrational mode. As described in detail by \\citet{oom03}, the peak position of the bands is a function of vibrational energy (i.~e. cross\\textendash anharmonic shifts) and quantum number (i.~e. anharmonic shifts) of the emitting vibrational mode. Knowledge of these molecular parameters, together with the distributions in energy and quantum number (i.~e. Figs.~\\ref{histo_bisanthene_n} to \\ref{quantahisto_bisanthene_c}) would enable one to derive the effect of anharmonicity on band shape for each band.  Furthermore, intramolecular energy redistribution is extremely effective at high excitation energies and causes strong lifetime broadening. The rate of radiationless transitions causing intramolecular energy redistribution is known to be a steep function of molecular vibrational excitation energy. Thus again, in principle, if this function were known this broadening might be estimated from the distribution of energies of the emitting molecules. In the case of the lowest frequency modes, lifetime broadening becomes negligible, since they are emitted preferentially below or only slightly above the decoupling energy $E_\\mathrm{dec}$.  We performed a completely detailed calculation of anharmonic and lifetime broadening effects on band shape for the C\\textendash H stretching modes near 3.3~$\\mu$m for neutral anthracene, phenanthrene and pyrene in the Red Rectangle \\citep{mul06}, using experimental data \\citep{job95b,pec02}. Since the molecular parameters required for this are only available for few bands of few species, we did not include this effect in our present calculations.  \\begin{figure}[t!] \\includegraphics[width=\\hsize]{quantahists_bisanthene_n.eps} \\caption{Distribution of the photons emitted by neutral bisanthene in three RFs and three vibrational bands as a function of the vibrational excitation quanta in the mode.} \\label{quantahisto_bisanthene_n}  \\end{figure} \\begin{figure}[t!] \\includegraphics[width=\\hsize]{quantahists_bisanthene_c.eps} \\caption{As Fig.~\\ref{quantahisto_bisanthene_n} for the bisanthene cation.} \\label{quantahisto_bisanthene_c} \\end{figure}  Finally, far\\textendash IR bands are expected to show rotational profiles. The population of rotational levels, for isolated PAHs in space, is probably far from thermal equilibrium \\citep{rou92,lec98,lec99}. While our modelling approach can be extended to estimate this precisely, this implies a huge increase in computational costs, with a factor of the order of the number of rotational levels to be traced (i.~e. $\\sim10^5$). It also requires the detailed knowledge of more molecular parameters \\citep[e.~g. the matrix of vibration\\textendash rotation constants and anharmonic vibrational parameters,][]{bar05}, which are not readily available for our sample and far from trivial to calculate. We performed such a complete calculation for two of the smallest neutral PAHs in our sample, namely naphthalene and anthracene, as a test case \\citep{mul06b}. For the present paper, we simulated rotational profiles by including, for each band, two relatively broad gaussians, displaced by their width $\\sigma_\\mathrm{PR}$ on either side of the calculated band origin, representing P and R branches and, for perpendicular bands, a narrower gaussian of width $\\sigma_\\mathrm{Q}$ representing the central Q branch. A fraction of $\\sim20\\%$ of the flux in the band was assigned to the Q branch, when present, the remaining split evenly between the P and R branches. Real values will deviate possibly by a factor of $\\sim2$ from this. The population of levels between different values of the quantum number $J$ of the total angular momentum is qualitatively expected to follow a pseudo\\textendash thermal distribution \\citep{rou92}, its pseudo\\textendash temperature scaling with the average energy of the emitted IR photons. The parameter $\\sigma_\\mathrm{PR}$ in our approximate representation of the P and R branches scales with the square root of this pseudo\\textendash temperature, since they involve transitions between states of different $J$ \\citep[see e.~g.][]{cam04}.  A c\\textendash type\\footnote{IR\\textendash permitted bands are labelled according to the direction of the change of the electric dipole moment in the involved vibrational mode, in relation to the axes of the principal momenta of inertia (I$_i$, with $i=x,y,z$) of the molecule. These axes, and the corresponding rotational constants and transition types, are labelled in order of increasing I$_i$. A c\\textendash type band therefore corresponds to a transition in which the electric dipole moment changes along the direction of maximum momentum of inertia, which for planar molecules is always perpendicular to the symmetry plane \\citep{her91b}.} band of a planar molecule such as a PAH will generally exhibit a sharp central Q branch, since in the rigid rotor approximation all the transitions composing it are coincident in energy if the rotational constants do not change between the two vibrational states involved \\citep{her91b}. For this class of molecules, the fractional variation in rotational constants due to vibrational excitation is usually of the order of $\\sim10^{-4}$, arriving at most to $\\sim10^{-2}$ in the infrequent case of accidental near degeneracy, giving rise to very strong Coriolis perturbations \\citep{her91b,mul06b}. The transitions in the Q branch, by definition, connect states with the same $J$. The relative population of levels with the same $J$, due to Internal Vibration\\textendash Rotation Energy Transfer \\citep[IVRET][]{rou92} remains in thermal equilibrium with the bath of vibrational states down to the decoupling energy. This means that the relative population of levels with the same $J$ will reflect the excitation of the molecule when each photon gets emitted. Since low\\textendash energy photons are emitted near the end of the relaxation cascades, this latter distribution will be very ``cold'' for them. This ``cold'' relative population of levels and the almost unchanged rotational constants in the transition conspire to produce Q branches very much sharper than the corresponding P and R branches.  In the absence of reliable estimates in this respect, we had to assume a starting point from which to apply the above scaling relations. We took $0.7~\\mathrm{cm^{-1}} \\leq \\sigma_\\mathrm{PR} \\leq 5.5 ~\\mathrm{cm^{-1}}$ for neutral naphthalene, which amounts to pseudo\\textendash temperatures in $J$, calculated according to the recipe in \\citet{cam04}, between $\\sim$10~K (highly subthermal) and $\\sim$500~K (highly suprathermal). Any realistic value for the environments considered ought to fall in such a wide range.  For $\\sigma_\\mathrm{Q}$, the profile of Q branches will be blurred by anharmonic effects, which are expected to perturb the energy of the transition by a fractional amount of less than 1\\%, essentially of the same order as the empirical correction factors of (i~.e. $\\sim$0.96 in this case) commonly used to account for anharmonicity, bringing harmonic frequencies into near coincidence with experimental values.  \\begin{figure}  \\includegraphics[width=\\hsize]{rraibs.eps} \\caption{Continuum-subtracted AIBs spectrum of the Red Rectangle, from the online ISO spectral database.} \\label{rraibs}  \\end{figure}   \\begin{figure*}  \\includegraphics[width=\\hsize]{PlotweightedPAH_cont.eps} \\caption{Weighted sum of the spectra of all PAHs in our sample, calculated for the Red Rectangle halo (see text for its exact definition). The dotted line shows, for comparison, the estimated dust continuum in the same source. Different panels correspond to different assumptions of rotational excitations and anharmonic widths of the bands. The axes are in logarithmic scale, wavelength in abscissa and flux in ordinate.} \\label{allsums}  \\end{figure*}  \\begin{figure}  \\includegraphics[width=\\hsize]{rrcont.eps} \\caption{ISO SWS and LWS spectrum of the Red Rectangle, showing the extrapolated dust continuum to long wavelengths (dotted line).} \\label{rrcont}  \\end{figure}  For a direct comparison with present and forthcoming observations, we chose the normalisation factor $k$ in the definition of $\\mathrm{I}_\\lambda^\\mathrm{tot}$ (see Eq.~\\ref{sum}) by requiring that its integral between 3 and 14~$\\mu$m be equal to the integral, over the same wavelength interval, of the observed ISO\\textendash SWS spectrum of the AIBs in the Red Rectangle. The latter spectrum, taken from the online ISO database, is reproduced after continuum subtraction in Fig~\\ref{rraibs} for reference. Our synthetic summed spectrum is shown in Fig.~\\ref{allsums}. The same figure also shows, overplotted in dotted line for comparison, the estimated dust continuum in the Red Rectangle, obtained by extrapolating the long wavelength tail of the observed continuum in available ISO\\textendash LWS observations of the same source with a grey body spectrum in the Rayleigh\\textendash Jeans approximation (see Fig.~\\ref{rrcont}), given (in W~cm$^{-2}$~$\\mu$m$^{-1}$) by $\\mathrm{I}_\\lambda^\\mathrm{dust}~=~A\\,\\lambda^{-\\beta}$, with $A=8.14~10^{-12}$, $\\beta=3.31$ and $\\lambda$ in $\\mu$m. This extrapolation is consistent with the spectral energy distribution previously obtained by detailed modelling in \\citet{men02}.  The molecules in our sample are likely not to be the interstellar population, as the mid\\textendash IR spectra of the mixtures show some difference with the interstellar spectra. Nonetheless, this combined spectrum serves the intended purpose to get a visual estimate of the density of bands, and hence of probable spectral congestion, as a function of wavelength.  In almost all assumed conditions, P and R branches tend to blend in a structureless continuum, from which Q branches of perpendicular bands stand out more or less clearly depending on their assumed width. Indeed, the main obstacle for the detection and identification of these bands is probably not posed by their absolute fluxes, but rather by spectral congestion and spectral contrast against a strong background continuum.  We conclude that a\\textendash type and b\\textendash type bands of PAHs, which display no sharp Q branches, will probably remain unobservable for the foreseeable future. On the other hand, the Q branches of some c\\textendash type bands, for some specific molecules, might turn out to be observable, if molecular parameters conspire to make them narrow enough. More theoretical and experimental work will be needed to find the most promising candidates in the zoo of possible species.  The density of bands decreases with the wavelength. The search for such bands appears therefore favoured at lambda $>$100\\textendash200~$\\mu$m. The decrease of the dust continuum towards the millimeter range will make the detection of the bands easier. Indeed, Fig.~\\ref{allsums} shows that, under the most favourable combination of molecular parameters, most of the bands in our sample at wavelengths beyond $\\sim$120~$\\mu$m are calculated to have peak intensities over 10\\% of the estimated dust continuum, a few of them even exceeding it. Since our sample does not reproduce the actual population of interstallar PAHs, chemical diversity will dilute the peak intensity of individual bands by an unknown factor, especially if this population is evenly spread over a vast number of different species. Chemical diversity is produced by many concurrent causes, e.~g. different carbon skeletons, different hydrogenation and/or charge states, inclusion of heteroatoms, substitution of hydrogen atoms at the edges by different functional groups, isotopic substitutions. \\citet{lep01,lep03} showed that the population of a given PAH in any given environment is expected to be dominated by one or two ionisation and hydrogenation states, depending on its size and physical conditions. \\citet{mul03} showed the effect of the most likely isotopic substitutions on low\\textendash energy bands to be extremely small in the test case of ovalene, which makes it negligible in the present context. As to the other causes of diversity we mentioned above, chemical selection effects will favour the most stable species over the others, so that the most abundant of them \\emph{may} be detectable in the far\\textendash IR. Dedicated laboratory experiments, based on tentative identifications, will be needed to say the final word.  HSO could be the ideal tool for this difficult quest, with its combined sensitivity and spectral resolving power in the far\\textendash IR. The target of choice for such a search would be a source with the following properties: \\begin{itemize} \\item it should be a known powerful source of the AIBs; \\item it should be an environment with mild excitation conditions, i.e. with a not too strong far\\textendash UV RF (to avoid a ``hot'' dust continuum) but with enough near\\textendash UV photons to effectively excite PAH emission. \\end{itemize} The above description clearly fits carbon\\textendash rich circumstellar environments such as the Red Rectangle as well as some relatively bright photon\\textendash dominated regions (e.g. the Orion Bar). On the other hand, neither the diffuse ISM (too weak AIBs) nor Planetary Nebulae (too strong far\\textendash UV flux) would be among the best targets."
    },
    "0606/astro-ph0606502_arXiv.txt": {
        "abstract": "We present Gemini Multiobject Spectrograph integral field spectroscopy of the extended emission-line region associated with quasar 3C\\,249.1. The kinematics of the ionized gas measured from the [O\\,{\\sc iii}]\\ $\\lambda$5007 line is rather complex and cannot be explained globally by a simple dynamical model, but some clouds can be modeled individually as having locally linear velocity gradients. The temperatures of the ionized gas appear uniform (varying from $\\sim$12000 to 15000 K), while the densities vary from a few tens to a few hundreds cm$^{-3}$. The emission mechanism of all of the emission clouds, as indicated by the line-ratio diagnostics, is consistent both with ``shock + precursor\" and pure photoionization models. The total mass of the ionized gas is on the order of $10^9$ M$_\\odot$. We estimate the bulk kinetic energy and momentum of the extended emission-line region of $2.5\\times10^{57}$ ergs and $10^{50}$ dyne s, and a dynamical timescale of $\\sim$10 Myr.  By comparing the injection rates of kinetic energy and momentum of different galactic wind models with the observation, we argue that the emission-line clouds are most likely a direct result from the feedback of the quasar.  We also discuss the nature of the extended X-ray emission surrounding the quasar. ",
        "introduction": "About a third of low-redshift (z $\\leq$ 0.45) steep-spectrum radio-loud quasars are surrounded by luminous emission-line regions that often extend to several tens of kpc from the nucleus \\citep{sto87}. The morphologies of these extended emission-line regions (EELRs) are often complex and clearly display a non-equilibrium situation showing strings of knots and filaments straddling tens of kpc, and they are in general uncorrelated with either the distribution of stars in the host galaxy or the structure of the extended radio source, if present.  The origin of the gas comprising these EELRs remains uncertain. The apparent correlation between the occurrence of EELRs and evidence for strong interaction inspired the idea that the gas is tidally disrupted material from a disk \\citep{sto87}. But this is unlikely because (1) there is generally no correlation between stellar tidal features and the distribution of the gas, and (2) some confinement mechanism is required for the gas to retain the densities implied by the emission-line spectra for more than a very brief time ($\\sim10^4$--$10^6$ years; \\citealt{fab87,sto02}). It has also been suggested that the gas comes from a cooling flow in a hot halo surrounding the quasar \\citep{fab87}. This scenario has been ruled out by deep {\\it Chandra} X-ray observations of four quasars showing strong EELRs, since the hot (10$^8$ K) halo gas from which the warm emission-line gas is suggested to condense is not detected, indicating strong cooling is not taking place \\citep{sto06}. And, in any case, the standard cooling-flow paradigm seems to have been ruled out by recent high-resolution X-ray spectroscopy \\citep{pet03}.  If the gas cannot come from the intergalactic medium (IGM) through a cooling flow, then it must originate from the interstellar medium (ISM) of the quasar host galaxy. Despite the non-detection of hot gaseous halos by the {\\it Chandra} observations, highly structured X-ray emission regions are seen around two of the four quasars surveyed. One of the plausible explanations for the X-ray emission is that it is thermal bremsstrahlung from high-speed shocks. On the other hand, detailed photoionization modeling indicates that a low-density phase ($\\sim$2 \\cc) and a high-density phase ($\\sim$500 \\cc) are needed to explain the optical spectrum of one of the EELR clouds around 4C\\,37.43 (E1; \\citealt{sto02}). The two phases cannot be in pressure equilibrium because their temperatures are almost the same. A reasonable possibility is that the high densities in the gas are continuously regenerated by shocks propagating through the surrounding medium.  The high-speed shocks could be driven by a galactic wind, which itself results from the feedback from the quasar \\citep{diM05} or a vigorous starburst in the host galaxy (see \\citealt{vei05} for a review). This evidence of ongoing galactic winds combined with the morphological properties of EELRs thus suggest a scenario where the EELRs are a direct result of a superwind---the gas is originally the ISM in the host galaxy and has been blown out and shocked by a galactic wind; the gas is then photoionized by the UV continuum from the quasar and/or ionized by the shocks.  3C\\,249.1 ($z = 0.31$) is a powerful quasar, also known as PG\\,1100+772. Around the quasar, distinct X-ray emission regions \\citep{sto06} and one of the most luminous optical EELRs at $z < 0.5$ \\citep{sto87} are seen extending to radial distances of tens of kpc, making it an ideal candidate for an in-depth study.  In this paper, we take advantage of the simultaneous spatial and spectral coverage of integral field spectroscopy (IFS) to further explore the origin of the EELR and the extended X-ray emission around 3C\\,249.1. Throughout we assume a flat cosmological model with $H_0=70$ km s$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}= 0.7$. ",
        "conclusions": "\\subsection{The nature of the X-ray emission} In Fig.~\\ref{fig:xrayvla} we compare the residual {\\it Chandra} X-ray image of 3C\\,249.1 and the 0\\farcs35 resolution radio image. This figure shows convincingly that one of the X-ray emission region, $xb$, has an intimate relationship with the radio jet. The X-ray emission is coincident with a region protruding perpendicular to the jet direction, a feature that has been noted by several authors (\\eg\\ designated X1 in \\citealt{gil04}). Since there is a large offset ($\\sim$1\\farcs4 $\\approx$ 6.4 kpc) between $xb$ and the nearest hotspot (N2, following \\citealt{gil04}), it might not simply be synchrotron radiation from the jet, but possibly a result of a gas cloud being shocked by the jet fluid.  \\begin{figure}[!tb] \\epsscale{1.2} \\plotone{f7.eps} \\caption{PSF-subtracted {\\it Chandra} ACIS image of 3C\\,249.1 overlaid with contours of a VLA image at 5 GHz and 0\\farcs35 resolution \\citep{gil04}. Contours at $0.5\\times(1,2,8)$ mJy beam$^{-1}$. The two X-ray emission regions (xa \\& xb) and a radio hotspot (N2) are labeled. The peak of the radio continuum has been registered to that of the X-ray emission, which is indicated by a cross, before a scaled synthetic PSF was subtracted from the X-ray image (refer to \\citealt{sto06} for more details on the PSF subtraction).} \\label{fig:xrayvla} \\end{figure}   The other X-ray emission region, $xa$, is covered by the IFU FOV.  Gas kinematics, line-ratio diagnostics and its morphology all fail to provide convincing evidence of shocks, so we can not rule out the possibility that the X-ray emission is actually a result of the same recombination process as what causes the optical emission. If the X-ray is indeed from recombination lines, then the other high surface-brightness clouds (EELR-$a,b,c$) should also be bright in X-ray; but it is impossible to detect the emission from the {\\it Chandra} data, since at such small radial distances their X-ray emission, if any, is completely overwhelmed by the X-rays from the quasar.  \\subsection{Constraints on the power source of the outflow\\label{outflow}}  The dynamical timescale of the EELR can be estimated from their radial velocities and distances to the nuclei, $\\tau_{dyn} \\approx 10^7 D_{5kpc} \\bar{v}_{500}^{-1}$ yr. Assuming an average velocity of 500 \\kms\\ is appropriate here, because the initial velocity of a galactic wind is expected to be on the level of 1000 \\kms\\ and it has since been slowed to the observed a few hundreds \\kms\\ due to mass entraining and working against gravity. At the redshift of 3C\\,249.1, the angular scale is 4.6 kpc arcsec$^{-1}$, so 5 kpc is about the projected distances of the two brightest EELR clouds to the nucleus. The time scale is consistent with the age of typical luminous radio lobes, $\\sim10^7-10^8$ yr \\citep{blu00}. The rate of mass outflow can be estimated assuming it is constant over the dynamical age of the ionized gas,  $\\dot{M} \\simeq M / \\tau_{dyn} = 100 M_9 D_{5kpc}^{-1} \\bar{v}_{500}$ M$_{\\odot}$ yr$^{-1}$. Similarly the injection rate of kinetic energy is $\\dot{E}_{KE} \\simeq E_{KE} / \\tau_{dyn} = 8\\times10^{42} M_9 v_{500}^2 D_{5kpc}^{-1} \\bar{v}_{500}$ ergs s$^{-1}$, and the input momentum rate is $\\dot{p}_{kin} \\simeq p_{kin} / \\tau_{dyn} = 3\\times10^{35} M_9 v_{500} D_{5kpc}^{-1} \\bar{v}_{500}$ dyne. We emphasize that the rates can only be considered as lower limits to the actual wind energetics since the majority mass of the wind flow may not be detectable in the optical \\citep{vei05} and a substantial amount of wind-entrained ISM may be shielded from the ionizing flux of the quasar.  If the outflow is driven by a starburst, a star formation rate (SFR) of $\\sim$10 M$_{\\odot}$ yr$^{-1}$ is required to deposit enough kinetic energy into the clouds, or $\\sim$60 M$_{\\odot}$ yr$^{-1}$ to inject enough momentum into the ionized gas (Equations 2--3 in \\citealt{vei05}. Since a significant fraction of the originally injected kinetic energy may have been lost to radiation and counteracting the gravitational potential, the kinetic energy is not conserved. But the wind may still conserve momentum, so the SFR of $\\sim$60 M$_{\\odot}$ yr$^{-1}$ is probably more realistic. The SFR scales to an infra-red (IR) luminosity of $L_{IR}$(8-1000$\\mu$m) $\\approx 5.8\\times10^9 {\\rm SFR} L_{\\odot} = 3.5\\times10^{11} L_{\\odot}$ \\citep{ken98}. Is this luminosity consistent with mid-IR observations? 3C\\,249.1 was detected in 12, 25 and 60 $\\mu$m by IRAS \\citep{san89},  in 70 $\\mu$m by MIPS \\citep{shi05}, and in 100 $\\mu$m by ISO \\citep{haas03}. By using a model of two blackbody components to fit the mid-IR spectral energy distribution (SED), we obtain a total IR luminosity of $L_{IR}$(8-1000 $\\mu$m) $\\simeq 8.6\\times 10^{11}$ L$_{\\odot}$. Thus more than 40\\% of the total IR luminosity must be powered by star formation to explain the input momentum rate. However, the SED peaks near 12 $\\mu$m, so the IR emission is dominated by warm dust (T $\\sim$ $100 - 200$ K) presumably directly heated by the quasar \\citep{san89}. From the SED fitting, less than 23\\% of the total luminosity could have come from a ``starburst\" component with a characteristic dust temperature $T \\lesssim 65$ K. Apparently the IR data argues against an on-going starburst as strong as $\\sim$60 M$_{\\odot}$ yr$^{-1}$, though it doesn't formally rule out the possibility that there was an energetic starburst which had caused the outflow but ceased not too long ago ($< \\tau_{dyn} \\approx$ 10 Myr).  Note also that the mass outflow rate is almost an order of magnitude larger than the predicted mass injection rate from supernovae (SNe) for a SFR$\\sim$60 M$_{\\odot}$ yr$^{-1}$ ($\\dot{M}_{inj} \\simeq 16$ M$_{\\odot}$ yr$^{-1}$), which implies that the wind has entrained considerable amount of gas from the ISM. This is in contrary to the low efficiency for transporting mass out of the galaxy predicted by numerical simulations of starburst-driven galactic winds \\citep{str00}.  Now we consider the possibility that the outflow is directly driven by the quasar of 3C\\,249.1. Radiation from the quasar can couple to the surrounding gas via various processes such as electric scattering, photoionization, etc. The input momentum rate from radiation pressure is $\\dot{p} = 1.3\\times10^{35} (\\eta/0.1) \\dot{M}_{acc}$ dyne s, where $\\eta$ is the radiative efficiency and $\\dot{M}_{acc}$ is the mass accretion rate of the black hole in units of M$_{\\odot}$ yr$^{-1}$. Hence a mass accretion rate of 2.5 M$_{\\odot}$ yr$^{-1}$ is sufficient to inject enough momentum to the clouds in $10^7$ yr. This accretion rate is in agreement with the one implied by its bolometric luminosity, $\\dot{M}_{acc}$ = $8 (\\eta/0.1)^{-1}$ M$_{\\odot}$ yr$^{-1}$ (L$_{bol}$ = $8\\times10^{12}$ L$_{\\odot}$; Sanders \\etal\\ 1989).  In summary, although we cannot totally exclude a starburst origin, the available evidence favors the quasar as the power source of the outflow.  If this is indeed the case, the presence of a luminous EELR may tell us something about the quasar itself.  \\citet{sto87} found that roughly one-third to one-half of all powerful, steep-spectrum radio-loud quasars show luminous EELRs, but there are clearly some for which any extended emission is quite weak.  These latter cases may be those for which the quasar ignition occurred sufficiently long ago ($>$ a few $10^7$ years) that the expelled gas has dissipated to the point that its surface brightness falls below our detection limit."
    },
    "0606/astro-ph0606028_arXiv.txt": {
        "abstract": "I compute a renormalization group (RG) improvement to the standard beyond-linear-order Eulerian perturbation theory (PT) calculation of the power spectrum of large-scale density fluctuations in the Universe.  At $z=0$, for a power spectrum matching current observations, lowest order RGPT appears to be as accurate as one can test using existing numerical simulation-calibrated fitting formulas out to at least $k\\simeq 0.3\\ihmpc$; although inaccuracy is guaranteed at some level by approximations in the calculation (which can be improved in the future). In contrast, standard PT breaks down virtually as soon as beyond-linear corrections become non-negligible, on scales even larger than $k=0.1\\ihmpc$. This extension in range of validity could substantially enhance the usefulness of PT for interpreting baryonic acoustic oscillation surveys aimed at probing dark energy, for example.  I show that the predicted power spectrum converges at high $k$ to a power law with index given by the fixed-point solution of the RG equation.  I discuss many possible future directions for this line of work. The basic calculation of this paper should be easily understandable without any prior knowledge of RG methods, while a rich background of mathematical physics literature exists for the interested reader. ",
        "introduction": "Statistics of the large-scale density fluctuations in the Universe provide an invaluable source of information on the very early Universe, where the initial density perturbations that grow by gravitational instability were created, and on the material content and gravitational laws of the Universe, which determine the growth of structure and how it looks to us from a distance \\citep{2006JCAP...10..014S}. The standard practical toolbox we use for interpreting observations consists of linear theory on very large scales or at very early times \\citep{1996ApJ...469..437S}, and everywhere else numerical simulations which compute the fully non-linear evolution of a realization of the density field in a patch of Universe. Fitting formulas, motivated by physical intuition but ultimately heuristic, are often used to more efficiently interpolate between simulation results \\citep{2003MNRAS.341.1311S,2006ApJ...646..881W}.  Beyond linear order perturbation theory (PT) has been considered for a long time \\citep{1980lssu.book.....P,1981MNRAS.197..931J,1983MNRAS.203..345V, 1984ApJ...279..499F, 1986ApJ...311....6G,1994ApJ...431..495J,1996ApJ...473..620S}, but does not currently play an important role in our interpretation of precision observations (see \\cite{2002PhR...367....1B} for a thorough review of large-scale structure perturbation theory). The desire to measure dark energy properties using observations of baryonic acoustic oscillation features \\citep{1998ApJ...504L..57E, 1998ApJ...496..605E,2001ApJ...557L...7C,2003astro.ph..1623E, 2003ApJ...594..665B,2003PhRvD..68h3504L,2003ApJ...598..720S, 2004ApJ...615..573M,2005ApJ...631....1G,2005astro.ph..7457G, 2005MNRAS.357..429A,2005MNRAS.363.1329B,2006MNRAS.365..255B, 2006MNRAS.366..884D,2006astro.ph..7122M}, which fall more or less exactly in the range of scales where higher order perturbation theory could be most useful, gives this line of work much more pressing practical relevance than it has had in the past \\citep{2006ApJ...651..619J}. More traditional uses of galaxy clustering to measure cosmological parameters, long reliant on the assumption that galaxy power is simply proportional to linear mass power, have also reached the point where weakly non-linear effects are a substantial limitation \\citep{2004ApJ...606..702T,2006MNRAS.366..189S}. Additionally, other probes of large-scale structure, like the \\lyaf\\ \\citep{2005ApJ...635..761M,2006MNRAS.365..231V}, weak lensing \\citep{2006ApJ...647..116H}, galaxy cluster/Sunyaev-Zel'dovich effect (SZ) measurements \\citep{2005astro.ph.11060D}, and possibly future 21cm surveys \\citep{2005MNRAS.364..743N}, could all benefit from improved computational techniques.  The problem with higher order perturbation theory is that it does not work very well in exactly the regime where it could currently be most useful: at low redshift and wavenumber $k\\sim 0.1\\ihmpc$, where galaxy clustering is significantly but still only weakly non-linear \\citep{2004PhRvD..70h3007S}. Recently, \\cite{2006PhRvD..73f3520C,2006PhRvD..73f3519C} proposed a renormalization technique aimed at improving PT calculations.  They re-sum an infinite series of perturbation theory terms (Feynman diagrams) which determine the memory of perturbations to initial conditions as a function of scale (see also \\cite{2001NYASA.927...13S}). This method is apparently computationally difficult because they have not yet computed a power spectrum using it.  The renormalization group (RG) method I present here is different, but surely related. There has been some work toward using RG techniques to model large-scale structure \\citep{1997EL.....38..637B,1999A&A...344...27D, 2001IJMPA..16.2041G}. The methods used are different than mine, and have not yet seen much practical application for comparison with observations.  My implementation of the RG is inspired more by its use for dealing with secular divergences in perturbative solutions of differential equations \\citep{1994PhRvL..73.1311C,1996PhRvE..54..376C} than by the probably better known applications in high energy and statistical physics \\citep{1974PhR....12...75W, 1983RvMP...55..583W,1994RvMP...66..129S,2002PhR...363..223B}, although all of these uses are related. The basic idea is to identify a quantity in the problem that is not directly observable and modify it to remove breakdowns in the perturbation theory, leading to a well-behaved renormalized observable. In this paper I renormalize the initial (linear theory) power spectrum to remove the relative (not actually infinite for realistic power spectra) divergence of the lowest order correction to linear theory. More detailed mathematical clarification and explanation of the meaning of the method of \\cite{1994PhRvL..73.1311C,1996PhRvE..54..376C}, in terms of the theory of envelopes, can be found in \\cite{1995PThPh..94..503K,1997PThPh..97..179K}. Even more detailed treatments can be found in \\citep{2001PhR...352..219S, 2006JPhA...39.8061K}. These papers should be useful to the reader who is unsatisfied with my presentation, or wants to extend it. During the last several years, \\cite{1999PhRvD..60f5003B,1999PhRvD..59j5019B,2000PhRvD..61f5006B, 2000PhRvD..62j5026W,2002PhRvD..65d5014B,2003PhRvD..67f5022B, 2003AnPhy.307..335B} have extended and generalized this dynamical renormalization group method to compute the real time evolution of expectation values of quantum fields.  The rest of the paper is as follows: In \\S \\ref{seccalc} I derive a simple formula for the RG-improved power spectrum of mass density fluctuations.  In \\S \\ref{secresults} I show numerical results and compare them to standard PT and fitting formulas calibrated using simulations. Finally, in \\S \\ref{secdiscuss} I discuss possible future directions for this kind of work. ",
        "conclusions": "}  The central result of this paper is Eq.~(\\ref{eqbeta}), the RG equation for the renormalized power spectrum, demonstrated in Fig. \\ref{z0}(a) to give good results for the power spectrum in the quasi-linear regime at low redshift, where standard PT performs poorly. It is clear from Figs. \\ref{z0}-\\ref{z6} that the RG improvement to standard PT is generally helpful, in both the weakly and strongly non-linear regimes.  Except in the most challenging case of $z=0$ and high $k$, the simulation-calibrated fitting formula predictions of Peacock and Dodds \\cite{1996MNRAS.280L..19P} and HALOFIT \\cite{2003MNRAS.341.1311S} are sufficiently contradictory that it is not completely obvious that they are more reliable than RGPT.  It is possible that this method will turn out to be a computational curiosity, maybe leading to improved a posteriori understanding of what we see in simulations, but unable to achieve sufficient accuracy for practical problems, and thus leading to little fundamental change in how we carry out precision cosmology measurements. (It may appear that the method is guaranteed to be useful for describing weakly non-linear galaxy clustering, but this could still be foiled by bias and redshift-space distortions, which I discuss below.)  On the other hand, I hope that this line of work is only in its infancy, and will lead to a diversification of computational methods, beyond standard N-body simulations. I am distinguishing here between systematic approximation methods and heuristic ones like the halo model, which rely heavily on calibration by simulations. One of the implicit goals of this work has been to avoid any approximations where the path to improved accuracy is unclear. A good quality for an approximation to have is a clear limit in which it can be expected to return to the exact result. Note that standard simulations are themselves exactly the kind of systematic approximation that I would like to see more of.  Simulating a finite volume with a finite number of particles and limited force resolution are all approximations, but each can be tested for convergence as their controlling parameters are changed.  With the hope of stimulating further work, I now discuss a variety of possible directions that could be pursued, although I can't guarantee that all of these are good ideas:  The first priority should be to include the velocity-velocity and velocity-density power spectra as independent functions to be renormalized, as I discussed extensively in \\S \\ref{secexplanation}. The main requirement for doing this properly is a straightforward generalization of the standard PT results to allow for general initial conditions, including decaying modes. The $n\\simeq-1.4$ fixed-point power spectrum will be modified, hopefully to a more interesting, accurate fixed-point shape.  After that, brute force computation to higher order in perturbation theory should produce more accurate results.  The fact that the current calculation is a substantial improvement over linear theory at all $k$ (as opposed to diverging wildly at some point), well beyond the scale where it is no longer very precise, makes me hopeful that each additional order could extend the effective scale significantly.  The possibility of estimating the errors in the calculation by comparing results at different orders is an equally strong motivation for going to higher order.  [In the same sense, the present calculation should show where linear theory breaks down much more accurately than the common practice of guessing based on the value of $\\Delta^2(k)=k^3 P_L(k)/2 \\pi^2$.]  These calculations give clear motivation and targets for high accuracy simulation work.   RGPT, being a first principles method containing the possibility of internal error control, may in places lead the comparison rather than simply being tested itself.  If a higher order of PT is computed, and the results agree between different orders to some level of precision, it is reasonable to expect that the result is correct, independent of simulations.  (Missing higher moments of the particle velocity distribution function may be a loophole in this argument, as I discuss below.) With  the fixed-point concept in mind, it may be interesting to look more carefully at the evolution of the logarithmic slope of the power spectrum with time and scale in simulations.  Abandoning the first principles philosophy, RGPT results may provide a useful template for a new simulation fitting formula.  Because the errors in RGPT are not too large, and should be slowly varying functions of $k$ and cosmological parameters, the new formula could be just a very simple parameterization of corrections to RGPT computed using simulations.  One can obviously apply the techniques of this paper to higher order statistics such as the bispectrum \\citep{2005MNRAS.361..824G,2006A&A...456..421B,2006PhRvD..74b3522S}, trispectrum,  etc. \\citep{2006ApJ...649...48R}.  As discussed in \\S \\ref{secexplanation}, one can consider promoting the leading order for these statistics to the more generally possible $\\mathcal{O}(\\delta_1^3)$ and $\\mathcal{O}(\\delta_1^4)$ (for the bispectrum and trispectrum, respectively). Bispectrum results can be compared to the simulation fitting formula of \\cite{2001MNRAS.325.1312S}. If the trispectrum computation is sufficiently accurate, RGPT should be useful for computing expected statistical errors on the power spectrum, which is notoriously painful to do with simulations \\citep{2006MNRAS.371.1188H}.  Similarly, errors on higher order statistics may be computable.  Mostly for convenience, I have assumed that the only effect of deviations from Einstein-de Sitter background evolution is to change the linear growth factor. This is probably sufficient in the weakly non-linear regime, but if RGPT is to achieve much accuracy at higher $k$, this assumption inevitably must be relaxed.  This is clear, even though the corrections in standard PT are very small, because the dependence of the power spectrum on the equation of state of dark energy, $w$, at fixed observation-time linear theory power \\citep{2006MNRAS.366..547M}, can not be reproduced under the approximation used here.  There is simply no way the dependence quantified in \\cite{2006MNRAS.366..547M} can enter the current calculation.  Fortunately, even small changes in the right hand side of Eq.~(\\ref{eqbeta}) can lead to significant changes in the solution of this non-linear equation. It will also be necessary to move beyond the very simple time dependence in Eq.~(\\ref{Pspeceq}) if one wants to include any non-gravitational effects that break this form (e.g., gas pressure).  To describe direct tracers of gas, like the \\lyaf\\ or SZ, gas dynamics could be included in the calculation.  Even for weak lensing, which traces mass directly, gas physics has been shown to matter at a level problematic for high precision experiments \\citep{2006ApJ...640L.119J}.  Very rudimentary pressure approximations should be relatively straightforward to include \\citep{1998MNRAS.296...44G}.  While it may be that very sophisticated implementations of heating and cooling are hopeless, one should remember that there is some non-trivial inclusion of non-linearity in this calculation, so this is at least worth considering.  To describe clustering of galaxies, or other unavoidably inexact tracers of density, some concept of bias will need to be included in the calculation. I think this may be the area where RGPT is most likely to provide a real fundamental improvement over simulations, or at least complement to them.  With a large but probably achievable amount of computer power, one can simulate the power spectrum of dark matter and its dependence on a relatively small number of cosmological parameters more or less perfectly.  This will not be possible in the foreseeable future for galaxies, so robust, high precision cosmological measurements using them will rely on developing a comprehensive understanding of how relatively microscopic galaxy formation details translate into large-scale clustering patterns. While the separation of scales probably is not clear enough to make the analogy perfect, this is reminiscent of the classic uses of the renormalization group to understand how complex microscopic theories lead to simple macroscopic behavior \\citep{1983RvMP...55..583W}. The simplest way to introduce bias is to follow \\cite{1993ApJ...413..447F,1998MNRAS.301..797H} in writing the galaxy density fluctuation, $\\delta_g$, as a Taylor series in the mass density, $\\delta_g=\\sum_i b_i \\delta^i /i!$, which leads to an expression for the PT power spectrum of galaxies involving a few more terms than the mass power spectrum,  but not fundamentally more complicated.  This path has been followed by \\cite{2006PhRvD..74j3512M}. Much more ambitiously, one could write local formation rate equations for a galaxy density field and apply perturbation theory to them (something like this was discussed using linear theory in \\citep{1998ApJ...500L..79T}). For example, $\\dot{\\rho}_g=\\rho~f(\\rho)$, with $\\dot{\\rho}=-\\dot{\\rho}_g$ to represent depletion of gas, and with both components obeying the usual gravitational evolution equations.  The mean galaxy density as a function of time  would be an output of the calculation as well as fluctuations.  The goal here would not be to write down a realistic local model and compute exact predictions as much as to see what kind of relations between observables can be generically expected. An intriguing possibility is that one could show that, regardless of the small-scale galaxy formation details, bias can only take certain universal forms, described by a small number of free parameters.  Note that the validity of the scale-independent bias assumption has never really been proven even in the perfectly linear regime (see \\cite{1998ApJ...504..607S} for one attempt).  This entire discussion of bias of galaxies applies equally well to many other traces of density, e.g., the \\lyaf\\ \\citep{2003ApJ...585...34M}.  For many applications, redshift space distortions will need to be included. Like bias, these should be relatively straightforward to include in the present RG formulation using the extension of \\cite{1987MNRAS.227....1K}'s linear theory calculation described by \\cite{1998MNRAS.301..797H}. \\cite{2004PhRvD..70h3007S} discusses the imperfections of the \\cite{1987MNRAS.227....1K} approach, and proposes alternative ideas which may be useful.  It may be interesting, as a computational curiosity, to see if RGPT can give sensible results for $n\\geq -1$ power law initial conditions.  In this case the integrals giving the first correction in standard PT truly diverge, in the sense of being infinite rather than just large \\citep{1996ApJ...473..620S}. The integrals can be cut off at wavenumber $k_c$, but the results then depend on $k_c$, which is arbitrary. One may be able to show that, using the RG as described in this paper, the results will converge as $k_c$ is taken to infinity.  The small-$A_\\star$ integration of Eq.~(\\ref{eqbeta}) will need to be done increasingly carefully as $k_c$ is increased, but hopefully convergence to the $n=-1.4$ fixed point will quickly erase all trace of the cutoff. Note that $k_c$ could affect the normalization of the asymptotic power law, which would spoil this idea (e.g., it may only work over some limited range of $n$).  One potential limitation in RGPT as presented here is the absence of higher moments of the velocity distribution function, e.g., the dispersion in the velocities of particles around their mean velocity at a point. I wrote the Vlasov equation (\\ref{eqvlasov}) instead of cutting straight to the hydrodynamic equations to point out this problem, shared by standard PT (subsequently discussed by \\cite{2006astro.ph.10336A}). Dropping higher moments of the velocity distribution, also known as the single-stream approximation, is not really an added approximation in standard Eulerian perturbation theory. This is surely known to experts (e.g., \\cite{1998A&A...335..395B}), but may not be universally recognized. One can easily include these moments as variables in the calculation, for example multiplying the Vlasov equation by $p_i p_j$ and integrating over $p$ gives (after some substitution to remove bulk velocity terms) \\begin{equation} \\frac{\\partial \\sigma^{i j}}{\\partial \\tau}+2~ {\\mathcal H}~\\sigma^{i j}+ v^k~ \\partial_k \\sigma^{i j}+\\sigma^{j k}~ \\partial_k v^i+ \\sigma^{i k}~ \\partial_k v^j+ \\frac{\\partial_k\\left[\\left(1+\\delta\\right) q^{i j k}\\right]}{1+\\delta}=0 ~, \\end{equation} where $\\partial_k=\\partial/\\partial x^k$, $\\sigma^{i j}(\\vx,\\tau)\\equiv\\left<\\delta v^i~\\delta v^j\\right>$, $q^{i j k}(\\vx,\\tau)\\equiv\\left<\\delta v^i~\\delta v^j~\\delta v^k\\right>$, with $\\delta v^i$ the deviation of a particle's velocity from the local mean velocity (note that a $\\sigma^{i j}$ term would appear in Eq. \\ref{eqeuler} if I had not dropped it). Unfortunately, the lowest order evolution equations contain only the Hubble drag term, so any initial velocity dispersion will self-consistently disappear from the perturbation theory. Furthermore, higher order equations contain no source terms, meaning that if the velocity dispersion is initially zero it can never become non-zero. In fact, \\cite{2001A&A...379....8V} showed that even perturbation theory using the distribution function directly leads to the same result. This is a vexing problem because stream crossing is clearly ubiquitous in the real Universe; however, \\cite{2001NYASA.927...13S} estimated that it only becomes significant at $k\\gtrsim 1 \\ihmpc$, so it may not be fatal for intermediate scales. It seems likely that, if these variables really matter, there should be some way to recover them using a RG method.  The basic idea is that, no matter how small the effect is in the bare perturbation theory, if it is dynamically relevant it will grow through a RG equation like Eq.~(\\ref{eqbeta}). Note that even WIMP cold dark matter has a very small seed velocity dispersion \\citep{2004MNRAS.353L..23G}. The solution is probably to simply retain the dispersion variable in the calculation, even though it is decaying, in the same way I have argued above that one should retain the decaying combination of $\\delta$ and $\\theta$ -- because it can be rapidly fed by higher order terms. Vorticity, $\\vnabla \\times \\vv$, falls out of the standard PT calculation for the same reason velocity dispersion does \\citep{2002PhR...367....1B}. The standard calculation makes use only of the velocity divergence, $\\vnabla \\cdot \\vv$.  A method that successfully reintroduces velocity dispersion may also work for vorticity.  I started this line of work with the intention of using a significantly different type of renormalization group method, based on integrating out (averaging over) small-scale (high-$k$) modes while modifying the equations describing evolution of the large-scale modes in a way that preserves their statistics \\citep{1994RvMP...66..129S}. This method usually starts with a path integral formulation of the problem. \\cite{2001A&A...379....8V} presented a path integral approach to large-scale structure using the full distribution function, and it should not be too hard to write something similar for the usual moments of the distribution function. Velocity dispersion should arise inevitably in this approach, because smoothing a dispersionless field obviously produces dispersion.  It is less inevitable that this dispersion will actually fundamentally change the outcome of the calculation, because the effect of these velocities is already present in the standard PT calculation. Among other possibly interesting features, the noise discussed in previous RG approaches to large scale structure \\citep{1997EL.....38..637B,1999A&A...344...27D} may be generated naturally using this approach (i.e., if high-$k$ modes are eliminated, the time evolution of the remaining modes will no longer be perfectly deterministic).  Further work could elucidate the connection between this work and the re-summation method of \\cite{2006PhRvD..73f3520C,2006PhRvD..73f3519C}. Related renormalization group procedures, starting from general relativistic perturbation theory, are used by \\cite{2000PhRvD..62j4010N} and \\cite{2005astro.ph..6534K} to compute the back-reaction of inhomogeneities on the global evolution of the Universe.  Finally, the success of the RGPT calculation in this paper, which leads to what looks like a time evolution equation for the power spectrum, suggests that it might be useful to consider the exact time evolution equation for the power spectrum, $\\dot{P}(k)\\propto 2~{\\rm Re}\\left<\\dot{\\delta}_\\vk~\\delta_{-\\vk}\\right>$, where $\\dot{\\delta}_\\vk$ is obtained from the usual evolution equations. The usual argument against this would be that $\\left<\\dot{\\delta}_\\vk~\\delta_{-\\vk}\\right>$ involves higher order terms like the bispectrum, which then require their own time evolution equations, and so on, leading to an infinite hierarchy of equations. Firstly, this could be a useful way to look at perturbation theory, especially if we are going to be renormalizing statistics like the power spectrum instead of the field itself.  Second, one may be able to make progress by simply solving these equations numerically, in the same sense that simulations are usually used. One could, for example, truncate the hierarchy by replacing the connected spectra at some level by the appropriate functions of lower order spectra given by perturbation theory, and look for convergence in the results as higher order terms are added. The advantage of working with statistics of the density field rather than the density field itself is enormous.  One obtains the desired results directly rather than averaging over fluctuations in simulations, where there is usually an unavoidable conflict between the need for large boxes to limit finite volume effects like sample variance, and the need for high force and mass resolution to properly compute small-scale structure. The fact that we are working with continuous, relatively slowing varying functions, rather than wildly fluctuating fields, allows many fewer points to be computed, e.g., we can compute a power spectrum over many decades of dynamical range by interpolating between a few tens of logarithmically spaced computation points when millions or billions of particles would be needed to cover the same range with a simulation. Tasks that at first glance seem arduous, e.g., parameterizing the trispectrum or even higher spectra and setting up their evolution equations, should be considered in contrast to the decades of person-power expended on simulations."
    },
    "0606/astro-ph0606734_arXiv.txt": {
        "abstract": "We report the results of CCD $V$ and $R$ photometry of the RR Lyrae stars in M2. The periodicities of most variables are revised and new ephemerides are calculated. Light curve decomposition of the RR Lyrae stars was carried out and the corresponding mean physical parameters [Fe/H]= $-1.47$, $T_{\\rm eff}$= 6276 K, log $L= 1.63 ~L_\\odot$ and $M_V = 0.71$ from nine RRab and [Fe/H]= $-1.61$, $M= 0.54 ~M_{\\odot}$, $T_{\\rm eff}$ = 7215 K, log $L = 1.74 ~L_{\\odot}$ and $M_V= 0.71$ from two RRc stars were calculated. A comparison of the radii obtained from the above luminosity and temperature with predicted radii from nonlinear convective models is discussed. The estimated mean distance to the cluster is $10.49 \\pm 0.15$ kpc. These results place M2 correctly in the general globular cluster sequences Oosterhoff type, mass, luminosity and temperature, all as a function of the metallicity. Mean relationships for $M$, log $L/L_{\\odot}$, $T_{\\rm eff}$ and $M_V$ as a function of [Fe/H] for a family of globular clusters are offered. These trends are consistent with evolutionary and structural notions on the horizontal branch. Eight new variables are reported. ",
        "introduction": "RR Lyrae stars are of particular importance in the age and distance determination to ancient stellar systems, such as globular clusters. Their usefulness is based in the fact that their light curves are easily distinguishable and  their high intrinsic brightness allows their detection in systems of the Local Group. Their well defined absolute magnitude allows using them as standard candles in the cosmic distance scale (e.g. Alcock et al. 2004 for the LMC).  The RR Lyrae stars have been used in the determination of absolute ages of globular clusters, by measuring the magnitude difference between the main sequence turn off point and the horizontal branch (HB), on which the RR Lyrae stars reside. This is a useful reddening free parameter that helps in tracing the early stages of the formation of our Galaxy. This, however, requires a proper calibration of the absolute magnitude $M_V$ for the RR Lyraes.  The pulsational behaviour of RR Lyrae stars in globular clusters has been a subject of study for decades since it offers insight on stellar evolution in the HB stages. The mean periods of fundamental pulsators (type RRab) divide globular clusters into two groups, the Oosterhoff type I (OoI) with $<P_{ab}>$ smaller than 0.6 days and the Oosterhoff type II (OoII) with $<P_{ab}>$ larger than 0.6 days (Oosterhoff 1939).  Recent light curve Fourier decomposition calculations have clearly shown that RR Lyrae stars in OoII clusters are, on average, more luminous, more massive and cooler than in OoI clusters, and that these quantities are closely correlated with the mean cluster metallicity. For example these stars have longer periods in metal poor clusters than in metal rich clusters (e.g. Simon \\& Clement 1993; Clement \\& Shelton 1997; Kaluzny et al. 2000; Arellano Ferro et al. 2004; Arellano Ferro et al. 2006 (see their Tables 8 and 9)).  On the other hand the period distribution of fundamental periods ($P_{ab}$ for RRab stars), and the fundamental periods calculated from the first overtone period and the period ratio ($P_f$ for RRc stars), as well as the relative number of RRc and RRab stars, provide important insights relevant to the instability strip structure and HB structure and evolution (Castellani et al. 2003). Nevertheless, such studies are strongly limited by bona fide completeness of the sample of RR Lyrae pulsators in a cluster. It is evident from recent works that a substantial number of RR Lyrae stars can be discovered when new image-subtraction techniques (Alard 2000; Alard \\& Lupton 1998; Bond et al. 2001; Bramich et al. 2005) are applied to CCD images of globular clusters (e.g. Lee \\& Carney 1999a; Kaluzny et al. 2001; Clementini et al. 2004).  The latest CCD study of M2 (NGC 7089) was published by Lee \\& Carney (1999a) (hereinafter LC99). These authors reported 13 new RR Lyraes, for a total of 34 variables in the cluster, 30 of which are RR Lyrae stars. They calculated the ratio n(c)/n(ab+c) = 0.40, which approaches the mean typical value of 0.44 in OoII clusters. Hence, not many undiscovered RR Lyraes are expected. However, since these authors did not perform a light curve decomposition analysis, we decided to supplement their $BV$ data with new $VR$ observations and use the two data sets to refine the ephemerides, determine the mean physical parameters of the RR Lyrae stars from the Fourier decomposition technique, compare the results with those of other similarly studied OoI and OoII clusters and, in passing, search for new variables. The effort was not fruitless.  In Sect. 2 we describe the observations and data reductions. In Sect. 3 we calculate new ephemerides, new times of maximum, discuss individual objects and report new variables. In Sect. 4 we calculate the physical parameters using the light curve decomposition method. In Sect. 5 we discuss the results in the wider context of other globular clusters and in Sect. 6 we state our conclusions. ",
        "conclusions": "Physical parameters of astrophysical relevance have been derived for the RR Lyrae stars in M2 using the Fourier decomposition of their light curves. The estimates of the stellar radii for RRab stars from the above approach are about 5\\% smaller than radii derived from Period-Radius-Metallicity relations obtained from nonlinear convective models, i.e. a completely different set of pulsational models and approach. The agreement for the RRc is generally good but the dispersion is markedly larger.  Previously recognized trends between the cluster mean values of $log~(M/M_{\\odot})$, $log~(L/L_{\\odot})$, $log T_{\\rm eff}$ and $M_V$ against [Fe/H], obtained from RRc stars, have been established for a sample of 7 OoII and 5 OoI clusters. We extended the study of the trends for RRab stars in 4 OoII and 5 OoI clusters.   Our analysis, based on a larger data set than previous studies, allows one to quantify the relations (see eqs. 14-19 and Figs. 10 and 11). In particular, much has been discussed about the relation between luminosity (or  $M_V$) and metallicity for the RR Lyrae variables. Our eqs. 15, 17, and 19 are in very good agreement with the slopes of $M_V - [Fe/H]$ relations solidly determined and available in the literature. The zero points indicate however that the luminosities for the RRc and RRab stars, as calculated from eq. 3 (Simon \\& Clement 1993) and those derived from eqs. 7 and 9 (Kov\\'acs 1998; Kov\\'acs \\& Jurcsic 1996) respectively, differ by the equivalent to 0.2 mag. This suggests these calibrations need to be revised.  RR Lyrae Fourier light curve decomposition produces a zero point of the RR Lyrae distance scale consistent with those from other methods. When the light curve decomposition technique is used on families of OoI and OoII globular clusters, it provides physical parameters, as a function of the metal content of the cluster, that are consistent with general notions of stellar evolution in the HB and of early stages of galactic formation."
    },
    "0606/astro-ph0606444_arXiv.txt": {
        "abstract": "{} {We provide a library of some 7000 SEDs for the nuclei of starburst and ultra luminous galaxies.  Its purpose is to quickly obtain estimates of the basic parameters, such as luminosity, size and dust or gas mass and to predict the flux at yet unobserved wavelengths.  The procedure is simple and consists of finding an element in the library that matches the observations.  The objects may be in the local universe or at high $z$.} {We calculate the radiative transfer in spherical symmetry for a stellar cluster permeated by an interstellar medium with standard (Milky Way) dust properties. The cluster contains two stellar populations: old bulge stars and OB stars. Because the latter are young, a certain fraction of them will be embedded in compact clouds which constitute hot spots that determine the MIR fluxes.  } {We present SEDs for a broad range of luminosities, sizes and obscurations.  We argue that the assumption of spherical symmetry and the neglect of clumpiness of the medium are not severe shortcomings for computing the dust emission.  The validity of the approach is demonstrated by matching the SED of seven of the best studied galaxies, including M82 and Arp220, by library elements.  In all cases, one finds an element which fits the observed SED very well, and the parameters defining the element are in full accord with what is known about the galaxy from detailed studies.  We also compare our method of computing SEDs with other techniques described in the literature.} {} ",
        "introduction": "By definition, the rapid conversion of a large amount of gas into predominantly massive ($> 8 M_\\odot$) stars, or the result of such a conversion, is called a starburst.  Starburst galaxies constitute a unique class of extragalactic objects.  The phenomenon is of fundamental importance to the state and evolution of the universe, as outlined, for example, in the review by Heckman (1998). According to him, in the local universe starbursts are responsible for about 25\\% of the high--mass star formation rate and for 10\\% of the total luminosity; in the early universe, beyond $z \\sim 0.7$, IR luminous galaxies dominate the star forming activity (Flo'ch et al.~2005).  Starbursts are also cosmologically significant if one interprets the high bolometric luminosities of high redshift galaxies to be due to star formation (Hirashita et al. 2003) at a rate so high that it can only be maintained over a cosmologically short spell ($< 10^8$ yr). The study of nearby starbursts would then help us to understand the processes underlying the star formation history of the universe.   \\noindent Starbursts, we think, are triggered by the gravitational interaction between galaxies (Kennicutt et al. 1987), but they occur, as a result of mass and angular momentum transfer, predominantly in their nuclei, at the center of a massive and dynamically relaxed cluster of old stars (the bulge).  Although the region where OB stars form is relatively small (a few hundred parsec), its luminosity often exceeds that of the host galaxy. Starbursts are almost pure infrared objects, opaque to stellar photons.  Whereas, on average, in the local universe $\\sim$60\\% of the star formation is obscured by dust (Takeuchi et al., 2006), in starbursts the fraction is typically 90\\%.  To interpret infrared observations and to arrive at a self--consistent picture for the spatial distribution of stars and interstellar matter in the starburst nucleus and of the range of dust temperatures, one has to simulate the transfer of continuum radiation in a dusty medium.  Line emission is energetically negligible.  \\noindent A starburst has four basic parameters: total luminosity, $L$, dust or gas mass, $M_{\\rm d}$ or $M_{\\rm gas}$, visual extinction, $A_{\\rm V}$, and size.  Size, $A_{\\rm V}$ and $M_{\\rm d}$ are, of course, related, for a homogeneous density model, only two of them are independent. The luminosity follows observationally in a straight forward way by integrating the spectral energy distribution over frequency, $M_{\\rm d}$ is readily derived from a millimeter continuum data point, if available, and the outcome is almost independent of the internal structure or viewing angle of the starburst.  The size is best obtained from radio observations as it does not suffer extinction by dust.  It does not matter whether the radio emission is thermal or non-thermal, both are connected to high--mass stars.  \\noindent In this paper, we present a set of SEDs for starbursts covering a wide range of parameters.  Anyone with infrared data and interested in their interpretation can compare them with our models, find an SED that matches (after normalization of the distance) and thus constrain the properties of the starburst under investigation without having to perform a radiative transfer computation himself.    \\begin{figure*} [htb] \\includegraphics[width=18cm,height=19.3cm,angle=0]{./Fig/para4.ps} \\caption{Influence of starburst parameters on the SED for a distance of 50 Mpc.  The parameters which are kept constant are listed in square parentheses. \\ {\\bf a)} Total luminosity is varied between $L^{\\rm {tot}} = 10^{10}$ and $10^{12.7}$\\Lsun; [$R=3$kpc, $A_{\\rm V} \\sim 17$mag, $L_{\\rm {OB}} / L^{\\rm {tot}} = 0.6$ and $n^{\\rm hs} = 10^3$ cm$^{-3}$].  {\\bf b)} Here we vary two parameters: the radius of the nucleus from $R=0.35$ over 1 to 3kpc, and the luminosity ratio: $L_{\\rm {OB}} / L^{\\rm {tot}} = 0.4$ (full lines), 0.9 (dotted); [$L^{\\rm {tot}} = 10^{11.1}$\\Lsun, $A_{\\rm V} \\sim 4.5$mag and $n^{\\rm hs} = 10^4$ cm$^{-3}$].  {\\bf c)} Variation of the hot spot density: $n^{\\rm hs} = 10^2$, $10^3$ and $10^4$ cm$^{-3}$; [$L^{\\rm {tot}} = 10^{10.5}$\\Lsun, $R=3$kpc, $A_{\\rm V} \\sim 9$mag, $L_{\\rm {OB}} / L^{\\rm {tot}} = 0.9$].  {\\bf d)} Variation of the dust extinction: $A_{\\rm V} \\sim $2.2, 4.5, 6.7, 9, 18, 35, 70 and 125mag; [$L^{\\rm {tot}} = 10^{10.5}$\\Lsun, $R=3$kpc, $L_{\\rm {OB}} / L^{\\rm {tot}} = 0.9$ and $n^{\\rm hs} = 10^4$cm$^{-3}$]. \\label{para4.ps}} \\end{figure*} ",
        "conclusions": "We have computed in a self--consistent radiative transfer SEDs of spherical, dusty galactic nuclei over a wide range of their basic parameters such as luminosity, dust mass, size and obscuration.  The SEDs can be accessed in a public library \\footnote{ The SED library is available at: \\\\{\\tt http://www.eso.org/\\~ \\/rsiebenm/sb\\_models}}.  \\noindent Given a set of data points for a particular galaxy, there is a simple procedure, described in the README file, to select from the library those elements which best match them.  If the observations cover the full wavelengths band from a few $\\mu$m to about 1 mm, one usually finds only one library element that fits very well, as demonstrated for seven famous active galaxies.  If the data points are widely spaced, there may be a few elements of less fitting quality, but similar in their basic parameters.  \\noindent The library therefore allows one to constrain the fundamental properties of any nucleus which is powered by star formation and for which data exist, without any further modeling.  Two observed fluxes in the MIR plus one submm point are usually sufficient for a crude characterization of the nucleus.  If there are only two MIR points, for example, from Spitzer at 8 and 24$\\mu$m, one can still bracket the total luminosity within a factor of $\\sim 2$.    \\noindent In the UV, optical and NIR, it may be necessary to add to the library SED a low luminosity stellar component, at least, this was necessary for M82 and NGC6240. This component may be due to photons that escaped the nucleus without interaction because of clumping or it may be light from the galactic disk."
    },
    "0606/astro-ph0606391_arXiv.txt": {
        "abstract": "Using analytic arguments and a suite of very high resolution ($\\sim 10^3$\\,M$_\\odot$ per particle) cosmological hydro-dynamical simulations, we argue that high redshift, $z \\sim 10$, $M \\sim 10^8$\\,M$_\\odot$ halos, form the smallest `baryonic building block' (BBB) for galaxy formation. These halos are just massive enough to efficiently form stars through atomic line cooling and to hold onto their gas in the presence of supernovae winds and reionisation. These combined effects, in particular that of the supernovae feedback, create a sharp transition: over the mass range $3-10 \\times 10^7$M$_\\odot$, the BBBs drop two orders of magnitude in stellar mass. Below $\\sim 2 \\times 10^7$M$_\\odot$ galaxies will be dark with almost no stars and no gas. Above this scale is the smallest unit of galaxy formation: the BBB.  We show that the BBBs have stellar distributions which are spheroidal, of low rotational velocity, old and metal poor:  they resemble the dwarf spheroidal galaxies (dSphs) of the Local Group (LG). Unlike the LG dSphs, however, they contain significant gas fractions. We connect these high redshift BBBs to the smallest dwarf galaxies observed at $z=0$ using linear theory. A small fraction ($\\sim 100$) of these gas rich BBBs at high redshift fall in to a galaxy the size of the Milky Way. We suggest that ten percent of these survive to become the observed LG dwarf galaxies at the present epoch. This is consistent with recent numerical estimates. Those in-falling halos on benign orbits which keep them far away from the Milky Way or Andromeda manage to retain their gas and slowly form stars - these become the smallest dwarf irregular galaxies; those on more severe orbits lose their gas faster than they can form stars and become the dwarf spheroidals. The remaining 90\\% of the BBBs will be accreted. We show that this gives a metallicity and total stellar mass consistent with the Milky Way old stellar halo. ",
        "introduction": "\\label{sec:introduction}  The Local Group (LG) of galaxies provide a unique test-bed for galaxy formation theories and cosmology. Their close proximity allows individual stars to be resolved giving accurate kinematics, stellar populations and star formation histories (see e.g. \\bcite{2001ApJ...563L.115K} and \\bcite{2002MNRAS.332...91D}); their spatial distribution can be compared with cosmological predictions to give useful constraints \\citep{2005astro.ph.10370M}; and their large mass to light ratios can be used, through dynamical modelling, to place constraints on the nature of dark matter \\citep{2001ApJ...563L.115K}.  The LG dwarf galaxies are usually split into three types: dSph galaxies, which typically have old stellar populations, are spheroidal in morphology, lie close to their host galaxy\\footnote{We use the terminology `host galaxy' throughout this paper to refer to either the Milky Way (MW) or Andromeda (M31) depending on which of these is closer to the satellite being discussed.} and are devoid of HI gas; dIrr galaxies, which have younger stellar populations, irregular morphology, lie further away from their host galaxy and contain significant HI gas; and the transition galaxies, which are in between the dSph and dIrr types \\citep{1998ARA&A..36..435M}.  In our current `vanilla' cosmological paradigm ($\\Lambda$CDM - cold dark matter with a cosmological constant), all structure forms from the successive mergers of smaller substructures \\citep{1978MNRAS.183..341W}. While this theory has been tremendously successful on scales larger than $\\sim 1$\\,Mpc, on smaller scales it has fared less well (see e.g. \\bcite{2004ApJ...612..628D}). A now long-standing puzzle is the `missing satellites' problem: there appears to be an order of magnitude fewer satellite galaxies in the LG than would naively be predicted from the mass function of dark matter halos (see e.g. \\bcite{1993MNRAS.264..201K} and \\bcite{1999ApJ...522...82K}). A number of solutions to this puzzle have been presented, the two main threads being either to alter the nature of dark matter (\\bcite{2001ApJ...556...93B} and \\bcite{2001ApJ...559..516A}), or to invoke some form of feedback from supernovae explosions (see e.g. \\bcite{1974MNRAS.169..229L} and \\bcite{2000MNRAS.317..697E}), or photo-evaporation (see e.g. \\bcite{1996MNRAS.278L..49Q} and \\bcite{1999ApJ...523...54B}). In the feedback scenario, one might naively expect only the most massive substructure satellites, with the deepest potential wells, to form stars and remain visible in the LG at the present epoch \\citep{2002MNRAS.335L..84S}. However, this presents a problem since the most massive substructure dark matter halos predicted by $\\Lambda$CDM models have central stellar velocity dispersions which are factors of 2-3 larger than those observed in the LG satellites, even after extreme tidal stripping of shocking of these halos (\\bcite{2003ApJ...584..541H}, \\bcite{2004ApJ...608..663K} and \\bcite{2006MNRAS.tmp..153R}).  An alternative view, is that the LG satellites are fossil galaxies left over from reionisation\\footnote{The epoch of reionisation is caused by UV flux emitted by the first forming massive stars. Throughout this paper we suggest that reionisation occured at $z \\sim 10$. Observationally there are two bounds on this epoch. The old WMAP satellite data favour larger redshifts \\citep{2003ApJS..148..175S}; the data from Quasar absorption spectra favour smaller redshifts (\\bcite{2005astro.ph.12082F}; but see also \\bcite{2006NewAR..50...94B}); the new WMAP data favour our chosen redshift \\citep{2006astro.ph..3449S}.} (\\bcite{2000ApJ...539..517B}, \\bcite{2001ApJ...548...33B}, \\bcite{2002MNRAS.333..177B}, \\bcite{2004ApJ...609..482K}, \\bcite{2004ApJ...600....1S}, \\bcite{2005astro.ph..9402K}, \\bcite{2005ApJ...629..259R} and \\bcite{2006astro.ph..1401G}). In this scenario, only those rare over-dense peaks which collapse before redshift $z \\sim 10$ (the epoch of reionisation) and achieve a potential well deep enough to form stars remain visible at the present epoch; the remaining satellite galaxies have star formation quenched by the background UV flux from reionisation. In this model it is not the most massive substructure halos at $z=0$ which are the LG satellites, but rather the {\\it survivors} from halos which form stars at $z \\sim 10$. There is observational evidence for such a scenario, both indirectly from high redshift quasar absorption spectra \\citep{Wyithe:2006st}, and from the star formation histories of LG dIrrs,  which show a strong suppression in star formation up to $z \\sim 1$ \\citep{2005NewAR..49..453S}.  While this general model has been investigated by a number of authors in the literature, there are significant differences in the details. \\citet{2004ApJ...609..482K} argue that only a few of the LG satellites are genuine fossils; the majority were significantly more massive in the past, formed most of their stars after reionisation (at $z \\sim$ 3), and then subsequently lost their mass through tidal stripping and shocking. They find that the central velocity dispersions can be sufficiently lowered in their model once satellite-satellite interactions are taken into account alongside stripping and shocking from the Milky Way (MW) or Andromeda (M31). This model complements the `tidal model' proposed by  \\citet{2001ApJ...559..754M} and \\citet{2001ApJ...547L.123M} to explain the distance-morphology relation between dSphs and dIrrs. In these two papers, it is suggested that  all of the Local Group satellite galaxies started out looking more like the dIrrs with a disc-like morphology. Those satellites on orbits which brought them close to the MW or M31 then formed induced bars which buckled leaving a spheroidal remnant: a dSph.  \\citet{2005ApJ...629..259R} and \\citet{2006astro.ph..1401G} focus on the pure fossils left over from reionisation. They present a detailed cosmological model which includes radiative transfer and molecular cooling from H$_2$. These new key ingredients allow them to study star formation in mini halos with virial temperatures $T < 10^4$\\,K - the temperature at which hydrogen starts to become collisionally ionised. They find, contrary to previous studies, that such mini-halos can cool efficiently and form stars. They posit that such halos could then (if they survive) be the progenitors of dSph galaxies in the LG. They require more massive halos to then become the dIrr galaxies. However, they do naturally recover the spheroidal morphology, low gas fractions, low rotational velocity and old stellar populations observed in the LG dSphs (see e.g. \\bcite{1998ARA&A..36..435M}), without recourse to any tidal transformations.  \\citet{2005ApJ...629..259R} and \\citet{2006astro.ph..1401G} study in detail the effects of photo-ionising feedback from star formation, but they do not include the effects of feedback from supernovae winds. It is important to separate feedback from supernovae (which is included in the models by \\citet{2005ApJ...629..259R} and \\citet{2006astro.ph..1401G}) from feedback from supernovae {\\it winds}, which are not. In this paper, we implement and test the effect of both kinds of feedback. This is a key difference in the work we present here. Observationally, we can see that supernovae winds are driven in galaxies undergoing a phase of star formation and that such winds are important, especially on the scale of dwarf galaxies (see e.g. \\bcite{2005MNRAS.358.1453O}). Theoretically, \\citet{1999ApJ...513..142M} and \\citet{Marcolini:2006mk} have shown that there is a transition scale at about $M_{\\mathrm{crit}} \\sim 10^7 - 10^8$\\,M$_\\odot$ below which dwarf galaxies efficiently lose their gas from supernovae winds\\footnote{We note, however, that this conclusion is degenerate with the mechanical luminosity in the wind. For the mean plausible range, efficient mass loss occurs at $\\sim 10^7$M$_\\odot$.}; while \\citet{1986ApJ...303...39D} and \\citet{2003Astro-Ph..0210454} demonstrated that supernovae feedback could account for the global scaling relations of the LG dwarfs.  In this paper, we use a suite of very high resolution cosmological hydro-dynamical simulations, which include gas cooling, star formation, feedback from supernovae, galactic winds and reionisation, to study a new model for the formation of the LG dwarf galaxies. We do not include the effects of the detailed radiative transfer and cooling physics required to model star formation in mini-halos with mass $M \\simlt 10^7$\\,M$_\\odot$, since we are interested primarily in halos more massive than this. We suggest that the smallest dwarf galaxies of the LG share a common progenitor: rare, $\\sim 3 \\sigma$, $\\simgt 10^8$\\,M$_\\odot$, dark matter halos at $z \\sim 10$. These halos are just massive enough to efficiently form stars through atomic line cooling and to hold onto their gas in the presence of supernovae winds and reionisation. As a result they are the smallest  `baryonic building block' (BBB) available for galaxy formation.  Some of these gas-rich early forming galaxies fall-in late to the LG and survive as dwarf galaxies. Those in-falling halos on benign orbits which keep them far away from the MW or M31 manage to retain their gas and slowly form stars - these become the dIrrs; those on more severe orbits lose their gas faster than they can form stars and become the dSphs. This suggests that the dIrrs should also have an old extended spheroidal component of stars: a stellar halo. There is increasing observational evidence that this is indeed the case (\\bcite{1996ApJ...467L..13M}, \\bcite{2003Sci...301.1508M} and \\bcite{2000AJ....119..177A}; and see \\bcite{2005AJ....130.1593V} for a study of dIrrs outside of the LG). It also suggests that the star formation histories of the LG dwarfs should all show an early, pre-reionisation burst of star formation. This also appears to be the case (see e.g. \\bcite{2000MNRAS.317..831H} and \\bcite{2002MNRAS.332...91D}).  The idea of a smallest building block has a long history in the literature. \\citet{1953ApJ...118..513H} pointed out that gas cooling becomes efficient for ionised hydrogen at $10^4$\\,K; \\citet{1968ApJ...154..891P} presented a model with a smallest mass block of $10^5$M$_\\odot$ for star formation in the context of globular clusters; \\citet{1984ApJ...277..470P} updated this argument to include dark matter and showed that the relevant mass was $\\sim 10^8$M$_\\odot$, as suggested here; and, from an observational point of view, \\citet{1978ApJ...225..357S} showed that the galactic globular cluster abundances are not correlated with distance -- a result which has led to the hierarchical merging model being currently favoured over the then-popular model of monolithic collapse \\citep{1962ApJ...136..748E}. Here, we take the next logical step by investigating the morphology and kinematics of the stars, gas and dark matter in these high redshift BBBs, and making a link to the smallest galaxies observed in the Local Group at the present epoch.  To test our model, we use a small box size of 1\\,Mpc and stop the simulation at $z=10$ (to avoid simulations becoming non-linear on the scale of the box). With such a small box size we achieve an unprecedented mass resolution of $\\sim 10^3$\\,M$_\\odot$ per particle. This allows us to accurately track the kinematics and morphology of galaxies of total mass $M \\sim 10^8$\\,M$_\\odot$ and compare our results with observations from the LG dwarfs. Ideally, one would like to use a larger box and evolve all the way to redshift $z=0$. However, this is not technically feasible for a mass resolution of $\\sim 10^3$\\,M$_\\odot$ per particle, at the present time. Instead we have to compromise by stopping at redshift $z=10$ and making the link to redshift $z=0$ using a mixture of linear theory arguments (section \\ref{sec:motivation}) and results from other studies in the literature (\\bcite{2005astro.ph.10370M} and \\bcite{2005astro.ph..4277M}).  This paper is organised as follows: in section \\ref{sec:motivation} we present some theoretical motivation for our model. In section \\ref{sec:hydro} we describe the simulations. We used a control simulation with only dark matter, and five other simulations which explore the effect of the star formation prescription and supernovae winds of varying strength. All simulations were run with identical initial phase space distributions. In section \\ref{sec:results} we describe the results from our suite of high resolution simulations. In section \\ref{sec:discussion}, we discuss our results in the context of recent observations. Finally, in section \\ref{sec:conclusions}, we present our conclusions.  \\begin{figure} \\begin{center} \\includegraphics[width=8.25cm]{numberhalo.eps} \\caption{The number of halos, $N$, formed up to redshift $z = 15,10,5$ (solid, dotted and dashed lines), in a given mass range which are likely to have fallen into a larger halo of mass $M_0 = 2\\times 10^{12}$\\,M$_\\odot$ at $z_0=0$. Over-plotted are the number of dSph galaxies currently observed around the MW (horizontal line) and the mass scale $M_\\mathrm{crit}$ (vertical line).} \\label{fig:numberhalo} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  Using analytic arguments and a suite of very high resolution ($\\sim 10^3$\\,M$_\\odot$ per particle) cosmological hydro-dynamical simulations, we have argued that high redshift, $z \\sim 10$, $M \\sim 10^8$\\,M$_\\odot$ halos, form the smallest `baryonic building block' (BBB) for galaxy formation. These halos are just massive enough to efficiently form stars through atomic line cooling and to hold onto their gas in the presence of supernovae winds and reionisation. These combined effects, in particular that of the supernovae feedback, create a sharp transition: over the mass range $3-10 \\times 10^7$M$_\\odot$, the BBBs drop two orders of magnitude in stellar mass. Below $\\sim 2 \\times 10^7$M$_\\odot$ galaxies will be dark with almost no stars and no gas. Above this scale is the smallest unit of galaxy formation: the BBB.  We have shown that supernovae feedback is important for these smallest galaxies. Not because it ejects the gas, but because it keeps the gas hot and extended. We find that the details of such feedback is not critical. Whether implemented as a heating term for gas surrounding star forming regions, or as a galaxy-wide wind, the results are similar. However, the combination of supernovae heating {\\it and} supernovae driven galactic winds gives the best agreement with observations. Such feedback works by reducing the star formation efficiency in subhaloes. This keeps the surface density and rotational velocity of the stars which do form low. The smallest observed galaxies in the Local Group have very low surface brightness, as does the Milky Way old stellar halo. Without such supernovae feedback our model cannot reproduce these properties. Efficient cooling of the supernovae ejecta is prevented by the ionising background from reionisation.  We connected these BBBs to galaxies observed at $z=0$ using a mixture of linear theory arguments and results from other studies in the literature. In a galaxy the size of the Milky Way, O(100) such building blocks will be accreted. \\citet{2005astro.ph.10370M} have recently shown that, of these, $\\sim 10$\\% survive to form the lowest mass Local Group dwarf galaxies. The remainder form the bulk of the Milky Way old stellar halo. The survivors will slowly form stars and lose gas over a Hubble time. Since neither reionisation nor supernovae winds actually eject gas from these galaxies, we require some other mechanism for this. Ram pressure stripping is a likely candidate. In this case, those BBBs on benign orbits which keep them far away from the Milky Way or Andromeda manage to retain their gas and slowly form stars - these become the lowest mass dwarf irregular galaxies (dIrr); those on more severe orbits lose their gas faster than they can form stars and become the dwarf spheroidals (dSphs). Both galaxy types become more metal rich and of lower abundance than the progenitor building blocks due to their more extended star formation. In this picture, the dSphs are `naked stellar halos', while the dIrrs have both an old, metal poor, stellar halo and a younger, more centrally concentrated population of stars. There is increasing observational evidence that this is indeed the case (\\bcite{1996ApJ...467L..13M}, \\bcite{2003Sci...301.1508M} and \\bcite{2000AJ....119..177A}).  We have shown that the stars in our high redshift BBBs resemble in those in the faintest dSphs in the Local Group at the present epoch. This does not prove, however, that all of the Local Group dSphs formed in this way. \\citet{2004ApJ...609..482K} and \\citet{2005astro.ph..4277M} have suggested that some formed from quite different mechanisms and were much more massive in the past than at present.  There should be many galaxies in the Local Group with surface brightness an order or two orders of magnitude fainter than those already found. We predict that these galaxies will have very similar total mass ($2-10 \\times 10^7$\\,M$_\\odot$) to those satellites already discovered.  Finally, we have commented almost exclusively on the low-density Local Group, rather than the high-density cluster environment. This is mainly because the BBBs we form in our simulations are so small and faint that, at the moment, we can only hope to observe their $z=0$ counterparts in the very nearby universe. We expect, however, many such BBBs to survive to the present epoch, even in cluster environments. The continued accretion of such BBBs over a Hubble time should lead to an old metal-poor stellar halo being a ubiquitous feature of all large galaxies. However, it is important not to think of these BBBs as {\\it the} building block for galaxy formation. We stress that they are the {\\it smallest} building block and the major contributor to old stellar halos, not the bricks from which all galaxies are made."
    },
    "0606/astro-ph0606672_arXiv.txt": {
        "abstract": "We consider the Fundamental Plane of elliptical galaxies lensed by the gravitational field of a massive deflector (typically, a cluster of galaxies). We show that the Fundamental Plane relation provides a straightforward measurement of the projected mass distribution of the lens with a typical accuracy of $\\approx 0.15$ in the dimensionless column density $\\kappa$.  The proposed technique breaks the mass-sheet degeneracy completely and is thus expected to serve as an important complement to other lensing-based analyses. Moreover, its ability to measure directly the mass distribution on the small pencil beams that characterize the size of background galaxies may lead to crucial tests for current scenarios of structure formation. ",
        "introduction": "\\label{sec:introduction}  One of the most remarkable properties of early-type galaxies is the fact that their distribution in the three-dimensional parameter space of \\textrm{effective} radius $R_\\mathrm{e}$, central velocity dispersion $\\sigma_0$, and average (intrinsic) surface brightness $\\langle \\mathrm{SB} \\rangle_\\mathrm{e}$ within $R_\\mathrm{e}$ is extremely well localized.  In particular, if $\\sigma_0$ is measured in $\\mbox{km s}^{-1}$, $\\langle \\mathrm{SB} \\rangle_\\mathrm{e}$ in $\\mbox{mag arcsec}^{-2}$, and $R_\\mathrm{e}$ in $\\mbox{kpc}$, then the parameter space is mostly filled around a locus called the ``Fundamental Plane'' \\citep{1987ApJ...313...42D, 1987ApJ...313...59D} and described by the equation \\begin{equation} \\label{eq:1} \\log R_\\mathrm{e} = \\log r_\\mathrm{e} + \\log D_A(z) = \\alpha \\log \\sigma_0 + \\beta \\bigl\\langle \\mathrm{SB} \\bigr\\rangle_\\mathrm{e} + \\gamma \\; , \\end{equation} where $\\{\\alpha, \\beta, \\gamma\\}$ are three (possibly wavelength-dependent) constants, $r_\\mathrm{e}$ is the angular effective radius (taken to be measured in radians), and $D_A(z)$ is the angular diameter distance of the galaxy at redshift $z$ in $\\mbox{kpc}$.  In general, the parameters $r_\\mathrm{e}$ and $\\langle \\mathrm{SB} \\rangle_\\mathrm{e} \\equiv 42.0521 - 2.5 \\log \\bigl(L / 2 \\pi R_\\mathrm{e}^2 L_\\odot^B \\bigr)$ are operationally defined in terms of a best fit with the $R^{1/4}$ law \\citep{1948AnAp...11..247D}; moreover, in the evaluation of the \\textit{intrinsic\\/} surface brightness $\\langle \\mathrm{SB} \\rangle_\\mathrm{e}$ from the data, a number of factors (such as cosmological dimming, k-correction, and galactic extinction) need to be taken into account. Note that $r_\\mathrm{e}$ corresponds to the half-light radius of the galaxy only if the luminosity profile is exactly described by the de~Vaucouleurs law.  [Some concerns about the universality of the $R^{1/4}$ law and the structural homology of elliptical galaxies have been raised, but the small systematic deviations can be interpreted in terms of weak homology, see \\citealp{2002A&A...386..149B}.]\\@ Usually, the central velocity dispersion $\\sigma_0$ is defined as the luminosity weighted average dispersion inside an aperture radius of $r_\\mathrm{e} / 8$.  The tightness of the Fundamental Plane, which in the nearby universe is characterized by a relative scatter in $R_\\mathrm{e}$ of the order of $15\\%$ or below \\citep[e.g.][]{1996MNRAS.280..167J}, makes this scaling relation equivalent to the existence of a ``standard rod.'' Therefore, the Fundamental Plane has been very useful as a distance indicator (e.g.\\ \\citealp{2000ApJ...529..768K,1999MNRAS.304..595G, 1998ApJ...493..529B}; see also \\citealp{1995MNRAS.276.1255V}). Surprisingly, this scaling relation has been found to hold tightly also at redshifts of cosmological interest, so that it has allowed astronomers to investigate the evolution properties of early-type galaxies out to $z \\approx 1$ (\\citealp{1999MNRAS.308.1037T, 1999MNRAS.308..833J, 2003ApJ...595...29R, 2005ApJ...623..666R} and a number of related articles; see also the discussion reported in Sect.~3 below).  So far, in the context of gravitational lensing (and mostly in the case of strong lensing) the Fundamental Plane has been mainly employed to characterize the properties of the deflector \\citep[see][]{2000ApJ...543..131K}, a natural approach since the lenses are often early-type galaxies. In contrast to previous applications, in this Letter we investigate the use of the Fundamental Plane for the background elliptical galaxies observed through a gravitational lens.  Since the Fundamental Plane essentially provides a ``standard rod,'' that is an \\textit{absolute length-scale\\/} through $R_\\mathrm{e}$, we propose to employ it to measure the magnification due to an intervening gravitational lens.  This novel technique not only is expected to provide measurements of the projected mass density of the lens with relatively high signal-to-noise ratio, but also presents several unique characteristics that make it an invaluable diagnostics of the structure of gravitational lenses. The method is applicable, in principle, to a study of single lensed galaxies or, more naturally, to a small sample of lensed objects; for a given lens, it will find best applications as a complement to other lensing-based investigations. ",
        "conclusions": "\\label{sec:conclusions}  In this Letter we have considered the application of the Fundamental Plane of elliptical galaxies, as equivalent to a ``standard rod'', to investigate the mass distribution of gravitational lenses from their magnification of the intrinsic scale-length $R_\\mathrm{e}$. We have shown that the proposed technique can provide column density measurements on pencil beams with an accuracy of $0.15$ in the reduced dimensionless mass density $\\kappa$.  As explained above, our technique has several advantages with respect to standard lensing methods, but is also observationally challenging and can be applied to the limited number of early-type galaxies typically observed through a lens.  These considerations suggest that the lensing of the Fundamental Plane would be best employed in conjunction with other lensing analyses.  In fact, ground-based weak lensing observations are based on approximately $25 \\mbox{ galaxies arcmin}^{-2}$, with each galaxy typically giving an estimate of the local shear with $0.4$ error.  Hence, by averaging the various shear measurements over a square arcmin, an error of $\\approx 0.08$ on the shear can be achieved.  This value should be compared with the $\\approx 0.1$ error expected on the convergence $\\kappa$ from our method alone, or with the $\\approx 0.06$ error from a combined analysis.  In addition, as noted earlier in this Letter, a combined analysis would break the mass-sheet degeneracy of weak lensing.  In principle, an analogous technique could be based on the existence of a ``standard candle'', as provided by the Tully-Fisher \\citep{1977A&A....54..661T} relation for late-type galaxies, \\begin{equation} \\label{eq:9} L \\propto V^p \\; , \\end{equation} where $L$ is the total absolute luminosity of the galaxy, $V$ is its \\textit{maximum\\/} rotational velocity, and $p$ is a (wavelength-dependent) constant.  Similarly to the Fundamental Plane, the Tully-Fisher scaling relation is often used as a distance estimator or to analyse the evolution properties of spiral galaxies. In the nearby universe, the measured scatter in $L$ can be as small as $20\\%$ (e.g.\\ \\citealp{1988ApJ...330..579P, 1990ApJ...351L...5W, 1997AJ....113.2046R, 2001ApJ...563..694V}; see \\citealp{1995PhR...261..271S} for a review), suggesting that the Tully-Fisher relation might be used to measure with a similar accuracy the lens magnification $| \\det A |$, and would thus provide a direct estimate of $\\kappa$ with an rms error as small as $0.1$.  The practical use of the Tully-Fisher relation as a distance estimator has been limited to low redshifts.  Recently, it has been possible to investigate the Tully-Fisher relation at relatively high redshifts \\citep[e.g.][]{2006MNRAS.366..308B, 2004A&A...420...97B}. Unfortunately, a rather large scatter is observed in these studies (approximately $0.8$--$1 \\mbox{ mag}$), which severely limits applications of the type proposed in this Letter on the basis of the Fundamental Plane. Nevertheless, if the observed scatter is not (totally) intrinsic, next generation telescopes, such as the James Webb Space Telescope or the Extremely Large Telescopes, may open the way to this other, possibly powerful, investigation technique.  Type Ia supernovae (SNe Ia) also appear to be very good (calibrated) ``standard candles'' \\citep{1996ApJ...473...88R, 1998ARA&A..36...17B} and thus they could be suitable for the lensing analysis proposed in this Letter.  In fact, such a possibility has already been considered by various authors (see, e.g., \\citealp{2001ApJ...556L..71H, 2002A&A...393...25G, 2003ApJ...583..584O, 2003MNRAS.338L..25O}), especially in relation to the prospects offered by the \\textit{Supernova/Acceleration Probe\\/} (\\textit{SNAP\\/}).  This satellite is expected to provide $\\approx 8$ multiple imaged supernovae per year within its 20 fixed, $1 \\mbox{ deg}^2$ fields \\citep{2001ApJ...556L..71H}; of these, $\\approx 2$ are expected to be Type Ia SNe for which the \\textit{absolute\\/} magnification can be measured.  Since the probability of observing a strong lensing effect on a randomly selected area of the sky (the strong lensing optical depth) is $\\tau \\approx 0.001$ for sources out to redshift $1.5$, lensing applications of SNe will be limited to the study of the few clusters for which serendipitous observations become available."
    },
    "0606/astro-ph0606358_arXiv.txt": {
        "abstract": "Angular momentum transport must have occurred in the Sun's radiative zone to explain its current solid body rotation.  We survey the stability of the early Sun's radiative zone with respect to diffusive rotational instabilities, for a variety of plausible past configurations. We find that the (faster rotating) early Sun was prone to rotational instabilities even if only weak levels of radial differential rotation were present, while the current Sun is not. Stability domains are determined by approximate balance between dynamical and diffusive timescales, allowing generalizations to other stellar contexts. Depending on the strength and geometry of the weak magnetic field present, the fastest growing unstable mode can be hydrodynamic or magneto-hydrodynamic (MHD) in nature.  Our results suggest that diffusive MHD modes may be more efficient at transporting angular momentum than their hydrodynamic (``Goldreich-Schubert-Fricke'') counterparts because the minimum spatial scale required for magnetic tension to be destabilizing limits the otherwise very small scales favored by double-diffusive instabilities.  Diffusive magneto-rotational instabilities are thus attractive candidates for angular momentum transport in the early Sun's radiative zone. ",
        "introduction": "\\label{sec:intro}  The Sun's internal structure, which has been probed extensively with seismological techniques (Thompson et al. 2003), is a unique laboratory for the study of fluid dynamical processes on astronomical scales. In this context, it is rather significant that the physical process at the origin of the solid body rotation in the Sun's radiative zone has not yet been unambiguously identified. By analogy with young Sun-like stars, it is believed that the young Sun was rotating faster than it currently is (by a factor up to $\\sim 50$) and that it was subsequently spun down by magnetic torques associated with the solar wind (see, e.g., Sofia et al. 1991 for a review).  Although details are missing, the end result of this rotational evolution should be a differentially rotating interior, with a faster rotating core. Consequently, the well-established solid body rotation in the current Sun's radiative zone is evidence that an efficient angular momentum transport mechanism has operated in the past. It is also tempting to associate this transport mechanism with ``deep,'' large-scale mixing of elements (below the convection zone), as required to explain the much depleted lithium abundance at the solar photosphere (e.g. Chaboyer, Demarque \\& Pinsonneault 1995a, 1995b).   The motivation to identify a successful transport mechanism for the Sun's radiative zone goes in fact much beyond the Sun itself, as this mechanism is likely to also have consequences for the evolution of rotating massive stars (e.g., Woosley, Heger \\& Weaver 2002), the spin of newly-formed compact objects resulting from this evolution (e.g. Ott et al. 2006) and possibly the origin of gamma-ray bursts (e.g. MacFadyen \\& Woosley 1999). The mechanism could be operating in the radiative zones of all kinds of stars or only at some stages of their evolution, but until its nature is known, it is difficult to be more specific.    Efforts to understand transport in stellar radiative zones are not new (e.g. Tassoul 1978). It was clear early on that the microscopic kinematic viscosity in the Sun, $\\nu_\\odot$, would not be sufficient to reduce any global scale differential rotation present, since the corresponding viscous timescale, \\begin{equation} \\label{eq:tvis} \\tau_{\\rm vis} \\simeq \\frac{R_\\odot^2}{\\nu_\\odot} \\sim 10^{12-13}~{\\rm years}, \\end{equation} is prohibitively long (e.g. Goldreich \\& Schubert 1967). To explain the solid body rotation of the solar radiative zone which was later revealed by helioseismology, another process must have been operating.    Several candidate mechanisms have indeed been proposed. The list includes the so-called GSF instability (Goldreich \\& Schubert 1967; Fricke 1968), the ``secular hydrodynamical shear instability'' (Zahn 1974) and internal gravity waves (e.g. Kumar \\& Quataert 1997; Talon \\& Zahn 1998; Talon, Kumar \\& Zahn 2002; Charbonnel \\& Talon 2005). No consensus has been reached yet as to which of these mechanisms, if any, provides a satisfactory answer to the current state of solid body rotation. The GSF instability alone is not expected to bring a radiative zone to a state of solid body rotation because its marginal stability conditions still allow for substantial differential rotation across cylindrical radii (Goldreich \\& Schubert 1967). The existence of the secular shear instability has not been rigorously established but rather suggested on the basis of heuristic arguments (Zahn 1974; Schatzman 1991), which makes it a difficult subject of detailed study. The efficiency of internal gravity waves at transporting angular momentum has been much debated in the literature (e.g. Press 1981; Talon \\& Zahn 1998; Talon, Kumar \\& Zahn 2002; Charbonnel \\& Talon 2005; Rogers \\& Glatzmeier 2005, 2006). It is also a priori unclear whether waves, which by nature transport energy and momentum but not the fluid itself, can achieve a sufficient level of deep mixing of elements such as lithium, whose depleted photospheric abundance is arguably more easily attributed to instabilities and turbulent mixing (but see Garcia-Lopez \\& Spruit 1991; Schatzman 1993; Charbonnel \\& Talon 2005 for a different view).   In recent years, magnetic fields have taken central stage in this discussion (e.g. Turck-Chieze et al. 2005). The role of a magnetic field has been discussed in relation to that of internal gravity waves (e.g. Mathis \\& Zahn 2005; Rashba, Semikoz \\& Valle 2005). Solar evolutionary models which account for rotation and effects attributed to magnetic fields have been presented by Eggenberger, Maeder \\& Meynet (2006). Spruit (2002) and Braithwaite \\& Spruit (2006) have discussed the possibility of growing an azimuthal magnetic field via stretching by differential rotation, followed by an instability, in stably stratified radiative zones. Finally, Menou, Balbus \\& Spruit (2004) have generalized the diffusive stability analysis of GSF to the case when a weak magnetic field is present in the radiative zone. These diffusive rotational instabilities and their potential role for the early Sun's radiative zone are our main subjects of interest here.   The usual difficulty one faces when invoking fluid instabilities in stellar radiative zones is that the strong radial thermal stratification present stabilizes the fluid against a wide variety of (adiabatic) perturbations. This is the case for rotational instabilities in particular, even if strong radial differential rotation (across spherical shells)\\footnote{Any radiative zone with negative differential rotation within a spherical shell ($\\dd \\ln \\Omega / \\dd \\theta < 0$) is subject to the standard magneto-rotational instability because perturbations tangential to the spherical shell are blind to the radial stratification, as shown by Balbus \\& Hawley (1994). Here, we are focusing exclusively on the more challenging, ``orthogonal'' situation where negative differential rotation exists across spherical shells ($\\dd \\ln \\Omega / \\dd \\ln r < 0$; see Menou et al. 2004 for a discussion).} is present. For instance, one verifies readily that, for any reasonable amount of imposed radial differential rotation, the current Sun's radiative zone would be stable according to both the general Solberg-H\\o iland adiabatic criteria (e.g. Tassoul 1978) and their generalization to weakly-magnetized fluids (Balbus 1995).  A way around this difficulty has been proposed by Goldreich \\& Schubert (1967) and Fricke (1968), who have shown that rotational instabilities can still exist, provided perturbed fluid elements are allowed to exchange heat with their environment much faster than they exchange momentum. For such (diabatic) perturbations, the stabilizing role of thermal stratification is effectively neutralized when the perturbed fluid element reaches thermal equilibrium with its environment while its original momentum remains largely unchanged. In the solar interior, radiative heat diffusion is indeed orders of magnitude faster than viscous diffusion of momentum, thus motivating the double-diffusive analysis of Goldreich \\& Schubert and Fricke (hereafter GSF altogether). The basic double-diffusive mechanism invoked in these instabilities is related to that of ``salt-finger'' instabilities operating in the Earth's oceans, except that the destabilizing salinity stratification is replaced here by a destabilizing profile of angular momentum.   In Menou et al. (2004), we studied numerically the stability of the current Sun's radiative zone to diffusive magneto-rotational instabilities under a variety of imposed rates of radial differential rotation. It was found that the Sun is relatively stable, unless a fairly substantial rate of negative differential rotation is imposed. However, as we alluded before, this is probably not the right question to ask because angular momentum transport in the Sun's radiative zone must have occurred in the past, when it was rotating faster. One would generally expect this faster rotation to be favorable to rotational instabilities. Therefore, we reconsider the stability of the radiative zone of the (faster rotating, early) Sun for various plausible rates of radial differential rotation. We do find that, if the early Sun was rotating at least a few times faster than it currently is, it was indeed much more prone to diffusive hydro- and magneto-rotational instabilities.   We describe our method of solution in \\S\\ref{sec:met}. In \\S\\ref{sec:res}, we present the results of our detailed stability survey for the early Sun's radiative zone. In \\S\\ref{sec:conc} we summarize key results, discuss some of their consequences and possible extensions of this work. ",
        "conclusions": "\\label{sec:conc}  The main conclusions of our stability survey for the early Sun's radiative zone can be summarized as follows. Provided that the early Sun was rotating at least a few times faster than it currently is, we find that it was easily subject to diffusive hydro- and magneto-rotational instabilities. That is, rather weak levels of negative radial differential rotation (across spherical shells) were sufficient to trigger these linear instabilities. This result complements the one stating that any level of negative differential rotation within a spherical shell ($\\dd \\ln \\Omega / \\dd \\theta < 0$) triggers the standard (adiabatic) magneto-rotational instability within that shell (Balbus \\& Hawley 1994). An inspection of our survey results suggests that diffusive magneto-rotational instabilities may be more efficient at transporting angular momentum than their hydrodynamical (GSF) counterparts because fastest growing unstable modes typically have larger wavelengths in the magnetized case. A simple interpretation of this trend in terms of the scale favored by a destabilizing magnetic tension was offered. This leads us to conclude that diffusive magneto-rotational instabilities are attractive candidates to explain the angular momentum transport (and possibly the elemental mixing) that must have occurred in the Sun's radiative zone.  Additional work will be required to establish these promising trends on firmer grounds. Indeed, a number of questions raised by our investigation remain unanswered at this point. Chief among them is the issue of the primordial strength (and to some extent geometry) of the magnetic field present in the Sun's radiative zone.  Our study of the hydro-to-MHD transition shows that if the magnetic field initially present in the Sun's radiative zone was too weak, the fluid would have behaved largely as if it were unmagnetized, as far as rotational instabilities are concerned.\\footnote{We note that there is a well-known pathology if one attempts to recover the stability conditions of an unmagnetized fluid by taking the zero-field limit of the corresponding ideal MHD stability conditions (e.g. Balbus 1995). Our results are interesting in this respect, as they show how the proper unmagnetized limit can be recovered at small, but finite, field values, when diffusive effects are properly accounted for.}  The minimum magnetic field values needed for the fluid stability to be determined by MHD processes are not large ($B \\gsim 0.1$ G) and they do not violate current upper limits on the field strength in the Sun's radiative zone (e.g. Friedland \\& Gruzinov 2004 and references therein). However, if the Sun's radiative zone was initially seeded with only a very weak magnetic field, the growth of GSF instabilities would be favored and the resulting transport may then be weak. Diffusive magneto-rotational instabilities could eventually play a role if these GSF instabilities were able to grow the background seed field to a large enough value for MHD modes to take over. Although this hypothetical scenario shares some similarities with the kinematic dynamo problem, the likelihood of its success remains very unclear at this point.   In addition to this dependence on seed field conditions, the role of diffusive rotational instabilities is likely to be affected by the chemical structure and evolution of the radiative zone in which they operate. By analogy with the result of Goldreich \\& Schubert (1967) on the strongly stabilizing role of composition gradients, one would generally expect diffusive rotational instabilities to be quenched when substantial composition gradients start developing in a radiative zone. This may limit to very early times only their role for the innermost radiative core of the Sun and it points, more generally, to the possibility of a complex, coupled chemical-rotational evolution (since turbulent mixing could affect the magnitude of the stabilizing composition gradients being established). In regions where composition gradients end up being significant (like the current Sun inner radiative core), it appears likely that the dynamo and transport process advocated by Spruit (2002) and simulated by Braithwaite \\& Spruit (2006) would be dominant and free to operate (see also Heger, Woosley \\& Spruit 2005 for an application to massive star evolution). In regions with negligibly small composition gradients, however, diffusive modes, growing exponentially with time constants approaching the local value of the angular velocity (Figs.~\\ref{fig:one} and~\\ref{fig:two}), are likely to dominate over a linear growth by field stretching, which is invoked in the initial phase of the scenario proposed by Spruit (2002; see also Braithwaite \\& Spruit 2006).   Let us conclude by mentioning that the locations of stability domains, as found in our survey, are broadly consistent with a simple argument involving a balance between local diffusive and dynamical timescales. Acheson (1978) already emphasized that instability to double-diffusive GSF modes should be approximately determined by the criterion \\begin{equation} \\frac{\\nu}{\\xi} \\lsim - \\frac{\\kappa^2}{N^2}, \\end{equation} where $\\kappa$ is the epicyclic frequency (which must be imaginary for instability). This simply states that the stabilizing role of thermal stratification ($N^2$) must be neutralized by a strong enough heat diffusivity ($\\xi$), but that at the same time the destabilizing role of the angular momentum profile ($\\kappa^2 < 0$) must not be neutralized by momentum diffusion (i.e. viscosity, $\\nu$). This simple criterion appears to be approximately satisfied in our stability survey. Similarly, our survey results indicate that instability to double-diffusive MHD modes is roughly determined by a similar criterion, \\begin{equation} \\frac{\\eta}{\\xi} \\lsim -\\frac{\\dd \\Omega^2 / \\dd \\ln r}{N^2}, \\end{equation} where $\\eta$ is the resistivity. This explains why unstable MHD modes only appear for values of $\\Omega / \\Omega_\\odot$ which are a few times larger than for GSF modes, in approximate proportion with the ratio $\\sqrt{\\eta/\\nu}$. It would be interesting to apply these simple criteria to the radiative zones of other stars than the Sun, at different stages of their evolution, to determine more generally how relevant diffusive rotational instabilities may be for the evolution of rotating stars."
    },
    "0606/hep-ph0606311_arXiv.txt": {
        "abstract": "We deduce the cosmogenic neutrino flux by jointly analysing ultra high energy cosmic ray data from HiRes-I and II, AGASA and the Pierre Auger Observatory. We make two determinations of the neutrino flux by using a model-dependent method and a model-independent method. The former is well-known, and involves the use of a power-law injection spectrum. The latter is a regularized unfolding procedure. We then use neutrino flux bounds obtained by the RICE experiment to constrain the neutrino-nucleon inelastic cross section at energies inaccessible at colliders. The cross section bounds obtained using the cosmogenic fluxes derived by unfolding are the most model-independent bounds to date. ",
        "introduction": "Neutrino cross sections at high energies could herald new physics because several extensions of the Standard Model predict enhanced cross sections. Examples include electroweak instanton processes~\\cite{instanton}, black hole production~\\cite{banks} and exchange of towers of Kaluza-Klein gravitons~\\cite{gravitons} in low scale gravity models~\\cite{add}, and TeV-scale string excitations~\\cite{strings}. The highest center-of-mass energy at which the neutrino-nucleon cross section has been measured is about $300$ GeV at the HERA accelerator. The only practical way to probe this cross section at center-of-mass energies above 100 TeV is by detecting the interactions of neutrinos with energy above $10^{10}$ GeV incident on Earth. By {\\it ultra high energy} one often means energies above $10^9$ GeV, but this usage is not standard.  While neutrinos with energies above $\\cal{O}$($10^4$)~GeV have not been observed so far~\\cite{Ribordy:2005fi}, they are expected to accompany ultra high energy cosmic rays (that have been observed with energies exceeding $10^{10}$ GeV) because almost all potential cosmic ray sources are predicted to produce protons, neutrinos and gamma rays with comparable rates. Experiments indicate that the highest energy cosmic rays are primarily protons~\\cite{bird}. A guaranteed source of neutrinos, the cosmogenic neutrinos, arises from the inelastic interactions of cosmic ray protons on the cosmic microwave background~\\cite{gzk}, prominently $p\\gamma \\rightarrow \\Delta^+ \\rightarrow n \\pi^+$, followed by pion decay~\\cite{cosmicnu}. The threshold energy for this reaction is the Gresein-Zatsepin-Kuzmin (GZK) energy, $E_{GZK} \\sim 4\\times 10^{10}$ GeV~\\cite{gzk}. At such high energies, the gyroradius of a proton in the galactic magnetic field is larger than the size of the Galaxy and it is therefore expected that these cosmic ray protons are of extragalactic origin. Since the attenuation caused by the above reaction has a length scale of about 50 Mpc~\\cite{stecker,yoshida}, a strong suppression in the cosmic ray spectrum is expected above $E_{GZK}$. There is evidence for the GZK cutoff in Fly's Eye/HiRes~\\cite{hires,berg2} data but not in AGASA~\\cite{agasa,takeda} data. The Pierre Auger Observatory~\\cite{pao} (which we will refer to as Auger in what follows) is expected to resolve this conflict.  Estimates of the cosmogenic flux are very model-dependent and differ by about $2-3$ orders of magnitude~\\cite{yoshida,cosmogenic,cosmomodels}. As we discuss below, the uncertainty in the cosmogenic flux is further exacerbated by recent evidence from the HiRes experiment that protons dominate the flux at about $5\\times 10^{8}$ GeV~\\cite{morehires}, two orders of magnitude below $E_{GZK}$.  The uncertainty in the cosmogenic neutrino flux must be dealt with before progress can be made to extract or constrain the neutrino-nucleon inelastic cross section using cosmic neutrinos. We prefer to use cosmic ray data to infer the cosmogenic neutrino flux, thus side-stepping theoretical modelling. So long as a lower bound on the cosmogenic flux is established, it will be possible to place an upper bound on the neutrino-nucleon inelastic cross section.  In Refs~\\cite{ring,ahlers}, estimates of the cosmogenic neutrino flux have been made from separate analyses of AGASA and HiRes-II data combined with their predecessor collaborations, Akeno and Fly's Eye, respectively. These results have been used to constrain the neutrino-nucleon cross section in Ref.~\\cite{ring2}. Our approach is different.  It has been shown in previous work that the apparent disagreement between AGASA and HiRes data could be a result of large systematic uncertainties in the energy determinations of the two experiments~\\cite{olinto}. The energy scale uncertainty of the AGASA experiment is 30\\%~\\cite{takeda}, while that of the HiRes detectors is about 20\\% above $10^{10.5}$ GeV and about 25\\% above $10^{9.5}$ GeV~\\cite{berg}. The data from Auger have an energy uncertainty of about 25\\%~\\cite{sommers}.  Rather than perform separate analyses of these datasets, we take into account the large uncertainties in the energy measurements and perform a combined analysis of Auger~\\cite{sommers}, HiRes-I (sample collected between June 1997 and May 2005) and II (sample collected between December 1999 and May 2003)~\\cite{berg2}, and AGASA data~\\cite{agasa}. We expect the double-counting of events that are common to the HiRes-I and HiRes-II datasets to have a negligible effect on our results. We suppose that the energy uncertainties for HiRes-I and II are the same and allow a total of 3 floating energy scale parameters: one each for Auger, HiRes and AGASA.  It was thought that the ankle at $10^{10}$ GeV indicates the onset of the dominance of a higher energy flux of particles over a lower energy flux, such as an extragalactic component starting to dominate over the galactic component of the cosmic ray flux. The dip structure in the vicinity of the ankle can be explained as a consequence of pair production from extragalactic protons on the CMB~\\cite{ber}. The latter process has a threshold of $10^{8.6}$ GeV thus permitting the interpretation that protons dominate even below this energy. Recent evidence for a steepening of the spectrum at about the same energy~\\cite{moreagasa}, the second knee, suggests that a low energy transition is consistent with the data. Moreover, the HiRes collaboration finds a drastic change of composition across the second knee, from about 50\\% protons just below this knee to about 80\\% protons above~\\cite{morehires}. Since the end-point of the galactic flux is thought to be comprised of heavy nuclei, a changeover to protons is interpreted as the onset of the dominance of the extragalactic flux. We do not analyse events with energy below the second knee since they are probably not of cosmological origin. However, since composition measurements are strongly model-dependent (via the theoretical uncertainty in predicting electron and muon shower sizes and the depth of shower maximum for hadronic showers), the only statement about the transition energy that can be made with some confidence is that it lies in the interval $10^{8}$ GeV to $10^{10}$ GeV. ",
        "conclusions": "We obtained the proton injection spectrum from cosmic ray data using two methods. In the first, we performed a combined analysis of HiRes-I and II, AGASA and Auger data in the energy interval, $10^{9.6}$ GeV to $10^{11}$ GeV, using a power law injection spectrum. We found the quality of the fit to be unsatisfactory even though we did not include super-GZK events in the analysis. The primary reason for the poor fit is apparently that the initial Auger data are noisy. In the second method, we implemented a regularized unfolding of the injection spectrum for each dataset separately.  We found the resultant cosmogenic neutrino flux using the injection spectra from both methods to be in excellent agreement with eachother. This suggests that the injection spectrum is well-modeled as a power-law and that the poor fit mentioned above is due to unaccounted-for systematic uncertainties.  Using the model-independent limit on the neutrino flux from RICE, we constrained the neutrino-nucleon cross section under the assumption that new physics modifies the charged-current and neutral-current interactions by the same constant factor.  Our results are succintly summarized in Figs.~\\ref{fig:protons}--\\ref{fig:cross}."
    },
    "0606/astro-ph0606281_arXiv.txt": {
        "abstract": "We show that dark matter emerging from late decays ($z  \\lesssim  1000$) produces a linear power spectrum identical to that of Cold Dark Matter (CDM) on all observationally relavant scales ($\\gtrsim 0.1$ Mpc), and simultaneously generates observable constant-density cores  in small dark matter halos. We refer to this class of models as {\\it meta}-Cold Dark Matter ({\\it m}CDM), because it is born with non-relativistic velocities from the decays of cold thermal relics. The constant-density cores are a result of the low phase-space density of $m$CDM at birth. Warm dark matter cannot produce similar size phase-space limited cores without saturating  the Ly$\\alpha$ power spectrum bounds. Dark matter dominated galaxy rotation curves and stellar velocity dispersion profiles may provide the  best means to discriminate between $m$CDM and CDM. {\\it m}CDM candidates are motivated by the particle spectrum of supersymmetric and extra dimensional extensions to the standard model of particle physics. ",
        "introduction": "In the standard cosmological model, the bulk of the matter in the universe is in the form of Cold Dark Matter (CDM). Structure in CDM is built up hierarchically, from initial density perturbations that are adiabatic and nearly  scale-invariant  \\cite{adiabatic}. CDM provides a compelling description  of the universe on large scales, but it has not conclusively passed all observational tests on small scales. Two well-known  shortcomings are  that simulations predict  Milky Way-sized    galaxies   should have nearly an order of magnitude more dark matter substructures than observed satellites \\cite{missing}, and density profiles of low-mass galaxies tend to be less cuspy than those seen in simulations \\cite{Simon05,McGaugh05,Zackrisson06,KuziodeNaray06} and are often well-fit by constant-density cores. While the former can be alleviated both by  astrophysical \\cite{astro} and cosmological \\cite{Kamionkowski99,ZB03} methods,  the  latter, if ultimately true, may require a more drastic modification of the CDM paradigm \\cite{drastic}.  Warm dark matter (WDM) provides an alternative that may alleviate the problems on small scales \\cite{WDM,HD}. Indeed, we {\\em expect} constant density cores in the centers of WDM halos because the initial phase-space maximum can never be exceeded \\cite{TG}.   An additional consequence of the large WDM particle velocities is that its free-streaming length is large. This suppresses the power spectrum on small scales.  One problem with WDM as an alternative to CDM, however, is that models with velocities large  enough to produce observationally-relevant cores are in conflict with measurement of the Lyman-$\\alpha$ forest power spectrum \\cite{Strigari:2006ue}. Indeed,  Lyman-$\\alpha$  forest  analyses  suggest that the linear density fluctuation spectrum  is nearly identical to the CDM prediction on scales larger than $\\sim 1$ Mpc \\cite{Narayanan00,Viel05,Abazajian05,Seljak06,Viel06}.  It is thus interesting to consider if any modification to CDM can reproduce the pattern of hierarchical structure formation, and also create large phase-space limited cores in the halos of low-mass galaxies. In this paper  we investigate such a dark matter model, which we refer to as ``{\\it meta}-Cold Dark Matter''  ($m$CDM). We define $m$CDM as particles that emerge relatively late in cosmic  time  ($z \\lesssim  1000$)    and are  born  non-relativistic from  the decays  of   cold particles. As we will show, the $m$CDM model provides both large phase-space cores {\\em  and}   CDM-like   power spectra    on  $\\gtrsim  0.1$ Mpc scales, consistent with the most stringent measurements of the Ly$\\alpha$ forest power spectrum \\cite{Seljak06,Viel06}. Specifically, $m$CDM is born late enough that the free-streaming scale is small, but the velocity dispersion is large enough to give rise to a reduced phase-space density. Even though $m$CDM velocities are non-relativistic, the redshifted velocities today are larger than the corresponding velocities for WDM. Previous studies of similar models considered decays with shorter lifetimes $\\sim  M_{pl}^2/M_{weak}^3$ (of order   a month) to {\\em relativistic} daughter particles \\cite{Kaplinghat05,Cembranos05}. In  these models the relationship between the phase-space density and cut-off scale in the power spectrum is similar to that of WDM.  Dark matter in the $m$CDM class arises in many extensions  to  the standard model of particle physics. For  example  in  supersymmetric   models we can  consider   a sneutrino decaying into a neutrino and gravitino with a lifetime $\\tau \\simeq  3.6 \\times    10^8$  s  $   (100   \\, {\\rm  GeV}/\\Delta   m)^4 m_{\\tilde{G}}/{\\rm  TeV}$  \\cite{Feng03,Ellis04}, where $\\Delta m$ is the mass difference between the sneutrino and gravitino, and $m_{\\tilde{G}}$ is the gravitino mass.  A  similar relation holds for Kaluza-Klein WIMPs   decaying to gravitons in theories  with extra  dimensions. For example, we can consider a 10 TeV gravitino produced from sneutrino decays with a lifetime $\\tau = 5 \\times  10^{12}$ s, where the velocity imparted to the gravitino is $v/c \\sim 10^{-3}$. ",
        "conclusions": "For the models considered, the net effect of the injected relativistic energy during decay may be phrased in terms of the effective number of light neutrinos, $\\Delta N_\\nu =  1.8  \\times    10^{-2} \\sqrt{\\tau/{\\rm yrs}} (p_{cm}/m)$. For the $m_{\\tilde G} = 10$ TeV, $\\tau = 5 \\times 10^{12}$ s model described above, we have $\\Delta N_{\\nu}  \\simeq 0.01$ which would be hard to detect.  Neutrinos  and    photons  produced  along   with  $m$CDM  will arrive unscattered in the diffuse radiation backgrounds today.  The strongest constraint  comes from  the  neutrinos produced  directly in  two-body decays. The relevant low-energy  neutrino flux limit is   $< 1.2$ cm$^{-2}$  s$^{-1}$ from  Super-Kamiokande for energies $18-82$ MeV \\cite{SK}. The  above sneutrino model is consistent  with this limit and in the sensitivity range of future experiments with  reduced    energy thresholds \\cite{gadzooks}. Photons may also be produced in three-body decays and subsequent hadronization into neutral pions. Using a very conservative branching  fraction of 10$^{-3}$   into photons \\cite{Feng04}, we find that the photon flux is consistent with observations \\cite{Kinzer97}.  In  conclusion, we have  shown  if dark  matter  is produced from the late decay of cold relics, it will give rise to large  cores in low-mass  galaxies {\\em and}  alleviate  the dwarf satellite problem,   without  destroying  agreement with   the observed Lyman-$\\alpha$ forest power spectrum.  This  differs from  standard WDM models, which,  in   order  to remain consistent   with   constraints from the Lyman-$\\alpha$ power  spectrum,  cannot produce sizeable cores in small galaxies. Dynamical studies of nearby galaxies may provide the best means to test the $m$CDM scenario and compare it to CDM predictions.   We  thank J. Beacom, J. Cooke, S. Kazantzidis, S. Koushiappas, and R. Wechsler for  discussions. LES is supported in part by a Gary McCue Postdoctroral Fellowship through the  Center for Cosmology at UC Irvine.  We acknowledge the use of CMBfast \\cite{cmbfast}.   \\vspace*{-0.25cm}"
    },
    "0606/astro-ph0606248_arXiv.txt": {
        "abstract": "We present the full set of power spectra of cosmic microwave background (CMB) temperature and polarization anisotropies due to the coupling between quintessence and pseudoscalar of electromagnetism. This coupling induces the rotation of the polarization plane of CMB photons, thus resulting in non-vanishing $B$ mode and parity-violating $TB$ and $EB$ modes. Using the 2003 flight of BOOMERANG (B03) data, we give the most stringent constraint on the coupling strength. In some cases, the rotation-induced $B$ mode can confuse the hunting for the gravitational lensing-induced $B$ mode. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606412_arXiv.txt": {
        "abstract": "Several different processes could be changing the density in the core of a neutron star, leading to a departure from $\\beta$ equilibrium, quantified by the chemical potential difference $\\delta\\mu\\equiv\\mu_n-\\mu_p-\\mu_e$. The evolution of this quantity is coupled to that of the star's interior temperature $T$ by two functions that quantify the rate at which neutrino-emitting reactions proceed: the net reaction rate (difference between $\\beta$ decay and capture rates), $\\Gamma_{\\rm net}(T,\\delta\\mu)$, and the total emissivity (total energy emission rate in the form of neutrinos and antineutrinos), $\\epsilon_{\\rm tot}(T,\\delta\\mu)$. Here, we present a simple and general relationship between these variables, ${\\partial\\epsilon_{\\rm tot}/\\partial\\delta\\mu=3\\Gamma_{\\rm net}}$, and show that it holds even in the case of superfluid nucleons. This relation may simplify the numerical calculation of these quantities, including superfluid reduction factors. ",
        "introduction": "The simplest weak interaction process that could proceed in the core of a neutron star is the so-called direct Urca (or {\\it Durca}) process, which consists of the two successive reactions, $\\beta$ decay and capture. \\begin{equation} \\label{beta}n\\rightarrow p+e+\\bar{\\nu}_e\\qquad\\qquad p+e\\rightarrow n+\\nu_e. \\end{equation} It is the most powerful of the neutrino processes potentially leading to the cooling of the neutron star. However, these reactions are allowed only if each of the Fermi momenta of neutrons ($n$), protons ($p$), and electrons ($e$) is smaller than the sum of the two others (triangle condition). This implies a (highly uncertain) threshold in the matter density  for the direct Urca processes. Another reaction, which overcomes this restriction, is the so-called modified Urca (or {\\it Murca}) process, which involves an additional spectator nucleon: \\begin{equation} \\label{MURCA} n+N\\rightarrow p+e+\\bar{\\nu}_e+N\\qquad\\qquad p+e+N\\rightarrow n+\\nu_e+N, \\end{equation} where the additional nucleon $N$ can be either a neutron or a proton. The matter is transparent to neutrinos, which escape freely, transporting their energy away from the star.  Furthermore, these two reactions bring nucleons into the state of beta (or chemical) equilibrium, which determines the concentration of neutrons and protons. The beta equilibrium condition $\\mu_n=\\mu_p+\\mu_e$ involving the chemical potentials of the constituent particles implies the equality of $\\bar{\\nu}_e$ and $\\nu_e$ emission rates (or the net reaction rate set equal to zero). If a slight departure from equilibrium $\\delta\\mu=\\mu_n-\\mu_p-\\mu_e$ is produced by any external or macroscopic phenomenon that changes the density of a fluid element, one of the two reactions in equations (\\ref{beta}) or (\\ref{MURCA}) becomes more intense and changes the fraction of protons and neutrons towards the new equilibrium values (Le Ch\\^atelier's principle). In the subsequent evolution, the chemical imbalance affects the stellar interior temperature $T$ by increasing the phase space available to the products of the neutrino-emitting reactions and by converting chemical energy into thermal energy, and the temperature affects the chemical imbalance by also determining the rate at which reactions proceed. Several authors have investigated external processes that induce non-equilibrium $\\beta$ reactions, such as radial pulsation \\citep{finzi65,finzi_wolf}, gravitational collapse \\citep{haensel92,Gourgou93}, a changing rotation rate \\citep{andreas95,andreas97,fernand05,andreas06}, and a hypothetical time-variation of the gravitational constant \\citep{jofre06}. However, none of these has considered an adequate model for the effects of the likely Cooper pairing (superfluidity) of nucleons (see \\citealt{andreas97} for a very rough estimate of their likely importance).  Two basic rates are relevant in order to follow the coupled evolution of $T$ and $\\delta\\mu$ (see, e.~g., \\citealt{andreas95,fernand05}): the total (neutrino and antineutrino) emissivity (energy per unit volume per unit time), $\\epsilon_{\\rmn{tot}}=\\epsilon_{\\bar{\\nu}}+\\epsilon_\\nu$, and the net reaction rate (number of reactions or emitted lepton number per unit volume per unit time), $\\Gamma_{\\rmn{net}}=\\Gamma_{\\bar{\\nu}}-\\Gamma_\\nu$. Recently, \\citet{villain05} analyzed the effect of nucleon superfluidity on the net reaction rates, $\\Gamma_{\\rm net}$, calculating superfluid reduction factors for the direct and modified Urca processes by means of sophisticated numerical methods. Nevertheless, they did not evaluate the total emissivity of neutrinos, $\\epsilon_{\\rm tot}$, which is also required in order to compute the eventual time-evolution of $T$ and $\\delta\\mu$.  In this work, we present a curious relationship between $\\epsilon_{\\rmn{tot}}$ and $\\Gamma_{\\rmn{net}}$, both considered as functions of the chemical imbalance parameter $\\dmu$, that holds even in the superfluid case: \\begin{equation} \\label{3sig}\\frac{\\partial\\epsilon_{\\rmn{tot}}}{\\partial\\dmu}=3\\Gamma_{\\rmn{net}}. \\end{equation} The advantage of this relation is that, having computed one of the functions numerically (as done by \\citealt{villain05}), the other is obtained very easily, saving time in the calculation.  In the non-superfluid case, \\citet{andreas95} calculated $\\epsilon_{\\rmn{tot}}$ and $\\Gamma_{\\rmn{net}}$ analytically. For Durca processes, \\begin{equation} \\label{R1} \\epsilon_{\\rmn{tot}}^{\\rmn{Durca}}(T,\\dmu)=\\epsilon_0^D\\left(1+\\frac{1071u^2+315u^4+21u^6}{457}\\right), \\end{equation} and \\begin{equation} \\label{R2} \\Gamma_{\\rmn{net}}^{\\rmn{Durca}}(T,\\dmu)\\dmu=\\epsilon_0^D\\frac{714u^2+420u^4+42u^6}{457}, \\end{equation} where $u\\equiv\\dmu/(\\pi T)$ is a dimensionless parameter, and $\\epsilon_0^D\\equiv\\epsilon_{\\rmn{tot}}^{\\rmn{Durca}}(T,0)\\propto T^6$ is the equilibrium emissivity. For Murca processes, \\begin{equation} \\label{R3} \\epsilon_{\\rmn{tot}}^{\\rmn{Murca}}(T,\\dmu)=\\epsilon_0^M\\left(1+\\frac{22020u^2+5670u^4+420u^6+9u^8}{11513}\\right), \\end{equation} and \\begin{equation} \\label{R4} \\Gamma_{\\rmn{net}}^{\\rmn{Murca}}(T,\\dmu)\\dmu=\\epsilon_0^M\\frac{14680u^2+7560u^4+840u^6+24u^8}{11513}, \\end{equation} with $\\epsilon_0^M\\equiv\\epsilon_{\\rmn{tot}}^{\\rmn{Murca}}(T,0)\\propto T^8$ (see \\citealt{haensel92} for precise estimates). In these two pairs of polynomial expressions, it is easy to verify our proposed relation (eq.~\\ref{3sig}). The purpose of this paper is to show that it is valid beyond these simple cases, encompassing, for example, the cases of superfluid neutrons and/or protons. A brief and clear discussion of superfluidity in neutron star was given by \\citet{yako2001b}, and a rigorous and deep analysis can be found in \\citet{lombardo}. ",
        "conclusions": "\\label{sec:concl} We have proven a simple, general relationship (eq.~\\ref{3sig}) between the main rates characterizing non-equilibrium $\\beta$ processes in both superfluid and non-superfluid neutron star matter. This relation could simplify the evaluation of these quantities in superfluid neutron star models, complementing numerical calculations such as those of \\citet{villain05}.  We thank Francisco Claro for very helpful suggestions. We are also grateful to Olivier Espinosa, Claudio Dib, Miguel Kiwi, Paula Jofr\\'e, and Elena Kantor for discussions that benefited the present paper. S.~F.-T. was supported by a CONICYT PhD fellowship and A.~R. by FONDECYT grants 1020840 and 1060644."
    },
    "0606/astro-ph0606138_arXiv.txt": {
        "abstract": "A sophisticated SBF analysis package has been developed, designed to measure distances of early-type galaxies by means of surface brightness fluctuations of unresolved stars. This suite of programs called SAPAC is made readily available to the astronomical community for extensive testing with the long-term goal to provide the necessary tools for systematic distance surveys of early-type galaxies using modern optical/NIR telescopes equipped with wide-field cameras. We discuss the technical and scientific concepts of SAPAC and demonstrate its capabilities by analysing deep $B$ and $R$-band CCD images of 10 dwarf elliptical (dE) galaxy candidates in the Fornax cluster obtained with FORS1 at the VLT. All candidates are confirmed as cluster members. We then turn our attention to the innermost region of the Fornax cluster. A total of 29 early-type galaxies closer than three cluster core radii ($2^\\circ$) from the central galaxy NGC1399 have radial velocities and SBF distances.  Their Hubble diagram exhibits a pronounced S-shaped infall pattern suggesting a picture that Fornax is still in the process of formation during the present epoch through a general collapse and possible accretion of distinct groups of galaxies. From fitting a model we estimate the cluster mass within 720\\,kpc projected distance from NGC1399 which is $2.3\\pm0.3\\times 10^{14}\\mathcal{M}_\\odot$. The associated collapse time is $t_{\\rm collapse}=2.9{+1.6 \\atop -0.9}$\\,Gyr. After cleansing our galaxy sample from a few kinematical outliers the true distance of the Fornax cluster core is determined at 20.13$\\pm$0.40\\,Mpc [$(m-M)_0=31.51\\pm 0.04$\\,mag]. Applying a bootstrap resampling technique on the distance distribution with individual distance errors taken into account further reveals a small intrinsic cluster depth of $\\sigma_{\\rm int}=0.74{+0.52 \\atop -0.74}$\\,Mpc in best agreement with the cluster's linear extension in the sky: $\\sigma_{\\rm R.A}=\\sigma_{\\rm DEC}\\approx0.5$\\,Mpc. We conclude that the early-type galaxy population in the Fornax cluster must be spatially well constrained with no evidence of elongation along the line-of-sight in contrast to the Virgo cluster. Moreover, we find marginal evidence for substructure, a results that is consistent with the young evolutionary state of the cluster and the overall galaxy infall. Combining the kinematically defined cluster distance with the mean cosmological velocity for the central cluster galaxy sample yields a Hubble constant of $H_0 = 63\\pm 5$ \\kms Mpc$^{-1}$. ",
        "introduction": "Finding low star density dwarf galaxies in the local Universe is, by nature, a challenging task. Over the last two decades, numerous surveys \\citep[e.g.][]{B85,J97,F90,JBF00,K95,CA98} have produced hundreds of candidates in nearby groups and clusters based on morphological properties. But due to the lack of any distance information large numbers still remain only dwarf {\\em candidates} today.  To separate the wheat from the chaff, i.e., genuine dwarf galaxies from similar looking background spirals and giant ellipticals, distances are required. The importance of distances, however, extends far beyond the unambiguous identification and spatial location of new dwarf galaxies. The careful construction of highly complete, volume-limited galaxy samples like the 10\\,Mpc sphere \\citep{K79,S92,K04} yields the exact number and distribution of dwarfs relative to larger galaxies and provides crucial empirical input to better constrain theories of galaxy formation, including the faint-end of the galaxy luminosity function and the merging rate \\citep{TS02,JTT04}. Furthermore, accurate distances permit studies of the dwarf galaxy distribution in 3-D and, if combined with kinematical and stellar population data, provide insight into the importance of environmentally driven over self-induced physical processes that govern dwarf galaxy evolution \\citep[e.g.][]{PP01}.  Unlike the scarcely evolved, low density environment populating dwarf irregular galaxies, dwarf elliptical (dE) galaxies are lacking significant amounts of neutral hydrogen gas preventing a relatively simple distance estimate via 21cm\\,redshift. Furthermore, little help is coming from empirical scaling relations between galaxy luminosity and surface brightness \\citep{BC91} or S\\'ersic shape parameter \\citep{B98} as the large intrinsic scatter turns out to be unsuitable for individual distance measurements.  That is why in the last decade a formidable observational effort has been devoted to deriving distances to dEs by using the tip magnitude of the red giant branch (TRGB), a powerful Pop II distance indicator \\citep[][and references therein]{K03,D90}. However, the requirement of accurate photometry for individual stars 1--2 magnitudes below the tip [$m_I\\geq 25.8+5\\log(D/5 {\\rm Mpc})$] at distance $D$[Mpc] sets a practical application limit at $\\approx$ 5\\,Mpc beyond which TRGB measurements quickly become very costly and time consuming. This was demonstrated by \\cite{H98} who spent nine hours with the HST to obtain the TRGB distance for a single dE galaxy (IC3388) in  the Virgo cluster.  Because of this limitation, a new distance indicator, the so-called surface brightness fluctuation (SBF) method, has received increasing attention in recent years. The method quantifies the statistical pixel-to-pixel variation of star counts across a galaxy image with the major technical advantage of working on unresolved stellar populations. Ground-based CCD images with relatively short exposure times can be used instead of high resolution space-based observations making the SBF method extremely efficient. Distance measurements for early-type galaxies far beyond the reach of any other distance indicator become possible.  The theoretical framework of the SBF method was developed by \\cite{T88} to study the distribution of luminous giant ellipticals in the local Universe \\citep[e.g.][]{T01}. While SBF applications were initially focused on high surface brightness systems \\citep{A94,J96,T97}, Jerjen and collaborators \\citep{J98, J00, J01,J03} demonstrated that the method works equally well for dEs with stellar surface densities as low as $\\mu_{\\rm B_{eff}} =26$\\,mag\\,arcsec$^{-2}$ and luminosities comparable to Local Group dwarfs like Leo\\,I, And\\,II, and Sculptor ($M_B= -10$\\,mag). The SBF method has since been successfully applied to dwarf candidates in the distance range from 10--40\\,Mpc, employing 2-m ground-based telescopes \\citep{R05}, the 8-m ESO-VLT \\citep{M03,J04}, and the HST \\citep{Jo05, M05}.  Noticing that dE galaxies are by far the most numerous galaxy type at the current cosmological epoch \\citep{F94}, the SBF technique, in combination with modern wide-field imaging at optical and near-IR telescopes like SUBARU, VISTA, and LSST, has the potential to play a major role in future {\\em distance} surveys of early-type galaxies in the local Universe. Using extended source catalogues from SDSS and Pan-STARRS in the North and SkyMapper in the South there is the prospect of carrying out SBF distance surveys in a similar fashion as redshift surveys are done today. For the efficient processing of such large amount of imaging data we are currently engaged in a project of developing an automated SBF Analysis Pipeline for the Astronomical Community called SAPAC.  In \\S\\ref{sapac} we will introduce SAPAC, discuss the data requirements for the pipeline and give an overview of the SBF method.  We will then demonstrate its capabilities in \\S\\ref{test} by processing VLT+FORS1 CCD images for 10 Fornax cluster dE candidates of similar quality as existing and future survey telescopes are expected to produce. In \\S\\ref{discuss} we endeavor to extract clues about the 3-D cluster structure from the spatial distribution of early-type galaxies in the central region of Fornax. Combining galaxy velocities and distances and analysing the Hubble diagram enables the measurement of various cluster parameters such as the total mass and dynamical collapse time, as well as a value for the Hubble constant. We summarize our results in \\S\\ref{conc}. ",
        "conclusions": "\\label{discuss} \\subsection{Distance distribution and Substructures}\\label{distsub} With the availability of significant numbers of distances for cluster galaxies true 3-D substructure studies in clusters have finally come of age. Most recent distance studies of the Virgo cluster have found extensive substructure \\citep{W00,S02,J03}. \\cite{W00} used SBF distance measurements from \\cite{T01} to 14 elliptical galaxies in the northern part of the Virgo cluster to show a strong collinear arrangement in three dimensions, defining a principal axis between NGC4387 and NGC4660, otherwise known as cluster A \\citep{B87}.  Cluster B is centered around M49 and these two major concentrations of the Virgo Cluster are thought to be surrounded by three galaxy clouds. \\cite{J04} also used the SBF method to determine distances to 16 dE galaxies in the Virgo cluster and found first evidence for line-of-sight bimodality in cluster A, when combined with 24 ellipticals with SBF distances from \\cite{T01} and \\cite{N00}.  They concluded that the northern cluster A consists in fact of two dynamically distinct systems, a small group around M86 falling into the M87 subcluster from the back.  The Fornax cluster is less rich than Virgo and considered structurally to be less complex. It hosts about 350 likely cluster members \\citep{F89}. They are essentially concentrated in one major component that is dominated by early-type galaxies (see Fig.~\\ref{distribution}). The cluster has a King profile core radius of $R_{\\rm c}=0.67^\\circ$ ($\\simeq 0.24$\\,Mpc) and a center that almost exactly coincides with the location of the giant elliptical NGC1399 \\citep{F89}. Approximately three degrees to the southwest lies a loose aggregate that is dominated by the peculiar Sa galaxy Fornax A (NGC1316) and composed of mainly late-type galaxies. In Ferguson's initial study, cluster membership and background objects were distinguished for the most part based on morphology. Subsequent spectroscopic surveys have reassigned a few members from likely membership to background status and vice versa and have also determined some previously unknown redshifts \\citep{D98,SD01,D01}. \\cite{M99} measured a relatively smooth Gaussian velocity distribution with very little substructure compared to Virgo. \\cite{DG01} analysed redshifts of 108 Fornax cluster members, including 55 dwarfs, and found kinematical evidence for a large main concentration, centered around the velocity of NGC1399. The small second component around NGC1316 has a slightly higher ($\\approx 100$\\,km\\,s$^{-1}$) mean velocity suggesting an infall towards the  cluster center from the near side. But all this substructure evidence is based on velocities only and thus does not provide a genuine three dimensional picture of the Fornax cluster.  Fornax was also subject to some initial distance work. The Cepheid distances of three spiral galaxies were measured, NGC1365 at 18.3$\\pm2.3$\\,Mpc, NGC1326A at 18.7$\\pm1.9$\\,Mpc \\citep{S99}, and NGC1425 at 22.2$\\pm1.4$\\,Mpc \\citep{M00}, only one of them (NGC1365) inside the central area as defined in Fig.~\\ref{distribution}. These results provided some first insight into the spatial depth of the cluster. In order to get a better understanding of the galaxy distribution in the cluster core, \\cite{J03} determined SBF distances to eight dEs and found a mean distance of $(m-M)_0=31.54\\pm0.07$\\,mag, or 20.3$\\pm$0.7\\,Mpc and a cluster depth of $\\sigma_{\\rm int}=1.4{+0.5 \\atop -0.8}$\\,Mpc. The former value was in good agreement with previous SBF work on 26 Fornax E/S0, Sa giants \\citep[31.50$\\pm0.04$\\,mag,][]{T01} spread over a larger cluster area.  With the availability of SBF distances for a total of 43 early-type galaxies (including the 10 dEs from this study) we have the first-time luxury to conduct a 3-D Fornax cluster analysis right in the cluster center. Fig.~\\ref{ddist} shows three different representations of the distance distribution of the 31 dE+E galaxies that are in projection closer than three core radii ($2^\\circ$) from NGC1399. This sample size is equivalent to 35 percent of all early-type galaxies in the Fornax Cluster Catalog brighter than $-13.5$\\,mag located in that area. The top panel is the histogram, binned in 0.1\\,mag bins.  The second panel uses the adaptive kernel method \\citep{V94} with a Gaussian kernel.  Interestingly, both graphs show evidence for two or even three concentrations, the two major ones centered at a distance modulus of $\\approx$ 31.3 and at just under 31.6. The bottom panel is probably the closest to reality showing the distribution using the adaptive kernel method and adopting an average SBF distance uncertainty of 0.2\\,mag.  The process completely straightens out the substructure indicating that the number of galaxies and the accuracy in the distance measurements are not sufficient to establish the reality of this effect. Nevertheless, the concentrations in the distribution of the individual galaxies still stand out and should be tested in future studies. The selected 31 dE+E galaxies give a preliminary distance of $19.95\\pm0.37$\\,Mpc [$(m-M)_0=31.49\\pm0.04$\\,mag] to the Fornax cluster core. A final value will be given in \\S\\ref{fdist}.  Next we wish to investigate the line-of-sight cluster depth to further reduce the uncertainties in the \\citet{J03} result. The observed dispersion in the distance distribution of our sample galaxies, $\\sigma_{(m-M)}=0.24$\\,mag is influenced  by two factors: the distance errors and the intrinsic cluster depth. The depth can be calculated via: $\\sigma_{\\rm los}^2=\\sigma_{(m-M)}^2-\\sigma_{\\Delta(m-M)}^2$, where $\\sigma_{\\Delta(m-M)}$ must be a representative estimate of the distance errors. Because the distance accuracy of our sample shows quite some variation we generated 100,000 bootstrap samples \\citep{E86} of distances and errors from the original data set. For each sample we calculated its mean distance and mean error and used the previous equation to determine $\\sigma_{\\rm los}$. Histograms were then produced for all three parameters, each well approximated by a Gaussian. The $\\sigma_{\\rm los}^2$ distribution was found to peak at 0.0075 with a 1$\\sigma$ scatter of 0.0105.  This translates into a true line-of-sight cluster depth of $\\sigma_{\\rm int}=0.74{+0.52 \\atop -0.74}$\\,Mpc  which is in good agreement with the small linear extension of Fornax in the sky as estimated from the positions of early-type galaxies in the FCC: $\\sigma_{\\rm R.A.}=0.55\\pm0.06$\\,Mpc; $\\sigma_{\\rm DEC}=0.46\\pm0.05$\\,Mpc. We conclude that the early-type galaxy population in the central region of the Fornax cluster is spatially well constrained with no evidence of elongation along the line-of-sight.  \\subsection{The Hubble diagram} \\label{hub} Unlike previous work \\citep[e.g.][]{DG01} that concentrated on measuring global kinematical properties of the Fornax cluster we now have velocities and distances available for 29 early-type galaxies within a 2 degree (720\\,kpc) radius from NGC1399, a data set highly suitable for analysing gravitationally induced motion right in the core of the Fornax cluster. The velocity-distance diagram for these galaxies is shown in Fig.\\ref{hubbledia}.  On the near side, almost all galaxies (9 out of 12) with distance estimates closer than the tentative cluster center are redshifted relative to the cosmological expansion. Only the three galaxies NGC1386, FCC222, and FCC278 stand out as they appear isolated in the velocity-distance space. They are blueshifted by $\\approx 360$\\,km\\,s$^{-1}$ relative to the Hubble flow and $\\approx 700$\\,km\\,s$^{-1}$ to the the other sample galaxies at comparable distances. The main galaxy in the ensemble, NGC1386, whose SBF distance was measured by \\cite{T01}, is a Sa galaxy hosting a Seyfert 2 nucleus \\citep{W91}. FCC222 was classified as a SAB(s)0 galaxy but no evidence for a disk or spiral features was found in our high resolution FORS1 images. Hence, the morphological type ``dE0,N\" as initially given in the Fornax cluster catalogue \\citep{F89} must be considered to be a more accurate description.  The situation in the Hubble diagram is equally pronounced on the far side of the cluster. With only two and possibly a third exception all galaxies are blueshifted to the systemic velocity. Those are NGC1404 together with NGC1382 and the dwarf FCC100. NGC1404 is a well known object in the X-ray community \\citep{JS97,MD05}. The X-ray emission of NGC1404 shows a distorted envelope towards NGC1399, suggesting that NGC1404, which has a similar distance as NGC1399 but a very different velocity, experiences ram pressure stripping of the galaxy gas, caused by its relative motion of $\\Delta v=544$\\,km\\,s$^{-1}$ through the surrounding intracluster medium. Furthermore, \\cite{DG01} speculated that another galaxy, the irregular NGC1427A at $03^h 40^m09.3^s$ $-35^d37^m28^s$ (J2000) with a velocity $v_\\odot=2028$\\,km\\,s$^{-1}$ may be associated to NGC1404 but no distance is available to clarify this to date.  Overall the bulk of sample galaxies exhibit a prominent S-shaped infall pattern suggesting that Fornax is still in the process of formation during the present epoch through a general collapse. The Hubble line intercepts the kinematically defined curve at a distance of $\\approx 19.9$\\,Mpc right at the location, i.e.~distance and velocity, of NGC1399 thereby highlighting the central role of this galaxy in the cluster. Based on the prescription of \\citet{T84} we used a point mass model to estimate the total cluster mass. Adopting 13.7\\,Gyr  as the age of the Universe \\citep[WMAP, ][]{S03} we fitted a mass model at the data points. The best-fitting envelope for the line of sight angle of two degrees from the Fornax cluster center is indicated as solid curve in Fig.~\\ref{hubbledia} mapping the extreme model velocity. From the amplitude of $735$\\,km\\,s$^{-1}$ we get a cluster mass of $2.3\\pm0.3\\times 10^{14}\\mathcal{M}_\\odot$ within 720\\,kpc projected distance from NGC1399. This result is two times larger than the previously estimated Fornax mass by \\citet{DG01} based solely on redshifts for early-type galaxies, and approximately 40 times the mass of NGC1399 \\citep[$0.6\\times 10^{13}\\mathcal{M}_\\odot$, ][]{JS97}. The estimated collapse time for the system is $t_{\\rm collapse}=2.9{+1.6 \\atop -0.9}$\\,Gyr.  \\subsection{Fornax cluster distance and the Hubble constant} \\label{fdist} To determine the physically most meaningful distance to the Fornax cluster we need to remove the six outlying galaxies. The remaining selection of 23 early-type galaxies within a 2 degree radius of NGC1399 that follow the overall infall pattern of the cluster yields a true cluster distance of $20.13\\pm0.40$\\,Mpc [$(m-M)_0=31.51\\pm0.04$\\,mag]. This value shall now be used to estimate the Hubble constant. The standard procedure is to transform the mean velocity for the sample galaxies, $1447\\pm 66$\\,km\\,s$^{-1}$, to the frame of the cosmic microwave background. This required two corrections: (i) a correction of $-115\\pm 18$\\,km\\,s$^{-1}$ from heliocentric to the barycenter of the Local Group \\citep{C99}, and (ii) a projected Virgocentric flow correction of $-54\\pm 9$\\,km\\,s$^{-1}$ in the direction of Fornax. The latter is based on the distance geometry and a $1/R$ flow-field characteristics of the Local Group-Virgo-Fornax system as outlined in \\citet{M99} and the assumption of an infall velocity of $240\\pm40$\\,km\\,s$^{-1}$ of the Local Group into the Virgo cluster \\citep{J93}. Finally we use the derived cosmological expansion velocity $v_{CMB}^{Fornax}=1278\\pm69$\\,km\\,s$^{-1}$ to compute the Hubble constant of $63\\pm5$km s$^{-1}$ Mpc$^{-1}$.  We find good agreement with the values derived from various independent techniques such as Cepheid stars by the HST ``Key Project'' team \\citep[$72\\pm8$\\,km s$^{-1}$ Mpc$^{-1}$, ][]{F01}, the Cosmic Microwave Background by the WMAP team ($73\\pm3$\\,km s$^{-1}$ Mpc$^{-1}$), and galaxy clusters in the local Universe \\citep[$69\\pm5$km s$^{-1}$ Mpc$^{-1}$, ][]{G97}."
    },
    "0606/astro-ph0606624_arXiv.txt": {
        "abstract": "An understanding of stellar acoustic oscillations is attempted, using the local wave concept in semi-analytical calculations. The local homogeneity approximation allows to obtain simplified equations that can nevertheless describe the wave behavior down to the central region, as the gravitational potential perturbation is not neglected. Acoustic modes are calculated as classical standing waves in a cavity, by determining the cavity limits and the wave phases at these limits. The internal boundary condition is determined by a fitting with an Airy function. The external boundary condition is defined as the limit where the spatial variation of the background is important compared to the wavelength. This overall procedure is in accordance with the JWKB approximation. When comparing the results with numerical calculations, some drawbacks of the isothermal atmosphere approximation are revealed. When comparing with seismic observations of the Sun, possible improvements at the surface of solar models are suggested. The present semi-analytical method can potentially predict eigenfrequencies at the precision of $\\pm$3 $\\mu$Hz in the range 800-5600 $\\mu$Hz, for the degrees $l=0-10$. A numerical calculation using the same type of external boundary conditions could reach a global agreement with observations better than 1 $\\mu$Hz. This approach could contribute to better determine the absolute values of eigenfrequencies for asteroseismology. ",
        "introduction": "\\label{Introduction} Asteroseismology describes mainly the propagation of matter waves and their eigenmodes in stars (or planets) that are considered as cavities. By an inverse procedure one can then deduce the physical properties of the layers crossed by the waves. It is a very effective tool, and often the unique one to study the internal structure of stars.  Since several decades, many intensive works were undertaken, especially for the Sun, to calculate precisely the oscillation modes and to understand the physical mechanisms involved. Different reviews are dedicated to this subject \\citep{Cox_1980, Unno_etal_1989, JCD&Berthomieu_1991, Turck-Chieze_etal_1993, JCD_cours_1998}.  Semi-analytical approaches have been developed, essentially based on various asymptotic approximations. \\citet{Lamb_1909, Lamb_1932}, \\citet{Cowling_1941}, \\citet{Tassoul_1980}, \\citet{Leibacher&Stein_1981}, \\citet{Duvall_1982}, and \\citet{Deubner&Gough_1984}, have explored the first-order asymptotic approximation, and have helped to interpret the high-order modes measured for the Sun. Then with the availability of low-order modes, \\citet{Vorontsov_1991}, \\citet{Lopes&Turck-Chieze_1994},\\citet{Roxburgh&Vorontsov_2000}, have taken into account the effect of gravity in second-order theories, allowing to better describe the propagation in central regions of stars. On another hand, numerical calculations are required to reach the high precision in frequency determination, needed for inversion of stellar internal layer physics. Several high accuracy oscillation codes already exist, offering the computation of eigenfrequencies and eigenfunctions \\citep{JCD_1982, Unno_etal_1989}. They have been widely used to interpret the large amount of solar modes observed from Earth or space.  Just at the beginning of the spatial-asteroseismology era, it is worth assessing the progress of these theoretical efforts in the prediction of observations. It is interesting to note that the agreement between numerical results and observations is remarkable in the range 1000 - 2400 $\\mu$Hz where $\\nu_{num} - \\nu_{obs}$ is contained in $\\pm$1 $\\mu$Hz. At larger frequencies however, the disagreement increases gradually, up to 20 $\\mu$Hz at 5000 $\\mu$Hz. Semi-analytical results agree well with numerical ones only in that high-frequency range, while their dispersion at lower frequencies stays important even when they have been reduced by the second order approximation. Moreover, the acoustic cut-off frequency is theoretically predicted around 5000 $\\mu$Hz, while from observations it is rather around 5600 $\\mu$Hz \\citep{Fossat_etal_1992, Garcia_etal_1998}. Those difficulties can be overcome in the case of the Sun thanks to the great number of eigenmodes accessible by observations. But such difference could be more problematic in asteroseismology as only low-degree modes ($l \\leq 3$) are available and the stellar radius and mass are less known.  The discrepancy between calculation and observation is generally attributed to an incorrect description of the outer layers of the current solar models at the surface. Several studies have tried to introduce a more relevant model of convection via the introduction of a turbulent pressure \\citep{Kosovichev_1995}. That reduces the discrepancy at high frequencies, but the excellent agreement at low frequencies is also deteriorated \\citep{Rosenthal_1995, Rosenthal_1999}. The role of the chromosphere was also invoked, but no real improvement is yet obtained \\citep{Wright&Thompson_1992, Vanlommel&Goossens_1999, Dzhalilov_etal_2000}.  With the regular progress of observational instruments, these remaining problems become more crucial: the theory-observation discrepancies at high frequencies and the inaccuracy of semi-analytical methods at low frequencies are up to 100 times larger than observational errors. So, we suggest to search the reasons not only in the solar model at the surface, \\emph{but also} in the eigenmode calculation itself, especially about the wave description in the solar interior and the boundary conditions at the surface. Concepts closer to classical wave physics remain to be further explored. It can be considered as a continuation of the works of \\citet{Leibacher&Stein_1981} and \\citet{Duvall_1982}. Some first acoustic-frequency calculations was attempted in \\citet{Nghiem_2000}, and the effect of gravitational potential was taken into account in \\citet{Nghiem_2003a}. In this paper, we aim to further formalize the concepts, by putting a special accent on the internal and external boundary conditions. This work is organized as follows. In \\S 2, a new p-mode calculation is developed. Its approximations and its boundary conditions are proved to be in accordance with the JWKB method. In \\S 3, the obtained results are first compared to present numerical calculations, and some drawbacks of the isothermal atmosphere approximation are pointed out. Then in \\S 4 comparisons with observations suggest a way to improve solar surface descriptions, and at the same time to evaluate the potential of the proposed method. Concluding remarks are presented in \\S 5.  In this paper, the solar model used is the seismic model \\citep{Turck-Chieze_etal_2001, Couvidat_etal_2003}, computed with the CESAM code \\citep{Morel_1997}. The numerical results for solar oscillations are obtained with the ADIPLS code \\citep{JCD_code_1997}. Acoustic mode observations of the Sun come from the SoHO spacecraft (MDI data from \\citet{Rhodes_etal_1997}; GOLF data from \\citet{Bertello_etal_2000}, and from \\citet{Garcia_etal_2001}.The graphic representation for modes of degrees $l = 0$ to 10 follows the symbols indicated in Figure \\ref{fig_legende}.  \\begin{figure}[htb] \\centering \\plotone {f1.eps} \\caption{Symbols used for the graphic representation of the $l=0$ to 10 modes in all this paper.} \\label{fig_legende} \\end{figure} ",
        "conclusions": "The local wave approximation has been fully and coherently applied to stellar acoustic modes. Conceptually, it is in total accordance with the JWKB approach. Pratically, it consists in clearly determining the cavity limits and the reflection conditions. The standing wave nature of the modes are thus firmly claimed. In this simple context, a particularly reduced number of equations is necessary.  The wave propagation in the innner part is nevertheless correctly represented, as the gravity effect is directly taken into account. No development into higher orders, no artificial change of the horizontal eigenvalue $l(l+1)$ are needed. It turns out that the $l=0$ modes do not touch the centre. Not surprisingly, the gravity plays an important role for low-$l$ modes, being able to shift frequencies by tens of $\\mu$Hz. The main result is that the present analytical calculation is equally valid for low than high frequencies and degrees, only the coupling with gravity modes is not considered.  At the surface, a new boundary condition is proposed, directly based on the propagation condition of an acoustic wave, which corresponds to the JWKB validity criteria. Although that returning point should perhaps be refined, some first conclusions can already be drawn. A coefficient like $a_v$ seems to be sufficient to characterize the validity of a JWKB solution. But the main result is that an alternating boundary condition leads to frequency shifts of the same range than present theory-observation discrepancies. To explain the latter, a search for a more appropriate boundary condition cannot be excluded. The method proposed here has the advantage to use the atmosphere model as it is, without any modification. Comparisons with observations should allow to directly recover the real atmosphere structure. Examples are given to find out the atmosphere used by numerical codes in the context of isothermal atmosphere approximation, and to suggest a way to deduce physical properties of the solar atmosphere when observations are considered.  Despite the simplicity of the proposed semi-analytical formalism, it seems that main physical phenomenons involved are taken into account. That is why for the Sun, a precision of $\\pm 3 \\mu$Hz can be reached in the range 800-5600 $\\mu$Hz, and for $l=0-10$. Some calculations for other stars show a comparable behaviour. If now a similar type of external boundary conditions is used in numerical calculations, a potential precision well better than 1$\\mu$Hz can be expected for the whole frequency range. That could represent a step forward a global agreement between the three actors involved in helio-, astero-seismology: semi-analytical calculation bringing the physical understanding by predicting the good trend for the results, numerical calculation bringing the high precision by avoiding most of the approximations, and observation allowing to infer real stellar structures when searching to cancel theory-observation discrepancies."
    },
    "0606/astro-ph0606678.txt": {
        "abstract": "Accretion processes in quasars and active galactic nuclei are still poorly understood, especially as far as the connection between observed spectral properties and physical parameters is concerned. Quasars show an additional degree of complexity compared to stars that is related to anisotropic emission/obscuration influencing the observed properties in most spectral ranges. This complicating factor has hampered efforts to define the equivalent of an Hertzsprung-Russel diagram for quasars. Even if it has recently become possible to estimate black hole mass and Eddington ratio for sources using  optical and UV broad emission lines, the results are still plagued by large uncertainties. Nevertheless, robust trends are emerging from multivariate analysis of large spectral datasets of quasars. A firm observational basis is being laid out by accurate measurements of broad emission line properties especially when the source rest-frame is known. We consider the most widely discussed correlations (i.e. the so-called ``eigenvector 1 parameter space\" and the ``Baldwin effect\") and analyze how they can be explained in terms of accretion properties, broad line region structure, and source evolution. We critically review recent estimates of black hole mass, accretion rate, spin and possible orientation indicators, stressing that any improvement in these parameters will provide a much better understanding of the physics and dynamics of the region producing the optical and UV broad emission lines. More accurate measurements of Eddington ratio and black hole mass may have a significant impact on our ideas about evolution of quasar properties with redshift and luminosity as well as on broader cosmological issues.  %Possible techniques and minimum quantitative requirements needed %to achieve this goal are discussed. ",
        "introduction": "Tremendous data gathering advancements make now possible to see a distant solution for most of the deepest conundrums concerning quasar research. There has been an exponential growth in the number of known quasars lasting since the 1970s. Presently, the 11$^{th}$\\ edition of the V\\'eron-Cetty \\& V\\'eron Catalogue of AGNs and quasars \\cite{veroncettyveron03} lists $\\approx$50000 quasars.  The {\\it Sloan Digital Sky Survey} (SDSS) plans to compile a sample of $\\sim$ 100000 quasars. The SDSS Third Data Release provides {\\it Charge Coupled Device} (CCD) spectra for 51000 quasars found over 4200 deg$^2$\\ \\cite{sdss}. The improvement in observational capabilities has been possible by the introduction and spread of CCDs and other linear, high detective quantum efficiency devices as detectors for astronomical observations since the late 1980s. At a second stance comes the increase in access to large light-gathering power telescopes of aperture $\\simgt$ 4m,  crucial for spectroscopic observations. The ability to carry out multi-frequency observations with increasing spectral resolution and sensitivity (most notably provided by {\\it Hubble Space Telescope} (HST) and by the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) for the optical/UV) through the 1990s and early 2000s has been a third factor of relevance. The data gathering improvements -- which are the foundation of every astronomical advancement -- have led to the discovery of systematic trends in quasar properties. These suggest that quasars are {\\em not} well described by an average spectrum. We can amplify this statement to say that the spectra of quasars seen at a fixed viewing angle  are also not the same -- the basic tenet of Unification Schemes \\cite{antonucci93}. Yet, data have not been fully digested in physical terms. Observational constraints have been only very recently organized in meaningful ways. We still lack the ability to derive important physical information from observational parameters on an object-by-object basis. The first aim of this paper is to set the point of our present knowledge on the diversity of quasar spectral properties, especially in the optical and UV spectral ranges (\\S\\ \\ref{e1}). We then discuss much needed improvement to ensure accurate measurements of the main physical parameters (\\S\\ \\ref{mass}) as well as structural constraints on the line emitting regions (\\S\\ \\ref{blr}) and the physical basis of quasar diversity (\\S\\ \\ref{spin}, \\S\\ \\ref{edd}, \\S\\ \\ref{end}).  \\subsection{Basic Accretion Parameters}  We regard quasars and active galactic nuclei as systems accreting matter onto a massive ($\\simgt 10^5$ \\msol) black hole (see \\S \\ref{arebhneeded} for possible caveats). The main parameters are, as for any accreting system, black hole mass (\\mbh, \\S \\ref{mass}), Eddington ratio (\\S \\ref{edd}), and mass accretion rate $\\dot{M}$ in \\msoly. The Eddington ratio \\lledd\\ is defined as the ratio between the AGN bolometric luminosity \\lbol\\ and the Eddington luminosity \\ledd, i.e., the limiting luminosity beyond which radiation pressure overcomes gravitational attraction if accretion is spherical \\cite{peterson97}. The Eddington luminosity is directly proportional to  \\mbh, and can be written as  $$L_{\\rm Edd} \\approx 1.3 \\times 10^{38} (M_{\\rm BH}/M_\\odot)~ {\\rm ergs~ s}^{-1}.$$  Since the power emitted by an active nucleus through conversion of mass into energy can be expressed as \\lbol$ = \\eta \\dot{M} c^2$, it is possible to define an Eddington accretion rate $\\dot{M}_{\\rm Edd} = L_{\\rm Edd}/(\\eta c^2)$, and hence a dimensionless accretion rate $\\dot{m}={\\dot{M}}/\\dot{M}_{\\rm Edd}$.  We stress that \\lledd $\\propto$ \\lm, and that \\lm\\ is a quantity that can be derived from observations. On the other hand, the value of the efficiency $\\eta$ ($\\sim 0.1$) depends on the accretion mode which may not be the same for all sources: if matter is confined in an accretion disk as it is customarily assumed \\cite{shakurasunyaev73}, $\\eta$\\ depends  on the geometry and radiative properties of the disk, which in turn may depend on $\\dot{M}$). Therefore we will not always confuse the dimensionless accretion rate $\\dot{m}$ and the Eddington ratio as it is frequently done in literature: for a fixed efficiency, there could be well super-Eddington accretion even if the source is radiating at, or below, the Eddington limit \\cite{collinetal02}. Black hole spin (\\S \\ref{spin}) and an orientation angle (defined as the angle between the line of sight and the axis of the accretion disk around the black hole $\\theta$, \\S \\ref{theta})  probably matter in the context of UV/optical properties although they have turned out to be very elusive to measure.  \\subsection{Nomenclature and Samples \\label{samples}}  We use the word quasar here in a generic fashion which means all classes of extragalactic objects that show {\\em broad} (full width at half maximum, FWHM $\\simgt$ 1000 \\kms) optical and UV emission lines. This includes the nuclei of Seyfert and broad-line radio galaxies (BLRG) as well as radio-quiet (RQ) and radio-loud (RL) optically unresolved sources. They are often referred to under the umbrella of {\\it Active Galactic Nuclei} (AGNs) which reenforces the paradigm that they are driven by the same physics differing only in their redshift-implied distances and, hence, luminosity. At the same time, there is no divide in AGN occurrence at the canonical boundary (absolute B magnitude \\mb\\ $\\approx$ 22.5 for $H_0 \\approx$ 65 \\kms Mpc$^{-1}$) that separates Seyfert nuclei from quasars \\cite{veroncettyveron03}, so that nomenclature may be kept luminosity-independent. For instance, we keep using the word Narrow Line Seyfert 1 (\\S \\ref{e1extremes}) for sources that are actually luminous quasars.  Most sources considered in this review were ``classically\" selected through color criteria. Major color-based surveys include the {\\it Palomar-Green} (PG), the {\\it Large Bright Quasar Survey} (LBQS), the {\\it ESO-Hamburg} (HE) quasar survey and the SDSS. The SDSS photometric system was designed to allow quasars at 0 $\\simlt z \\simlt$ 6 to be identified with multicolor selection techniques \\cite{schneideretal02}. Other high-quality data-sets, even if  more heterogeneous and less complete, are also considered \\cite{borosongreen92,corbinboroson96,grupeetal99,marzianietal03a,shangetal04,willsetal93}, especially if spectral resolution $\\lambda/\\Delta\\lambda \\sim 1000$ and continuum $S/N \\simgt 20$. These samples often include $\\sim 100$ objects of ``good\" spectra; studies on the LBQS involve $\\sim 10^3$\\ sources, while the SDSS accounts for another order of magnitude leap in sample size, with $\\sim 10^4$\\ sources.  No consideration to obscured AGNs will be given since the very possibility of accretion parameter estimation from the observed optical/UV quasar spectra turns out to be related to the measurement of quasar broad emission line shifts, profile widths, and equivalent widths. Quasars showing broad lines are often referred to as ``type-1\" AGNs to distinguish them from Seyfert 2 nuclei and type 2 quasars (see e.g., Refs. \\cite{antonucci93,padovanietal04}) which will not be considered here.  There are more cumbersome, borderline AGNs which could be grouped into three classes:  \\begin{enumerate} \\item sources that show very large broad Balmer intensity decrement \\ha/\\hb $\\simgt$ 7 \\cite{dongetal04}. Objects of this kind often show dramatic differences between the broad component of \\ha\\ and \\hb, and are  not always found in color-based surveys. RL AGNs hosted in early-type galaxies like Pictor A are among them \\cite{sulenticetal95}.  \\item  An analysis of spectra of quasars in the {\\it Half-Jansky Parkes Survey} \\cite{francisetal01} suggests that the wide majority ($\\approx$ 80\\%) are sources whose continuum is dominated by a pure power-law spectrum ($f_\\nu \\propto \\nu^{0 \\div -2}$) ascribed to synchrotron radiation. These sources show broad lines and are bona-fide type-1 sources although their equivalent width is somewhat lower than that expected for color-selected quasars (see also \\S \\ref{edd}). Sources with rather low W(\\hbbc) and very broad \\heii\\ (most notably NGC 1275, indeed a core-dominated BRLG) may be seen as somewhat peculiar with respect to the other AGNs, especially if an important goal of the spectral analysis is \\mbh\\ measurements.  \\item A third class of sources may show contamination in the nuclear spectrum of emission components associated to strong star formation. At the very least, the bolometric luminosity can be significantly affected as in the case of Mkn 231 \\cite{braitoetal04,crescietal04,sulenticetal05}. Circumnuclear starbursts do occur in Seyfert 1 sources and quasars \\cite{guetal01a,imanishiwada04,daviesetal04}, and are usually {\\em not} resolved. The emission line spectra of luminous Seyfert 1 nuclei may well be affected by the absorption/emission features observed in strong wind galaxies (e.g. NGC 4691; \\cite{garciabarretoetal95}).  {\\it Broad Absorption Lines} (BAL) observed in quasars could be associated to nova stars \\cite{shields96}. \\end{enumerate}  We mention these three AGN optical/UV spectral typologies since they may be extremes of AGNs accretion parameters, especially of \\lledd\\ (\\S \\ref{mass}, \\S \\ref{edd}): sources in the third class are likely to be ``young\" quasars radiating at high \\lledd; the first class may include ``dying\"  quasars radiating at very low \\lledd\\   with little reservoir of matter (\\S \\ref{evolution}).  \\subsection{Rationale for a Black Hole in Active Galactic Nuclei \\label{arebhneeded}}  Do we have a definitive proof that the central massive object at the center of quasar is a black hole?  From the definition of event horizon (\\rg $ = 2 G M /c^2$\\ for a non-rotating black hole, where $G$\\ is the gravitational constant), the mass of the black hole must satisfy the condition \\mbh/\\rg $\\simgt 6.7 \\cdot 10^{27}$ g cm$^{-1}$. Although this criterion has not been satisfied yet by observations of extragalactic sources, circumstantial evidence in favor of a black hole is now considered overwhelming \\cite{blandford90}. From the time delay in the response of emission lines to continuum variations (``reverberation mapping\") it is possible to estimate the central mass assuming that the line broadening is due to Doppler effect, and to test whether the emitting gas motion is virial. In luminous Seyfert 1 galaxies, the central mass can be $\\sim 10^8$ \\msol\\ within the emitting region of the broad lines (the ``Broad Line Region\", BLR) size, of the order of 1 light month. The resulting \\mbh/\\rb $\\sim$ 10$^{24}$ g cm$^{-1}$\\ is far less than the \\mbh/$r$\\ requirement for a black hole. However, time responses for lines of different width are roughly consistent with a Keplerian trend i.e., with a ``point-like\" mass concentration with respect to the BLR size \\cite{krolik01}. It is therefore legitimate to talk at the very least about ``compact massive objects\" at the center of AGNs. This would be a proper, observation-bounded definition.   We will use the term black hole with the reassurance of a large body of cirumstantial evidence as well as with the perspective of decisive tests. In addition, most of the reasoning presented in this paper is largely independent on the exact nature of the ``compact massive object,\" provided that it is really compact and massive. ",
        "conclusions": "The main purpose of this paper was to point out how it is becoming possible to measure accretion  parameters (especially \\mbh\\ and \\lm) of AGNs.  We highlighted major results and sources of uncertainty, as well as the connections between \\mbh\\ and \\lledd, the structure of the broad line emitting region, and physical evolution. If we compare our present knowledge to the one of about 10 years ago, we can see an impressive progress. \\mbh\\ determinations through the virial assumption have become widely popular, and have been applied to quasars at all known $z$, up to 6.4, even if several sources of uncertainty limit strongly the accuracy by which \\mbh\\ is known for individual objects. It has been possible to show that many line properties and trends (including the Baldwin effect) correlate with \\lledd, and the the BLR structure is strongly affected by \\lledd\\ through E1. There is no doubt that we have a much better framework with a strong observational support to understand the BLR structure, minority objects like ``double-peakers\" and BAL QSOs, and source evolution. Improvements on \\mbh\\ and \\lm\\ estimates are possible through new multivariate analysis including orientation sensitive parameters, or through disk and wind BLR models accounting for the combined effects on  line parameters of orientation and physics. We may be able to achieve a satisfactory vision of quasar structural evolution (not to mention boundaries on cosmography) well before the emitting regions may be resolved through space interferometry missions (e.g., Ref. \\cite{unwinetal02}).  %\\newpage \\bigskip  %"
    },
    "0606/astro-ph0606142_arXiv.txt": {
        "abstract": "SKA\\footnote{The Square Kilometer Array, a next generation interferometric instrument for centimeter-wave radio astronomy. See http://www.skatelescope.org for a detailed description.}  will be operating at the same time with several new large optical, X--ray and $\\uGamma$--ray facilities currently under construction or planned. Fostering synergies in astrophysical research made across different spectral bands presents a compelling argument for designing the SKA such that it would offer imaging capabilities similar to those of other future telescopes. Imaging capabilities of the SKA are compared here with those of the major future astrophysical facilities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606374_arXiv.txt": {
        "abstract": "{We discuss the synthesis of some short-lived isotopes and of $^{23}$Na in thermally pulsing AGB stars with initial mass of 2 M$_{\\odot}$ and two different metallicities ($Z$=1.5$\\times$10$^{-2}$, corresponding to the metal amount in the present Sun, and $Z$=10$^{-4}$), representative of disk and halo stars, respectively.\\\\ The different nucleosynthesis channels are illustrated in some details. As previously found, the $^{13}$C formed after each third dredge up episode is usually completely consumed by $\\alpha$ captures before the onset of the subsequent thermal pulse, releasing neutrons. This is the most efficient neutron source in low mass AGB stars, and the resulting s-process nucleosynthesis is at the origin of the solar main component. However, in the solar metallicity model, we find that the temperature of the first formed $^{13}$C pocket remains too low during the interpulse and the $^{13}$C is not completely burnt, being partially engulfed in the convective zone generated by the following thermal pulse. Due to the rapid convective mixing in this zone, the $^{13}$C is exposed to a larger temperature and a nucleosynthesis characterized by a relatively high neutron density develops. The main effect is the strong enhancement of isotopes located beyond some critical branching in the neutron-capture path, like $^{60}$Fe, otherwise only marginally produced during a {\\it standard} s-process nucleosynthesis. ",
        "introduction": "Low mass AGB Stars (1$<$$M$/M$_{\\odot}$$<$3) are among the most important polluters of the Milky Way, because of the strong winds eroding their chemically enriched envelopes. In the AGB phase, the material processed during the alternating series of H and He burnings is recurrently mixed within the convective zones generated by thermal pulses (TPs) \\cite{sh65} and then partially mixed with the surface by the so called third dredge up (TDU) episodes. As a matter of fact, MS, S, C(N) and some post-AGB stars are enriched in C and s-process elements and the spectroscopic detection of unstable Tc demonstrates that the s process is at work in these stars \\cite{merr52}. \\\\ \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{roneu.ps}} \\caption{Peak neutron density evolution after the ingestion of the $^{13}$C-pocket formed after the first TDU episode (solar metallicity model).} \\label{Fig1} \\end{figure} Besides $^{99}$Tc, a number of short-lived isotopes are produced during the s-process, among which $^{41}$Ca, $^{60}$Fe and $^{107}$Pd. Moreover, $^{26}$Al is an important product of the H-shell burning, by proton capture on $^{25}$Mg. In this work we want to study the impact of an improved version of the FRANEC code on the production of short-lived radioactivities, and discuss the challenging  hypothesis  that the protosolar nebula was polluted by the winds of a close-by AGB star (see \\cite{was06} for a detailed review of this argument). A major improvement in the stellar evolution code was the introduction of a physical algorithm for the treatment of the convective/radiative interfaces \\cite{cris01} (see Section \\ref{c13conv}). In such a way we succeeded to overcome the discontinuity in the thermal gradients that otherwise would result at the border of the H-rich convective envelope and the inner radiative C and He-rich intershell region. Furthermore, the partial diffusion of protons in the top layers of the He shell gives naturally rise to the subsequent formation of a primary $^{13}$C pocket, partially overlapped with an outer $^{14}$N pocket, followed by a further minor $^{23}$Na pocket, both of primary origin. The new models presented here are for AGB stars of initial mass $M $ = 2 M$_{\\odot}$ and two metallicities (Z=1.5$\\times$10$^{-2}$, Z=10$^{-4}$), representative of disk and halo stars, respectively. In Section \\ref{c13conv} we discuss some key characteristics of the self-consistent formation of the $^{13}$C pocket and of the $^{14}$N pocket along the AGB phase. In Section \\ref{radioact} we examine in detail the impact of the new models on the production of short-lived radioactivities, while in Section \\ref{sodium} we reanalyze the production of $^{23}$Na in AGB stars. ",
        "conclusions": "In this paper we study how short-lived isotopes are created in a low mass ($M$=2 M$_{\\odot}$) AGB star, showing the strong effects a single convective $^{13}$C burning episode implies. A possible explanation for the ESS short-lived isotopes has been proposed, this however implying strong {\\it ad hoc} assumptions.\\\\ In addition, we present detailed nucleosynthetic results concerning the production of $^{23}$Na in low mass AGB stars at different metallicities (Z=1.5$\\times$10$^{-2}$ and Z=10$^{-4}$), representative of disk and halo stars respectively. Different processes responsible for the $^{23}$Na nucleosynthesis have been analyzed in detail, stressing the importance of varying key nuclear reaction rates."
    },
    "0606/astro-ph0606004_arXiv.txt": {
        "abstract": "{We present a discovery and observation of an extraordinarily bright prompt optical emission of the GRB\\,060117 obtained by a wide-field camera atop the robotic telescope FRAM of the Pierre Auger Observatory from 2 to 10 minutes after the GRB. We found rapid average temporal flux decay of $\\alpha = -1.7 \\pm 0.1$ and a peak brightness $R = 10.1$ mag. Later observations by other instruments set a strong limit on the optical and radio transient fluxes, unveiling an unexpectedly rapid further decay. We present an interpretation featuring a relatively steep electron-distribution parameter $p \\simeq 3.0$ and providing a straightforward solution for the overall fast decay of this optical transient as a transition between reverse and forward shock.}   \\offprints{ \\hbox{Martin Jel\\'{\\i}nek, e-mail:\\,{\\tt mates@iaa.es}}}  \\date{Received  / Accepted } ",
        "introduction": "After the establishment of the GRB Coordinates Network (GCN) \\citep{bacodine} in 1995, the technical advances have enabled a fast and reliable dissemination of satellite gamma-ray burst (GRB) data to ground-based observers. Subsequently, wide use of the sophisticated robotic follow-up systems has led to the first insight into the very early phases (less than 5 minutes after trigger) of the optical transients (OTs) accompanying some GRBs. However, despite the extended efforts, optical data at the very early stages are still sparse, so the whole picture remains unclear.  The definition of the prompt optical emission was given by \\cite{piranold} as the optical emission arising during the $\\gamma$-emission period. This early emission is usually explained in terms of either a reverse or an internal shock \\citep{saripiran99}. With a certain delay after the GRB, the afterglow -- emission due to the forward shocks -- starts to dominate the steeply decaying prompt emission. This transition flattens the lightcurve. The original rapid decay of the prompt emission slows down to a more modest decay due to the afterglow. It is generally accepted that the physical mechanisms of the prompt emission and the afterglow are distinct. The distinction is probably not absolute, some observations suggest \\citep[cf.][]{chincarini} that at least the $X$-ray intensity of the early emission phases can be extrapolated from the GRB emission itself.  In this letter we present the observation of a very bright optical transient associated with GRB\\,060117 observed by the robotic telescope FRAM.  \\begin{figure}[b!] \\begin{center} \\label{fig1} \\resizebox{\\hsize}{!}{\\includegraphics{060117-fig1.eps}} \\caption{The optical light curve of GRB\\,060117 in comparison with other well-covered early GRB optical emissions: GRB\\,990123 \\citep{Akerlof99}, GRB\\,021211 \\citep{Li03}, GRB\\,050502a \\citep{Guidorzi05}, GRB\\,050525a \\citep{Blustin06}, and GRB\\,050801 \\citep{Rykoff06}. In this figure we use $z \\approx 1.0$ for GRB 060117. Observed $R$-band magnitudes are shown, except for GRB 050525a, where the $V$-band values are plotted.} \\end{center} \\end{figure} ",
        "conclusions": "\\label{Conclusions}  The GRB\\,060117 is the most intense (in terms of peak flux) GRB detected so far by Swift. With the maximum brightness of $R = 10.1$ mag, FRAM has discovered one of the optically brightest prompt optical emissions ever detected. The initial optical decay was found to be one of the steepest of an early GRB optical afterglow observed.  We have presented 3 scenarios for explaining the lightcurve. The apparent change in its slope is neglected in two simple scenarios, where we suppose the observed lightcurve to only be the trace of a forward, resp. reverse, shock. In the third (preferred) scenario, the shape of the lightcurve is explained as a transition between reverse and forward shock emission. The forward-shock-only interpretation is flawless regarding the PROMPT limit, but shows rather spurious value of $p$. The other two scenarios (i.e. those involving reverse shock) result in a relatively slowly decaying forward shock, and later limits require an early jet break $t_{\\mathrm j} \\sim 0.2$\\,d. The detailed analysis of this problem is beyond the scope of this letter.  Progress in further study of the particular case of GRB 060117 depends on the measurement of its distance. Therefore the follow-up and identification of host galaxy with a large telescope is of high importance.  A larger sample of GRB rapid follow-ups is needed to decide whether this kind of transition, already suggested for other bursts, is common."
    },
    "0606/astro-ph0606232_arXiv.txt": {
        "abstract": "We report the results of spectroscopic mapping observations carried out toward the Herbig-Haro objects HH7--11 and HH54 over the $5.2 - 37\\rm\\,\\mu m$ region using the Infrared Spectrograph of the {\\it Spitzer Space Telescope}.  These observations have led to the detection and mapping of the S(0) -- S(7) pure rotational lines of molecular hydrogen, together with emissions in fine structure transitions of Ne$^+$, Si$^+$, S, and Fe$^+$.  The H$_2$ rotational emissions indicate the presence of warm gas with a mixture of temperatures in the range $400 - 1200$~K -- consistent with the expected temperature behind nondissociative shocks of velocity $\\rm \\sim 10 - 20\\, km\\, s^{-1}$ -- while the fine structure emissions originate in faster shocks of velocity $\\sim 35 - 90\\, \\rm km \\, s^{-1}$ that are dissociative and ionizing.  Maps of the H$_2$ line ratios reveal little spatial variation in the typical admixture of gas temperatures in the mapped regions, but show that the H$_2$ ortho-to-para ratio is quite variable, typically falling substantially below the equilibrium value of 3 attained at the measured gas temperatures.  The non-equilibrium ortho-to-para ratios are characteristic of temperatures as low as $\\sim 50$~K, and are a remnant of an earlier epoch, before the gas temperature was elevated by the passage of a shock.  Correlations between the gas temperature and H$_2$ ortho-to-para ratio show that ortho-to-para ratios $< 0.8$ are attained only at gas temperatures below $\\sim 900$~K; this behavior is consistent with theoretical models in which the conversion of para- to ortho-H$_2$ behind the shock is driven by reactive collisions with atomic hydrogen, a process which possesses a substantial activation energy barrier ($E_A/k \\sim 4000$~K) and is therefore very inefficient at low temperature.  The lowest observed ortho-to-para ratios of only $\\sim 0.25$ suggest that the shocks in HH54 and HH7 are propagating into cold clouds of temperature $\\simlt 50$~K in which the H$_2$ ortho-to-para ratio is close to equilibrium. ",
        "introduction": "Stars have a profound effect upon the interstellar medium (ISM) that surrounds them.  Stellar radiation heats -- and sometimes photoionizes and photodissociates -- the interstellar gas, while supersonic outflows associated with stellar birth (in protostellar outflows) and death (in supernova remnants) send shock waves propagating through the ISM.  In regions of high mass star-formation -- such as the Orion Molecular Cloud, for example -- many of these processes operate simultaneously, making their separate effects hard to disentangle and elucidate.  In Herbig-Haro objects and supernova remnants, by contrast, the effects of shock waves can be studied in relative isolation in an environment where the stellar radiation field is not greatly enhanced.  Over more than three decades, the physics and chemistry of interstellar shock waves have been the subject of intensive study, both theoretical and observational, the results of which have been summarized in several review articles (e.g. Walmsley \\& Pineau Des For{\\^e}ts 2005; Draine \\& McKee 1993).  When shock waves propagate within molecular clouds, the shock velocity and magnetic field determine the fate of the gas.  Fast shocks (with propagation velocities $v_s \\simgt 40 \\,\\rm km \\, s^{-1}$ given interstellar magnetic fields of typical strength) are dissociative: they destroy molecules and may ionize the resultant atoms.  Slow shocks ($v_s \\simlt 40 \\,\\rm km \\, s^{-1}$), by contrast, are non-dissociative: they heat and compress the molecular gas through which they pass, and may modify its chemical composition, but molecules survive their passage.  The chemical changes wrought by non-dissociative shocks typically result from elevated gas temperatures or the phenomenon of ion-neutral drift (crucially important in non-dissociative shocks), which can drive chemical processes that are negligibly slow in cold clouds because they are endothermic or possess activation energy barriers.  The result of these processes depends critically upon the timescale for key chemical reactions relative to the period of time for which shock elevates the temperature and ion-neutral drift velocities. In a previous study of HH54 carried out with the {\\it Infrared Space Observatory (ISO)}, Neufeld et al.\\ (1998; hereafter NMH98) argued that the H$_2$ ortho-to-para ratio, observed to lie significantly below its equilibrium value ($\\sim 3$) in the warm ($T \\sim 650\\,$K) shocked gas in HH54, provided a valuable probe of the relevant timescales.  The Infrared Spectrograph (IRS) on board the {\\it Spitzer Space Telescope} is a powerful tool for the study of shocked interstellar gas.  With a wavelength coverage of 5.2 -- 37 $\\mu$m, {\\it Spitzer}/IRS provides access to the S(0) through S(7) pure rotational lines of H$_2$, as well as fine structure emissions of several atoms and ions produced in fast, dissociative shocks: Ne$^+$, Si$^+$, S, and Fe$^+$.  The great sensitivity of {\\it Spitzer}/IRS, together with the multiplex advantage inherent in long slit spectroscopy, allows large regions to be mapped rapidly.  In this paper we report the results of spectral line mapping in HH54 and HH$7-11$, two well-studied Herbig-Haro objects known to contain warm molecular gas.  HH 7--11 is a chain of Herbig-Haro objects in the Perseus molecular cloud, associated with NGC 1333 (Snell \\& Edwards 1981).  These objects lie in the blue-shifted part of a CO outflow that originates near the protostar SVS 13.  HH 7--11 has been studied extensively through observations of optical emission lines (Hartigan, Curiel, \\& Raymond 1989), near-IR emission lines of [Fe~II] (Stapelfeldt et al 1991; Gredel 1996), H$_2$ vibrational emissions (Noriega-Crespo et al. 2002, Khanzadyan et al.\\ 2003), far-infrared atomic and molecular lines (Molinari et al. 2000); and through high resolution studies of CO emission (Bachiller et al. 2000).   One key result to emerge from previous studies is that the hot H$_2$ emission exhibits a classic bow shaped morphology in HH~7, with an emission peak offset from the peak emission of atomic fine structure lines.   Observations of the HH7--11 flow have proven valuable in constraining detailed models for shock excited gas  (e.g. Hartigan, Raymond, \\& Hartmann 1987; Smith, Khanzadyan, \\& Davis 2003).  The distance to HH7-11 is a matter of debate; we adopt the value of 250 pc used by Enoch et al. (2006; see discussion in that paper which refers to \\u{C}ernis 1993; \\u{C}ernis \\& Strai\\u{z}ys 2003; Belikov et al. 2002). HH54 is located in the Chamaeleon II molecular cloud at an assumed distance of 200 pc (Hughes \\& Hartigan 1992), and is known to have a large column of warm ($>$ 2000 K) shocked molecular and atomic gas (Schwarz \\& Dopita 1980; Gredel 1994).  While it is less well studied than HH7-11, HH54 is notable because of an unusually large abundance of water vapor (Liseau et al. 1996) and H$_2$ ortho-to-para ratios that are below the equilibrium value expected at the observed gas temperature.  Our observations and data reduction methods are discussed in \\S 2 and Appendix A, and the results are presented in \\S 3.  The implications of these results for our understanding of interstellar shock waves are discussed in \\S 4.  A brief summary follows in $\\S 5$. ",
        "conclusions": "\\subsection{Fast and slow shocks}  In both sources that we observed, clear offsets between the fine-structure emissions and the molecular hydrogen emissions are evident.  Similar offsets have also been observed toward Herbig-Haro objects in the GGD37 (Cepheus A West) region  (Raines, 2000).  The fine structure emissions in HH7 are offset by $\\sim 5^{\\prime\\prime}$ from the bow shock, corresponding to a projected distance of $\\sim 2 \\times 10^{16} \\rm \\, cm$ for an assumed distance to the source of 250~pc. Such an offset has been observed previously in HH7 (Staplefeldt et al. 1991), where near-infrared emissions from [FeII] and H$_2$ have been mapped at subarcsecond resolution by Noriega-Crespo et al.\\ (2002) using the {\\it Hubble Space Telescope}.  Noriega-Crespo et al.\\ interpreted the [FeII] emissions as emerging from a ``Mach disk'', in which the collimated supersonic outflow is decelerated and heated in a fast shock.  In this picture, the fast shock is offset from a slower non-dissociative bow shock that is driven into the ambient molecular medium, the latter giving rise to the molecular hydrogen emissions observed from HH7.  Our mid-infrared observations indicate that the [SiII] and [SI] fine structure emissions -- like the [FeII] emissions -- also emerge from a region offset from the bow shock.  The morphology of the [SiII] map is most similar to the [FeII] map.  These observations support the two-shock picture in which the [FeII], [SI] and [SiII] fine structure emissions emerge from fast shock.  Our observations of HH54 indicate a similar offset ($\\sim 10^{\\prime\\prime}$, corresponding to a projected distance of $\\sim 3 \\times 10^{16} \\rm \\, cm$ for a source distance of 200~pc) between the fine structure and molecular hydrogen emissions.  As in HH7, the centroid of the fine structure emissions lies closer to the presumed source of the outflow than the H$_2$ emissions, consistent with the two-shock picture in which the fine structure emissions trace a fast shock that decelerates the collimated outflow.  We have compared the observed fine structure line intensities listed in Table 2 with predictions of the fast shock model of Hollenbach \\& McKee (1989, hereafter HM89).  Because the beam filling factor of shocked material is not known, line ratios provide a more valuable diagnostic than absolute intensities.   Although the predictions of the HM89 shock models are undoubtedly somewhat uncertain -- given substantial uncertainties in the collisional excitation rates for the various transitions and in the gas-phase abundances of the various elements -- the predicted [NeII]~12.8$\\,\\mu$m/[FeII]~26.0~$\\,\\mu$m line ratio is a strongly increasing function of the shock velocity, values $\\simgt \\rm 60\\,km\\,s^{-1}$ being needed to ionize Ne significantly.  The [NeII]~12.8$\\,\\mu$m/[FeII]~26.0$\\,\\mu$m line ratio of $\\sim 0.6$ observed in HH54 requires a shock velocity $\\sim \\rm 80 \\, km \\,s^{-1}$.  In HH7, by contrast, the absence of detectable [NeII]~12.8$\\,\\mu$m emission places an upper limit $v_s < \\rm 70\\,km\\,s^{-1}$, whilst our measurement of the [FeII]~17.9$\\,\\mu$m transition requires a shock velocity\\footnote{Considerably smaller shock velocities would be obtained in a comparison with the model predictions of Hartigan et al.\\ (1989), which apply to shocks in low density clouds that are assumed to be ionized prior to the passage of the shock.} of at least $35\\,\\rm km\\,s^{-1}$.  The HM89 shock models predict that the [FeII] and [SiII] line intensities show a very similar dependence upon the shock velocity, whilst the [SI]/[FeII] and [SI]/[SiII] line ratios are a decreasing function of shock velocity. This prediction provides a natural explanation for the fact that the [SiII] map is most similar to the [FeII] map: insofar as the fine structure emissions may arise from a superposition of shocks at varying velocities, the [SI] emissions would be expected to show a spatial distribution that is somewhat different from [SiII] and [FeII].  An independent check of the two-shock picture is provided by our detection of multiple [FeII] lines. % The mimimum temperature derived from the [FeII] line ratios ($\\sim 5000$~K) is larger than the maximum temperature at which H$_2$ can survive for any appreciable time period -- and is a factor $\\sim 5$ larger than that inferred from the H$_2$ line ratios (see \\S 4.2 below) -- lending strong support to the theory that both dissociative and nondissociative shocks are present in the source.  \\subsection{The temperature and column density of shocked molecular hydrogen}  The H$_2$ S(0) -- S(7) emissions provide a valuable probe of the temperature distribution in the shocked gas in which H$_2$ is the dominant constituent.  The observed line intensities for these H$_2$ quadrupole transitions provide a direct measurement of the column density, $N(J_u)$ in each upper state, $J_u$, because all the lines involved are optically-thin.  As in NMH98, we applied extinction corrections based upon the interstellar extinction curves of Weingartner \\& Draine (2001; ``Milky Way, $R_V=3.1$\").  For HH54, we adopted Gredel's (1994) estimate of 0.3 for $E(J-K)$, obtained from observations of the [FeII] 1.257 and 1.644 $\\mu$m transitions toward HH54E+K, whilst for HH7 we adopted the estimate $E(J-K) = 0.7$ obtained by Gredel (1996) for the nearby source HH8.  In Figure 11, based upon the line intensities tabulated in Table 2, we present rotational diagrams for each of the four circular regions for which we computed an average spectrum (\\S 3.2).  Two features are common to all observed positions.  First, as discussed further in \\S 4.4 below, the rotational diagrams all show the ``zigzag\" patterns characteristic of non-equilibrium ortho-to-para ratios, in which a fit to the para- (even-$J$) states lies above that for the ortho- (odd-$J$ states).  Second, the rotational diagrams all exhibit a degree of curvature, with the slope of the rotational diagram decreasing with increasing $J$.  This curvature indicates that each beam is sampling material at a range of temperatures.  Furthermore, the absence of any turnover in the curvature at high $J$ suggests that the rotational level populations are close to local thermodynamic equilibrium (LTE) even for $J$ as high as 9.  For each circular region in Table 2, we determined an excitation temperature, $T_{XY}$, from the ratio of the $J=X$ and $J=Y$ column densities and defined according to $$T_{XY} = {E_{Y} - E_{X}  \\over k\\, {\\rm ln}\\,[N_{X} g_{Y}/ N_{Y} g_{X}]}$$ where $E_J$ and $g_J$ are the energy and degeneracy of state $J$. Because the ortho-para ratio shows departures from equilibrium, the quantity $T_{XY}$ is most useful when $X$ and $Y$ are either both odd or both even. The excitation temperatures $T_{42}$, $T_{53}$, $T_{64}$, $T_{75}$, $T_{86}$ and $T_{97}$ are given in Table 3;  the inequalities $T_{42} < T_{64} < T_{86}$ and $T_{53} < T_{75}  < T_{97}$ observed in every case are a simple manifestation of the curvature apparent in the rotational diagrams.  We have obtained two-component fits to the rotational diagrams, in which a least squares fit was obtained for a combination of warm and hot gas components, assumed to be in LTE except with non-LTE ortho-to-para ratios, with temperatures  $T_w$ and $T_h$ (typically found to be $\\sim 400$~K and $\\sim 1000$~K, respectively). The solid lines in Figure 11 show the two-component fits to the rotational diagrams, while the points indicate the measured values.  The two-component fits provide an excellent fit to the data, although the solution is clearly non-unique.  Indeed, there is no specific evidence for a bimodal temperature distribution, and the data could be fit equally well by models that invoke a continuous range of gas temperatures. The assumption of LTE is clearly valid for the lower rotational states, but the data would allow modest departures from LTE for the highest states that we have observed. Indeed, calculations by Le Bourlot et al.\\ (1999) indicate that the critical density for $J=9$ (at which the probabilities of collisional and radiative de-excitation are equal) is $\\sim 4 \\times 10^5 \\rm \\, cm^{-3}$, a factor of several larger than the densities inferred from observations of CO in some of these sources (see discussion below). Departures from LTE, if present, would be an increasing function of the rotational quantum number, $J$, and would therefore require higher gas temperatures than the values inferred under the assumption of LTE.  The best-fit temperatures, computed under the assumption of LTE, together with the beam-averaged column densities ($N_w$ and $N_h$) and the H$_2$ ortho-to-para ratios (OPR$_w$ and OPR$_h$), are listed in Table 3.  Table 3 also shows the sum of the $N(4)$ though $N(9)$ column densities, $N_{4-9}$, all of which are derived directly from observations with the Short-Low module; not surprising, $N_{4-9}$ falls a factor of several below $(N_w + N_h)$ because it only includes states with $J \\ge 4$. The estimated gas temperatures, $T_w$ and $T_h$, are characteristic of low-velocity nondissociative shocks (which are of ``C\"-type in the designation of Draine (1980); see also Appendix B).  By means of a simple analytic treatment of energy conservation in such shocks (Appendix B), constrained by the results of the detailed shock models of Kaufman \\& Neufeld (1996; hereafter KN96), we find the characteristic gas temperature within such shocks to be given by $$T_s = 375\\, b^{-0.36} v_{s6}^{1.35}\\, \\rm K$$ (eqn.\\ B8), where $10\\,v_{s6}\\, \\rm km \\,s^{-1}$ is the shock velocity and $b (n_0 / {\\rm cm}^{-3})^{1/2} \\, \\rm \\mu G$ is the assumed preshock magnetic field for a shock propagating in material of preshock H$_2$ density $n_0$.  Given our standard assumption about the magnetic field, $b=1$, the warm and hot gas components typically correspond to shock velocities of $\\sim 10$ and $\\sim 20$ $\\rm km \\, s^{-1}$ respectively.  The H$_2$ rotational diagrams predicted by KN96 and other theoretical shock models (e.g. Wilgenbus et al.\\ 2000) all exhibit less curvature than our observed rotational diagrams, implying the presence of a range of shock velocities within each of the circular apertures demarked in Figure 7.  The preshock H$_2$ density, $n_0$, is less well constrained by the H$_2$ line fluxes.  The best estimates for the density of the warm shocked gas have been obtained by observations of high-lying rotational lines of CO, for which non-LTE line ratios provide a valuable probe.  Based upon {\\it ISO} observations of such transitions toward HH54, Liseau et al.\\ (1996) estimated the H$_2$ density in the warm shocked gas as $1 - 4 \\times 10^5 \\rm \\, cm^{-3}$.  Since most of the high-$J$ CO emission occurs before the compression of the neutral species becomes large, the corresponding preshock density is only a factor $\\sim 1.5$ smaller\\footnote{Liseau et al.\\ significantly overestimated the compression of the neutral material, thereby obtaining a significant underestimate of the preshock density required to match the CO line ratios.} (Appendix B, eqn.\\ B6): $n_0 = 6 - 25 \\times 10^4 \\rm \\, cm^{-3}$.  Similar  values were obtained for HH7--11, based upon measurements of CO line ratios (Molinari et al.\\ 2000), although measurements of the CO/H$_2$ line ratios argued for smaller values ($n_0 \\sim 10^4 \\rm \\, cm^{-3}$) toward HH7.  The observed column density of shocked H$_2$ is predicted to be $$N_s = 4 \\times 10^{20} \\, b\\, n_{04}^{0.5}  \\, v_{s6}^{-0.75}\\, \\Phi_s,$$ (Appendix B, eqn.\\ B7), where $\\Phi_s$ is a geometric factor defined to be the ratio of the projected beam size to the surface area of the shock. In principle, the quantity $\\Phi_s$ can be either larger or smaller than unity.  For a planar shock with a beam covering factor of unity that is viewed at angle $\\theta$ to the shock normal, the observed value of $\\Phi_s$ is ${\\rm sec}\\,\\theta \\ge 1$.  On the other hand, $\\Phi_s$ can be smaller than unity when the beam is incompletely covered by shocked material.  For our two component fits to the circular apertures presented in Table 3, the values of $\\Phi_s$ required for the warm and hot components are typically $0.2\\,n_{04}^{-0.5}$ and $0.05\\,n_{04}^{-0.5}$, respectively, implying that the shocked material has a covering factor considerably smaller than unity and possesses substructure that is not resolved by these observations.  In HH7, the subarcsecond resolution maps of near-IR H$_2$ vibrational emissions presented by Noreiga-Crespo et al.\\ (2002) clearly indicate that $\\Phi_s \\ll 1$ for observations with the IRS instrument.  The spectral line maps obtained for HH54 and HH7 also permit us to obtain a rotational diagram at each spatial position, and thereby to derive detailed maps of the gas temperature, ortho-to-para ratio, and column density of the shocked gas. While the average spectra presented in Figure 8 are of sufficient signal-to-noise ratio (SNR) to permit the two-component fits described above, the lower SNR obtained in a single spectrum does not warrant a two-component fit to the rotational diagram.  Accordingly, in constructing maps of the gas temperature, we obtained a least-squares fit to the $N(4)$ through $N(7)$ column densities (all measured with the SL module) for a {\\it single} gas component with a temperature and H$_2$ ortho-to-para ratio, maps of which are presented in Figure 12 (middle and lower panels respectively).  Here, the white contours show the H$_2$ S(3) line intensity, and regions of low SNR have been masked and appear in black.  The upper panel in Figure 12 shows the sum of the $N(4)$ though $N(9)$ column densities, $N_{4-9}$. Figure 12 shows that the fitted gas temperature is remarkably constant over the mapped region, lying generally within the range $900 \\pm 200$~K, and indicating that a similar admixture of shock velocities is present along each line-of-sight.  The ortho-to-para ratio, however, shows substantial variations -- as noted previously by Wilgenbus et al.\\ (2001) -- and is discussed further in \\S 4.4   \\subsection{Upper limits on the H$_2$O abundance}  The increase in temperature and density behind a shock front can alter significantly the composition of the gas in its wake. Specifically, chemical reactions not favored at temperatures characteristic of the preshock quiescent material (i.e., $T\\,\\simlt\\,$50~K) can proceed rapidly in the warm post-shocked gas.  One such example are the reactions that lead to the production of gas-phase \\water.  In the cold preshock gas, cosmic-ray driven ion-neutral reactions are capable of converting O and H into water yielding a peak water abundance relative to H$_2$, $x({\\rm H_2O})\\,\\equiv \\,n({\\rm H_2O})/n({\\rm H_2})$, of $\\sim\\,$10$^{-6}$ in approximately 10$^4$ years (Bergin et al.\\ 1998). When the gas temperature exceeds $\\sim\\,$300~K, a series of neutral-neutral reactions, $\\rm H_2 + O \\rightarrow OH + H$ and $\\rm H_2 + OH \\rightarrow H_2O + H$, dominate and chemical models (e.g., Elitzur \\& de Jong 1978; Neufeld et al.\\ 1995) predict that most of the gas-phase oxygen not bound as CO will be processed into water vapor. These reactions are rapid in warm gas, producing $x({\\rm H_2O})\\,\\sim\\,$2$\\times$10$^{-4}$ in less than a few hundred years (Bergin et al.\\ 1998), but are negligibly slow at low temperatures because they possess significant activation barriers.  Observations of both HH54 and HH7 have been conducted using ESA's {\\em Infrared Space Observatory (ISO)}.  {\\it ISO}'s Long Wavelength Spectrometer (LWS) was used in its grating mode to acquire low resolution ($\\lambda/\\Delta\\lambda \\sim\\,$200) spectra between 43 and 197\\um\\ within an 80$^{\\prime\\prime}$ diameter beam.  Toward HH54, Liseau \\etal\\ (1996) detected a number of low-lying (i.e., $E/k\\,\\simlt\\,$300~K above the ground state) \\water\\ and OH lines, along with CO emission lines from the rotational upper level $J\\,=\\,$14 to 19.  These data are best fit by a model that assumes collisional excitation behind a 10-15~km~s$^{-1}$ nondissociative $C$-type shock and $x({\\rm H_2O})\\,=\\,$10$^{-5}$.  Toward HH7, Molinari et al.\\ (2000) find evidence for the presence of both fast dissociative $J$-type and slower $C$-type shocks.  In particular, Molinari et al.\\ find that a 15-20 km~s$^{-1}$ $C$-type shock accounts well for the observed strengths of the CO $J\\,=\\,$14--13 to $J\\,=\\,$17--16 transitions plus the H$_2$ (0-0) S(0)--S(5) transitions, which were measured with the {\\it ISO} Short Wavelength Spectrometer.  Assuming that the {\\it ISO}-detected \\water\\ 179.5\\um\\ 2$_{12}-$1$_{01}$ and 174.6\\um\\ 3$_{03}-$2$_{12}$ lines arise predominately from behind the same $C$-type shock, Molinari \\etal\\ derive $x({\\rm H_2O})\\,\\leq\\,$6$\\times$10$^{-6}$.  Molinari et al.\\ suggest that the lower-than-expected post-shock \\water\\ abundance is due to the rapid condensation of water onto dust grains in the post-shocked gas.  A similar effect is observed toward IC$\\:$443, in which a low water abundance (2$\\times$10$^{-8}\\,<\\,x({\\rm H_2O})\\,<\\,$3$\\times$10$^{-7}$) is inferred\\footnote{Snell et al.\\ expressed their findings in terms of the ratio $n({\\rm o-H_2O})/n({\\rm CO})$.  Here we make the assumption that $n({\\rm CO})/n({\\rm H_2}) = 10^{-4}$.} in the presence of both $C$- and $J$-type shocks (Snell et al.\\ 2005). Snell et al.\\ ascribe the low water abundance to a combination of two effects. First, more than half of the preshock oxygen not in CO remains trapped as water-ice on grain surfaces despite the passage of a weak $C$-type shock, suppressing the production of gas-phase \\water\\ in the post-shocked gas.  Second, ultraviolet radiation generated by the nearby fast $J$-type shock photodissociates much of the gas-phase water produced behind the $C$-type shock.  Eight H$_2$O transitions with upper state energies $E_U/k < 1000$~K lie within the wavelength range accessible to {\\it Spitzer}/IRS. Prior to the measurements reported here, none of these transitions had been observationally investigated toward either HH54 or HH7. To assess the extent to which the upper limits to the \\water\\ line strengths set here constrain the water abundance, the equilibrium level populations of all ortho and para rotational levels of the ground H$_2^{\\;16}$O vibrational state with energies {\\em E/k} up to 7700~K have been calculated using an escape probability method described by Neufeld \\& Melnick (1991).  Only two updates have been made in the calculations undertaken here: (1) for collisionally induced transitions among the lowest 45 levels of ortho- and para-water, the rate coefficients calculated by Green, Maluendes, \\& McLean (1993) for He-H$_2$O collisions were used (multiplied by 1.348 to account for the different reduced mass when H$_2$ is the collision partner), and (2) the expression for the photon escape probability used by Neufeld \\& Kaufman (1993) was adopted.  We obtain level populations for a plane-parallel emitting region in which the limit of large velocity gradient applies.   A two-temperature component model is used, consistent with the results of our H$_2$ measurements (which are summarized in Table~3).  In addition, preshock densities of 10$^4$ and 10$^5$~\\cmc\\ are considered.  The H$_2$ column densities and velocity gradients are assumed to be those consistent with the simple shock model presented in Appendix~B, i.e.:  \\vspace{2.5mm}  \\begin{displaymath} N(H_2) = 7.9 \\times 10^{20}\\;\\left[\\frac{n(H_2)}{10^5}\\right]^{0.5}\\, \\left[\\frac{T_{gas}}{1000}\\right]^{-0.555}~~~{\\rm cm^{-2}} \\end{displaymath}  \\vspace{-2mm}  and  \\vspace{-2mm}  \\begin{displaymath} \\frac{dv}{dz} = 1.6 \\times 10^4\\;\\left[\\frac{n(H_2)}{10^5}\\right]^{0.5}\\, \\left[\\frac{T_{gas}}{1000}\\right]^{1.30}~~~{\\rm km~s^{-1}~pc^{-1}}, \\end{displaymath}  \\vspace{2.5mm} \\noindent where $n(H_2)$ is the preshock gas density and $T_{gas}$ is the gas temperature.  Finally, we assume that the fraction of the beam filled by the warm and hot \\water\\ is given by the ratio of the measured values of $N({\\rm H}_2)$ (given in Table~3) to those derived above.  For these conditions, the strongest water lines within the {\\em Spitzer}/IRS wavelength range are predicted to be the 29.837\\um\\ 7$_{25}$-6$_{16}$ and 30.899\\um\\ 6$_{34}$-5$_{05}$ ortho-\\water\\ transitions.  Table~4 lists the {\\em Spitzer}-measured 3$\\sigma$ upper limits to the line intensities for each transition along with the derived upper limits to the gas-phase water abundance.   Several things are apparent.  First, the upper limits to the \\water\\ abundance given in Table~4 are higher than, but nonetheless consistent with the water abundances derived previously from the {\\it ISO} data.   In particular, for a preshock density of 10$^5$~\\cmc\\ -- favored by the H$_2$ data toward HH54 (Liseau \\etal\\ 1996) -- the inferred abundance limits are less than would be expected from the full conversion of elemental oxygen not locked in CO into \\water.  Unfortunately, the present 29.837 and 30.899\\um\\ line flux limits are insufficient to allow us to comment further upon the previous HH54 models. Second, the \\water\\ 174 and 179\\um\\ line fluxes measured by Liseau \\etal\\ provide an additional constraint on $x({\\rm H_2O})$ within the warm and hot components discussed here.  In particular, the {\\it ISO}-measured flux in the \\water\\ 174\\um\\ line, $\\sim\\,$3$\\times$10$^{-20}$ W~cm$^{-2}$, requires that $x({\\rm H_2O})\\,<\\,$2$\\times$10$^{-5}$ in the warm and hot H$_2$ components or these components will produce a 174\\um\\ line flux in excess of that measured. This result is in agreement with the water abundance determined by Liseau et al.\\ of 10$^{-5}$ and suggests that the size of the \\water\\ 174 and 179\\um\\ emitting region is larger than that inferred for the {\\em Spitzer}-measured H$_2$ emission (or that assumed for the mid-infrared \\water\\ emission).  This is plausible since the 174 and 179\\um\\ lines probe cooler gas than either the H$_2$ or mid-infrared \\water\\ lines.  Toward HH7, Molinari et al.\\ estimate the preshock density to be $\\sim\\,$10$^4$ cm$^{-3}$. For this density, rather high H$_2$O abundances are required to reproduce the 29.837 and 30.899\\um\\ line intensity limits. However, the strength of the \\water\\ 179\\um\\ line observed by Molinari et al.\\ does set a more stringent upper limit of $x({\\rm H_2O})\\,<\\,$6$\\times$10$^{-4}$ within the warm and hot components measured here.  Unfortunately, this limit remains significantly greater than the water abundance inferred from the {\\it ISO} measurements and, thus, does not permit a test of their model.     \\subsection{The H$_2$ ortho-to-para ratio}  The H$_2$ ortho-to-para ratios (Table 3) derived for the two-component fits, OPR$_w$ and OPR$_h$, are uniformly less than the local thermodynamic ratio, $\\rm OPR_{\\rm LTE}\\sim 3$ at the temperatures of relevance.  This non-equilibrium behavior, first observed in the atmospheres of giant planets (e.g. Conrath \\& Gierasch 1983), was noted in {\\it ISO} observations of HH54 (NMH98; Wilgenbus et al.\\ 2001) and subsequently observed in several other interstellar regions (e.g.\\ Fuente et al.\\ 1999, Rodr{\\'{\\i}}guez-Fern{\\'a}ndez et al.\\ 2000).  As discussed in NMH98, the H$_2$ ortho-to-para ratio equilibrates very slowly, exhibiting ``memory effects\" in which the OPR can reflect the temperature of a previous epoch in its thermal history\\footnote{Such effects are a well-known problem in the industrial production and storage of liquid hydrogen (LH2).  The rapid refrigeration of molecular hydrogen leads to the production of LH2 with an ortho-to-para ratio $\\sim 3$, reflecting the ``frozen-in\" equilibrium value at room temperature, and greatly exceeding the value of 0.02 obtained at equilibrium at 21~K, the boiling point of LH2.  Equilibration occurs over a timescale of 6.5 days, resulting in the eventual release of a total heat content that exceeds the latent heat of vaporization of LH2: without continued refrigeration, the LH2 simply boils away!  The solution to this problem involves the introduction of a paramagnetic catalyst to facilitate ortho-to-para conversion during the original liquification process and thereby prevent the room temperature OPR from being frozen-in.}.  Thus NMH98 argued that the OPR of $1.0 \\pm 0.4$ observed toward HH54 indicated (1) that the warm ($T \\sim 650$~K) shocked H$_2$ that they observed in HH54 must have previously acquired an OPR characteristic of equilibrium at a temperature $\\le 90$~K and (2) that the observed gas has not been warm long enough for a high-temperature OPR to have been established.  The equilibration of ortho- and para-H$_2$ has been considered in the detailed shock models of Timmermann (1998) and Wilgenbus et al.\\ (2000).  Both studies concluded that reactive collisions with atomic hydrogen are the dominant para-to-ortho conversion process in non-dissociative molecular shocks: $$\\rm \\hbox{para-H}_2 + H \\rightarrow H + \\hbox{ortho-H}_2 \\eqno(R1)$$ Because reaction (R1) has a substantial activitation energy barrier ($\\Delta E_A/k \\sim 3900$~K), the efficiency of para-to-ortho conversion is a strong function of temperature (and, consequently, shock velocity).  This strong temperature dependence is entirely consistent with the results in Table 3, which show that for each region the hot gas component has a much higher OPR than the warm gas component.  This behavior is also apparent in the H$_2$ rotational diagrams (Figure 11), in which the departures from LTE (i.e.\\ the degree of zigzag) are plainly larger for small $J$ than for large $J$.  In Figure 13, the ortho-to para ratios, OPR$_w$ and OPR$_h$, for the two-component fits in Table 3, are plotted as a function of the fitted gas temperatures, $T_w$ and $T_h$ (open triangles).  We have also obtained two-component fits to the gas temperature and OPR elsewhere in the mapped regions.  Averaging the observed spectra in a set of $5^{\\prime\\prime} \\times 5^{\\prime\\prime}$ subregions to attain adequate SNR to yield meaningful two-component fits, we determined OPR$_w$, OPR$_h$, $T_w$ and $T_h$ estimates for each such subregion.  Green and red squares in Figure 13 indicate respectively the ($T_w, $OPR$_w$) and ($T_h$, OPR$_h$) values obtained in subregions within the HH54 and HH7 maps.  Green and red crosses apply to the discontiguous regions northwest of HH7 (Figure 12; these are close to HH8).  The orange curve denotes the ortho-to-para ratio expected in LTE.  The green and red symbols in Figure 13 confirm that the ortho-to-para ratio tends to increase with gas temperature.  The black curves show the expected behavior for a simple model in which all the observed gas started with the same initial ortho-to-para ratio -- OPR$_0 = 0.4$ (corresponding to equilibrium at temperature 55~K) for HH54 and OPR$_0 = 0.25$ (corresponding to equilibrium at temperature 48~K) for HH7 -- and has been at the currently-observed temperature $T$ for the same time period, $\\tau$.  In that case, the current OPR is given by $${\\rm OPR(\\tau) \\over 1+OPR(\\tau)} =  {\\rm OPR_0 \\over 1+OPR_0} \\, e^{-n({\\rm H})\\, \\tau \\, k_{\\rm PO}} +  {\\rm OPR_{LTE} \\over 1+OPR_{LTE}} \\, \\biggl( 1- e^{-n(H)\\, \\tau \\, k_{\\rm PO} }\\, \\biggr),$$ where $n({\\rm H})$ is the atomic hydrogen density and $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). The black curve applies to the case $n_2 ({\\rm H}) \\tau = 120$~yr, where $n_2 ({\\rm H}) = n({\\rm H})/10^2 \\rm  \\, cm^{-3}$. If, following NMH98, we assume that H atoms are produced primarily by reaction of O with H$_2$ to form OH, followed by reactions of OH with H$_2$ to form water, we expect an $n({\\rm H})/n({\\rm H}_2)$ ratio $\\sim 10^{-3}$, corresponding to $n_2 ({\\rm H}) \\sim 1$ for a preshock H$_2$ density $n_0 = 10^5 \\rm \\, cm^{-3}$. The steep rise in the black curve at $T \\sim 1000$~K is a consequence of the strong temperature-dependence of $k_{\\rm PO}$.  To within the likely errors, all the points lie between the black and orange curves in Figure 13.  This behavior then implies that all the shocked gas had an initial OPR between 0.4 and 3 and has been warm for at least $150\\,n_2({\\rm H})^{-1}$~yr.  The foregoing analysis is premised on the assumption that the rotational temperature inferred for the hot gas component is equal to the kinetic temperature.  As discussed in \\S4.2, above, subthermal excitation of the S(7) transition is allowed by the data and would mean that the actual gas temperature is higher than the rotational temperature for the hot component, $T_h$.  Such effects, in turn, would decrease the required timescale for ortho-para conversion.  For a postshock atomic H abundance $\\sim 10^{-3}$, the timescale of $150\\,n_2({\\rm H})^{-1}$~yr inferred above is broadly consistent with the time period for which gas temperature is elevated upon passing through a nondissociative shock.  For $n_0 = 10^5 \\rm \\, cm^{-3}$ and $v_s = 15\\,\\rm  km\\,s^{-1}$, the simple treatment of Appendix B implies a shock column density, $N_s \\sim 10^{21} b\\, \\lambda_{in}\\rm cm^{-3}$, corresponding to a timescale $N_s/(n_0 \\, v_s) =  200\\,b\\, \\lambda_{in}$~yr.  However, the OPR values plotted in Figure 13 are entirely inconsistent with a shock model with a large H abundance, as has been invoked by Smith et al.\\ (2003), who proposed an $n({\\rm H})/n({\\rm H}_2)$ ratio of 3 for the bow shock in HH7.  Such a large H abundance would lead to very rapid para-to-ortho conversion, producing OPR values significantly in excess of those observed here.  We may also compare the results plotted in Figure 13 with the predictions of detailed theoretical models for para-to-ortho conversion in non-dissociative molecular shocks.  The cyan loci in Figure 13 represent the predictions of Wilgenbus et al.\\ (2000) for a shock propagating in gas with an initial OPR of 0.01. The loci correspond to preshock H$_2$ densities of $10^3$, $10^4$, 10$^5$ and $10^6\\, \\rm cm^{-3}$, respectively from left to right, and the asterisks apply to shock velocities of 30 and 40 $\\rm km \\, s^{-1}$ ($10^3 \\, \\rm cm^{-3}$ case, from lower left to upper right); 20, 30, and 40 $\\rm km \\, s^{-1}$ ($10^4$ and $10^5 \\, \\rm cm^{-3}$ cases, from lower left to upper right); or 10, 15, 20, 25, and 30 $\\rm km \\, s^{-1}$ ($10^6 \\, \\rm cm^{-3}$ case, again from lower left to upper right).  As in the case of the toy model (black curve), the curves for 10$^5$ or $10^6 \\, \\rm cm^{-3}$ are a reasonable approximation to the lower envelope of the observed data points.  The absence of sight-lines with high temperature and low OPR is consistent with shock models in which para-to-ortho conversion is driven by reactive collisions with atomic hydrogen.  On the other hand, the presence of sight-lines with {\\it low} temperature and {\\it high} OPR suggests that some of the shocked material has a high {\\it initial} OPR.  Such material is clearly present for the sight-lines represented by the crosses in Figure 13.  These regions lie behind the ``working surface\" of the jet in HH7--11; they may, perhaps, have been heated by relatively fast shocks during a previous epoch of shock activity, thereby acquiring a larger OPR, and are currently being reheated by a slower shock.  The bow shock in HH7, by contrast, appears to be propagating into fresh material in which the OPR is $\\simlt 0.25$, corresponding to equilibrium at temperature $\\sim 50$~K."
    },
    "0606/astro-ph0606226_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ Future neutrino detector at Megaton mass might enlarge the neutrino telescope thresholds revealing cosmic supernova background and largest solar flares neutrino. Indeed the solar energetic ($ E_p> 100 MeVs $) flare particles (protons, $\\alpha$) while scattering among themselves  on Solar corona atmosphere must produce prompt charged pions, whose chain decays are source of  a solar (electron-muon) neutrino \"flare\" (at tens or hundreds MeV energy). These brief (minutes) neutrino \"burst\" at largest flare peak may overcome by three to five order of magnitude the steady atmospheric neutrino noise on the Earth, possibly leading to their detection above detection thresholds (in a full mixed  three flavor state). Moreover the birth of $anti-neutrinos$ at a few tens MeVs is well loudly flaring above a $null$ thermal \"hep\" anti-neutrino solar background and also above a tiny supernova relic and atmospheric noise. The largest prompt solar anti-neutrino \"burst\" may be well detected in future Super Kamikande (Gadolinium implemented) anti-neutrino $\\overline{\\nu_e}$ signatures mostly in inverse Beta decay ($\\overline{\\nu_e} + p \\rightarrow n + e^+$) and rarely in higher energy $muon$ , or even rarest $tau$ neutrino leptons. Our estimate for the recent and exceptional October - November $2003$ solar flares and January $20th$ 2005 exceptional flare might lead to a few events above or near unity for existing Super-Kamiokande and above unity for Megaton detectors. The $\\nu$ spectra may reflect in a subtle way the neutrino flavor oscillations and mixing in flight.  A comparison of the solar neutrino flare (at their birth place on Sun and after oscillation on the arrival on the Earth) with other neutrino foreground is estimated: it offers an independent track to disentangle the neutrino flavor puzzles and its most secret mixing angles. The sharpest noise-free anti-neutrino imprint maybe its first clean voice. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606010_arXiv.txt": {
        "abstract": "We outline a novel satellite mission concept, DEMON, aimed at advancing our comprehension of both dark matter and dark energy, taking full advantage of two complementary methods: weak lensing and the statistics of galaxy clusters. We intend to carry out a 5000 $\\rm deg^{2}$ combined IR, optical and X-ray survey with galaxies up to a redshift of z $\\sim$ 2 in order to determine the shear correlation function. We will also find $\\sim 100000$ galaxy clusters, making it the largest survey of this type to date.  The DEMON spacecraft will comprise one IR/optical and eight X-ray telescopes, coupled to multiple cameras operating at different frequency bands. To a great extent, the technology employed has already been partially tested on ongoing missions, therefore ensuring improved reliability. ",
        "introduction": "\\label{sec:introduction}  The open question concerning the nature of dark matter and dark energy is certainly one of the most compelling issues of modern astrophysics, that demands new and better observations in order to be answered. Until the physical properties of dark matter can be determined experimentally (something that cannot be done in the laboratory up to now), trying to improve our knowledge about these properties implies improving the current measurement of the cosmological parameters. In particular, in the case of dark energy, the first crucial step is to try to distinguish a cosmological constant model from time dependent quintessence models.  The current best constraints on the cosmological parameters come from the combined results of three different techniques with complementary degeneracies: SNIa experiments, CMB measurements, and the study of galaxy clusters.  In addition to these well known methods, weak lensing is emerging as one of the most powerful and robust probes of the large scale structure of the universe and the nature of dark matter and dark energy\\cite{mellier, bartelmann, refrieger}.  Several existing and proposed wide-field weak lensing surveys, ground-based\\footnote{VISTA (Visible \\& Infrared Survey Telescope for Astronomy), Pan-STARRS (Panoramic Survey Telescope And Rapid Response System), LSST (Large Synoptic Survey Telescope), CFHTLS (Canada-France-Hawaii Telescope Legacy Survey), KIDS (Kilo-Degree Survey) and DES (Dark Energy Survey)} and space borne\\footnote{JDEM (Joint Dark Energy Mission) from NASA-DOE and DUNE (Dark UNiverse Explorer) from CNES, also proposed for ESA's Cosmic Vision} are currently being developed. In order to contribute to the design of possible space missions we present DEMON (Dark Energy and Matter Observational Nexus), a novel concept satellite conceived to study the properties of dark matter and dark energy combining two complementary techniques: weak lensing and clusters statistics. ",
        "conclusions": "\\label{sec:conclusions}  This paper presents a new concept mission for a satellite aimed at provide better constraints on the cosmological parameters and, in particular, studying the evolution of the dark energy equation of state with time.  The innovate concept of the DEMON spacecraft is to carry out a combined optical/IR and X-ray survey of 5000 $\\rm deg^{2}$ to perform weak lensing and cluster statistics studies. The satellite is meant to be complemented by one of the ground-based wide-field surveys currently under development, to provide the best accuracy in the photometric redshift measurements required.  For the optical/IR part of the survey, we investigated telescope design similar to the one proposed by SNAP.  To sample the FWHM PSF with $2\\times2$ pixels (something crucial for weak lensing studies), the proposed telescope requires a pixel size of 5 $\\mu m$. This technology is currently under development and is expected to be available within 5 years (J. W. Beletic, director of sensor systems of Rockwell Scientific, private communication).  An alternative possibility consist on use micro-dithering in the pointing strategy and include an additional magnification of 1.5 in the optical system. This configuration allows in principle to sample the FWHM PSF with $2\\times2$ pixels using CCDs of 10 $\\mu$m. However, the derived consequences in the technical telescope design and image quality should be investigated in detail to determine the feasibility of this option.  For the X-ray part of the mission, we explore the possibility to include 8 X-ray telescopes in the spacecraft, taking advantage of the technology developed for the eROSITA mission.  The result is an X-ray survey with a flux limit of $8\\times10^{-15}$ erg cm$^{-2}$ s$^{-1}$ (10 times more sensitive than DUO\\cite{duo} and 5 times more sensitive than eROSITA) that will identify all the clusters with a mass $\\gtrsim$ $2\\times 10^{14}$ $\\rm M_\\odot$ up to $z \\sim 1.5$. Therefore, the DEMON $5000 \\rm deg^{2}$ X-ray survey is expected to produce a catalog of $\\sim 100000$ clusters with a mean redshift of 0.75.   The final spacecraft is 7.1 m high and 3.75 m wide (at its broadest parts). It has an estimated mass of 2500 kg, which is well within the transport capabilities of a Soyuz rocket into the desired LEO orbit. The estimated cost of the spacecraft is 510 M euro~(see Table~\\ref{table:costs} for details).  The DEMON concept design shows that combining an optical and X-ray survey in the same spacecraft is technically possible and very interesting in terms of costs. In addition, it is able to provide high quality data for both weak lensing and cluster statistics, something that is not possible with the weak lensing missions proposed by SNAP\\cite{snap} and JEDI\\cite{jedi} due to their poor sampling of the FWHM PSF.  We also note that, although DEMON is optimized for a weak lensing survey, the high quality of the observations and the wide sky coverage of the survey will provide valuable data for other studies, like strong lensing systems, baryonic wiggles, galactic structure or detailed studies of clusters of galaxies. Due to the characteristics of the X-ray instruments, DEMON will also contribute to the study of the intracluster medium and galaxies with active nuclei (AGNs). Furthermore, a number of faint objects will be identified as future targets for narrow field instruments performing pointed imaging observations, like XEUS.  Finally, the planned far-infrared surveys based on the Sunyaev-Zel'dovich (SZ) effect (e.g. the ESA Planck mission and South Pole Telescope) will produce the redshift dependent cluster mass function and the power spectrum with different systematics compared to the DEMON X-ray survey, permitting for a very important cross-check of the two cosmological techniques to be made."
    },
    "0606/astro-ph0606683_arXiv.txt": {
        "abstract": "We present a self-contained description of everything needed to write a program that calculates the CMB power spectrum for the standard model of cosmology ($\\lcdm$). This includes the equations used, assumptions and approximations imposed on their solutions, and most importantly the algorithms and programming tricks needed to make the code actually work. The resulting program is compared to CMBFAST and typically agrees to within $0.1 \\,\\%$ -- $0.4 \\,\\%$. It includes both helium, reionization, neutrinos and the polarization power spectrum. The methods presented here could serve as a starting point for people wanting to write their own CMB program from scratch, for instance to look at more exotic cosmological models where CMBFAST or the other standard programs can't be used directly. ",
        "introduction": "Since the first detection of anisotropies in the Cosmic Microwave Background (CMB) by the Cosmic Background Explorer (COBE) satellite in the early 90's \\cite{COBE}, there has been considerable activity in this field of cosmology throughout the world. With the more accurate measurements of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP_firstyear, WMAP_threeyear} and the future Planck satellite \\cite{Planck}, we are entering the era of precision cosmology. Results from different kinds of experiments seem to converge to what is being referred to as the Standard Model of Cosmology \\cite{StandardModel}, describing the history of the universe from inflation through Big Bang Nucleosynthesis to the release of the CMB radiation during recombination. With the growing precision of the observational data comes also the need for fast and accurate theoretical calculations of CMB power spectra. Often one seeks the best fit to observations of a model with several parameters, requiring typically hundreds of spectra to be calculated and compared to the data.  Of course, several standard computer programs that calculate CMB power spectra are already available. The most commonly used include CMBFAST \\cite{CMBFAST, CMBFAST_web}, CMBEASY\\- \\cite{CMBEASY}, and CAMB \\cite{CAMB}. These are all excellent programs when calculating power spectra for the $\\lcdm$ model, including also some extensions like a quintessence field, hot dark matter or a simple fifth dimension. One may therefore wonder what the point is writing a new program from scratch, a task which obviously requires quite a lot of work. The need may arise when considering more exotic cosmological models which are not included in any of the standard programs. This could for instance include quintessence with a non-trivial coupling to other fields, extra dimensions with non-trivial geometry, or a model with varying constants of nature. In these cases one has the choice of either extending existing code, or writing a completely new program. The former can certainly be the easiest solution in some situations, but not necessarily all. The problem with updating existing code is that one probably doesn't know exactly how it works, and it is therefore difficult to make changes without doing something wrong. By instead writing new code, starting with a simple model and comparing the results to existing programs, it will be much easier to later extend the program since one then knows precisely what every line of code is doing. By knowing the code in detail, more confidence can also be put in the results. Obviously a lot of work is needed to write a CMB program from scratch. Fortunately this work is not a complete waste of time, since a lot of new insight about the underlying physics can be obtained in the process.  The purpose of this text is to provide a self-contained collection of all the ingredients needed to write a program that calculates the CMB power spectrum. This includes both the equations governing the physics, any assumptions or approximations used in their solutions, and pointing out what algorithms and tricks to use when implementing the equations in a computer program. The main focus will be the practical computer implementation of the equations, not their derivations, for which we instead point the reader to the references. Only new or less readily available derivations will be included. We are only assuming that the reader has some basic knowledge of CMB physics, and some experience with a high-level programming language~\\footnote{We will not show any code written in a spesific programming language, only the general algorithms used. For the record, we have used Delphi for Windows \\cite{Delphi} in this work, but C or Fortran (or similar languages) are equally well suited.}. Most of the presentation will follow the notation and conventions of Dodelson~\\cite{Dodelson}.  In section \\ref{cha:background} we start by going through the background cosmology, which includes background geometry and the recombination history of the universe, and in section \\ref{cha:perturbations} we introduce perturbations to this background. The equations of motion for the perturbations are given by the various Boltzmann equations. We then state the initial conditions used, and the approximations used during tight coupling. In section \\ref{cha:cmbspectrum} we go from the perturbations of the CMB temperature to the spectrum of $C_l$'s. The most important programming techniques are mentioned in section \\ref{cha:programming}, including the cutoff scheme for the Boltzmann hierarchies, how to integrate the various equations numerically, and the normalization of the spectrum. The resulting spectra are compared to CMBFAST in section \\ref{cha:results}. In section \\ref{cha:more_ingredients} the program is extended to include a few more effects, including helium, a simple model of reionization, massless neutrinos, and the (E-mode) polarization power spectrum. Finally we conclude in section~\\ref{cha:conclusion}. ",
        "conclusions": "\\label{cha:conclusion}  We have here presented all the main steps required in writing a program that calculates the CMB anisotropy power spectrum. Our focus has been on the computer-technical side of the problem, by including all the small details that make the program actually work, something which is often left out in the literature. We have consentrated on the $\\lcdm$ model, where the program achieves an accuracy comparable to CMBFAST~\\footnote{The accuracy of CMBFAST is of order $0.1 \\,\\%$~\\cite{CMBFAST_accuracy}.} ($\\sim \\, 0.1 \\,\\%$ -- $0.4 \\,\\%$) over a range of cosmological parameters. The program runs in a couple of minutes on a mid-range personal computer (as of 2006). While certainly not as good as CMBFAST, this is still acceptable considering that the code hasn't really been optimized for speed.  The purpose of this work has been to give a running start to those needing to calculate the CMB power spectrum for some exotic cosmological model where the standard programs can't be used. With the growing precision of the observed spectrum, a calculation to within the $1 \\,\\%$ level is often what distinguishes the models and makes it possible to rule out some of them. We hope this work will encourage others by showing that writing a program from scratch to within this accuracy is not really as difficult or time-consuming as one may think.   \\textbf{Acknowledgement:} I want to thank Tomi Koivisto for very useful discussions about the various programming techniques involved in writing the program. This work has been supported by grant no. NFR 153577/432 from the Research Council of Norway."
    },
    "0606/astro-ph0606360_arXiv.txt": {
        "abstract": "We show by means of a high-resolution N-body simulation how the mass assembly histories of galaxy--size cold dark matter (CDM) halos depend on environment. Halos in high density environments form earlier and a higher fraction of their mass is assembled in major mergers, compared to low density environments. The distribution of the present--day specific mass aggregation rate is strongly dependent on environment. While in low density environments only $\\sim 20\\%$ of the halos are not accreting mass at the present epoch, this fraction rises to $\\sim 80\\%$ at high densities. At $z=1$ the median of the specific aggregation rate is $\\sim 4$ times larger than at $z=0$ and almost independent on environment. All the dependences on environment found here are critically enhanced by local processes associated to subhalos because the fraction of subhalos increases as the environment gets denser. The distribution of the halo specific mass aggregation rate as well as its dependence on environment resemble the relations for the specific star formation rate distribution of galaxies. An analogue of the morphology--density relation is also present at the level of CDM halos, being driven by the halo major merging history. Nevertheless, baryonic processes are necessary in order to explain further details and the evolution of the star formation rate--, color-- and morphology--environment relations. ",
        "introduction": "The build--up of galactic dark matter halos is a crucial ingredient of galaxy formation in the context of the popular hierarchical $\\Lambda$ Cold Dark Matter (\\lcdm) scenario. Many of the present--day galaxy properties are expected to be tightly related to their halo mass assembly process. A simple way to characterize this process is by the halo mass--aggregation history (MAH). The MAH implies both the mass growth by violent major mergers and by quiescent accretion. The influence of the MAHs on the halo and galaxy properties remains imprinted for example on the halo concentration \\citep{AFH98, FA00, W02}, the gas infall and star formation rates in disk galaxies \\citep[e.g.,][]{AF00,vdB02,Murali02} and the major merging rate and therefore, the morphology of galaxies \\citep{KWG93,Baugh96,Rachel99,Cole00,Springel01,Granato01, Stein02, Maller05}.  The continuous mass growth of isolated halos is a generic process in the hierarchical CDM scenario \\citep[e.g.,][]{Gunn82,KWG93,LC93}. Semi--analytic models and numerical simulations have been used to predict the MAHs of halos of a given present--day mass, starting from a primordial Gaussian density fluctuation field characterized by a power spectrum. On one hand, the results have shown that, due to the stochastic nature of the primordial density field, the MAHs of halos of a given mass at $z=0$ span a wide range of ``tracks'', influencing galaxy features such as the scatter of the Tully--Fisher, halo concentration--mass and color--magnitude relations \\citep{EL96,AFH98,FA00,Bullock01,Eke01,Berlind05}. On the other hand, the MAHs depend on mass: less massive halos tend to build up a given fraction of their present--day mass on average earlier than the massive ones. This dependence is at the basis of the halo concentration--mass relation \\citep{NFW97,FA00,Eke01,vdB02,W02,zhao03b}.  To study the general behavior of the MAHs as a function of mass and environment, we introduce the average MAH (AMAH) as $\\tilde{M}(a)\\equiv \\frac {<M(a)>}{M_0}$, where $a$ is the scale factor and $<M(a)>$ is the average mass of the most massive progenitors (MMP) of present--day halos of mass $M_0$ at the epoch $a$. It was found that the smooth AMAHs can be approximated by simple (universal) functions, where (1) the main function parameter is related to some typical formation epoch of the halo, and (2) this epoch depends on $M_0$ \\citep{Bullock01,vdB02, W02}. On the basis of N--body cosmological simulations, it was also suggested that all the MAHs, independent of $M_0$, present an early phase of fast mass aggregation (mainly by major mergers) and a late phase of slow aggregation (mainly by smooth mass accretion) \\citep{zhao03a,Sole05,Lin05}. These authors also show that the depth of the potential well of a present--day halo, traced by its virial velocity, is set mainly at the end of the fast, major--merging driven, growth phase.   The role that the environment plays on the CDM halo assembly has not been explicitely explored in the literature. Yet, it is well known that some galaxy properties vary as a function of environment. Present-day galaxy properties as well as their evolution have been studied by some authors by combining cosmological N--body simulations and semi--analytic modeling of the baryon galaxy processes, a technique pioneered by \\citet{Kaufetal99} and \\citet{Springel01}. Although in the papers that used this technique the galaxy dependencies on environment are intrinsically taken into account, the impact the MAH dependence on environment has on the galaxy properties is not clearly established.   The main observable dependencies with environment are seen for the morphological mix of elliptical and spiral galaxies \\citep[e.g.,][]{Dressler80, PG84,Whitmore93,Dom02,Goto03}, color and specific star formation rate \\citep[e.g.,][]{Ba98,Balogh_ecol04, Ko01, Tran01, Pimbblet02, Lewis02, Blanton03, Gomez03, Kauffmann04, Hogg04, Tanaka04,Croton05}.  Besides, there is some evidence that these dependencies evolve, becoming weaker at higher redshifts \\citep{Treu03,Goto04,Bell04,sm05,Postman05}. As mentioned above, the morphology of galaxies as well as their colors and star formation rates are certainly related to the assembly history of their halos. Astrophysical external mechanisms, acting mainly in high--density environments, are also important. {\\it To disentangle the role of one from the other and to understand which are the drivers of the changes of galaxy properties with environment, it is important to explore and quantify how the CDM halo assembly history depends on environment.}  N--body cosmological simulations are required for this endeavor.   Based on N--body simulations, \\citet{LK99} concluded that only the halo mass function varies with environment. No significant dependence of any other halo property on environment was found by these authors. \\citet{Gott01} found that the major merging rate histories of CDM halos vary as a function of environment. More recently, \\citet{AR05} found that some properties of halos of {\\it similar masses} (for example, the mean concentration and spin parameter) actually vary systematically between voids and clusters. Also, recently, it was shown that halos of a given mass form statistically earlier in denser regions \\citep{ST04, Harker05,Gao05,W05}. Nonetheless, the question of how the mass assembly history of CDM halos of similar masses depend on environment has not been yet explored in detail. Furthermore, we would like to know to what extent this dependence is able to explain the observed dependencies of galaxy morphology, color, star formation rate, etc. on environment. We would also like to know if the bimodality of the star formation rate (SFR) and color distributions, found in large statistical galaxy samples, could be accounted for, at least in part, by the physics of dark matter halos only.   In this Paper we construct the MAHs of $\\sim 4700$ halos with present--day masses larger than $10^{11}\\msunh$ from a $50\\mpch$ box simulation with high--mass resolution: $m_p=7.75 \\times 10^7 \\msunh$ (\\S 2).  We find that several quantities that characterize the halo MAH (the specific mass aggregation rate and merging mass fraction histories, formation times, etc.) change significantly with environment for halos of similar present--day masses (\\S 3). We also explore possible systematical dependences of observational--related quantities on intermediate--scale density, and discuss to what extent these dependences are able to explain the observed galaxy property--density relations (\\S 4). Finally, we highlight the main conclusions of our work (\\S 5). ",
        "conclusions": "In previous works it was shown that the present--day mass function \\citep{LK99} and internal properties \\citep{AR05} of  \\lcdm\\ halos change systematically with environment. In the present paper we investigated environmental dependences of the mass assembly process of \\lcdm\\ halos. Our results suggest that a significant part of the relations observed between galaxy properties and environment can be explained by the environmental dependences of the halo assembling process. Our main findings are as follows:  $\\bullet$  The MAHs of the \\lcdm\\ halos change with intermediate--scale environment: the lower the density contrast $\\delta_4$, the later on average a given fraction of the maximum halo mass is assembled. The averages of the specific present--day mass aggregation rate, \\mdotp, are 4--5 times higher for halos in low--density environments ($\\delta_4<0$) than for halos in high--density environments ($\\delta_4>5$). These differences are smaller at higher redshifts, at $z\\sim 1$ disappear, and for $z\\gsim 1$ the trends are reversed. The average MAHs and mass aggregation rates also depend on halo mass, but to a much lesser degree than on environment.  $\\bullet$ The distribution of \\mdotp\\  has two parts: $\\approx 40\\%$ of all the halos in the mass range $11.5\\!\\le\\!\\log M_0/\\msunh\\!<\\!12.5$ do not aggregate mass, i.e (\\mdot)$_0 \\le 0$, while the rest has a distribution of \\mdotp\\ peaked at $\\sim 0.03-0.04$ Gyr$^{-1}$. The distribution changes little with mass but dramatically with environment: the fraction of halos with (\\mdot)$_0 \\le 0$ increases to $\\approx 81\\%$ and reduces to $\\approx 22\\%$ for the high-- and low--density environments, respectively. The distribution of the halo formation time,  also changes with environment: for instance, halos of $\\sim 10^{12}\\msunh$ assemble half of their maximum mass approximately 1.5 Gyrs earlier in the high--density regions than in the low--density ones.   $\\bullet$ Using only isolated halos, the dependences of halo MAHs on environment are weak, showing that the intermediate--scale density around halos affects only partially their mass assembling process. The strongest effects of environment appear when subhalos are included. Although the effect due to subhalos is local, the fraction of halos with masses above $10^{11}$ \\msunh\\ that become subhalos increases as the environment gets denser. The two main effects that a halo suffers when it becomes subhalo of a larger host virialized structure are: (i) its MAH is suddenly truncated (the mass growth is even reversed to mass loss due to tidal stripping), and (ii) the major merging probability drops drastically due to the high velocity dispersions inside the virialized host halo.  $\\bullet$ Present--day halos in high--density environments suffered more major mergers on average and assembled a larger fraction of their mass in major mergers than halos in low--density regions. Halos in dense regions become subhalos at $z\\lsim 1$ much more frequently than those in low--density regions. For these halos the (high) major merging mass fractions and (large) last major merger mass-to-current mass ratios, \\bt,  become defined at the epoch of their infall. On the other hand,  halos in low--density environments continue growing mainly by mass accretion, thus, their major merging mass fraction and \\bt\\ ratio continuously decrease.  We have shown that the distribution of \\mdotp\\ and the strong dependence this distribution presents with environment could explain partly the observed bimodal distributions of specific SFR and color for galaxies and their systematical dependences on environment. The main difficulty for this direct halo--galaxy connection would be to explain the existence of the population of red luminous galaxies in low--density environments, and the high fraction of blue star forming galaxies among low mass galaxies. Astrophysical processes related to baryonic matter should be invoked here.  The morphology--density relation can be partly explained if the halo \\bt\\ ratio is used as measure of the bulge--to--total mass ratio (morphology). However, even close to the cluster center, the fraction of halos with \\bt\\ as high as to harbour elliptical galaxies is too small at $z=0$. This suggests the need of astrophysical processes able to avoid gas accretion onto the early formed spheroids. The main trends of the observed morphology-- density and cluster-centric radius relations at $z=0$ agree with the trends of the halo \\bt\\ ratio with environment reported here.   Future comparisons of observations with models that include the astrophysical processes of galaxy formation in highly resolved evolving \\lcdm~ halos will tell us whether or not the \\lcdm\\ paradigm agrees with observations, and what are the key ingredients at the basis of the morphology-, color-, and star formation--density correlations. Our results suggest that the halo assembly process and its dependence on environment should be one main ingredient."
    },
    "0606/astro-ph0606156_arXiv.txt": {
        "abstract": "We report the discovery of a medium-strength ($\\sim$0.5~kG) magnetic field on the young, massive star $\\tau$~Sco (B0.2$\\;$V), which becomes the third-hottest magnetic star known. Circularly polarized Zeeman signatures are clearly detected in observations collected mostly with the ESPaDOnS spectro{\\-}polarimeter, recently installed on the 3.6-m Canada--France--Hawaii Telescope; temporal variability is also clearly established in the polarimetry, and can be unambiguously attributed to rotational modulation with a period close to 41~d.  Archival UV spectra confirm that this modulation repeats over timescales of decades, and refine the rotation period to $41.033\\pm0.002$~d.  Despite the slow rotation rate of $\\tau$~Sco, we nonetheless succeed in reconstructing the large-scale structure of its magnetic topology. We find that the magnetic structure is unusually complex for a hot star, with significant power in spherical-harmonic modes of degree up to 5.  The surface topology is dominated by a potential field, although a moderate toroidal component is probably present.  We fail to detect {\\em intrinsic} temporal variability of the magnetic structure over the 1.5-yr period of our spectro{\\-}polarimetric observations (in agreement with the stable temporal variations of the UV spectra), and infer that any differential surface rotation must be very small.  The topology of the extended magnetic field that we derive from the photo{\\-}spheric magnetic maps is also more complex than a global dipole, and features in particular a significantly warped torus of closed magnetic loops encircling the star (tilted at about 90\\degr\\ to the rotation axis), with additional, smaller, networks of closed field lines. This topology appears to be consistent with the exceptional X-ray properties of $\\tau$~Sco and also provides a natural explanation of the variability observed in wind-formed UV lines. Although we cannot completely rule out the possibility that the field is produced through dynamo processes of an exotic kind, we conclude that its magnetic field is most probably a fossil remnant from the star-formation stage. ",
        "introduction": "Magnetic fields in hot, high-mass stars of spectral types O and early~B may have a drastic influence on the physics of the stellar interiors \\citep[e.g.,][]{Spruit02} and atmospheres \\citep[e.g.,][]{Babel97}.  As a consequence, they can also significantly modify these stars' long-term evolution, and in particular their rotational history \\citep[e.g.,][]{Maeder03, Maeder04, Maeder05}.  However, quantifying these effects requires that we know the basic properties of such fields.  From a theoretical point of view, several quite different mechanisms have been proposed as potential means of generating large-scale magnetic fields in very hot stars.  For example, such fields may be fossil remnants of the star-formation stage, either as relics of the field that pervaded the inter{\\-}stellar medium from which the star formed, or as left{\\-}overs from dynamo action in the convective Hayashi phase.  This hypothesis was initially proposed for the lower-mass, chemically peculiar, magnetic Ap and Bp stars \\citep[e.g.,][]{Moss01} and has recently been rediscussed for the particular case of high-mass O and early-B stars by \\citet{Ferrario05, Ferrario06}; in this scenario, very hot magnetic stars are expected to be the progenitors of magnetic neutron stars.  A second possibility is that fields may be produced by continuing dynamo action; this option, initially suggested about two decades ago \\citep[e.g.,][]{Moss82}, has not been closely examined until quite recently, but is now attracting increasing interest from theoreticians.  A dynamo action could operate in the convective core \\citep{Charbonneau01, Brun05}, but the outer radiative envelope could also be involved in genuine dynamo processes, in a subsurface shear layer \\citep{Tout95, Lignieres96} or even throughout the whole envelope \\citep{Spruit02, Macdonald04, Mullan05, Maeder05, Braithwaite06}.  Unfortunately, all available theoretical options still suffer significant problems.  Fossil-field theories have yet to demonstrate that sufficient magnetic flux can survive the accelerated decay and expulsion associated with Hayashi convection\\footnote{This difficulty may not be as problematical as it first seemed, as discussed in \\citet{Moss03}, and should not concern stars with masses larger than 10~\\msun.}, while core-dynamo theories are still lacking detailed models linking the field produced in the core with that emerging at the surface \\citep{Moss01}.  Being based on a very new idea, radiative-zone dynamo theories are still in their infancy and need to establish their potential validity, by demonstrating that, in particular, stellar radiative zones are capable of sustaining differential rotation throughout their whole volume.  In addition to these issues, existing theories face a number of problems when compared to observations.  At some point, all theories involve the coexistence of large-scale toroidal fields in stellar interiors and possibly even close to the surface \\citep[e.g.,][]{BraithSpruit04, BraithNord06}, which observations of magnetic Ap and Bp stars do not yet confirm \\citep{Moss01}.  Moreover, dynamo theories predict magnetic topologies that are expected to vary on short timescales and to depend strongly on stellar rotation rates, which again is not observed in intermediate-mass magnetic stars \\citep{Moss01}.  Last but not least, dynamos theories should in principle apply to most hot stars (as they do for most cool stars), making it hard to understand why only a small fraction of them (about 10\\% in the case of Ap/Bp stars) apparently hosts magnetic fields.  Indirect evidence for the presence of magnetic fields in high-mass stars is regularly reported in the literature, these being postulated as a potential explanation for many otherwise enigmatic phenomena, including unanticipated X-ray line profiles and high X-ray temperatures \\citep[e.g.,][]{Cohen97, Robinson02, Cohen03, Smith04}.  However, with {\\em direct} detections of magnetic fields in only two O stars to date \\citep{Donati02, Donati06}, together with less than a handful of early-B stars with masses larger than 10~\\msun\\ \\citep[e.g.,][]{Donati01, Neiner03}, very little is known reliably about magnetic strengths and topologies from a statistical point of view.  One reason for this is that absorption lines of these stars are not only relatively few in number in the optical, but are also generally rather broad (due to rotation, or to some other type of as yet unknown macroscopic mechanism; e.g., \\citealt{Howarth97}), decreasing dramatically the size of the Zeeman signatures that their putative fields can induce.  In these respects, $\\tau$~Sco (HR~6165, HD~149438, HIP~81266) appears as a very promising target for magnetic-field studies; thanks to its brightness and unusually narrow absorption lines (among the sharpest known for stars more massive than 10~\\msun), $\\tau$~Sco is an obvious candidate for accurate spectro{\\-}polarimetric experiments. Moreover, its strong, hard X-ray emission (with $\\log L_{\\rm X}/L_{\\rm Bol}\\simeq-6.5$, \\citealt{Cohen97}) poses a severe challenge to the standard picture of O-star wind-shock models, leading some authors \\citep[e.g.,][]{Cohen03} to speculate that it displays the presence of a magnetically confined wind, as proposed by \\citet{Stahl96} and \\cite{Babel97}, and as detected on the very young O~star $\\theta^1$~Ori~C \\citep{Donati02, Gagne05, Gagne05err}.  We therefore included $\\tau$~Sco in a list of candidates for observation with ESPaDOnS, a high-efficiency spectro{\\-}polarimeter installed in 2004/5 on the 3.6-m Canada--France--Hawaii Telescope (CFHT; Donati et al., 2006, in preparation).  In this paper, we first present our new spectro{\\-}polarimetric observations and Zeeman detections obtained for $\\tau$~Sco, and briefly summarize the UV spectra (Sec.~\\ref{sec:obs}).  We then carry out a basic analysis of these data to establish the rotation period (Sec.~\\ref{sec:rot}), and review the fundamental stellar parameters in the light of our result (Sec.~\\ref{sec:params}).  We perform detailed modelling of the magnetic topology of $\\tau$~Sco, by direct fitting to the observed Zeeman signatures (Sec.~\\ref{sec:mod}).  Finally, we discuss the implications of our results for models of the X-ray emitting magneto{\\-}sphere of $\\tau$~Sco, as well for as theories of large-scale magnetic-field generation in high-mass stars (Secs.~\\ref{sec:disc} and \\ref{sec:disc2}).  \\begin{table*} \\caption[]{Journal of spectropolarimetric observations.  Columns 1--6 list the UT date \\& time,   heliocentric Julian date (all at mid-exposure), observing site,  exposure time, and peak signal to noise ratio (per 2.6~\\kms\\ velocity bin) of each observation.  Columns 7 and 8 list the rms noise level (per 1.8~\\kms\\ velocity bin, relative to the unpolarized continuum level $I_{\\rm c}$) in the circular-polarization profile produced by Least-Squares Deconvolution (see Sec.~\\ref{sec:obs}) and the estimated longitudinal field \\Bl\\ (with corresponding 1$\\sigma$ error bars). The rotational cycle $E$ from the ephemeris of eqn.~\\ref{eq:ephem} is given in column~9.} \\begin{tabular}{lccclrcrr} \\hline \\multicolumn{2}{c}{UT}                & HJD          & Obs  & \\multicolumn{1}{c}{$t_{\\rm exp}$} & \\multicolumn{1}{c}{\\sn}  & $\\sigma_{\\rm LSD}$ & \\multicolumn{1}{c}{\\Bl} & \\multicolumn{1}{c}{Cycle} \\\\ \\multicolumn{1}{c}{Date}   & (h:m:s)  & (2,453,000+) &      & \\multicolumn{1}{c}{(s)}         &       &   (\\ptt)  & \\multicolumn{1}{c}{(G)} &  \\\\ \\hline 2004 Sep.\\ 04 & 05:58:31 & 252.7475     & CFHT & $4\\times30$  & 1440 & 0.70 & $-24.4\\pm3.3$ & 1.456 \\\\ 2004 Sep.\\ 04 & 06:23:17 & 252.7647     &      & $4\\times30$  & 1280 & 0.75 & $-18.3\\pm3.4$ & 1.457 \\\\ \\hline 2004 Sep.\\ 25 & 05:03:15 & 273.7075     & CFHT & $4\\times30$  & 1060 & 0.96 & $+36.2\\pm4.5$ & 1.967 \\\\ 2004 Sep.\\ 25 & 05:10:24 & 273.7124     &      & $2\\times60$  & 1100 & 0.92 & $+47.5\\pm4.3$ & 1.967 \\\\ \\hline 2004 Sep.\\ 26 & 08:51:15 & 274.8689     & AAT  & $4\\times120$ &  890 & 1.03 & $+64.0\\pm4.2$ & 1.995 \\\\ 2004 Sep.\\ 27 & 08:52:49 & 275.8700     &      & $4\\times120$ &  860 & 1.10 & $+78.4\\pm4.5$ & 2.020 \\\\ 2004 Sep.\\ 28 & 08:58:25 & 276.8739     &      & $4\\times120$ &  670 & 1.36 & $+80.5\\pm5.6$ & 2.044 \\\\ \\hline 2005 May 23   & 09:14:25 & 513.8903     & CFHT & $4\\times300$ & 1700 & 0.63 & $-47.9\\pm3.0$ & 7.820 \\\\ 2005 May 24   & 08:32:52 & 514.8614     &      & $4\\times60$  &  900 & 1.17 & $-50.3\\pm5.6$ & 7.844 \\\\ 2005 May 25   & 08:41:35 & 515.8675     &      & $4\\times60$  &  880 & 1.19 & $-41.4\\pm5.5$ & 7.868 \\\\ \\hline 2005 June  19 & 07:19:19 & 540.8096     & CFHT & $4\\times120$ & 1200 & 0.85 & $-19.2\\pm4.3$ & 8.476 \\\\ 2005 June  20 & 06:44:15 & 541.7852     &      & $4\\times120$ & 1550 & 0.66 & $-20.2\\pm3.3$ & 8.500 \\\\ 2005 June  21 & 06:51:16 & 542.7900     &      & $4\\times120$ & 1650 & 0.62 & $-19.8\\pm2.9$ & 8.525 \\\\ 2005 June  22 & 08:42:17 & 543.8671     &      & $4\\times120$ & 1440 & 0.71 & $-12.8\\pm3.6$ & 8.551 \\\\ 2005 June  23 & 06:07:24 & 544.7595     &      & $4\\times120$ & 1670 & 0.61 & $-5.9\\pm3.1$ & 8.572 \\\\ 2005 June  24 & 06:22:32 & 545.7699     &      & $4\\times120$ & 1710 & 0.60 & $+1.3\\pm3.1$ & 8.597 \\\\ 2005 June  25 & 06:05:50 & 546.7583     &      & $4\\times120$ & 1570 & 0.68 & $+4.7\\pm3.4$ & 8.621 \\\\ 2005 June  26 & 06:00:31 & 547.7545     &      & $4\\times120$ & 1590 & 0.66 & $+0.9\\pm3.3$ & 8.646 \\\\ 2005 June  26 & 11:02:47 & 547.9644     &      & $4\\times120$ & 1370 & 0.82 & $-5.1\\pm4.2$ & 8.651 \\\\ \\hline 2005 Aug.\\ 19 & 05:20:15 & 601.7224     & CFHT & $4\\times30$  & 1250 & 0.77 & $+44.5\\pm3.7$ & 9.961 \\\\ 2005 Aug.\\ 21 & 05:21:43 & 603.7232     &      & $4\\times30$  & 1200 & 0.81 & $+80.3\\pm4.0$ & 10.010 \\\\ 2005 Aug.\\ 23 & 05:11:45 & 605.7161     &      & $4\\times30$  & 1150 & 0.85 & $+87.8\\pm4.2$ & 10.058 \\\\ \\hline 2005 Sep.\\ 19 & 05:03:10 & 632.7079     & CFHT & $4\\times30$  & 1180 & 0.83 & $-21.8\\pm3.7$ & 10.716 \\\\ 2005 Sep.\\ 20 & 05:01:56 & 633.7070     &      & $4\\times30$  & 1280 & 0.76 & $-28.9\\pm3.6$ & 10.740 \\\\ 2005 Sep.\\ 24 & 04:55:44 & 637.7023     &      & $2\\times50$  &  680 & 1.48 & $-45.9\\pm6.8$ & 10.838 \\\\ 2005 Sep.\\ 25 & 04:48:19 & 638.6971     &      & $4\\times30$  & 1200 & 0.83 & $-51.9\\pm5.3$ & 10.862 \\\\ 2005 Sep.\\ 25 & 05:01:33 & 638.7063     &      & $4\\times30$  & 1120 & 0.89 & $-43.4\\pm4.2$ & 10.862 \\\\ \\hline 2006 Feb.\\ 07 & 15:03:44 & 774.1269    & CFHT & $4\\times30$  & 1090 & 0.88 & $+44.9\\pm4.1$ & 14.162 \\\\ 2006 Feb.\\ 08 & 16:22:33 & 775.1817    &      & $4\\times30$  &  910 & 1.04 & $+44.3\\pm4.8$ & 14.188 \\\\ 2006 Feb.\\ 09 & 16:16:56 & 776.1779    &      & $4\\times30$  &  620 & 1.58 & $+38.9\\pm7.0$ & 14.212 \\\\ 2006 Feb.\\ 10 & 16:33:25 & 777.1895    &      & $4\\times30$  & 1250 & 0.76 & $+32.8\\pm3.4$ & 14.237 \\\\ 2006 Feb.\\ 11 & 16:18:20 & 778.1791    &      & $4\\times30$  & 1020 & 0.94 & $+25.8\\pm4.2$ & 14.261 \\\\ 2006 Feb.\\ 13 & 16:15:23 & 780.1772    &      & $4\\times30$  &  490 & 2.06 & $+35.4\\pm9.4$ & 14.310 \\\\ 2006 Feb.\\ 14 & 16:23:12 & 781.1827    &      & $4\\times30$  &  840 & 1.13 & $+0.6\\pm5.6$ & 14.334 \\\\ 2006 Feb.\\ 15 & 14:07:12 & 782.0884    &      & $4\\times30$  & 1060 & 0.92 & $+8.0\\pm4.5$ & 14.357 \\\\ \\hline \\end{tabular} \\label{tab:log} \\end{table*} ",
        "conclusions": "We have reported the detection of a magnetic field on the massive B0.2$\\;$V star $\\tau$~Sco, using data obtained mostly with ESPaDOnS, the new high-resolution stellar spectro{\\-}polarimeter recently installed at CFHT.  From the Zeeman signatures and their temporal variability, we were able to identify the rotation period of $\\tau$~Sco and to reconstruct the large-scale topology of its photo{\\-}spheric field.  Archival IUE spectra confirm that the rotational modulation is stable on timescales of decades, with $P_{\\rm rot} \\simeq 41.03$~d.  We find that the surface magnetic topology is unusually complex (judged by the small sample of massive-star results) and is mostly potential.  It also includes a moderate toroidal component; in particular, the strength of this toroidal component (relative to that of the poloidal component) is much lower than that found in partly-convective cool stars hosting dynamo-generated magnetic fields.  No temporal variability of the magnetic structure is detected over the 1.5-yr period of our observations; we thus conclude that any surface differential rotation of $\\tau$~Sco is at least 20 times weaker than that of the Sun.  We determine that the large-scale magneto{\\-}spheric structure of $\\tau$~Sco is significantly more complex than a global dipole;  it features in particular a significantly warped torus of closed magnetic loops encircling the star, tilted at about 90\\degr\\ to the rotation axis, as well as additional (smaller) networks of closed field lines. The extended magnetic topology we derive from extrapolations of the photo{\\-}spheric magnetic maps is apparently compatible with the published X-ray luminosity and spectral characteristics of $\\tau$~Sco. Our model is also compatible with the observed modulation of UV spectral lines.  We predict that $\\tau$~Sco should exhibit a clear rotational modulation of its X-ray emission.  From these results, we conclude that its magnetic field is most probably a fossil remnant from the formation stage. We cannot yet completely rule out the possiblity that the field is produced through one of the recently-elaborated dynamo processes that may operate in the radiative zones of hot stars, but our findings already indicate that this option is rather unlikely."
    },
    "0606/astro-ph0606630_arXiv.txt": {
        "abstract": "The high-frequency--peaked BL Lacertae object Markarian~180 (Mrk~180) was observed to have an optical outburst in 2006 March, triggering a Target of Opportunity observation with the MAGIC telescope. The source was observed for 12.4 hr, and very high energy $\\gamma$-ray emission was detected with a significance of 5.5~$\\sigma$.  An integral flux above 200~GeV of $(2.3\\pm0.7)\\times10^{-11}\\, \\mbox{cm}^{-2}\\,\\mbox{s}^{-1}$ was measured, corresponding to 11\\% of the Crab Nebula flux. A rather soft spectrum with a photon index of $-3.3\\pm0.7$ has been determined. No significant flux variation was found. ",
        "introduction": "The search for very high energy (VHE, defined as $E\\gtrsim 100$~GeV) $\\gamma$-ray emission from active galactic nuclei (AGNs) is one of the major goals for ground-based $\\gamma$-ray astronomy. New detections open the possibility of a phenomenological study of the physics inside the relativistic jets in AGNs, in particular, to understand both the origin of the VHE $\\gamma$-rays as well as the correlations between photons of different energy ranges (from radio to VHE). The number of reported VHE $\\gamma$-ray--emitting AGNs has been slowly increasing and is currently 12 (2006 June).  Six of them have been seen by MAGIC: Mrk~421 \\citep{magic421}, Mrk~501 \\citep{mazin}, 1ES1959+650 \\citep{magic1959}, 1ES2344+514 \\citep{mazin}, 1ES1218+304 \\citep{magic1218}, and PG1553+113 \\citep{magic1553}.   The known VHE $\\gamma$-ray--emitting AGNs are variable in flux in all wave bands. Correlations between X-ray and $\\gamma$-ray emission have been found (e.g., \\citet{fossati04}), although the relationship has proven to be rather complicated, with $\\gamma$-ray flares also being detected in the absence of X-ray flares \\citep{holder,krawczynski} and vice versa \\citep{rebillot}. The optical-TeV correlation has yet to be studied, but the optical-GeV correlations seen in 3C~279 \\citep{hartman01} suggest that at least in some sources, such correlations exist. Using this as a guideline, the MAGIC collaboration has been performing Target of Opportunity observations whenever alerted that sources were in a high flux state in the optical and/or X-ray band.   The AGN Mrk~180 (1ES 1133+704) is a well-known high-frequency--peaked BL Lac (HBL) object at a redshift of $z=0.045$ \\citep{falco}. The spectral energy distribution (SED) of HBL objects exhibits a generic two-bump structure: one peak with a maximum in the X-ray band and the other peak located in the GeV-TeV band. The radiation is produced in a highly beamed plasma jet, which is almost aligned with the observer's line of sight. A double-peaked SED is normally attributed to a population of relativistic electrons, where one peak is due to synchrotron emission in the magnetic field of the jet and where the second peak is caused by inverse Compton (IC) scattering of low-energy photons. The low-energy photons can be external to the jet (external Compton scattering; \\citet{dermer}) or are produced within the jet via synchrotron radiation (synchrotron self-Compton [SSC] scattering; \\citet{maraschi}). Models based on the acceleration of hadrons can also sufficiently describe the observed SEDs and light curves \\citep{mannheim91,massaro2}. All of the up to now known AGNs with strong GeV/TeV $\\gamma$-emission belong to the HBL object class.   In many cases, the second peak of the SED is not observable because of low sensitivity above a few times 100~MeV of satellite-borne detectors and a too high energy threshold of ground-based $\\gamma$-ray detectors. In case of Mrk~180, HEGRA \\citep{54AGN} and Whipple \\citep{horan} observed this object but were only able to derive flux upper limits, and EGRET did not detect the source \\citep{fichtel, hartman}.  In the optical, Mrk~180 is characterized by a bright host galaxy $R=14.17\\pm0.02$ mag and a much fainter (variable) core $R=15.79\\pm0.02$ mag \\citep{nilsson}. The  variations in total optical brightness are therefore small. The source is observed regularly as part of the Tuorla Observatory Blazar monitoring program\\footnote{See http://users.utu.fi/kani/1m/.} with the Tuorla 1~m telescope and the KVA 35~cm telescope\\footnote{See http://tur3.tur.iac.es/.}.  To determine the core flux, we had to subtract the flux of the host galaxy and a nearby star within the $5''$ aperture radius (together R=14.96 mag, or 3.2 mJy; \\linebreak Nilsson, private communication).  After Mrk~180 had been detected in the X-rays by {\\it HEAO-1} \\citep{HEAO}, it was observed by various satellites with fluxes ranging from 6.3 to $22\\times 10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ \\citep{donato} at around 1~keV and from $5.0$ to $9.8\\times10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ in the 2--10~keV band \\citep{perlman,donato05}. The source is monitored by the all-sky monitor (ASM) on board the {\\it Rossi X-Ray Timing Explorer}.  In this Letter we present the first detection of VHE $\\gamma$-ray emission from Mrk~180. ",
        "conclusions": "In this Letter we have presented the first detection of VHE $\\gamma$-ray emission from Mrk~180. The discovery was triggered by an optical flare, but no significant variations in the VHE regime were found. The short observation period and the small signal prevent us from carrying out detailed studies. It is therefore impossible to judge whether the detected VHE flux level represents a flaring or a quiescent state of the AGN.  Earlier observations by other experiments have only set upper limits on the VHE $\\gamma$-ray emission from Mrk~180. HEGRA observed the source for 9.8 hr, resulting in an upper limit (99\\% c.l.) at 1.5~TeV of 12\\% of the Crab Nebula flux \\citep{54AGN}.  Whipple observed Mrk~180 during three observation periods for a total of 26.8 hr, resulting in a flux upper limit at 300~GeV of 10.8\\% of the Crab Nebula flux \\citep{horan}.  Both upper limits are above the flux presented in this Letter.  Fig.~\\ref{fig_SED} shows the SED of Mrk~180: the radio, optical and X-ray data are from UMRAO, KVA, ASM, and the NED database, and the VHE data are from this analysis and various upper limits.  Simultaneous data are noted in black, while historical data are noted in gray.  The predicted flux by the phenomenological model of \\citet{fossati} ({\\it dashed line}) is too high, as already emphasized by the non detection measurements \\citep{54AGN,horan}.  A more detailed SSC model from \\citet{costamante} ({\\it solid line}) seems to describe the data more closely.  The rather steep slope of the VHE spectrum suggests an IC peak position well below 200~GeV, while the non detection by EGRET gives a lower limit of $\\sim1$~GeV for the peak position. This is in agreement with the SSC model of \\citet{costamante} suggesting that the IC peak position is at around 10~GeV.  The overall IC luminosity, however, is underestimated in this model: the observed integral flux above 300~GeV is a factor 30 larger than predicted.  An explanation for this discrepancy could be the model's underlying assumption of a quiescent-state synchrotron spectrum to obtain the IC flux.  This could indeed suggest that our measurement was made during a high state.  \\begin{figure}[t] \\begin{center} \\begin{minipage}[b]{0.45\\textwidth} \\includegraphics*[width=0.9\\textwidth,angle=0,clip]{f5.eps} \\caption{\\label{fig_SED} SED of   Mrk~180. Simultaneous data (UMRAO, KVA, ASM, and MAGIC) are represented as black circles. The gray circles represent historical data \\citep{giommi,perlman}. The arrows denote the upper limits from ASM, EGRET \\citep{fichtel}, Whipple, and HEGRA.  The solid line is from \\citet{costamante} and the dashed line is from \\citet{fossati} (see text).} \\end{minipage} \\end{center} \\end{figure}  The discovery of VHE emission from Mrk~180 during an optical outburst makes it very tempting to speculate about the connection between optical activity and increased VHE emission.  Since Mrk~180 has not been observed prior to the outburst with MAGIC and since the upper limits from other experiments are above the observed flux level, further observations are needed.   {\\textit{Acknowledgements:}} We would like to thank the IAC for the excellent working conditions at the ORM in La Palma. The support of the German BMBF and MPG, the Italian INFN, the Spanish CICYT, ETH research grant TH~34/04~3, and the Polish MNiI grant 1P03D01028 is gratefully acknowledged.  We also thank H. D. and M. F. Aller, at the University of Michigan Radio Observatory, for providing us with the UMRAO data."
    },
    "0606/astro-ph0606406_arXiv.txt": {
        "abstract": "We calculate the yield of high energy neutrinos produced in astrophysical sources with negligible magnetic fields varying their interaction depth from nearly transparent to opaque. We take into account the scattering of secondaries on background photons as well as the direct production of neutrinos in decays of charm mesons. If multiple scattering of nucleons becomes important, the neutrino spectra from meson and muon decays are strongly modified with respect to transparent sources. Characteristic for neutrino sources containing photons as scattering targets is a strong energy-dependence of the ratio $R^0$ of $\\nu_\\mu$ and $\\nu_e$ fluxes at the sources, ranging from $R^0=\\phi_\\mu/\\phi _e\\sim  0$ below threshold to $R^0\\sim 4$ close to the energy where the decay length of charged pions and kaons equals their interaction length on target photons. Above this energy, the neutrino flux is strongly suppressed and depends mainly on charm production. ",
        "introduction": "Experimental high energy neutrino physics has become one of the most active areas of astroparticle physics, offering among others the prospect of identifying the sources of ultra-high energy cosmic rays~\\cite{reviews}. High energy neutrinos from astrophysical sources are the decay products of secondary mesons produced by scattered  high energy protons on background protons or photons.  The classic example are the so-called cosmogenic or GZK neutrinos produced in scatterings of extragalactic ultra-high energy cosmic rays on cosmic microwave photons during propagation~\\cite{GZKnu}. Two different kind of bounds on high energy neutrino fluxes exist: The cascade or EGRET limit uses bounds on the diffuse MeV-GeV photon background to limit the energy transferred to electromagnetically interacting particles that are produced unavoidably together with neutrinos~\\cite{casc}. The cosmic ray upper bounds of, e.g., Refs.~\\cite{WB,MPR} use the observed ultra-high energy cosmic ray flux to limit possible neutrino fluxes. The latter limit assumes that all neutrino sources are transparent to hadronic interactions and thus at least neutrons can escape from the source region without interactions. Dropping the assumption of transparent sources and hidding the acceleration region by sufficient material absorbing ultra-high energy cosmic rays allows one to avoid the cosmic ray limits~\\cite{TZ77,BG81,St,Berezinsky:2000bq}. One might therefore speculate that large neutrino fluxes at high and ultra-high energies might be produced in opaque sources, as they are needed, e.g., to perform neutrino absorption spectroscopy~\\cite{spec}.   In this work, we calculate the flux of high energy neutrinos produced as secondaries in astrophysical sources. We consider sources with interaction depth ranging from nearly transparent to opaque. In contrast to most earlier investigations~\\cite{Y,earlyI,earlyII}, we put emphasis on sources with such high densities that magnetic fields can be neglected and multiple scatterings are important. As a result, the neutrino spectra from meson and muon decays are strongly modified with respect to transparent sources. Since in most astrophysical environments the depth for $p\\gamma$ is above threshold much larger than for $pp$ interactions, we consider photons in the source as target material. In Sec. 2, we discuss the relevant production processes and present the neutrino yields from a single source. We discuss also briefly the consequences of the derived energy dependence of the neutrino yields on the expected high energy neutrino fluxes from astrophysical sources. In Sec.~3, we examine the neutrino flavor composition at the source as well as the effect of neutrinos oscillations on the flavor composition at the detector and find significant deviations from the canonical flavor ratio expected from pion decay. Finally, we summarize our results in  Sec.~4. ",
        "conclusions": "The high-energy cosmic ray flux from opaque or hidden neutrino sources is suppressed  and these source are therefore among the most promising candidates for powerful neutrino sources. We have calculated the yield of high energy neutrinos from sources with photons as target material, paying especially attention to opaque sources. We have found that the neutrino spectra from meson and muon decays are strongly modified with respect to transparent sources as soon as multiple scattering of nucleons becomes important. The main consequence is a strong suppression of the neutrino flux from pion and kaon decays at energies above a critical energy $E_{\\rm cr}$, defined as the energy at which the scattering length of mesons equals the decay length. Above this energy both pions and kaons scatter before decay and thus the main contribution to the neutrino flux is supplied by the decay of charm mesons. As a result, the parameter range where opaque sources can produce ultra-high energy neutrinos is rather restricted. An opaque source with $E_{\\rm cr}\\gsim 10^{21}$~eV requires basically a large extension with low densities.   Both neutrino telescopes and extensive air shower experiments can not only detect high energy neutrinos but also they have some flavor discrimination possibilities. Since the determination of the neutrino flavor ratio can shed light on the properties of both sources and neutrinos, we have calculated  also the expected flavor ratio for opaque sources containing photons as scattering targets. The main characteristic is a strong energy-dependence of the ratio $R^0$ of $\\nu_\\mu$ and $\\nu_e$ fluxes at the sources. A generic  prediction for neutrinos from hadron-photon scattering is the flavor ratio $R^0\\sim 0$  below the threshold, i.e.\\ a flux dominated by $\\bar\\nu_e$ at the source. In the case of opaque sources the flavor ratio presents a bump close to the energy where the decay length of charged pions and kaons equals their interaction length on target photons. The neutrino spectra at the source are modulated by oscillation. Therefore the expected flavor ratio $R_\\mu$ at the Earth will be different from the original one $R^0$: it significantly depends on the neutrino mixing parameters, and especially on $\\theta_{23}$. Therefore the observation of a strong energy dependence of $R_\\mu$ will not only be a hint on the kind of source but also allow obtaining complementary information on the neutrino mixing angles."
    },
    "0606/astro-ph0606540_arXiv.txt": {
        "abstract": "{} {We investigate the radio emitting structures of radio--quiet active galactic nuclei with an emphasis on radio--quiet quasars to study their connection to Seyfert galaxies.} {We present and analyse high--sensitivity VLA radio continuum images of 14 radio--quiet quasars and six Seyfert galaxies.} { Many of the low redshift radio--quiet quasars show radio structures that can be interpreted as jet--like outflows. However, the detection rate of extended radio structures on arcsecond scales among our sample decreases with increasing redshift and luminosity, most likely due to a lack of resolution. The morphologies of the detected radio emission indicate strong interactions of the jets with the surrounding medium. We also compare the radio data of seven quasars with corresponding HST images of the [\\ion{O}{iii}] emitting narrow--line region (NLR). We find that the scenario of interaction between the radio jet and the NLR gas is confirmed in two sources by structures in the NLR gas distribution as previously known for Seyfert galaxies. The extended radio structures of radio--quiet quasars at sub--arcsecond resolution are by no means different from that of Seyferts. Among the luminosities studied here, the morphological features found are similar in both types of objects while the overall size of the radio structures increases with luminosity. This supports the picture where radio--quiet quasars are the scaled--up versions of Seyfert galaxies. In addition to known luminosity relations we find a correlation of the NLR size and the radio size shared by quasars and Seyferts. } {} ",
        "introduction": "The population of active galactic nuclei (AGN) shows a well known wide spread in radio luminosity (Miller et al. \\cite{miller90}; Kellermann et al. \\cite{kellermann89},\\cite{kellermann94}; Miller et al. \\cite{miller93}; Kukula et al. \\cite{kukula98}; Xu et al. \\cite{xu99}). Introducing the R parameter as the ratio of radio flux at 5 GHz and optical flux at 4400\\,\\AA~ (Kellermann et al. \\cite{kellermann89}; Falcke et al. \\cite{falcke96b}), the radio--loud objects are separated  from the radio--quiet objects: Usually, the radio--loud quasars have R\\,$>$\\,100 and 5 GHz luminosities greater than $10^{26}\\,{\\rm W\\,Hz^{-1}}$. The radio--quiet ones typically have R values between 0.1 and 10 and luminosities less than $10^{24}\\,{\\rm W\\,Hz^{-1}}$ at 5 GHz. Thus, the separation between the classes is nearly three orders of magnitude in radio luminosity (Kellermann et al. \\cite{kellermann89},\\cite{kellermann94}). Radio--loud quasars (RLQs) and radio--galaxies are usually dominated by prominent double--lobe structures extending to tenth or hundreds of kpc. These extended structures in radio--loud objects are also dominating in luminosity by outshining the emission from the active central region by factors of 10 to 100, resulting in a very low core--to--lobe flux ratio. On the other hand, radio--quiet quasars (RQQs) only show small scale structures of typically a few kpc with most of the flux concentrated in the nuclear region (e.g. Miller et al. \\cite{miller90}; Kellermann et al. \\cite{kellermann94}). Nevertheless, some RQQs are known to show extended structures reminiscent of radio--loud objects, but on very low flux--density scales (Kellermann et al. \\cite{kellermann94}; Blundell \\& Rawlings \\cite{blundell01}). However, systematic high--dynamic range radio imaging studies of RQQs are still missing.\\\\\\indent At lower (radio) luminosities, Seyfert galaxies also show compact structures, often dominated by a central unresolved component, sometimes accompanied by jet--like features (Ulvestad et al. \\cite{ulvestad81}; Kukula et al. \\cite{kukula95}; Schmitt et al. \\cite{schmitt01}; Ho \\& Ulvestad \\cite{ho01}). Comparing the snapshot data of Seyfert galaxies and RQQs, there is no significant difference in the radio structure in both types of radio--quiet AGN. \\\\\\indent However, in high sensitivity data, the diversity in radio structure of Seyfert galaxies is obvious (Ulvestad \\& Wilson \\cite{ulvestad89}; Capetti et al. \\cite{capetti96}; Falcke et al. \\cite{falcke98}). The morphologies can comprise linearly arranged multiple components (interpreted as emission knots in a jet), one-- or two--sided diffuse jets or complicated structures that are, nevertheless, often well collimated. These radio structures coincide in shape and/or direction with the line--emitting gas of the narrow--line region (NLR), indicating that the radio ejecta are intimately connected to the NLR gas (Falcke et al. \\cite{falcke98}).\\\\\\indent In the ongoing debate of the evolution of active galaxies with luminosity, these findings of structured NLRs in Seyfert galaxies lead Bennert et al. (\\cite{bennert02}) to investigate the NLR structure of RQQs with high--resolution. They find resolved and structured quasar NLRs that closely resemble that of Seyfert galaxies. Moreover they discovered a close correlation between the size and the luminosity of the NLR, connecting Seyfert galaxies and RQQs over a few orders of magnitude in luminosity. These results strongly support the idea that, amongst radio--quiet AGN, RQQs are the high--power siblings of Seyfert galaxies. Since the close connection of Seyferts and RQQs is thus supported in the optical/NLR, we here explore the question if this can in fact be extended to the radio regime. We investigate if the interaction of radio ejecta and NLR gas, a typical phenomenon of active galaxies, persists with increasing luminosity. Therefore, we combine high--resolution, high--sensitivity radio and optical data for a sample of RQQs. We present radio continuum maps for 14 RQQs and some Seyfert galaxies to analyse the similarities and differences in both types of objects.\\\\\\indent In $\\S$2 and $\\S$3, observations, data reduction, and data analysis are summarised. In $\\S$4, we present the results on an object--by--object basis. In $\\S$5, we discuss the results and their implications for the role of radio jets. We close with the concluding remarks in $\\S$6. \\begin{table*} \\caption[]{\\sc VLA Observation Log \\label{sample}} \\begin{center} \\begin{tabular}{cllllccc} \\hline\\hline \\\\[-2.5ex] & & \\multicolumn{2}{c}{optical position} & & & Frequency & on--source \\\\[0.5ex]\\cline{3-4} & \\multicolumn{1}{c}{Object} & \\multicolumn{1}{c}{\\raisebox{0.2ex}[-0.2ex]{$\\alpha_{2000}$}} & \\multicolumn{1}{c}{\\raisebox{0.2ex}[-0.2ex]{$\\delta_{2000}$}} & \\hspace*{+1mm}z$_{\\rm hel}$ & R & (GHz) & time \\\\ & \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\hspace*{+1mm}(4) & (5) & (6) & (7)\\\\*[0.1cm]\\hline & PG0026+129    & 00 29 13.73  & +13 16 04.3    & 0.142    & 1.10 & 8.4 & 2:02 \\\\[0.3ex] & PG0052+251    & 00 54 52.16  & +25 25 38.7    & 0.155    & 0.24 & 4.8 & 3:57 \\\\[0.3ex] & PG0157+001    & 01 59 50.21  & +00 23 40.6    & 0.163    & 2.10 & 8.4 & 2:22 \\\\[0.3ex] \\multicolumn{1}{c}{\\rotatebox{0}{Quasars}} & PG0953+414    & 09 56 52.50  & +41 15 40.6    & 0.2341   & 0.44 & 4.8 & 3:45 \\\\[0.3ex] & PG1012+008    & 10 14 54.88  & +00 33 36.8    & 0.1874   & 0.50 & 8.4 & 2:46 \\\\[0.3ex] & PG1049$-$005  & 10 51 51.50  & $-$00 51 16.6  & 0.3599   & 0.25 & 4.8 & 3:34 \\\\[0.3ex] & PG1307+085    & 13 09 47.03  & +08 19 49.3    & 0.155    & 0.10 & 4.8 & 3:07 \\\\[0.3ex]\\hline & PG0003+199    & 00 06 19.193 & +20 12 10.51   & 0.0258   & 0.27 & 4.8 & 2:59 \\\\[0.3ex] & PG1119+120    & 11 21 47.102 & +11 44 18.28   & 0.0502   & 0.15 & 4.8 & 2:33 \\\\[0.3ex] & PG1149$-$110  & 11 52 03.505 & $-$11 22 23.87 & 0.049    & 0.88 & 4.8 & 3:05 \\\\[0.3ex] \\multicolumn{1}{c}{\\rotatebox{0}{Quasars}} & PG1351+640    & 13 53 15.847 & +63 45 45.72   & 0.0882   & 4.30 & 4.8 & 3:54 \\\\[0.3ex] & PG1534+580    & 15 35 52.358 & +57 54 09.17   & 0.0296   & 0.70 & 4.8 & 3:47 \\\\[0.3ex] & PG1612+261    & 16 14 13.214 & +26 04 16.40   & 0.131    & 2.80 & 4.8 & 3:17 \\\\[0.3ex] & PG2130+099    & 21 32 27.814 & +10 08 19.49   & 0.063    & 0.32 & 4.8 & 2:12 \\\\[0.3ex]\\hline & Mrk 612       & 03 30 40.9   & $-$03 08 16    & 0.0203   & & 8.4 & 1:30 \\\\[0.3ex] & ESO 428$-$G14 & 07 16 31.2   & $-$29 19 28    & 0.0054   & & 8.4 & 1:42 + 1:05 \\\\[0.3ex] & NGC 2639      & 08 43 38.1   & +50 12 20      & 0.0111   & & 8.4 & 1:06 \\\\[0.3ex] \\multicolumn{1}{c}{\\raisebox{1.8ex}[-1.8ex]{Seyferts}} & NGC 2992      & 09 45 42.0   & $-$14 19 35    & 0.0077   & & 8.4 & 1:06 \\\\[0.3ex] & Mrk 266       & 13 38 17.5   & +48 16 37      & 0.0279   & & 8.4 & 1:39 \\\\[0.3ex] & NGC 5643      & 14 32 40.8   & $-$44 10 29    & 0.004    & & 8.4 & 1:47 + 2:19 \\\\*[0.1cm]\\hline \\end{tabular} \\end{center} \\begin{list}{}{} \\item[\\emph{Note:}] The first seven RQQs are also included in the HST sample of Bennert et al. (\\cite{bennert02}). \\item[(2)--(3)] The optical position was taken from Miller et al. \\cite{miller93} and calculated for J2000 (in $^{\\rm h}\\,^{\\rm min}\\,^{\\rm sec}$ and $^{\\circ}\\,^{\\prime}\\,^{\\prime\\prime}$). \\item[(4)] heliocentric redshift as provided by NED \\item[(5)] R parameter taken from Kellermann et al. (\\cite{kellermann94}) \\item[(7)] on--source integration time in hrs:min \\end{list} \\end{table*} ",
        "conclusions": "We present deep radio images of 14 radio--quiet quasars and six Seyfert galaxies that show extended structures which can be interpreted as jets. Especially for the RQQs, aligned multiple components are detected that are generally interpreted as termination points of radio jets. The radio structures of RQQs are larger than those of Seyfert galaxies but exhibit the same morphology. Thus, RQQs can be regarded as scaled--up versions of Seyfert galaxies, keeping in mind distance and resolution effects when interpreting the images. The similarity can also be seen in the NLR gas distribution where the radio jet of the RQQs leaves its imprint on the NLR creating structures that are well known from Seyfert NLRs. This includes sweeping--up material and the creation of bow-shocks (PG0157+001 vs. Mrk 573) as well as one--sided jets which interact with one--sided NLRs (PG1012+008 vs. ).\\\\\\indent Comparing our results with (radio) snapshot surveys of Seyfert galaxies and quasars, we have shown that these surveys can miss significant structures. The apparent majority of unresolved RQQs in contrast to the Seyferts in our dataset can be attributed to selection and resolution effects.\\\\\\indent Including literature data of Seyfert NLRs and their radio emission, we find a clear correlation between the size of the NLR and the size of the radio emission. We conclude that there is no significant morphological difference between Seyfert galaxies and radio--quiet quasars on scales of typical NLR sizes (hundreds of pc to kpc), neither in the radio nor in the [\\ion{O}{iii}] emitting region. Moreover, the interaction of radio--ejecta with the NLR gas seem to be equally important in both types of sources for shaping the structure of the NLR {\\sl and} radio emission."
    },
    "0606/astro-ph0606589_arXiv.txt": {
        "abstract": "We present a new time-stepping criterion for $N$-body simulations that is based on the true dynamical time of a particle. This allows us to follow the orbits of particles correctly in all environments since it has better adaptivity than previous time-stepping criteria used in $N$-body simulations. Furthermore, it requires far fewer force evaluations in low density regions of the simulation and has no dependence on artificial parameters such as, for example, the softening length. This can be orders of magnitude faster than conventional ad-hoc methods that employ combinations of acceleration and softening and is ideally suited for hard problems, such as obtaining the correct dynamics in the very central regions of dark matter haloes. We also derive an eccentricity correction for a general leapfrog integration scheme that can follow gravitational scattering events for orbits with eccentricity $e\\rightarrow1$ with high precision. These new approaches allow us to study a range of problems in collisionless and collisional dynamics from few body problems to cosmological structure formation. We present tests of the time-stepping scheme in $N$-body simulations of 2-body orbits with eccentricity $e\\rightarrow1$ (elliptic and hyperbolic), equilibrium haloes and a hierarchical cosmological structure formation run. ",
        "introduction": "Achieving spatial adaptivity in the evaluation of forces in $N$-body simulations is a well-studied problem with many effective approaches based on the use of tree structures and multipole expansions or nested grids and FFT techniques. Such adaptivity in space also comes with a desire to achieve adaptivity in the time integration of these simulations since a large dynamic range in density implies a large dynamic range in time-scales for self gravitating systems ($T \\sim 1/\\sqrt{G \\rho}~$). Two very different problems present themselves when trying to achieve this. First, there are no practical (explicit), general purpose (applicable to a wide range of astrophysical problems), adaptive integration techniques known for the $N$-body problem which are symplectic. By this we mean that the numerical integration is an exact Hamiltonian phase flow very close to the phase flow of the continuous system under study. This is a very desirable property for following systems for longer than a single dynamical time. Such exact preservation of the geometrical properties of the dynamical system is possible for fixed time-step schemes. For general $N$-body simulations with adaptive time-stepping we have to resort to approximate symplectic behaviour or preservation of time symmetry (a property which is known to lead to very good integration methods). The second problem is that of continuously determining the appropriate time-step for each particle in the simulation so that the error in the integration remains within tolerance while performing the fewest possible force evaluations and minimising the computational overhead. Resolving this second problem is the main focus of this paper and can be considered independently from methods symmetrizing the time-stepping scheme such as presented in \\cite{2001PhDT........21S} and \\cite{2006NewA...12..124M}.  There are several known time-step criteria based on different properties of the simulation (e.g. local density $\\rho(r)$, potential $\\Phi(r)$, softening $\\epsilon$, acceleration $a$, jerk $\\dot{a}$ or even the velocity $v$ of the particle) that are used in state-of-the-art numerical codes \\citep{1997astro.ph.10043Q, 2001PhDT........21S, 2003gnbs.book.....A, 2003MNRAS.338...14P, 2005MNRAS.364.1105S}. Some of them have a physical motivation, others are just a clever combination of physical properties in order to obtain a criterion which has the physical unit of time. All are an attempt to find an inexpensive way of determining an appropriate time-step for each particle in the simulation. For cosmological simulations, the ad-hoc time-step criterion based on the acceleration and the gravitational softening of the particle ($T \\sim \\sqrt{\\epsilon / a}~$) has proven very successful, despite the fact that it is not directly related to the dynamical time in these simulations. One reason why this time-step criterion is thought to work well is that it results in a very tight time-step distribution with very infrequent changes in the time-step of a particle in block time-stepping schemes (power of 2 step sizes; see also section \\ref{chap:implementation}). This hides the evils due to the first problem since the behaviour is more like that of a fixed time-stepping scheme than for other criteria. The price, however, is that more time-steps are taken in lower density regions of the simulation than would seem to be necessary. Furthermore, while still adequate for the type of simulations that have been performed up to now, which adapt the softening and mass of particles in the highest density regions and thus reduce the time-step somewhat artificially in those regions, this time-stepping scheme is no longer effective in simulations covering a much larger dynamic range.  In state-of-the-art computer simulations, structures can be resolved by $\\Nvir \\approx O(10^7)$ which results in a resolution scale of $r_{\\mathrm{res}} \\approx O(10^{-3}~\\rvir)$. By comparing the dynamical scales at $\\rvir$ and $r_{\\mathrm{res}}$, we get $T_{\\mathrm{dyn}}(\\rvir)/T_{\\mathrm{dyn}}(r_{\\mathrm{res}}) \\approx O (10^3)$ or even larger for future high resolution simulations. This large dynamical range in time-scales demands an acutely adaptive criteria so that the dense, dynamical active regions are resolved correctly while the simulations remain fast.  However, in a general simulation, such as those used for cosmological structure formation, it is not straightforward to determine the dynamical time of a given particle. This depends on the dominant structure responsible for the orbit of a particle which needs to be quickly determined at each time-step as the particle is advanced. In this paper we develop a new fast method of determining each particle's true dynamical time using information directly computed in the force evaluation stage of the simulation. This is quite different from using the local dynamical time which fails dramatically under many circumstances, e.g. consider using the local density near the earth to estimate its time-step!  We present in Section \\ref{chap:TS} the basic ideas and our implementation of the new adaptive dynamical time-stepping criterion into a tree-code. Detailed and extensive tests are presented in Section \\ref{chap:TSTest}. In Section \\ref{chap:conclusions} we conclude. ",
        "conclusions": "\\label{chap:conclusions}  We have developed a physically motivated time-stepping scheme that is based on the true dynamical time of the particle. We also derive an eccentricity correction for a general leapfrog integration scheme. The combination of these schemes allows us to follow quite general dynamical systems that may contain a mixture of collisionless and collisional interacting components. Compared to the standard time-stepping scheme used in many $N$-body codes it has the following advantages:  \\begin{enumerate} \\item It does not depend directly on ad-hoc parameters such as the softening length $\\epsilon$. \\item It gives physically correct time-steps in dark matter haloes with arbitrary central cusp slopes. \\item It is faster in high resolution simulations. \\item It allows orbits with eccentricity $e\\rightarrow1$ to be followed correctly. \\item It allows us to follow complex dynamical systems where scattering events may be important. \\end{enumerate}  The main conclusion is that one should use a time-step criterion that is based on the dynamical time. This scheme shows the optimum scaling with the number of particles and always gives a physically motivated time-step."
    },
    "0606/hep-th0606031_arXiv.txt": {
        "abstract": "{ We present a general set-up for inflation in string theory where the inflaton field corresponds to Wilson lines in compact space in the presence of magnetic fluxes.  $T$-dualities and limits on the value of the magnetic fluxes relate this system to the standard D-brane inflation scenarios, such as brane-antibrane inflation, D3/D7 brane inflation and different configurations of branes at angles. This can then be seen as a generalised approach to inflation from open string modes. Inflation ends when the Wilson lines achieve a critical value and an open string mode becomes tachyonic. Then hybrid-like inflation, including its cosmic string remnants, is realized in string theory beyond the brane annihilation picture. Our formalism can be incorporated within flux-induced moduli stabilisation mechanisms in type IIB strings. Also, contrary to the standard D-brane separation, Wilson lines can  be considered in heterotic string models. We provide explicit examples to illustrate similarities and differences of our mechanism  to D-brane inflation. In particular we present an example in which the $\\eta$ problem present in brane inflation models is absent in our case. We have examples with both blue and red tilted spectral index and remnant  cosmic string tension $G\\mu \\lesssim 10^{-7}$.}   \\preprint{DAMTP-2006-22\\\\  hep-th/0606031} ",
        "introduction": "There are several classes of moduli fields in string compactifications. Closed string moduli include the geometric and dilaton fields, with the geometric ones separated into the deformations of the  K\\\"ahler structure and deformations of the complex structure for Calabi-Yau and related compactifications. Open string moduli include the location of branes as well as Wilson lines which correspond to background values of gauge fields, usually coming from the open string sector of the type I/II string theory. In heterotic string however,  Wilson lines are closed string modes.  Most of these fields have been proposed as inflatons in terms of explicit mechanisms to realize cosmological inflation in string theory. There are now concrete models for which the inflaton is a K\\\"ahler modulus \\cite{racetrack, conlon} and several scenarios in which the modulus is a brane separation \\cite{dvalitye,burgess, d3d7, rabadan, bcsq, dbi}. A general property of brane inflation models is that inflation terminates by means of a tachyon condensation mechanism similar to hybrid inflation. This open string tachyon itself has been also proposed to be the inflaton field in other scenarios \\cite{cqs}. Each of these  scenarios has distinct physical implications which can be put to test by observations, regarding remnants of inflation such as cosmic strings, the value of the spectral index, etc.  Obtaining several realizations of inflation within string theory is a major achievement of the past five years. It would be highly desirable to have a general set-up that includes at least general classes of scenarios and try to extract as much model independent predictions. At the moment we are still at the stage of exploring the different possible realizations of inflation within string theory. For instance, there is not yet a concrete proposal for having the complex structure moduli as inflatons, something that could be desirable. Also Wilson lines have not been yet used as inflatons in an explicit string construction\\fn{Notice however that  the authors of \\cite{berg} computed loop corrections to a variation of the KKLMMT model in terms of Wilson lines. Also, the possibility of using Wilson lines as inflatons was considered in \\cite{creminelli} in a five dimensional effective field theory setup.}. The purpose of this paper is to explore this possibility.  Wilson lines correspond to background values of gauge fields that can be turned on in spaces of non-trivial homotopy. They provide most of the moduli in toroidal compactifications of the heterotic string parametrising the Narain lattices \\cite{nsw}. Discrete and continuous Wilson lines have also been used to break the gauge symmetries in phenomenological attempts to realistic models \\cite{GSW, inq, aiq}. Furthermore, in the context of D-branes, Wilson lines have been playing a crucial role to understand dualities in type I/II strings \\cite{polchi}. A $T$-duality transforms branes at angles in such a way that the angles are mapped to magnetic fluxes and D-brane separations are mapped to Wilson lines in the dual model. Since brane separations have been successfully used as inflatons, it is natural to study the Wilson lines as inflatons. Furthermore, this formalism is the one that can naturally extend to the heterotic case in which the D-brane separation has no direct analogue.  We present our results as follows. Section 2 is dedicated to introduce the basic ideas and derive, in the case of toroidal compactifications, the string amplitude induced when magnetic fields and Wilson lines are turned on. We also comment on the fact that all potentials used in the literature as starting points for brane/antibrane inflation, D3/D7 inflation and D-branes at angles can be obtained from the amplitude we consider after performing $T$-dualities in different directions, as well as taking limits on the values of the magnetic fields. This means that all toroidal D-brane inflationary configurations can be obtained starting only with D9 branes with magnetic fluxes and Wilson lines in Type I string theory. The next section is dedicated to explore the Wilson line potential extracted from this amplitude for inflation and a concrete model of inflation is presented. Section 4 discusses the supergravity description of this set-up and compares with the KKLMMT approach to brane antibrane inflation in the context of moduli stabilization. We present our conclusions in section 5. We have an appendix explaining in detail the duality between Wilson lines and brane separation as well as the structure of Wilson lines in general. A further appendix discusses an explicit model in which the dilaton, complex structure and the location of D7 branes are fixed by RR and NS-NS fluxes but Wilson lines remain flat, leaving them as the right candidates for the inflaton field. ",
        "conclusions": "We have presented a  general set-up for inflation in string theory with open string modes as inflatons. Wilson lines can play the role of inflaton fields in a way similar to brane separations. In fact Wilson lines are T-dual to brane separations and therefore it is expected that they play a similar role. The physics of both systems is the same but in different regimes (large against small volume). Therefore if we are only interested on effective field theory descriptions then we would need only the large internal volume potential on each formalism and they are certainly not equivalent.  It is interesting that starting with a particular configuration of D-branes, fluxes and Wilson lines, all known proposals for open string inflation can be included by T-duality and limits on the value of the magnetic fluxes. Furthermore it is reassuring to know that the end of inflation, with corresponding reheating and topological remnants, such as cosmic strings, can go by the standard tachyon condensation mechanism.  The explicit models presented in section 3 provide examples of more possibilities than previously considered, such as the inclusion of the location of the D7 branes and the number of  parameters involved, like magnetic fluxes, that can parametrise a fully realistic treatment of inflation. It is worth pointing out that both blue and red tilted values of the spectral index can be obtained and that the string scale tends to be of order the GUT scale leading to the remnant cosmic string tension to be of order $G\\mu \\lesssim 10^{-7}$.  For a discussion with moduli fixing a la KKLT, our formalism is in general very similar to the brane/antibrane case, including the need to confront the $\\eta$ problem. But again there are more parameters that can be varied to contrast with experimental signatures such as the number of efoldings, the spectral index and the COBE normalized $\\delta\\rho/\\rho$. Remarkably, we found that the $\\varphi$ dependence in $V_0$ does cancel for the example of section 3 after including moduli stabilisation, as discussed in section 4, and therefore there is no $\\eta$ problem, so the phenomenological results of section 3 hold after moduli stabilisation. Notice that even though there is no $\\eta$ problem in these models, inflation is certainly not generic, as we needed to have the moduli, both geometric and D7 positions, as well as the magnetic fluxes, in particular ranges in order to satisfy the requirements of slow-roll, tachyon condensation and COBE normalisation. Still, it is encouraging that the fine-tuning required by the $\\eta$ problem generic in D-brane inflation models is not present. It would be interesting to understand how general this cancellation is. For this it would be worth exploring more general Calabi-Yau compactifications with four cycles having non-trivial two-cycles carrying either magnetic fluxes or Wilson lines.    Notice that without a moduli fixing scheme, branes at angles provided a more flexible approach to inflation than the brane/antibrane system. However in the extension to include moduli fixing, only the brane/antibrane system has been considered so far. % From the discussion above we can see that this omission is corrected here but  in terms of the dual version with Wilson lines, instead of brane separation,  as inflatons and magnetic fluxes representing the brane angles.  There are further open questions that may be worth exploring following our results. An explicit calculation on the heterotic string would be interesting, to have a concrete realization of inflation in the heterotic case similar to brane inflation in type II. For this nonsupersymmetric compactifications with Wilson lines would naturally provide both, the constant  and interaction terms as in the case we discussed in section 3. Also, it would be interesting to check how the same cancellation we found to obviate the $\\eta$ problem holds in T-dual configurations in terms of branes at angles. Finally, as stressed in the appendix, Wilson lines will be present whenever the homotopy of the submanifold wrapped by a given D-brane is not trivial, so that our mechanism is not restricted to the case of the torus. There are known examples of compactification manifolds richer than tori, which allow for cycles with non-trivial 1-homology. It would be interesting to see how our mechanism applies for these kind of manifolds. In particular, a complete discussion of our mechanism including warped geometries would be desirable."
    },
    "0606/hep-ph0606075_arXiv.txt": {
        "abstract": "\\noindent We examine non-thermal dark matter production from a late-decaying scalar field, with a particular attention on non-renormalizable operators of $D=5$ through which the scalar field decays into the standard model particles and their superpartners. We show that almost the same number of superparticles as that of particles are generally produced from the decay.  To avoid the gravitino overproduction problem, the decay is favored to proceed via interactions with an intermediate cut-off scale $M \\ll M_P$.  This should be contrasted to the conventional scenario using the modulus decay. The bosonic supersymmetry partner of the axion, i.e., saxion, is proposed as a natural candidate for such late-decaying scalar fields. We find that a right amount of the wino/higgsino dark matter with a mass of $O(100)$ GeV is obtained for the saxion mass around the weak scale and axion decay constant, $F_\\A = O(10^{9-12})$ GeV. ",
        "introduction": "\\label{sec:1}  The lightest superparticle (LSP) is one of the most plausible candidates for the dark matter (DM) in supersymmetric (SUSY) models with $R$-parity conservation. In the SUSY standard model, the lightest neutralino, composed of superpartners of the gauge and Higgs bosons,  is usually the LSP. From both the theoretical and experimental/observational points of view, the wino/higgsino DM is interesting, and it has been studied extensively so far~\\cite{AM,Mirage,recent}. In this paper, we will concentrate on scenarios of the wino/higgsino DM production.  The present DM abundance is \\begin{eqnarray} \\label{eq:dm} \\Omega_{\\rm DM} h^2 \\;=\\; 0.11 \\pm 0.01, \\end{eqnarray} from the latest WMAP three year data~\\cite{Spergel:2006hy}, where $h$ is the present Hubble parameter in units of 100km/s/Mpc. The relic abundance of the LSP must fall in this range to account for the DM abundance. The relic LSP abundance depends on the thermal history of the universe. Once the universe is reheated up to high temperature after inflation, the LSP is thermally produced by particle scatterings and reaches thermal equilibrium. However, it is known that the thermal processes cannot provide a sufficient amount of the wino/higgsino LSP with a mass, $m_\\chi = O(100)$ GeV, because of the large annihilation rates. Therefore, non-thermal production must take place below the decoupling temperature, $T_\\chi \\sim m_\\chi/20$, to realize the observed DM abundance.  The non-thermal DM production at such a low temperature may be realized by a late-time decay of a scalar field, $\\phi$, with a mass $m_\\phi$. Renormalizable interactions generically induce a too rapid decay into the standard model (SM) particles and their superpartners unless the couplings are suppressed by some symmetries or mechanisms. Thus the decay should proceed via non-renormalizable operators suppressed by a large cut-off scale. In this paper, we focus on the $D=5$ operators. We will show that the branching ratio of the scalar decay into the SM superparticles is comparable to that into the SM particles.  A modulus is one of such late-time decaying fields~\\cite{Moroi:1999zb}, because the interactions of the modulus field are suppressed by the Planck scale, $M_P = 2.4 \\times 10^{18}$ GeV.  The decay temperature of the modulus field is likely lower than the decoupling temperature, leading to  the non-thermal production of the wino/higgsino LSP.  However, according to the recent works~\\cite{Endo:2006zj,Kawasaki:2006gs, Asaka:2006bv, Dine:2006ii,Endo:2006tf}, the late-time scalar decay is generically plagued with the gravitino overproduction.  Thus the cosmological scenario with a modulus field is strongly constrained, for example, from the big--bang nucleosynthesis (BBN).  We point out that a natural source of the non-thermal DM production is related to {\\it new physics at an intermediate scale}, $M \\ll M_P$. The gravitino overproduction problem is then relaxed. This is because the total decay rate is enhanced as $\\Gamma \\sim m_\\phi^3 / M^2$, reducing the branching ratio of the gravitino production. Note also that, since a light scalar field becomes acceptable, the decay into the gravitinos may be kinematically forbidden.  In the following sections we will show that a right amount of the wino/higgsino DM is obtained by the scalar decay at an intermediate scale.  We propose the saxion field, which is the bosonic SUSY partner of the axion, as a natural candidate for such a scalar field.  Before closing the introduction, let us clarify differences from the works in the past, which discussed the non-thermal DM production from late-time decaying scalar fields. The decay from a heavy scalar with Planck-suppressed couplings, i.e., a modulus field, has been extensively discussed in Ref.~\\cite{Moroi:1999zb}. Also, Ref.~\\cite{Khalil:2002mu} analyzed a similar subject with a lower cut-off, although they regarded the branching ratios of the superparticle as a free parameter. In this paper, we explicitly show that the branching ratio of the SM superparticle production is generically large, and that the effective number of the LSP produced by one scalar decay, $N_\\chi$, is close to unity. This has a great impact on the non-thermal DM production scenarios.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:2} we estimate the branching fraction of the SM (super)particle production for the scalar decay through the non-renormalizable operators of $D=5$. We reexamine the conventional scenario using the modulus decay for the non-thermal DM production in Sec.~\\ref{sec:3}. It is shown that the saxion decay can account for the present DM abundance in Sec.~\\ref{sec:4}. The last section is devoted to conclusion. ",
        "conclusions": "\\label{sec:5}  In this paper we have pointed out that an intermediate scale physics is necessary for a successful non-thermal production of the wino/higgsino DM, to avoid the moduli-induced gravitino/LSP problem recently pointed out in Refs.~\\cite{Endo:2006zj}. This is because gravitinos and/or LSPs would be overproduced otherwise.  The conventional scenario using the modulus decay to produce the wino/higgsino DM non-thermally is disfavored from this point of view. Instead we have proposed an alternative candidate: the bosonic SUSY partner of the axion, i.e., the saxion. It is particularly interesting that the right amount of the DM can be realized for the saxion mass around the weak scale when we consider the axion decay constant within the cosmologically allowed range, $F_\\A = 10^{9-12}$ GeV.  Although our analyses  focused on the case of the saxion, we would like to stress that they are rather generic and can be applied to any scalar fields that decay via the non-renormalizable operators of $D=5$ with an intermediate cut-off scale.  It is also worth noting that the saxion decay is applicable for a source of the harmless entropy production. As was stressed in this paper, the late-time decay of the scalar field generally suffers from the gravitino/LSP overproduction. In contrast, the scalar field at an intermediate scale can reheat the universe avoiding overproduction of the gravitino/superparticles.  For instance, a scalar field at $M \\ll M_P$ may decay before BBN begins, while satisfying $m_\\phi < 2 m_{3/2}, 2 m_{\\rm LSP}$ to kinematically block the gravitino/LSP production channels."
    },
    "0606/astro-ph0606494_arXiv.txt": {
        "abstract": "We present optical observations of the low-luminosity Seyfert 1 nucleus of NGC\\,4395, as part of a multiwavelength reverberation-mapping program. Observations were carried out over two nights in 2004 April at Lick, Wise, and Kitt Peak Observatories. We obtained $V$-band and $B$-band photometry, and spectra over the range 3500--6800 \\AA. Simultaneous {\\it Hubble Space Telescope} UV and {\\it Chandra} X-ray observations are presented in companion papers. Even though NGC\\,4395 was in an extremely low state of activity, we detect significant continuum variability of 2--10\\%, increasing toward shorter wavelengths. The continuum light curves, both spectroscopic and photometric, are qualitatively similar to the simultaneous UV and X-ray light curves. Inter-band cross-correlations suggest that the optical continuum emission lags behind the UV continuum emission by $24^{+7}_{-9}$ min, and that the optical continuum emission lags behind the X-ray continuum emission by $44\\pm13$ min, consistent with a reprocessing model for active galactic nucleus emission. There are also hints of Balmer emission lines lagging behind the optical continuum by an amount slightly larger than the emission-line lag detected in the UV. These results are all similar to those of other Seyfert 1 nuclei. The emission-line lag yields a mass measurement of the central black hole, which although not very significant, is consistent with the value derived from the simultaneous UV data. ",
        "introduction": "The nucleus of NGC\\,4395 is a ``dwarf'' type 1 Seyfert active galactic nucleus (AGN), the least luminous known in its class \\citep{FS89, FHS93}. Its presence in an essentially bulgeless, extremely late-type galaxy is surprising; such AGNs (and quiescent black holes) are usually found in bulge-dominated early-type systems \\citep{HFS97}, with black hole mass and bulge mass strongly correlated \\citep{Magorrian98, Laor98, MH03, HR04}. Despite this, the nucleus of NGC\\,4395 exhibits all of the hallmarks of a type 1 AGN: it is detected as a variable X-ray point source \\citep{Lira99, Moran99, Moran05, Iwasawa00, Ho01, SIF03}, it has a high brightness temperature, highly compact radio core \\citep{Moran99, HU01, WFH01}, and its spectral energy distribution from X-ray to radio wavelengths is similar to that of other type 1 AGNs \\citep{Moran99}.  NGC\\,4395 offers a chance to probe the low end of the supermassive black hole (SMBH) mass distribution. Knowing the mass of the central object will help constrain the $M_{\\mathrm{BH}}-\\sigma_*$ relation \\citep{Gebhardt00,Ferrarese01,Tremaine02}, which suggests a strong link between the growth of the SMBH and evolution of the host galaxy.  The SMBH mass can be measured through emission-line reverberation mapping \\citep{BM82,Peterson93,NP97}.  The relations of \\citet{Kaspi00,Kaspi05} between continuum luminosity and radius of the broad-line region (BLR) suggest that the size of the Balmer line-emitting region in NGC\\,4395 is of order light hours.  We have used the reverberation method successfully in the ultraviolet (UV) with the {\\it Hubble Space Telescope} ({\\it HST}) to determine the size of the BLR and the mass of the central object, measured to be $M_{\\mathrm{BH}} = (3.6 \\pm 1.1) \\times 10^5 \\ \\Msun$ \\citep[hereafter Paper I]{Peterson05}. We present here simultaneous optical observations that complement the UV results. Simultaneous X-ray observations with {\\it Chandra} are presented by \\citet[hereafter Paper II]{ONeill06}. The observations and data reduction are outlined in Section~\\ref{sec:obs}, while the data analysis is described in Section~\\ref{sec:analysis}. In Section~\\ref{sec:disc}, we compare our results with those published for other AGNs and estimate the SMBH mass. We summarize our results in Section~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  We observed NGC\\,4395, the least luminous known Seyfert 1 galaxy, from five different ground-based optical telescopes over the course of 2 nights in 2004 April, as part of a coordinated multiwavelength reverberation-mapping campaign designed to measure the mass of the central black hole. The results of our observations are as follows.  \\begin{enumerate}  \\item NGC\\,4395 was in a very low state of activity during this period, with optical continuum flux levels, variability amplitudes, and broad emission-line fluxes all diminished compared with earlier epochs.  \\item Despite this, we detect intrinsic continuum variability, from both spectra and broad-band photometry, ranging from about 2\\% to 10\\% ($F_{\\mathrm{var}}$). In general, we clearly see evidence for increased variability toward shorter wavelengths, consistent with other typical AGNs. This trend extends into the simultaneous UV and X-ray bands. The overall shape of the optical continuum light curves agrees with the simultaneous UV and X-ray light curves. Variability was greatest on April 10, while April 11 exhibited a slightly shallower dependence of variability on wavelength.  \\item Cross-correlation of the various continuum light curves suggests that the optical emission lags behind the UV emission by $24^{+7}_{-9}$ min, and the optical emission lags behind the X-ray emission by $44\\pm13$ min. There is no formal detection of any lag between the UV and X-ray emission, as described in Paper II, and between the $B$ and $V$ bands. These results are consistent with the reprocessing model, in which X-ray photons are reprocessed into UV and optical photons. Our high sampling frequency has allowed us to make these measurements, which have previously been very difficult to achieve.  \\item We only formally detect a lag between the \\ha \\ light curve and the $V$-band light curve. Although all other emission-line light curves are consistent with zero lag, the trend as a whole is quite suggestive.  In particular, averaging the Balmer-line correlations with the optical continuum yields an emission-line lag of $80\\pm68$ min behind the continuum. While not very significant, this result is consistent with the relations of \\citet{Kaspi00} for the size of the BLR. The optical emission-line lag is also about 1.5 times larger than the UV emission-line lag behind the UV continuum. We have seen this behavior in other AGNs, which indicates NGC\\,4395 is behaving typically, albeit faintly, for an AGN.  \\item Using the Balmer emission-line lag and the FWHM of the variable broad-line region, we estimate the mass of the SMBH to be $M_{\\mathrm{BH}} \\approx 3\\times 10^5 \\Msun$. Again, this measurement is not very significant, and is meant only to compare with the value of $M_{\\mathrm{BH}} = (3.6 \\pm 1.1) \\times 10^5 \\ \\Msun$ derived in Paper I. The results are consistent.  \\end{enumerate}"
    },
    "0606/astro-ph0606388_arXiv.txt": {
        "abstract": "We propose  a practical method  to calculate Zernike  aberrations from analysis  of  a  single  long-exposure defocused  stellar  image.   It consists  in  fitting  the  aberration coefficients  and  seeing  blur directly to  a  realistic  image  binned  into  detector pixels. This \"donut\"  method is different from curvature  sensing in that it does not make the usual  approximation of  linearity.  We calculate  the sensitivity  of  this  technique  to  detector and  photon  noise  and determine optimal  parameters for some  representative cases. Aliasing of high-order un-modeled  aberrations is evaluated and  shown to be similar  to a  low-order Shack-Hartmann  sensor. The  method  has been tested  with  real data  from  the   SOAR and Blanco 4m telescopes. \\\\ ",
        "introduction": "An experienced optician can detect low-order aberrations by looking at the defocused  image of a  point source, and  it is trivial  to obtain defocused images  with modern telescopes equipped  with CCD detectors. Yet, measurements  of low-order  aberrations including focus  are still made  by  indirect techniques,  or  using  special  equipment such  as Shack-Hartmann (S-H)  sensors.  Astronomers spend  significant time in acquiring  ``focus sequences''  of  stellar images,  then fitting  the image  half-width  vs.  focus  curve  with  a  parabola  to  find  the best-focus position.  The appeal of estimating aberrations directly from defocused images is evident.   No  special equipment  is  needed  apart  from a  regular imager.   The  aberrations in  the  true  science  beam are  measured, including all optics of  the instrument but excluding additional optics of a wave-front  sensor.  The amount of defocus  is easily adjustable, providing flexibility.  It  was  recognized  since long  time  that optical  aberrations  cannot  be retrieved   from  a   focused  image   of  a   point   source  without ambiguity. However, combining {\\em two} images with a known difference of aberration provides a solution  to this problem, even for non-point sources.  The method of {\\it phase diversity} which exploits this idea has been used since the beginning of the 80-s \\citep{Thelen99}.  Phase diversity  works well  when the  image is  sampled to  the diffraction limit, e.g. in adaptive optics  \\citep{Hartung}.  This is not the case for conventional astronomical imagery with a pixel size matched to the seeing.    Yet  another   method  for   extracting   aberrations  from well-sampled focused images  by means of a trained  neural network was suggested by Sandler and later tried by \\citet{LH92}. The authors note that their method is  extremely computationally intensive and has some subtleties. To our knowledge, this method is not in use nowadays.  The relation of the intensity distribution in a defocused image to the local  wavefront curvature  is described  by the  so-called irradiance transport equation  \\citep{Roddier90}. This relation is  basic to {\\it curvature sensing}  as used  in adaptive optics  \\citep{Roddier99}.  A commercial software package for telescope aberration analysis based on the   same  principle  has   been  developed   by  Northcott\\footnote{ Northcott,  M.J., The  {\\it  ef} wavefront  reduction package.   1996, Laplacian  Optics  Inc.}  and  is  used  at  some observatories.  This method, however, is not very  practical because it requires two images with relatively large and equal defocus of opposite sign.  The need  of two images for  curvature sensing has  been questioned by \\citet{Hickson94}  who shows  that  even in  the  context of  adaptive optics a single extra-focal image  is sufficient and provides a better signal-to-noise  ratio with  a  CCD detector  and  faint guide  stars, despite   scintillation   noise.    One   image  is   sufficient   for non-ambiguous  aberration retrieval  as long  as the  rays originating from different  parts of  the aperture do  not cross each  other, i.e. for a  sufficiently large defocus  that avoids caustics.   The minimum defocus is proportional to  the amplitude of higher-order aberrations. \\citet{Ragazzoni} have used this technique in their experiment.  The intensity  transport equation  is not valid  for a  small defocus, where physical  optics must be used  instead. However, this  is not an obstacle for sensing low-order aberrations,  as long as they are small enough,  so that  a relation  between aberration  and  image intensity remains  linear.   \\citet{Bharmal}  develop  such  near-focus  sensing technique  for low-order  adaptive  optics, providing  in their  paper several valuable  insights into  this problem.  However,  their method still requires two images, intra- and extra-focal.  Here we present a quantitative method of measuring optical aberrations from  a single  defocused image.   Such images  often  resemble donuts (because of the shadow at the center caused by the central obscuration in a Cassegrain telescope), so we call this technique ``donut''.  This work  is primarily  motivated  by  the need  for  a simple  wave-front sensing     method    for    the     SOAR    telescope     in    Chile \\citep{Sebring98,Krab04}.  All numerical  examples in the article were computed for a telescope diameter $D=4.1$~m with a central obscuration 0.24,  appropriate  for SOAR.   The  proposed  technique is  primarily intended for  active optics, it  is too slow for  real-time correction of turbulence.  The donut method is different  from standard curvature sensing. We use physical optics and directly fit a model of the aberrated image to the measured  ``donut''. The  initial approximation  is obtained  from the second  moments   of  the  intensity  distribution   as  described  in Sect.~\\ref{sec:mom}.  Then an iterative fitting algorithm presented in Sect.~\\ref{sec:fit}, with further details  in the Appendix, is used to refine   the   model   including   higher   order   aberrations.    In Sect.~\\ref{sec:perf} we evaluate the errors of aberrations measured by this method and compare it  to a low-order Shack-Hartmann sensor while examples of  actual  performance are  given in  Sect.~\\ref{sec:exa}. Finally we present our conclusions in Sect.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We   have  shown  that   focus  and   astigmatism  can   be  evaluated quantitatively from the second moments of defocused images. One useful application  of this  analysis will  be a  fast and  accurate focusing procedure for  classical imaging, suggested  here as a  replacement of traditional  focusing  sequences.  Furthermore,  donut  images can  be fitted directly to a set of Zernike coefficients (complemented with an additional  parameter, seeing),  offering a  practical way  to measure aberrations and to tune the optics of ground-based telescopes.  The  donut  method  proposed  here  is  different  from  the  standard curvature sensing in several aspects.  First, only one defocused image is needed.   Second, no simplifying  assumption of linearity  is made, hence the  defocus may be  quite small while measuring  aberrations of significant amplitude -- comparable to the defocus itself. Third, we do not use the intensity  transport equation \\citep{Roddier90} but rather compute  the image  model by  a full  Fraunhofer  diffraction integral using an FFT.  Finite  detector pixel size and additional  blur caused by the  seeing are  explicitly  taken into  account.   These two  effects usually  wash  out  any  traces  of  diffraction,  so  the  calculated monochromatic image is a good model of a wide-band image as well.  The  down-side  of the  full  diffraction  image  modeling is  a  slow calculation time (a few  seconds for a 4-m telescope)  and a restriction of the  modeled field of  view.  The donut  method will work  best for small defocus  values and for measuring low-order  aberrations. On the other  hand, classical curvature  sensing would  be probably  a better choice  for high-resolution  sensing, where  a wave-front  map (rather than Zernike coefficients) is sought.  We plan to apply the donut technique to the closed-loop control of the SOAR  active  optics  and   to  optical  characterization  of  other telescopes at CTIO. The method seems to be simple and mature enough to be offered to other interested users. So far, it is implemented in the IDL language."
    },
    "0606/astro-ph0606611_arXiv.txt": {
        "abstract": "{This is the first paper of a series describing our measurement of weak lensing by large-scale structure, also termed ``cosmic shear'', using archival observations from the Advanced Camera for Surveys (ACS) on board the Hubble Space Telescope (HST).} { In this work we present results from a pilot study testing the capabilities of the ACS for cosmic shear measurements with early parallel observations and presenting a re-analysis of HST/ACS data from the GEMS survey and the GOODS observations of the \\textit{Chandra} Deep Field South (CDFS).} {We describe the data reduction and, in particular, a new correction scheme for the time-dependent ACS point-spread-function (PSF) based on observations of stellar fields. This is currently the only technique which takes the full time variation of the PSF between individual ACS exposures into account. We estimate that our PSF correction scheme reduces the systematic contribution to the shear correlation functions due to PSF distortions to \\mbox{$< 2 \\times 10^{-6}$} for galaxy fields containing at least 10 stars, which corresponds to $\\lesssim 5\\%$ of the cosmological signal expected on scales of a single ACS field. } {We perform a number of diagnostic tests indicating that the remaining level of systematics is consistent with zero for the GEMS and GOODS data confirming the success of our PSF correction scheme. For the parallel data we detect a low level of remaining systematics which we interpret to be caused by a lack of sufficient dithering of the data. Combining the shear estimate of the GEMS and GOODS observations using $96$ galaxies $\\mathrm{arcmin}^{-2}$ with the photometric redshift catalogue of the GOODS-MUSIC sample, we determine a \\textit{local single field estimate} for the mass power spectrum normalisation $\\sigma_{8,\\mathrm{CDFS}}=0.52^{+0.11}_{-0.15}\\mathrm{(stat)}\\pm0.07\\mathrm{(sys)}$ (68\\% confidence assuming Gaussian cosmic variance) at a fixed matter density $\\Omega_\\mathrm{m}=0.3$ for a $\\Lambda$CDM cosmology marginalising over the uncertainty of the Hubble parameter and the redshift distribution. We interpret this exceptionally low estimate to be due to a local under-density of the foreground structures in the CDFS. } {} ",
        "introduction": "Cosmic shear, the weak gravitational lensing effect of the large-scale structure, provides a powerful tool to constrain the total matter power spectrum without any assumptions on the relation between luminous and dark matter. Due to the weakness of the effect, it is challenging to measure, with the first detections only reported six years ago \\citep{bre00,kwl00,wme00,wtk00}. Since then cosmic shear has developed into a flourishing field of cosmology yielding not only constraints on the matter content $\\Omega_\\mathrm{m}$ and the normalisation of the power spectrum $\\sigma_8$ (\\citealt{mwm01,wmr01,wmp02,wmh05,hyg02,rrg02,bmr03,btb03,jbf03,hms03,hbh04,hbb05,rrc04,mrb05,hss06}; \\citealt{mhb07}), but recently also on the equation of state parameter $w$ of dark energy (\\citealt{jjb06,hmw06,smw06}; \\citealt{kht07}), see \\citet{bas01,mwb02,hyg02a,ref03} for reviews.  Due to the weakness of cosmological gravitational shear, proper correction for systematics, first of all the image point-spread-function (PSF) is indispensable. Within the Shear TEsting Programme (STEP) a number of algorithms to measure the shear from faint and PSF-distorted galaxy images are currently tested using a blind analysis of image simulations aiming to improve and quantify the accuracy of the different methods; see \\citet{hwb06}; \\citet{mhb07} for first results.  The majority of the previous cosmic shear measurements have been made with ground-based wide-field imaging data, mainly probing the matter power spectrum on linear to moderately non-linear angular scales from several degrees down to several arcminutes. In order to probe the highly non-linear power spectrum at arcminute and sub-arcminute scales, a high number density of usable background galaxies is required. While deep ground-based surveys are typically limited to $\\simeq 30$ galaxies/$\\mathrm{arcmin}^2$ due to seeing, significantly higher number densities of resolved galaxies can be obtained from space-based images. Cosmic shear studies have already been carried out with the HST cameras WFPC2 \\citep{rrg01,rrg02,crg03} and STIS \\citep{hms02,rrc04,meh05}. With the installation of the \\textit{Advanced Camera for Surveys} (ACS) \\textit{Wide-Field Channel} (WFC) detector, a camera combining improved sensitivity ($48\\%$ total throughput at 660 nm) and a relatively large field-of-view \\mbox{($\\sim 3\\farcm3 \\times 3\\farcm3$)} with good sampling (0\\farcs05 per pixel) \\citep{fch03,smc04}, the possibilities to measure weak lensing with HST have been further improved substantially. ACS has already been used for weak lensing measurements of galaxy clusters (\\citealt{jwb05,jwf05,jwf06,lrb05,cbg06,bcg06}; \\citealt{lgh07}) and galaxy-galaxy lensing \\citep{hbr06,gtr07}. The first cosmic shear analysis with ACS was presented by \\citet{hbb05}, \\citetalias{hbb05} henceforth, for the GEMS survey \\citep{rbb04}, a \\mbox{$\\sim 28^\\prime \\times 28^\\prime$} mosaic incorporating the HST/ACS GOODS observations of the \\textit{Chandra} Deep Field South \\citep[CDFS, ][]{gfk04}. Recently \\citet{lmk07} and \\citet{mrl07} presented a cosmological weak lensing analysis for the ACS COSMOS\\footnote{\\url{http://www.astro.caltech.edu/~cosmos}} field.   In this work we present results from a pilot cosmic shear study using early data from the ACS Parallel Cosmic Shear Survey (proposals 9480, 9984; PI J. Rhodes). Parallel observations provide many independent lines of sight reducing the impact of cosmic variance. With a separation of several arcminutes from the primary target (e.g. $\\sim 6^\\prime$ for WFPC2) they provide nearly random pointings for most classes of primary targets. This is important when measuring a statistical quantity like the shear. Still, for primary observations pointing at particularly over-dense regions, such as galaxy clusters, a significant selection bias might be introduced as they influence the shear field many arcminutes around them, which has to be checked carefully.   Analysing stellar fields we detect short-term variations of the ACS PSF, which are interpreted as focus changes due to thermal breathing of the telescope (see also \\citealt{kri03,rma05,ank06}; \\citealt{rma07}). Whereas earlier studies with other HST cameras assumed temporal stability of the PSF, a fully time-dependent PSF correction is required for ACS due to these detected variations. Given that only a low number of stars are present in high galactic latitude ACS fields \\mbox{($\\sim 10-30$)}, the correction cannot be determined from a simple interpolation across the field-of-view, but requires prior knowledge about possible PSF patterns. In this work we apply such a correction scheme, which is based on PSF models derived from stellar fields. Our method takes the full PSF variation between individual exposures into account and can be applied for arbitrary dither patterns and rotations. \\citet{rma05,rma07} propose a different correction scheme, in which they fit co-added frames with theoretical single-focus PSF models created with a modified version of \\texttt{TinyTim}\\footnote{\\url{http://www.stsci.edu/software/tinytim/tinytim.html}}.  \\citetalias{hbb05} use a semi-time-dependent model to correct for the image PSF, based on the combined stars for each of the two GEMS observation epochs. The tests for systematics presented by \\citetalias{hbb05} indicate zero contamination with systematics for the GEMS only data, but remaining systematics at small scales if the GOODS data are included. In order to test our fully time-dependent PSF correction we present a re-analysis of the GEMS and GOODS data, yielding a level of systematics consistent with zero for the combined dataset. We present a cosmological parameter estimation from the re-analysed GEMS and GOODS data in this work, whereas a parameter estimation from the parallel data will be provided in a future paper on the basis of the complete ACS Parallel Cosmic Shear Survey.  The paper is organised as follows: After summarising the weak lensing formalism applied in Sect.\\thinspace\\ref{se:method}, we describe the data and data reduction in Sect.\\thinspace\\ref{se:data}. We present our analysis of the ACS PSF and the correction scheme in Sect.\\thinspace\\ref{se:psf}. Next we elaborate on the galaxy selection and determined redshift distribution (Sect.\\thinspace\\ref{se:galaxy_select_redshift}) and compute several estimators for the shear and systematics in Sect.\\thinspace\\ref{se:preliminary_cosmic_shear}. After presenting the results of the cosmological re-analysis of the GEMS and GOODS data in Sect.\\thinspace\\ref{se:cosmo_para_estimate}, we conclude in Sect.\\thinspace\\ref{se:conclusion}. ",
        "conclusions": "\\label{se:conclusion}  We have presented a cosmic shear analysis of a first set of HST/ACS pure parallel observations and the combined GEMS and GOODS data of the CDFS. We estimate that our new correction scheme for the temporally variable ACS PSF reduces the systematic contribution to the shear correlation functions due to PSF distortions to \\mbox{$< 2 \\times 10^{-6}$} for galaxy fields containing at least 10 stars. This is currently the only technique taking the full time variation of the PSF between individual ACS exposures into account. In the GEMS and GOODS data the success of the PSF correction is confirmed by a number of diagnostic tests indicating that the remaining level of systematics is consistent with zero. For the parallel data we detect a low level of remaining systematics manifesting in a slight average alignment of the measured galaxy ellipticities in the $y-$direction, which we interpret to be due to a lack of proper dithering. We are currently further investigating this effect and exploring ways to correct for it, which will be necessary for the cosmic shear analysis of the complete set of ACS parallel observations. Although the degradation of the ACS charge-transfer-efficiency has not been found to be a problem for the early data analysed in this work, an in-depth analysis and correction will probably be required for the complete data set \\citep[see also][]{rma05,rma07}. Furthermore the parallel data are rather inhomogeneous regarding depth and extinction, raising the need for a well calibrated field-dependent redshift distribution. It will also be necessary to carefully exclude any selection bias which might arise for certain classes of primary targets, particularly galaxy clusters. Once these remaining obstacles are overcome, it will be possible to measure cosmic shear at small angular scales with unprecedented accuracy from the complete ACS Parallel Survey, with a strong reduction both of the shape noise and cosmic variance error due to many independent pointings. The main limitation of the cosmological interpretation of the data might then arise from the current accuracy of theoretical predictions for the non-linear power spectrum at small scales. An interesting comparison will be possible with the ACS COSMOS data \\citep{mrl07}, from which cosmic shear can be measured on a wide range of angular scales.  Given the high demands concerning the control over systematics for cosmic shear measurements with ACS, the derived technical expertise (see also \\citetalias{hbb05}; \\citealt{jwb05}; \\citealt{rma05}; \\citealt{rma07}; \\citealt{lmk07}) will also be of benefit for other weak lensing studies with the instrument, and possibly also other research fields requiring accurate PSF modelling.  Due to the weakness of the shear signal on the one hand, and the strong impact of poorly understood systematics on the other hand, an analysis of identical datasets with more than one independent pipeline is of great value to check the reliability of the algorithms employed. In this work we have independently re-analysed the ACS observations of the GEMS and GOODS fields. If we assume the same redshift parametrisation, our shear estimates are in excellent agreement with the earlier results found by \\citetalias{hbb05} indicating the reliability of both lensing pipelines. Such an independent comparison will also be highly desired both for the complete ACS Parallel Survey (Rhodes et al. in prep.) and the ACS COSMOS field \\citep{mrl07}. These comparisons, together with the results from the STEP project, will aid the preparations of future space-based cosmic shear survey such as DUNE or SNAP, which will reach a very high statistical accuracy \\citep{rmr04} requiring the continued advancement of improved algorithms such as shapelets (\\citealt{bej02,reb03,mar05,kui06}; \\citealt{nab07}).  Finally, we want to stress the possible impact of the field selection on a cosmic shear analysis: The \\textit{Chandra} Deep Field South was originally selected in a patch of the sky characterised by a low Galactic neutral hydrogen column density \\mbox{($N_\\mathrm{H}=8\\times 10^{19} \\mathrm{cm}^{-2}$)} and a lack of bright stars \\citep{grt01}. Additionally, it neither contained known relevant extragalactic foreground sources nor X-ray sources from the ROSAT ALL-Sky Survey Catalogue\\footnote{see \\url{http://www.mpe.mpg.de/~mainieri/cdfs_pub/index.html}} excluding e.g. the presence of a low redshift galaxy cluster. \\citet{ami05} present a detailed analysis of compact structures in the CDFS showing the presence of a chain-like structure at $z=0.66$, a massive group at $z=0.735$ embedded into a galaxy wall extending beyond the $21^\\prime \\times 21^\\prime$ field covered by the Vimos VLT Deep Survey \\citep{fvp04}, and a further massive group at $z=1.098$ \\citep[see also][]{gcd03,sbh04,vcd06}. \\citet{wmk04} identify a strong galaxy over-density at \\mbox{$z\\sim 0.15$}, which is too close to produce a significant lensing signal. Given the lack of massive structures at lower redshifts $0.3\\lesssim z \\lesssim 0.6$ with high lensing efficiency, one would expect to measure a shear signal biased to lower values in this field as a result of strong sampling variance. Therefore it is not surprising that our local single field estimate of \\mbox{$\\sigma_{8,\\mathrm{CDFS}} (\\Omega_\\mathrm{m}=0.3)=0.52^{+0.11}_{-0.15}\\mathrm{(stat)}\\pm0.07\\mathrm{(sys)}$} based on a source redshift distribution derived from the GOODS-MUSIC sample \\citep{gfs06}, is incompatible at the \\mbox{$\\sim 3 \\sigma$-}level assuming Gaussian cosmic variance with recent results of other weak lensing studies \\citep[e.g.][]{hmw06,smw06}, which probe much larger regions on the sky. \\citet{kis05} and \\citet{swh07} investigate the impact of non-Gaussianity on cosmic shear covariances. From their results we broadly determine an under-estimation of the cosmic variance contribution to our error on $\\sigma_8$ by a factor $\\approx 1.7$, indicating that the CDFS is still an exceptionally under-dense field, but with a lower significance (\\mbox{$\\sim 2 \\sigma$}) than under the Gaussian assumption. Our $\\sigma_8$ estimate is also significantly lower than the \\citetalias{hbb05} results of \\mbox{$\\sigma_8(\\Omega_\\mathrm{m}/0.3)^{0.65}=0.68\\pm0.13$} due to the deeper redshift distribution found in our analysis with a median source redshift \\mbox{$z_\\mathrm{m}=1.46 \\pm 0.12$}. Recently \\citet{pwp06} found a strong deficiency of faint red galaxies in the CDFS for the redshift range \\mbox{$0.25 \\lesssim z \\lesssim 0.4$} indicating a substantial under-density, which is in excellent agreement with the low shear signal found in our analysis.  We believe that the CDFS represents a somewhat extreme case. However, also other cosmic shear studies which observe a low number of small ``empty fields'' could be slightly biased just due to this prior selection. Such a bias can of course be eliminated either with the observation of sufficiently large fields or truly random pointings, which are realized in good approximation for a large fraction of the fields in the ACS Parallel Survey.   We plan to further investigate the peculiarity of the CDFS based on a shear tomography analysis with photometric redshifts derived for the full GEMS field, also using deep ground-based optical images from the MPG/ESO 2.2m telescope \\citep{hed06} in combination with infrared images from the ESO 3.5m NTT \\citep{omc06a,omc06b}. If the low estimate for $\\sigma_{8,\\mathrm{CDFS}}$ indeed stems from an under-density of foreground structures we would expect an increased shear signal for a high-redshift sample of source galaxies due to the spectroscopically confirmed structures at $z=0.735$ and $z=1.098$. Comparing the results with ray-tracing through N-body simulations we aim to further quantify the rarity of such an under-dense foreground field."
    },
    "0606/astro-ph0606561_arXiv.txt": {
        "abstract": "Dissipationless N-body models of rotating galaxies, iso-energetic to a non-rotating model, are examined as regards the mass in regular and in chaotic motion. Iso-energetic means that they have the same mass and the same binding energy and they are near the same scalar virial equilibrium, but their total amount of angular momentum is different. The values of their spin parameters $\\lambda$ are near the value $\\lambda=0.22$ of our Galaxy.  We distinguish particles moving in regular and in chaotic orbits and we show that the spatial distribution of these two sets of particles is much different. The rotating models are characterized by larger fractions of mass in chaotic motion (up to the level of $\\thickapprox 65\\%$) compared with the fraction of mass in chaotic motion in the non-rotating iso-energetic model (which is on the level of $\\thickapprox 32\\%$). Furthermore, the Lyapunov numbers of the chaotic orbits in the rotating models become by about one order of magnitude larger than in the non-rotating model. This impressive enhancement of chaos is produced, partly by the more complicated distribution of mass, induced by the rotation, but mainly by the resonant effects near corotation. Chaotic orbits are concentrated preferably in values of the Jacobi integral around the value of the effective potential at the corotation radius.  We find that density waves form a central rotating bar embedded in a thin and a thick disc with exponential mean radial profile of the surface density. A surprising new result is that long living spiral arms are excited on the disc, composed almost completely by chaotic orbits.  The bar excites an $m=2$ mode of spiral waves on the surface density distribution of the disc, emanating from the corotation radius. The bar goes temporarily out of phase with respect to an excited spiral wave, but it comes in phase again in less that a period of rotation. As a consequence, spiral arms show an intermittent behavior. They are deformed, fade, or disappear temporarily, but they grow again re-forming a well developed spiral pattern. Spiral arms are discernible up to 20 or 30 rotations of the bar (lasting for about a Hubble time). The relative power of the spiral $m=2$ mode with respect to all other fluctuations on the surface density is initially about $50\\%$, but it is reduced by a factor of about 2 or 3 at the end of the Hubble time. ",
        "introduction": "In previous papers some cases of non-rotating N-body models of elliptical galaxies have been investigated as regards their mass components in ordered and in chaotic motion (Voglis, Kalapotharakos \\& Stavropoulos 2002; Kalapotharakos, Voglis \\& Contopoulos 2004; Kalapotharakos \\& Voglis 2005). In these studies the fraction of mass in chaotic motion and its consequences were found to be important.  In particular, in Voglis et al. (2002) two non-rotating triaxial equilibrium models, called C and Q, resulting respectively  from clumpy (C) and quiet (Q) cosmological initial conditions after dissipationless collapse and relaxation, are examined. Both models have smooth centers, i.e. the density near the center flattens inside a radius of about 10\\% of the half mass radius. The C model is more close to spherical than the Q model having a maximum ellipticity E4. The mass in chaotic motion in C is found to be about $24\\%$ of the total mass. The Lyapunov numbers of the orbits are found in general small, so that only $2\\%$ of the total mass is able to develop chaotic diffusion in a Hubble time. The Q model has a high maximum ellipticity E7. The fraction of mass in chaotic motion in Q is found to be about $32\\%$ of the total mass and the values of the Lyapunov numbers are found higher than in C, so that a fraction of mass about $8\\%$ is able to develop chaotic diffusion in a Hubble time.  We have also found that the spatial distribution of the mass in chaotic motion is different than the spatial distribution of the mass in ordered motion, so that this difference can form an observable hump in the profile of the surface density.  The fact that chaos in models of non-rotating galaxies with smooth centers is not negligible was also found in Muzzio, Carpintero \\& Wachlin (2005). These authors found a fraction of mass in chaotic motion of about $53\\%$ in their non-rotating smooth center model, i.e. even larger than the fraction found in our models. One reason that explains this difference is the fact that their model is more cuspy, i.e. the density near the center flattens in a smaller radius than in our models. Another reason is that they have used a lower threshold of the Lyapunov numbers for distinguishing between regular and chaotic orbits. The two results are in agreement, in the sense that the fraction of mass in chaotic motion cannot be negligible (it is of the order of some tens percent) even in models with smooth centers that are expected to be mostly in favor of ordered motion. Another agreement with our results is that the mass in chaotic motion has a different spatial distribution than the mass in regular motion. However, the Lyapunov times of the vast majority of chaotic orbits are large compared with the mean dynamical times of galaxies (e.g. more than 10 dynamical times).   Models of non-rotating galaxies are the simplest possible models of galaxies, especially those with smooth centers. Rotation of galaxies introduces new parameters and nonlinear interactions that lead to more complicated dynamics. In rotating models density waves, able to form rotating bars or spiral arms, can grow.  The mechanism that is responsible for the formation and the stability of spiral arms in galaxies is a complicated problem that has not been very well understood yet. All researchers agree that the main driving force leading to this effect is gravity. Dissipative effects can also play a non-negligible role in the stability of the spiral arms, but gravity remains the main driving mechanism.  The frequent appearance of spiral arms in galaxies is in favor of the assumption that this is probably a quasi-stationary effect, i.e. spiral arms are long living features in galaxies. This assumption is known as the Lin-Shu Hypothesis (Lin \\& Shu 1964,1966; Binney \\& Tremaine 1987).  The dynamical evolution of galaxies can be studied by combining theoretical analysis with the results obtained by N-body simulations.  N-body simulations show that spiral arms appear as a transient effect, in the sense that they grow as density waves, but they are temporarily dissolved or fade and recreated again. This process may be repeated several times until they finally disappear.  A necessary ingredient for spiral structure to appear is the amount of angular momentum a galaxy possesses, as well as the way this angular momentum is distributed along the mass of the galaxy.  It is well known that resonances between the frequencies of stellar orbits and the angular velocity of the pattern $\\Omega_p$ play a significant role on the stability and the dynamical evolution of rotating discs. In particular the Inner Lindblad Resonance (ILR), Corotation (CR) and the Outer Lindblad Resonance (OLR) are the most important resonances. Lynden-Bell \\& Kalnajs (1972) have shown that in a rotating disc angular momentum is transferred outwards. This transference is due to the interaction of a bar-like perturbation of the gravitational field with stars moving in resonant orbits particularly near ILR and CR. Tremaine \\& Weinberg (1984) and Weinberg (1985), extended this analysis showing that angular momentum transference can be interpreted as an effect caused by a particular type of dynamical friction.  N-body simulations indicate that angular momentum transference outwards is always  present in rotating galactic models. This effect takes place not only between the bar and the disc, but it is extended to the halo particles (e.g. Debattista \\& Sellwood 1998, 2000; Athanassoula 2002, 2003).  In particular, it has been found that the material in the bar and in the inner disc emits angular momentum to resonant material in the outer disc, or in the halo and in the bulge. The evolution of bars in isolated disc galaxies is driven by the redistribution of angular momentum. Transference of angular momentum outwards in rotating galaxies produces secular dynamical evolution in which the bar grows and slows down.  Several questions arise regarding the role of regular motion and the role of chaotic motion in rotating galaxies and their evolution, examples are as follows.  \\noindent i) What is the effect of rotation on the degree of chaos in such systems, in comparison to the degree of chaos in non-rotating systems?  \\noindent ii) What is the typical level of the Lyapunov Characteristic Numbers  and what is the amount of mass that can develop efficient chaotic diffusion in a Hubble time?  \\noindent iii) What is the role of chaos on the secular evolution and how secular evolution affects chaos?  \\noindent iv) How the two components of mass in chaotic or in regular motion are distributed in space?  \\noindent v)How the pattern is affected by chaos? Can such systems develop spiral arms? For how long?   In order to investigate these questions, we use one of the non-rotating N-body models mentioned above, namely the Q-model, and we construct from it a series of iso-energetic rotating models. Iso-energetic means that they have the same total mass $M$, the same total binding energy $E$ and they are close to the same scalar virial equilibrium, but they carry different total amounts of angular momentum. In this way, a direct comparison can be made as regards the effect of their angular momentum on the behavior of the models. The total amount of angular momentum in these models is large, so that their values of the spin parameter $\\lambda$ are comparable to the value of $\\lambda$ of our Galaxy.  In this investigation we use the same code and the same number of particles as in Voglis et al. (2002), i.e. an improved version of the conservative technique code by Allen Palmer \\& Papaloizou (1990) with $N \\thickapprox 1.3 \\times 10^5$ particles.  We use as a scaling unit of length the half mass radius $R_h$ in every model. The half mass crossing time $T_{hmct}=(2R_h^3/GM)^{1/2}$ is defined as the unit of time. An important time scale is the radial period $T_{hmr}$ of an orbit with binding energy $E_{hmr}$ equal to the mean value of the potential at the half mass radius. If we assume that in a typical galaxy a star with binding energy $E_{hmr}$ describes about 50 circular periods $T_{cp}=2 \\pi T_{hmct}$ in a Hubble time, then it describes about 100 radial periods in this time (since the 2:1 resonance at this energy is typical). Therefore, a Hubble time is $t_{Hub} \\approx 50 T_{cp} \\approx 100 T_{hmr} \\approx 300 T_{hmct}$.    In Section 2 the method of derivation of the initial conditions of our models is described. In Section 3 various features of the models along their self-consistent run are examined, regarding their morphology, the radial profiles of the surface density, the rotation curves, the evolution of the pattern speed and the corotation radius, the longevity of the spiral structure and the secular evolution of the systems. In Section 4 the components of mass in regular and in chaotic motion for every model are estimated by the method introduced in Voglis et al. (2002). In Section 5 the spatial distribution of these two components is demonstrated. Our conclusions are summarized and a discussion is given in Section 6. ",
        "conclusions": "We have presented the results of several N-body simulations of iso-energetic models of galaxies, having different amounts of angular momenta, but the same scalar virial equilibrium. We have examined a number of morphological and dynamical properties of these models and we have found the sets of particles that move in regular or in chaotic orbits in every model. We have compared the spatial distribution of these two sets. Our main conclusions can be summarized as follows:  1) Density waves form a rotating central bar embedded in a rotating thin and a thick disc. The size of the bar decreases and the density in the region of the bar increases with increasing angular momentum of the models.  2) Beyond the ends of the bar an exponential radial profile of the surface density is formed, resembling the observed corresponding profiles of disc galaxies.  3) Spiral arms are spontaneously formed on the disc, due to an $m=2$ spiral wave fluctuation mode on the surface density, excited by the bar. Spiral arms show an intermittent behavior. They are deformed, fade or disappear temporarily, but they grow again, in a time scale less than the rotation period of the bar. The $m=2$ mode preserves its spiral structure outside corotation for a Hubble time, during which the bar describes about 20 to 30 rotations (depending on the angular momentum of the model). The relative amplitude of this mode with respect to higher modes and noise decreases in time. However, in the region of spiral arms the amplitude of $A_2(r,t)$ at the end of a Hubble time is on the same level as it is given from observational data.  4) The dynamical rotational velocity curve shows a relatively small variability in a Hubble time indicating a slow redistribution of mass. However, the rotational velocity curve, evaluated from the mean angular momentum of particles inside successive cylindrical rings around the rotation axis, shows a remarkable evolution in time. Namely, this rotation curve decreases in the inner parts, particularly inside one half mass radius, due to the transference of angular momentum outwards.  5) In faster rotating models (as in QR2, QR3, QR4) the central bar loses angular momentum and slows down. The pattern speed $\\Omega_p$ of the bar decreases in time and the respective corotation radius increases by about $50\\%$ at the end of a Hubble time. This result is in good agreement with observational data comparing the linear sizes of young and old bars.  6) Rotation enhances chaos substantially, both by increasing the fraction of mass in chaotic motion and by shifting the Lyapunov numbers to larger values (by about one order of magnitude with respect to the non-rotating model). Two effects collaborate for this enhancement of chaos, both due to rotation. The deformation of the distribution of mass relative to the non-rotating model and the resonant effects, particularly near corotation. The mass that can develop effective chaotic diffusion in a Hubble time increases significantly. In our examples, this mass rises (from the level of 8\\% of the total mass in the non-rotating Q model) to the level of 45\\% - 55\\% in the rotating models.  7) The energies of chaotic orbits are concentrated around the value of the energy near the corotation radius (based on the pattern of the central bar), while the energies of regular orbits are mainly in the deepest part of the potential well.  8) A most remarkable result in this work is that spiral arms are almost completely formed by mass in chaotic motion. This is clearly demonstrated by examining the spatial distribution of the mass in regular and in chaotic motion.   It is significant that in all our models the $m=2$ mode preserves a spiral pattern up to the end of a Hubble time. This mode is gradually obscured by excitation of higher modes and noise. At least part of the power of the higher modes and the noise is expected to be a real effect, but numerical effects are expected to be present as well. We should keep in mind that integrating orbits in a time varying potential is a demanding problem, mainly as regards the accuracy of reproducing the self-consistent potential and its evolution.  In the level of the available accuracy in the dissipationless simulations presented in this paper, our investigation shows only what gravity, combined with rotation, can do. The relative strength of the $m=2$ wave could be better preserved  by the end of a Hubble time under the presence of dissipation. Dissipation could absorb noisy components and reinforce the robustness of the $m=2$ mode. Furthermore, gaseous material could be sunk and trapped well inside the corotation radius, as e.g. in Bournaud \\& Combes (2002); Kormendy \\& Kennicutt (2004). This process enriches the bar with angular momentum, competing the angular momentum transference outwards.  Other parameters, not involved in this study, could play a role as well. E.g. the presence of a massive central black hole could moderate the strength of the bar (or even destroy the bar, see e.g. Shen \\& Sellwood 2004). Furthermore, the rate of redistribution of angular momentum can be modified by the presence of a halo with low values of $\\lambda$. However, chaos can still play an important role in the dynamics and morphology of these systems.   \\bigskip"
    },
    "0606/astro-ph0606757_arXiv.txt": {
        "abstract": "{} {This paper presents a method to identify gravitational arcs or more generally elongated structures in a given image.} {The method is based on the computation of a local estimator of the elongation. The estimation of the local elongation proceed in two steps: first the local orientation of the structure is computed, then in the next step, a rotation is performed, and the marginal distributions are used to compute the elongation. This procedure allows the computation of the local elongation at each point in the image. Then, using a threshold on the elongation map the elongated structures are identified and re-constructed using connectivity criteria.} {Finally a catalog of elongated structures is produced, and the properties of each object are computed, allowing the selection of potential arc candidates. The final selection of the arc candidates is performed by visual inspection of multi-color images of a small number of objects.} {This method is a general tool that may be applied not only to gravitational arcs, but to all problems related to the mapping and measurement of elongated structures, in an image, or a volume.} ",
        "introduction": "The new large scale surveys offers deep and well sampled views of large regions of the sky. These views contains such a large number of objects that automated method to find a specific type of objects is of great interest. In particular, the new release of the CFHTLS survey (see {\\it http://www.cfht.hawaii.edu/Science/CFHLS/} for a description of the project) presents, a large set of 1 square degree images, with a corresponding size of about $20000 \\times 20000$ pixels. Among a large number of interesting objects, this survey contains a number of gravitational arcs, typically one per image. Considering the huge number of pixels, finding theses arcs or arclets is a good task for an automated software. The identification of potential arcs in the image is mostly a search for elongated objects. Thus an automated method designed to find arcs should first produce a catalog of elongated objects and select possible arcs candidates among them. In the end the procedure should be concluded by a visual inspection of a small number of candidates. This method is a new approach, former methods relied on filtering, for instance, Lerzer, Scherzer \\& Schindler, (2004 \\& 2005), used anisotropic filtering, while Starck, Donoho \\& Candes (2003) used wavelet Filtering. Arc detections using catalogs from the Sextractor software (Bertin \\& Arnouts 1996), were also performed by Horesh {\\it et al.} 2005. ",
        "conclusions": ""
    },
    "0606/astro-ph0606269_arXiv.txt": {
        "abstract": "{} {We challenge the accepted distance of 3.2~kpc of GRO~J1655-40. We present VLT-UVES spectroscopic observations to estimate the absorption toward the source, and determine a maximum distance of GRO~J1655-40.} {We show that the accepted value of 3.2~kpc is taken for granted by many authors. We retrieved in the ESO archive UVES spectra taken in April 2004 when GRO~J1655-40 was in quiescence to determine the spectral type of the secondary star. For the first time we build a flux-calibrated mean (UVES) spectrum of GRO~J1655-40 and compare its observed flux to that of five nearby stars of similar spectral types. We strengthen our results with the traditional pair method, using published photometric data of the comparison stars.} {We show that the distance of 3.2~kpc is questionable. We determine a spectral type F6IV for the secondary star. We demonstrate in details that the distance of GRO~J1655-40 must be smaller than 1.7~kpc.} {The runaway black hole GRO~J1655-40 could be associated with the open cluster NGC~6242 which is located at 1.0$\\pm$0.1~kpc from the Sun. At $D \\leq$ 1.7~kpc the jets are not a superluminal, and GRO~J1655-40 becomes one of the closest known black holes to the Sun.} ",
        "introduction": "\\object{GRO~J1655-40} (a.k.a Nova Sco 94) is a Soft X-ray transient (SXT) in our Galaxy (l=345.0$^\\circ$, b=2.2$^\\circ$; R.A.=16$^h$54$^m$00$^s$, Dec.=$-39^\\circ$50$^m$45$^s$). It has been discovered on July 27, 1994 with \\emph{BATSE} on board the Compton Gamma Ray Observatory \\citep{Zhang-etal-1994}. Optical photometry of GRO~J1655-40 revealed a double-waved modulation with a period of 2.6 days \\citep{Bailyn-etal-1995b,vanderHooft-etal-1997}. Strong evidence that the compact object in GRO~J1655-40 is a black hole was presented by \\citet{Bailyn-etal-1995b} who established a spectroscopic period of 2.601 $\\pm$ 0.027 days. \\citet{Shahbaz-etal-1999} published the masses of the black hole and the secondary star: 5.5--7.9 $M_{\\odot}$ and 1.7--3.3 $M_{\\odot}$ respectively. Given a distance of about 3~kpc, the jets appeared to be superluminal: 1.5$\\pm$0.4~$c$ \\citep{Tingay-etal-1995}, 1.05~$c$ \\citep{Hjellming-Rupen-1995}, where $c$ is the speed of light. Various spectral types of the secondary star have been published. We mention in particular: F3-F6~IV \\citep{Orosz-Bailyn-1997}, and F7IV-F6III \\citep{Israelian-etal-1999}. The distance of GRO~J1655-40 that is used in the literature is 3.2$\\pm$0.2~kpc, although \\citet{Mirabel-etal-2002} pointed out that a distance of about 1~kpc cannot be ruled out with the current data. The aim of this paper is to show that none of the distance estimate methods used for GRO~J1655-40 is good enough to claim a distance of 3.2~kpc, and that this distance is actually based on one very uncertain interpretation. We present unpublished optical data of GRO~J1655-40 to show that the distance of the source is certainly smaller than 1.7~kpc. ",
        "conclusions": "We have shown that the distance of GRO~J1655-40 of 3.2~kpc quoted in the literature has been obtained through a refinement of a distance range, which was in turn not well established. We have determined a spectral type F6IV for the secondary star using UVES spectra. By flux-calibrating our spectra of GRO~J1655-40 during quiescence and comparing them with various stars of similar spectral types, we have shown that the distance of GRO~J1655-40 simply cannot be of 3.2~kpc, and is certainly smaller than 1.7~kpc. We consider the possibility that the source is associated with the open cluster NGC~6242 \\citep[see Fig.~1 of][]{Mirabel-etal-2002}, which lies at a distance of 1.0~kpc. With such distance, GRO~J1655-40 is not a superluminal source (the jets speed reaches $\\sim$ 0.37$c$). If the distance is confirmed (maybe by the future european satellite {\\it Gaia}), and assuming that the distance of other black-hole binaries are correct, GRO~J1655-40 would become one of the closest known black hole to the Sun."
    },
    "0606/astro-ph0606575_arXiv.txt": {
        "abstract": "We examine the effects of significant electron anti-neutrino fluxes on hydrogen burning.  Specifically, we find that the bottleneck weak nuclear reactions in the traditional pp-chain and the hot CNO cycle can be accelerated by anti-neutrino capture, increasing the energy generation rate.  We also discuss how anti-neutrino capture reactions can alter the conditions for break out into the $rp$-process.  We speculate on the impact of these considerations for the evolution and dynamics of collapsing very- and super- massive compact objects. ",
        "introduction": "Hydrogen burning involves the conversion of four protons into an alpha particle, two positrons, neutrinos and photons.  The principal bottleneck involved in this process is the weak interaction conversion of protons into neutrons.  For decades the primary mechanisms of hydrogen burning have been an astronomical staple.  \\citet{bet38} first elucidated the proton-proton chain (pp-chain) where the weak conversion is accomplished by two protons interacting to become a deuteron, $p ( p, \\nu_e e^+ ) d$.  \\citet{vwe37, vwe38} and \\citet{bet39} independently described the CNO cycle, where carbon is used as a catalyst in hydrogen burning, and the weak conversion of protons to neutrons occurs through the positron decay of isotopes of oxygen with half lives of about 100 seconds.  A large flux of electron anti-neutrinos ($\\bar\\nu_e$) could alter the hydrogen burning paradigm.  Anti-neutrino capture could perform the necessary conversion of protons to neutrons.  The $\\bar\\nu_e$-capture cross sections of relevance are very small, but depend strongly on neutrino energy.  The smallness of these cross sections allows energetic neutrinos to escape from deep within a compact object, where the temperature and other energy scales are high, and freely stream to where hydrogen burning is occurring.  Nevertheless, if $\\bar\\nu_e$-capture is to have a significant effect on hot hydrogen burning, a truly prodigious flux ($\\phi_{\\bar\\nu_e} \\gtrsim 10^{40} ~\\mbox{cm}^{-2} ~\\mbox{s}^{-1}$) and large neutrino energy ($\\langle E_{\\bar\\nu_e} \\rangle \\gtrsim \\mbox{a few MeV}$) would be necessary.  It should be kept in mind, however, that to affect hydrogen burning, the $\\bar\\nu_e$-capture rates need only be comparable to the corresponding positron decay rates.  The difficulty would be to find an environment capable of producing these fluxes of neutrinos, yet quiescent enough that simple hydrogen burning could be relevant, and the products of such burning could be ejected into space.  High entropy electron-positron plasmas are efficient engines for the production of neutrinos and anti-neutrinos of all flavors.  Possible environments that may merit future investigations into the effects of anti-neutrino capture on hydrogen burning include high mass accretion disks and collapsing very- and super- massive objects.  In this paper, we investigate the effects of a prodigious neutrino flux on hot hydrogen burning.  In section \\ref{burnmech} we point out the effects of anti-neutrino capture on the rate limiting steps in both the pp-chain and the $\\beta$-limited CNO cycle, and its implications for the energy generation rates.  In section \\ref{consequences}, we examine the consequences for the $rp$-process and energy generation mechanisms.  In section \\ref{smo}, we consider the case of a supermassive star collapsing on the general relativistic Feynman-Chandrasekhar instability, and the effects of its internal neutrino production on hydrogen burning in its envelope.  We give conclusions in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this paper we have examined the effects of a prodigious flux of electron anti-neutrinos on hydrogen burning.  We have found that the traditional positron decay bottlenecks in hydrogen burning can be removed and replaced by much faster $\\bar\\nu_e$-capture reactions under some conditions.  This would result in an increase of several orders of magnitude in the energy generation rate over what would be expected without such a flux.  Additionally, the $\\bar\\nu_e$-flux would alter the conditions necessary for break-out into the $rp$-process, increasing the temperature necessary to do so at a given density.  If conditions allow the break-out into the $rp$-process, we could expect an acceleration of the flow toward the iron-peak facilitated by and accelerated by $\\bar\\nu_e$-capture.  When applied to the neutrino flux emitted in the final stages of the collapse of a supermassive star, interesting changes from current simulations may occur on the lower end of the supermassive star mass spectrum.  Whether or not these effects are relevant, remains an open question that can only be answered by simulations that are able to include hydrogen burning during the final collapse of the star.  Important issues that remain open include finding an astrophysical environment where the effects discussed here could take place.  Accretion disks surrounding black holes may provide a combination of high accretion rates and hot, high entropy disks which could produce the necessary fluxes of electron anti-neutrinos.  (See for example \\citet{sur05} for a discussion of neutrino emission in lower mass accretion disks.) Supermassive stars may exhibit this effect, though there is uncertainty related to whether or not these objects ever existed.  Computer simulations would be useful to determine any changes in the expected nucleosynthetic yield, and the effects of their possible distribution into the surrounding IGM."
    },
    "0606/astro-ph0606096_arXiv.txt": {
        "abstract": "We have measured the infrared transit of the extrasolar planet HD\\,209458\\,b using the Spitzer Space Telescope.  We observed two primary eclipse events (one partial and one complete transit) using the 24~\\micron\\ array of the Multiband Imaging Photometer for Spitzer (MIPS).  We analyzed a total of 2392 individual images (10-second integrations) of the planetary system, recorded before, during, and after transit.  We perform optimal photometry on the images and use the local zodiacal light as a short-term flux reference.  At this long wavelength, the transit curve has a simple box-like shape, allowing robust solutions for the stellar and planetary radii independent of stellar limb darkening, which is negligible at 24~\\micron.  We derive a stellar radius of R$_*$ = 1.06 $\\pm$ 0.07~R$_\\sun$, a planetary radius of R$_p$ = 1.26 $\\pm$ 0.08~R$_{\\mbox{\\scriptsize J}}$, and a stellar mass of 1.17~M$_\\sun$.  Within the errors, our results agree with the measurements at visible wavelengths.  The 24~\\micron\\ radius of the planet therefore does not differ significantly compared to the visible result.  We point out the potential for deriving extrasolar transiting planet radii to high accuracy using transit photometry at slightly shorter IR wavelengths where greater photometric precision is possible. ",
        "introduction": "The transit of an extrasolar planet across its star allows us to measure the radius of the planet \\citep{charbonneau00, henry00, brown01a}.  Of the ten known transiting planets, HD\\,209458\\,b, with a radius (at visible wavelengths) of 1.320 $\\pm$ 0.025~R$_{\\mbox{\\scriptsize J}}$ \\citep{knutson06}, is inflated compared to the other known transiting planets and thus has a lower bulk density.  One explanation for the anomalous radius was inflation by the dissipation of tidal stress within the planet \\citep{bodenheimer01}.  However, the timing of the secondary eclipse as observed by the Spitzer Space Telescope \\citep{deming05b}, as well as improved radial velocity observations \\citep{laughlin05}, have ruled out a non-zero orbital eccentricity of the magnitude ($\\sim$0.03) needed by the tidal dissipation theory. \\citet{showman02} suggest that kinetic energy produced by atmospheric circulation and deposited in the planet's interior could account for the missing energy source and increase the planetary radius. Another popular proposed explanation is the possibility of obliquity tides \\citep{winn_holman05}, in which a non-zero obliquity (made possible by a spin-orbit resonance) could drive the tidal dissipation and provide the necessary energy to give the planet an inflated radius.  Our Spitzer program to measure the 24~\\micron\\ flux of HD\\,209458\\,b includes observations during transit (i.e., primary eclipse), revealing the infrared (IR) radius of the planet, which is reported in this paper. Section~\\ref{sec:motiv} further elaborates on why an IR radius measurement is of interest.  Section~\\ref{sec:obs} describes the observations; section~\\ref{sec:analysis} explains the photometric data analysis and radius fit.  Section~\\ref{sec:disc} concludes with results and discussion. ",
        "conclusions": "\\label{sec:disc} We have computed the stellar density directly from the observable quantities from the best-fit simple curve \\citep{seager_mallen-ornelas03}.  Assuming a stellar mass allows us to calculate the stellar radius.  This empirical mass-radius relation is shown in Figure~\\ref{fig:mstar}, where we have derived the radii for stellar masses from 0.9 to 1.3~M$_\\sun$. We break the degeneracy by intersecting the stellar radius curve with the mass-radius relation from \\citet{cody02}, which is derived by fitting stellar models to a constant luminosity.  This is shown as the dashed line in Figure~\\ref{fig:mstar}, following \\citet[][their Figure 5]{wittenmyer05}.  On this basis, we derive the stellar mass to be M = 1.171~M$_\\sun$, with R$_*$ = 1.06 $\\pm$ 0.07~R$_\\sun$ and R$_p$ = 1.26 $\\pm$ 0.08~R$_{\\mbox{\\scriptsize J}}$.  Our result for the planetary radius agrees with the updated visible radius of R$_p$ = 1.320 $\\pm$ 0.025~R$_{\\mbox{\\scriptsize J}}$ \\citep{knutson06}.  We are encouraged by the fact that our radius error is only four times larger than that obtained by \\citet{knutson06}, in spite of the fact that our infrared photometric precision is an order of magnitude poorer than the HST visible photometry. We attribute this to the character of the IR transit curve, where the lack of limb darkening produces a simple transit shape, from which radius information can be extracted with maximum efficiency.  We point out that photometry at other accessible Spitzer wavelengths such as 8 and 16~\\micron\\ would provide much higher photometric precision for bright transiting systems, because the stellar flux will be much higher, and the zodiacal background will not be a limiting factor.   Limb darkening remains sufficiently weak at these shorter IR wavelengths to maintain a relatively simple transit light curve shape.  Considering also that Spitzer's heliocentric orbit allows uninterrupted observations of complete transits, we suggest that IR transit photometry from Spitzer may be the optimal method for precise radius determination in bright transiting planet systems."
    },
    "0606/astro-ph0606605_arXiv.txt": {
        "abstract": "We present new results of \\FeII\\ $\\lambda$1.644~$\\micron$ spectroscopy toward the jets from HL~Tau and RW~Aur carried out with the Subaru Telescope combined with the adaptive optics system. We observed the regions within 2$^{''}$--3$^{''}$ from the stars with the sub-arcsecond resolutions of 0$\\rlap.^{''}$5 and 0$\\rlap.^{''}$2 for HL~Tau and RW~Aur, respectively. In addition to the strong, high velocity emission extended along each jet, we detected a blueshifted low velocity emission feature seen as a wing or shoulder of the high velocity emission at each stellar position. Detailed analysis shows that the low velocity emission of HL~Tau is peaked at 0$\\dotsec$077 $\\pm$ 0$\\dotsec$04 away from the stellar position, or 15~AU of actual distance from the star. We did not confirm such displacement for RW~Aur in its position velocity diagram (PVD). The PVDs of HL~Tau and RW~Aur show a characteristic similar to those of the cold disk wind and X-wind models in that the [\\ion{Fe}{2}] line width is broad in the vicinity of the stellar position and is narrower at the extended jet. A closer comparison suggests, however, that the disk wind model tends to have too large line width at the jet while the X-wind model has excess emission on the redshifted side at the stellar position. The narrow velocity width with symmetric line profiles of the observed high velocity emission supports an X-wind type model where the launching region is localized in a small radial range, while the low velocity emission located away from the star favors the presence of a disk wind. The \\FeII\\ emission from the HL~Tau jet shows a gap of 0$\\dotsec$8 between the redshifted jet and the star, indicating the presence of an optically thick disk of $\\sim$160 AU in radius. The \\FeII\\ emission from the RW~Aur jet shows a marked drop from the redshifted peak at $Y \\sim-$0$\\dotsec$2 toward the star, suggesting that its disk radius is smaller than 40~AU. The normalized Br12 emission of HL Tau shows a significantly large deviation from the normalized continuum in spatial profile along the blueshifted jet. It means that part of the Br12 emission originates in the blueshifted jet itself. ",
        "introduction": "Spectroscopy with high angular and spectral resolutions is essential for understanding the driving mechanism of the outflows from young stellar objects. In previous contributions, we demonstrated that such observations with the near-infrared \\FeII\\ $\\lambda$1.644~$\\micron$ line is a powerful tool to study the acceleration and collimation regions \\citep{pyo02, pyo03, pyo05}. In the position-velocity diagrams (PVDs) of L1551 IRS~5 and DG~Tau toward their blueshifted jets, the blueshifted \\FeII\\ emission showed two clearly distinct components in velocity and space: the high velocity component (HVC: $V=$~200--400~$\\kms$) is extended and has a narrow line width, while the low velocity component (LVC: $V=$~50--200~$\\kms$) is observed closer to the driving sources and shows a larger line width. Because of the distinct natures of these two velocity components, we suggested that they may be two physically different outflows: the HVC is a highly collimated jet, while the LVC is a widely opened disk wind.  HL Tau is a young star in a transitional phase from an embedded protostar to a exposed T~Tauri star, as has been suggested from its large amount of circumstellar material, strong outflows and optical nebulosity. Its ``flat'' spectral energy distribution from the near-infrared to far-infrared together with its strong optical veiling suggests that HL~Tau is still in the active accretion phase. This is also supported by the presence of an infalling envelope of $\\sim$1000~AU as a remnant of its placental molecular cloud core \\citep{HOM93}, although the infalling interpretation may not be unique because the flattened structure around HL~Tau overlaps with the wall of a bubble evacuated by XZ~Tau in a large scale CO map \\citep{Welch00}. Millimeter and sub-millimeter interferometry has revealed a circumstellar disk ($\\sim$100~AU) inside the envelope \\citep{lay97,mundy96,lay94}. Its jet was discovered by \\citet{MRB88}. The inclination angle of the disk with respect to the plane of the sky is $\\sim$42$^{\\circ}\\pm$5$^{\\circ}$ \\citep{lay97}. {\\it Hubble Space Telescope} observations revealed that the optical continuum emission from HL~Tau is not from the star itself but from a reflection nebula surrounding the blueshifted jet \\citep{Sta95}. \\citet{close97} showed in their adaptive optics observations that the near-infrared point source coincides with the star itself.  RW~Aur is a multiple star system located in the Taurus-Auriga molecular cloud. The brightest stellar component is called ``A,'' which is the source of the outflow called HH 229. RW~Aur~A itself may be a spectroscopic binary \\citep{gahm1999,petrov01}. The fainter component ``B'' is located 1$\\dotsec$5 away from A at a position angle of 258$\\degree$. Although RW~Aur~B was also known as a close binary system, its companion ``C'' was shown to be a false source \\citep{WG01}. RW~Aur~A and B are both classical T~Tauri stars (CTTS) with active accretion suggested by their strong H$\\alpha$ emission in the optical spectroscopy \\citep{Duchene99} and large veiling in the near-infrared wavelengths \\citep{FE99}. The inclination of the jet axis is $\\sim$46$^{\\circ}\\pm$3$^\\circ$ degrees with respect to the line of sight \\citep{LCD03}. Transverse velocity shifts across the jet axis were detected by \\citet{coffey04}, who interpreted it as jet rotation around its axis, although this rotation appears to be in the opposite sense to the rotation of the disk \\citep{cabrit05}.  In this article, we present new results of the \\FeII\\ $\\lambda$1.644~$\\micron$ line spectroscopy toward the outflows from HL~Tau and RW~Aur~A. The data were obtained with sub-arcsecond resolutions using the adaptive optics system, allowing us to study the spatial and radial velocity structure of the outflows in the vicinity of their driving stars. We assume the distance of 140~pc to HL~Tau and RW~Aur in this paper. ",
        "conclusions": "\\subsection{The Emission Source Locations around HL Tau}  The LVC of HL~Tau seems to be strongest at the stellar position with a weaker emission significantly extended to the northeast ($Y >$~0, see Figure~\\ref{fig2}). Because HL~Tau has a scattering envelope extended to the northeast \\citep{close97}, we must first examine whether the extended LVC emission comes from scattering.  Figure~\\ref{fig6} shows the normalized spatial profiles of intensity for the LVC, HVC, Br12 and continuum emissions. The intensities for the LVC and HVC are averaged over the velocity ranges of $-$100 to 0~$\\kms$ and $-$270 to $-$100~$\\kms$, respectively. The Br12 intensity is averaged over $-$700 to $-$280~$\\kms$ with respect to the rest wavelength of \\FeII\\ \\footnote{$-$190 to $+$231 $\\kms$ with respect to the rest velocity of Br12 ($\\lambda$1.6412 $\\mu$m)}, and the continuum emission is averaged over $+$400 to $+$800~$\\kms$ and $-$1200 to $-$1000~$\\kms$ where the spectrum is free from photospheric absorption lines.  The continuum, Br12 and LVC emissions show similarly skewed spatial profiles with sharp cutoffs on the redshifted side ($Y <$~0$''$) and gradual decreases on the blueshifted side ($Y >$~0$''$). The normalized spatial profile of Br12 exactly matches that of the continuum on the redshifted side ($Y <$~0$''$), reaching its half maximum at $Y = -$0$\\dotsec$28 consistent with the spatial resolution. This means that the dominant source of the Br12 emission coincides with that of the continuum, namely the star itself or its unresolved vicinity. We should note, however, that the peak position of the Br12 spatial profile is displaced by 0$\\dotsec$038 $\\pm$ 0$\\dotsec$02 with respect to that of the continuum in Figure~5. The small but finite amount of the observed spatial displacement between the Br12 and continuum peaks may mean that at least part of the Br12 emission at $Y\\sim$ 0$''$ arises in the unresolved outflow. This conclusion is consistent with the recent observations suggesting that infrared \\ion{H}{1} lines of classical T~Tauri stars or Class~I objects originate in the outflow acceleration region \\citep{FE01, NAG04, WRD04}, although existing outflow models as well as magnetospheric accretion models still have difficulties to quantitatively explain their line profile characteristics.  On the blueshifted side ($Y >$~0$''$), the gradual decrease of the continuum emission is caused by the scattering of the stellar light off the cavity located at the northeast of HL~Tau \\citep{close97}. Thus the spatial profiles of Br12 and LVC similar to that of the continuum suggest that their emissions in the northeast of HL~Tau arise mostly from scattering. However, the normalized spatial profile of Br12 is significantly larger than that of the continuum on the blueshifted side ($Y >$~0$''$). The Br12 emission along the blueshifted jet, therefore, does not entirely arise from scattering of the central unresolved light source, but part of it must originate in the extended blueshifted jet itself.  The normalized spatial profile of LVC is more interesting. Although it basically follows a skewed shape similar to those of Br12 and the continuum emission, the spatial profile of LVC in the southwest ($Y <$~0$''$) is offset by $\\Delta Y = +$0$\\dotsec$077 $\\pm$ 0$\\dotsec$04 with respect to that of the continuum. This means that the dominant emission of LVC arises not from the unresolved vicinity of the star but in the jet $\\sim$~15 AU away from the star. In the northeast ($Y >$~0$''$), the normalized intensity of LVC is larger than that of the continuum, although the excess is not much larger than the noise level of the LVC profile. Part of the extended LVC emission in the northeast thus also arises in the jet.  Figure~\\ref{fig4} shows that the LVC of RW~Aur also has its spatial maximum displaced to the blueshifted (southeast) side. However, this may be caused by the presence of the artificial dip around $Y=~-$0$\\dotsec$1 and V$_{\\rm LSR}=$~90~km~s$^{-1}$ and is not reliable.  \\subsection{The High and Low Velocity Components} \\label{diss_2v}  Our observations revealed that the \\FeII\\ emission lines show  both LVC and HVC in HL~Tau and RW~Aur. The presence and nature of the two velocity components are similar to the cases of L1551 IRS~5 and DG~Tau \\citep{pyo02, pyo03}. The LVCs of HL~Tau and RW~Aur are, however, not as distinct as those of L1551 IRS~5 and DG~Tau seen in their PVDs. The latter sources show that the LVC has a clearly separated peak located away from the stellar position in the PVD, while the former sources have the LVC observed only as a wing or shoulder of the HVC and its spatial location is close to, although not necessarily coincident with, the stellar position. For the jets from L1551 IRS~5 and DG~Tau, we suggested that the HVC is physically different from the LVC because of the largely different characteristics of the two components. In fact, it has not been easy to model the two distinct components with single outflow models \\citep{pesenti03, shang98}. In the cases of HL~Tau and RW~Aur, however, the LVC seems to be more physically related to the HVC from their appearance in the PVDs, and we will re-examine the relation between the two velocity components for these objects below\\footnote{Please refer to \\citet{pyo02,pyo03} for more details about two velocity components.}.  Up to date the magnetocentrifugal force is widely considered as the most comprehensive mechanism for launching outflows from the star-disk systems of young stellar objects \\citep[see reviews in][]{KP00,Shu00}. There are two kinds of popular models based on the magnetocentrifugal mechanism: the {\\it disk wind model} \\citep{KP00} and {\\it X-wind model} \\citep{Shu00}. The disk wind model assumes that the outflow launching region is located over a wider range of disk radii outside the inner disk edge, while the X-wind model postulates that the launching region is located at a very narrow disk radii close to the inner disk edge where the highly concentrated stellar magnetic field interacts with the disk. Synthetic PVDs of the forbidden emission lines predicted by their two models may be compared with observations \\citep[e.g.][]{mtakami06}.  In Figure~\\ref{fig7} we compare the PVDs of HL~Tau and RW~Aur with those of {\\it the cold disk wind models} \\citep{cabrit99, garcia01, pesenti03} and {\\it X-wind model} \\citep{shang98}. The two models show a similitude that their velocity widths are broad at the stellar position and become narrower when the position moves away from the star. The broad velocity width at the star is caused by the presence of lower velocity emission added to the radial velocity V $\\la$ 100 $\\kms$, while its spatial extent (20 -- 100 AU) depends on the adopted spatial resolutions of the models. This emission arises in the region where uncollimated outer stream lines have sufficient emissivity in these models.  The above noted characteristic of the model PVDs are also apparent in the observed PVDs of HL~Tau and RW~Aur. By analogy to the observed PVDs, we call the lower velocity emission at V $\\la$ 100 $\\kms$ the LVC and the spatially extended higher velocity emission at V $\\ga$ 100 $\\kms$ the HVC. The observed PVDs are, however, different from the model PVDs on the following points. \\begin{enumerate} \\item{The HVC of the {\\it disk wind models} shown in Figure~\\ref{fig7}$a$ and \\ref{fig7}$b$ shows a velocity width of 200 -- 500 $\\kms$, much larger than the observed width of $\\la$ 60 $\\kms$. On the other hand, the width of HVC for the {\\it X-wind model} shown in Figure~\\ref{fig7}$c$ is narrow, which may be consistent with the observed width of HVC and its symmetrical line profiles if internal velocity fluctuation in the HVC is taken into account. Recently \\citet{pesenti04} showed that a ``warm'' disk model can make narrower line widths, showing much better agreement with those of observed jets.} \\item{The LVC of the {\\it X-wind model} shows a significant redshifted emission arising from the streamlines on the far side. The presence of such redshifted emission coincident with the star is marginal in the observed PVDs because of the contamination by the photospheric \\ion{Fe}{1} absorption at $+$20 $\\kms$ (see \\S3.1). On the other hand, the LVC for the {\\it disk wind models} has little redshifted emission at the stellar position, which is consistent with the observed characteristics of LVC unless the marginal emission is real.} \\item{The HVC of the {\\it cold disk wind models} exhibits significantly higher peak velocity by a factor of 2--2.5 than the observed HVC does, while the HVC peak velocity of the {\\it X-wind model} is quantitatively consistent with the present observations. The {\\it disk wind models} can be accommodated to the observations if a smaller Alfv\\'en radius by a factor of 2-2.5 is adopted. \\citet{garcia01} also indicated that a denser and slower ``warm'' disk wind are favored to solve this problem.} \\end{enumerate}  With respect to the first point, material is launched from a large range of disk radii (0.07 -- 2 AU) in the disk wind models, where lower velocity material is ejected from the outer radii in the launch zone, while higher velocity material is ejected from the inner radii (the terminal velocity scales as the Keplerian speed at the jet base). The broad velocity width of HVC is hence caused by the outer streamlines having significantly lower radial velocities than the inner streamlines even when they are well collimated. On the other hand, all the streamlines of the X-wind have similar Alfv\\'en radii because they originate from the same launching point at the inner disk edge ($\\sim$ 0.07 AU), resulting in similar radial velocities when they are collimated. The narrow velocity width of the observed HVC thus means for the disk wind models that the outer streamlines with lower velocities are associated with little emission, with a rapid cutoff of emission from the inner to outer streamlines, or that all the streamlines originate from a narrow radial range, which is rather similar to the case of an X-wind. The presence of weak but spatially extended emission in the observed LVC, however, suggests that the extended LVC may correspond to the emission from the outer, low velocity streamlines that originate at larger radii than those of HVC.  Regarding the second point, the X-wind model has the streamlines at the base of the jet fanning out in a wide direction, toward and away from the observers. As a result, a very broad pedestal emission of $\\sim$ 300 $\\kms$ wide is predicted within 20 AU from the star at all inclinations, extending to $+$100 $\\kms$ on the redshifted side. The X-wind model may be accommodated to the observed PVDs if the outer streamlines with higher redshifted and blueshifted radial velocities give little emission, unless the marginal redshifted emission observed toward HL~Tau is real. Another possibility is that the redshifted emission is weak or hidden in the optical wavelength where the extinction across the outflow tends to be large near the disk surface.  In summary, none of the published model PVDs are fully consistent with the observed PVDs of \\FeII\\ emission from HL~Tau and RW~Aur.\\footnote{It is worth noting that the PVDs of \\SII\\ may differ from those of \\FeII\\ because the former line has a lower critical density than the latter one.} The narrow velocity width and symmetrical line profiles of the HVC seem to be more consistent with the X-wind model than the disk wind model, while the observed LVC emission located away from the star is in favor of the disk wind model. Given the difficulty to consistently explain the observed PVDs with a single model, we may need a hybrid type model, e.g. an X-wind type outflow emanating from the truncation radius combined with a disk wind in the outer region \\citep[see e.g.][]{pyo03}, although \\citet{FDC06} discussed recently that it may be difficult for an X-wind to co-exist with an extended disk wind.  \\subsection{The Truncation Radius of the HL Tau Disk}  For HL~Tau the peak intensity of the redshifted jet is 10 times smaller than that of the blueshifted jet when measured at the same projected distance of 1\\rlap.{$''$}4 or 200~AU from the star. Assuming the same intrinsic intensities for the blueshifted and redshifted jets, we estimate that the redshifted jet has an excess extinction of $\\Delta A_{H} = $2.5~mag, corresponding to $\\Delta A_{V} = $14~mag, with respect to the blueshifted jet. This means that a significant amount of envelope material is present around HL~Tau, as is consistent with its nature as being still embedded in a molecular cloud. In fact the optical \\SII\\ emission shows a larger gap of $\\sim$3$\\arcsec$ between the redshifted outflow and the star \\citep{solf89}. The optical gap size is similar to the size of the northeast scattering envelope observed by \\citet{Sta95} and \\citet{close97}, suggesting that the extinction is caused in the relatively spherical envelope surrounding HL~Tau and its disk.  The visual extinction of $A_{V}=$14~mag at the projected radius of 200~AU corresponds to the surface mass density of 0.05~g~cm$^{-2}$ or the average density of 2$\\times$10$^{6}$~cm$^{-3}$ if the excess extinction occurs along the line of sight length of 400~AU near the star. The surface mass density can be compared to that of the disk at the same projected radius. The disk has a surface mass density of 2 -- 4~g~cm$^{-2}$ at a radius of 100~AU \\citep{kita02}, giving an expected surface mass density of 0.4 -- 0.8~g~cm$^{-2}$ corrected for inclination at the projected radius of 200~AU when we follow the two models described by \\citet{kita02}. The expected surface density of the disk is still one order of magnitude larger than that derived from the excess extinction, suggesting that the disk should be truncated abruptly at some projected radius smaller than 200~AU from the star, most probably near the projected radius of 0$\\dotsec$8, which is the gap size between the redshifted outflow ($>$ 3$\\sigma$) and central star, corresponding to the actual radius of 160~AU.  \\subsection{Knots in the RW Aur Jet} The jet from RW~Aur has knots in optical forbidden line emissions \\citep{dougados00, woitas02, LCD03}, with the redshifted jet showing more knots than the blueshifted one. The current \\FeII\\ data also shows a hint of knots along the jets. In Figure~\\ref{fig5}$a$, the redshifted jet intensity shows four local maxima at $Y =$ $-$0$\\dotsec$25, $-$0$\\dotsec$73, $-$1$\\dotsec$20, and $-$2$\\dotsec$29, and the blueshifted jet intensity has two local maxima and one shoulder at $Y =$ 0$\\dotsec$17, 0$\\dotsec$48, and 1$\\dotsec$25, respectively. We compared these positions with the results by \\citet{LCD03}, who presented the proper motions of knots within $\\sim$12$^{''}$ from the star observed in \\SII\\ and \\OI\\ lines from 1997 to 2000. The results of the comparison are summarized as follows:  \\begin{enumerate}  \\item{ The local maxima at $Y =$ $-$0$\\dotsec$25 and 0$\\dotsec$17 may be stationary because knots were observed at similar positions in 1998 and 2000 in the optical forbidden lines.  } \\item{ The local maxima at $Y =$ $-$1$\\dotsec$20 and $-$2$\\dotsec$29 correspond to the knots R6 and R5, respectively, of \\citet{LCD03} when their proper motions of 0$\\dotsec$16~yr$^{-1}$ and 0$\\dotsec$24~yr$^{-1}$, respectively, are taken into account.} \\item{ The shoulder around $Y =$ 1$\\dotsec$25 $\\pm$ 0$\\dotsec$15 may correspond to the knot B4 if it has a proper motion of 0$\\dotsec$34$\\pm$0$\\dotsec$1~yr$^{-1}$.} \\item{ The local maxima at $Y =$ $-$0$\\dotsec$73 and 0$\\dotsec$48 may be newly ejected knots because they are close to the source and were not seen in \\citet{LCD03}.} \\end{enumerate}  In general, the local maxima observed in the \\FeII\\ intensity spatial profile show good agreement with the optical forbidden line knot positions if their proper motions are taken into consideration. Because \\FeII\\ has a larger critical density than [\\ion{S}{2}] \\citep{hamann94b,reipurth00}, \\FeII\\ traces denser regions of knots detected in [\\ion{S}{2}]."
    },
    "0606/astro-ph0606119_arXiv.txt": {
        "abstract": "In this paper, the first in a series, we present an overview of new {\\it Hubble Space Telescope} ({\\it HST}) imaging polarimetry of six nearby radio galaxies (3C 15, 3C 66B, 3C 78, 3C 264, 3C 346, and 3C 371) with optical jets. These observations triple  the number of  extragalactic jets with subarcsecond-resolution optical polarimetry. We discuss the polarization characteristics of each jet and,  as our Stokes {\\it I} images also represent by far the deepest optical images yet obtained of each of these jets,  we also discuss the morphology in total flux of each jet in detail. We find evidence of high optical polarization, averaging $20\\%$, but reaching upwards of $\\sim 50\\%$ in some objects, confirming that the optical emission is synchrotron, and that the components of the magnetic fields perpendicular to the line of sight are well ordered.  We find a wide range of polarization morphologies, with each jet having a somewhat different relationship between total intensity and polarized flux and the polarization position angle.  We find two trends in all of these jets.  First,   jet ``edges'' are very often associated with  high fractional optical polarizations, as also found in earlier radio observations of these and other radio jets.    In these regions, the magnetic field vectors appear to  track the jet direction, even at bends, where we see particularly high  fractional polarizations.  This indicates a strong link between the local magnetic field and jet dynamics.  Second, optical flux maximum regions are usually well separated from maxima in fractional polarization and often are associated with polarization minima. This trend is not found in radio data and was  found in our optical polarimetry of M87 with HST.   However, unlike in M87, we do not find a general trend for near-90$^\\circ$ rotations in the optical polarization vectors near flux maxima. We discuss possibilities for interpreting these trends, as well as implications for jet dynamics, magnetic field structure and  particle acceleration. ",
        "introduction": "Relativistic jets have been known for many years as a hallmark of powerful extragalactic radio sources. The compilation of Liu \\& Zhang (2002) listed over 660 sources with known, resolved radio jets, as of December 2000. The number of optical and X-ray jets is far smaller: about three dozen are known in each band (see the WWW pages maintained by Harris (XJET; Harris 2003) and Jester (Jester et al. 2004) respectively for lists of confirmed X-ray and optical jets). This does not signify a paucity of optical and X-ray emission, but instead, reflects the fact that few comprehensive surveys for jets have been done in those bands.  Indeed, until the advent of the {\\it Hubble} Space Telescope ({\\it HST}) and {\\it Chandra} X-ray Observatory, no survey had been done which  was reasonably optimized for the detection of jets.  Even now, there have been few {\\it HST} surveys for jets: the largest, the 3CR snapshot surveys (optical: de Koff et al. 1996, Martel et al. 1999; UV: Allen et al.  2002; near-IR: Madrid et al. 2006), includes more than 100 objects (not all of which have radio jets), but is complete only to $\\sim 21$ mag/ arcsec$^2$ due to the snapshot nature of the survey (typical integration time was $\\sim 300$ seconds with WFPC2). This is far less deep than required to see, for example, the knots in the jet of PKS 0637-752 (Chartas et al. 2000, Schwartz et al. 2000). Nevertheless, the WFPC2 survey achieved an overall success rate of 13\\%  for jet detection (Martel et al. 1999), and additional objects were discovered with NICMOS (e.g., Chiaberge et al. 2005, Floyd et al. 2006). More comprehensive and specifically targeted surveys have been  done in the X-rays, by Sambruna et al. (2002, 2004) and Marshall et al. (2005), with the former of these also having a deeper (full orbits with ACS) {\\it HST} component.  These surveys have  60-70\\% success rates of finding X-ray and optical jet emission from quasars with known radio jets.  Thus it is likely that a large proportion (perhaps even a majority) of jets generate significant levels of both optical and X-ray emission.  This explosion in the number of known optical and X-ray jets has been accompanied by a large increase in our knowledge about these objects. And yet, despite this, we are sorely lacking in information regarding the structures and composition of jets, as well as how these elements interact to produce the energetic emissions we see.  Part of the reason we lack such basic information is that it is difficult to obtain: since AGN jets are ionized flows, we lack many of the diagnostics that can be used to decipher physical conditions in other object classes. For sources that emit synchrotron radiation, polarimetry is a powerful tool that can yield direct information on the interactions between fields and particles.  The polarization vectors of synchrotron radiation indicate the magnetic field direction in the emission region, while the fractional polarization reflects the relative ordering of the magnetic field.  Since the electrons that emit synchrotron radiation in the optical and radio wavebands differ in Lorentz factor (and hence energy and cooling lengths) by factors of more than 100, comparing polarimetry in the radio and optical bands gives information on the magnetic field in regions occupied by particles in very different energy regimes.  Contrary to what one would expect in a homogeneous jet, polarimetry of the M87 jet shows different characteristics of polarization in the radio and optical bands (Perlman et al. 1999): in the optical, there is a strong anti-correlation between flux and polarization, and at the upstream ends of knots the optical magnetic field vectors become perpendicular to the jet. Much smaller changes are seen in the radio.  Thus the optical and radio emission in the M87 jet must occur in somewhat different regions.  In Perlman et al. (1999) we interpreted this as evidence that higher energy particles are concentrated closer to the center of the jet flow, with shocks that accelerate electrons and compress magnetic field lines located in the jet interior.  While {\\it HST} imaging observations now exist for virtually all known optical jets, it has been much more difficult to obtain {\\it HST} polarimetry for these objects.  Previous to this work, {\\it HST} polarimetry existed for only three optical jets:  M87 (Pre-COSTAR FOC: Boksenberg et al. 1992, Thomson et al. 1995, Capetti et al. 1997, WFPC2:  Perlman et al. 1999), 3C 273 (Pre-COSTAR FOC: Thomson et al. 1993) and 3C 293 (NICMOS:  Floyd et al. 2005).  Of these three, only in M87 were the observations of sufficient S/N ($\\gsim 30$) to track accurately both the polarization fraction and orientation throughout virtually the entire jet.  The other two objects had much lower S/N (typically $\\sim 10-15$ for 3C 293 and $\\lsim 4-20$ for 3C 273), making conclusions based on these quantities much less certain.   As a result very little is known about the energetic structures of most optical jets, including the extent to which the radio and optical emitting electron populations are co-spatial and encounter a common magnetic environment, as well as the role played by the magnetic field structure in the acceleration of high-energy particles.  We are engaged in a project to obtain {\\it HST} polarimetry of the nearest and brightest optical jets, in  order to remedy this situation.  In this paper, we present polarimetric observations of six nearby jets. These objects, namely 3C 15, 3C 66B, 3C 78, 3C 264, 3C 346 and 3C 371, constituted the six brightest optical jets as of 1999, when the project was first proposed for HST observations.  This paper presents the {\\it HST} observations in an atlas form, as well as a qualitative discussion of the polarization properties of individual jets and sample properties.  Upcoming papers will concentrate on individual jets and physical interpretation  (for example, 3C 15:  Dulwich et al. 2006; 3C 66B:  Perlman et al. in preparation; 3C 264:  Padgett et al., in preparation), including multiwavelength data as well as the relationship of polarized emissions to the spatially resolved optical and broadband structure and spectrum of each jet (e.g., for M87, Perlman et al. 1999, 2001, Perlman \\& Wilson 2005).  The paper is laid out as follows. In \\S 2 we discuss the observations and data reduction procedures. In \\S 3, we detail the results on each object.  In \\S 4, we compare them to the few objects for which previous {\\it HST} polarimetry has been done, and present the first discussion of the sample properties of jets in polarized optical light.  We close our discussion in \\S 5 with a summary. Throughout the paper we assume a cosmology with $H_0=70 {\\rm ~km ~s^{-1} ~Mpc^{-1}}$ (Riess et al. 2005 and references therein), $\\Omega_{m,0}=0.3$ and $\\Omega_{\\Lambda, 0}=0.7$ (Riess et al. 2004; Bennett et al. 2003; Perlmutter et al. 1999). ",
        "conclusions": ""
    },
    "0606/astro-ph0606433_arXiv.txt": {
        "abstract": "We introduce a \\MC\\ model of \\NL\\ diffusive shock acceleration allowing for the generation of large-amplitude magnetic turbulence, i.e., $\\Delta B \\gg B_0$, where $B_0$ is the ambient magnetic field.  The model is the first to include strong wave generation, efficient particle acceleration to \\rel\\ energies in \\nonrel\\ shocks, and thermal particle injection in an internally \\SC\\ manner.  We find that the upstream magnetic field $B_0$ can be amplified by large factors and show that this amplification depends strongly on the ambient \\alf\\ Mach number. We also show that in the nonlinear model large increases in $B$ do not necessarily translate into a large increase in the maximum particle momentum a particular shock can produce, a consequence of high momentum particles diffusing in the shock precursor where the large amplified field converges to the low ambient value. To deal with the field growth rate in the regime of strong fluctuations, we extend to strong turbulence a parameterization that is consistent with the resonant quasi-linear growth rate in the weak turbulence limit. We believe our parameterization spans the maximum and minimum range of the fluctuation growth and, within these limits, we show that the \\NL\\ shock structure, acceleration efficiency, and thermal particle injection rates depend strongly on the yet to be determined details of wave growth in strongly turbulent fields.  The most direct application of our results will be to estimate magnetic fields amplified by strong cosmic-ray modified shocks in supernova remnants. ",
        "introduction": "Recent observations and modeling of several young supernova remnants (SNRs) suggest the presence of magnetic fields at the forward shock (i.e., the outer blast wave) well in excess of what is expected from simple compression of the ambient circumstellar field, $\\Bcsm$.  These large fields are inferred from (i) spectral curvature in radio emission \\citep[e.g.,][]{RE92,BKP99}, (ii) broad-band fits of \\syn\\ emission between radio and non-thermal X-rays \\citep[e.g.,][]{BKV2003,VBK2005a} \\cite[see also][]{Cowsik80}, and (iii) sharp X-ray edges \\citep[e.g.,][]{VL2003,BambaEtal2003,VBK2005a,EC2005}. While these methods are all indirect, fields greater than 500~\\muG\\ are inferred in Cas A and values of at least several 100 \\muG\\ are estimated in Tycho, Kepler, SN1006, and G347.3-0.5. If $\\Bcsm \\sim 3-10$~\\muG, amplification factors of 100 or more may be required to explain the fields immediately behind the forward shocks and this is likely the result of a nonlinear amplification process associated with the efficient acceleration of cosmic-ray ions via diffusive shock acceleration (DSA). The magnetic field strength is a critical parameter in DSA and also strongly influences the \\syn\\ emission from shock accelerated electrons. Since shocks are expected to accelerate particles in diverse astrophysical environments and \\syn\\ emission is often an important emission process (e.g., radio jets), quantifying the magnetic field amplification has become an important problem in particle astrophysics and has relevance beyond cosmic-ray production in SNRs.  If \\CR\\ (CR) production is as efficient as expected in theories of nonlinear DSA \\citep[e.g.,][]{BE87,JE91,MD2001}, the CR pressure gradient in the shock precursor can do work on the incoming plasma and, in principle, place a large amount of energy in magnetic turbulence. In fact, if the DSA process is to work at all, magnetic turbulence must be self-generated on all resonant scale lengths to provide the scattering necessary to drive the particle distribution to isotropy. If the turbulence remains weak, i.e., if $\\Delta B/B \\ll 1$, the self-generation of turbulence by CR streaming in the shock precursor can be adequately described using quasi-linear theory \\citep[e.g.,][]{Skilling75a,MV82}. For large amplitude turbulence, however, this analytic description becomes questionable and less rigorous approximations must be made.  In principle, a complete description of \\NL\\ DSA, including \\MFA, is possible with particle-in-cell (PIC) simulations. In reality, however, these simulations are still far too computationally demanding to produce realistic models of astrophysical sources such as supernova remnants. Approximate methods describing \\NL\\ DSA must be used. Here, we have developed a method which incorporates a phenomenological model of particle diffusion and turbulence growth \\citep[similar to][]{BL2001,AB2006} in a fully \\NL\\ \\MC\\ model of DSA \\citep[e.g.,][]{JE91,EBJ96,BE99}.  While not a first-principles description of the plasma physics, the computational efficiency of our approach gives it a significant advantage over PIC simulations and \\SA\\ techniques based on the diffusion approximation in that we can determine the shock structure, the injection of thermal particles, the acceleration of these particles to \\rel\\ energies, and the magnetic turbulence, all \\SCly\\ in a fully \\NL\\ steady-state model.  An important advantage of the \\MC\\ simulation is that, in effect, it solves a more basic set of equations governing the shock structure and DSA than do techniques based on diffusion--convection equations \\citep[e.g.,][]{EE84}.  There is no assumption of isotropy for particle distributions and this allows an internally \\SC\\ treatment of thermal particle injection.  Injection is still phenomenological and depends on the assumptions made for the particle pitch-angle scattering, but these assumptions are applied equally to all particles. Since there is only one set of scattering assumptions, the \\MC\\ technique eliminates a free ``injection parameter'' which is present in all models based on the diffusion approximation to set the injection efficiency. As we show, the strong feedback between injection, shock structure, and \\MFA\\ makes this property of the \\MC\\ technique particularly important.  Our preliminary results show that magnetic field amplification results in an increase in the maximum CR momentum, $\\Pmax$, a given shock system can produce compared to the case without amplification. For the case we consider, acceleration truncated by a finite-size shock, this increase is not, however, as large as the increase in $B$ in the shocked region since fields on scales of the upstream diffusion length of the highest energy particles strongly influence $\\Pmax$. Furthermore, the fact that $B$ ranges smoothly from the unshocked value far upstream, $B_0$, to the amplified value, $B_2$, downstream from the shock will have consequences for electrons since radiation losses will be greatest in the downstream region where particles spend a large fraction of their time.\\footnote{We consistently use the subscript 0 for far upstream values, the subscript 2 for downstream values, and the subscript 1 for values immediately upstream from the subshock. Thus, the overall compression ratio is $\\Rtot = u_0/u_2$ and the subshock compression ratio is $\\Rsub=u_1/u_2$, where $u$ is the bulk flow speed.} Thus, the maximum momentum protons and electrons obtain may be determined by two very different field strengths.  We also show that the amplification factor, $B_2/B_0$, increases with \\alf\\ Mach number, that is, for a given shock speed and ambient density, a small $B_0$ will be amplified more than a larger $B_0$. The ability to dramatically increase very low ambient fields may make it possible for reverse shocks in young SNRs to accelerate particles and produce \\rel\\ electrons and radio \\syn\\ emission \\cite[see][for a discussion of DSA at reverse shocks is SNRs]{EDB2005}. The injection and acceleration of electrons in amplified fields will be considered in subsequent work, here, we consider only proton acceleration.  In this paper, we demonstrate the feasibility of applying \\MC\\ techniques to DSA with field amplification but have made a number of approximations regarding wave generation and have neglected wave damping.  An important advantage of the \\MC\\ method is that it can be generalized to include more realistic descriptions of non-resonant wave growth, linear and \\NL\\ damping, and the calculation of the momentum and space dependent diffusion coefficient from the turbulent energy density, and work along these lines is in progress. While a more accurate description of the plasma physics will influence all aspects of the acceleration process, we expect that $\\Pmax$ will be most strongly determined by the plasma physics details and thus will remain a critical problem for understanding the origin of galactic cosmic rays for some time. ",
        "conclusions": "We have introduced a model of diffusive shock acceleration which couples thermal particle injection, \\NL\\ shock structure, \\MFA, and the \\SC\\ determination of the maximum particle momentum. This is a first step toward a more complete solution and, in this preliminary work, we make a number of approximations dealing mainly with the plasma physics of wave growth. Keeping in mind that our results are subject to the validity of our approximations, we reach a number of interesting conclusions.  First, our calculations find that efficient shock acceleration can amplify ambient magnetic fields by large factors and are generally consistent with the large fields believed to exist at blast waves in young SNRs, although we have not attempted a detailed fit to SNR observations in this paper. While the numerical values we obtain depend on the particular parameters for our examples, and we will investigate in detail how the amplification depends on sonic Mach number, age, and size of the shock system, etc., in future work, amplification factors of several 100 are clearly possible.  More specifically, we find that the amplification, in terms of the downstream to far upstream field ratio $B_2/B_0$ is a strong function of \\alf\\ Mach number, with weak ambient fields being amplified more than strong ones. For the range of examples shown in Figure~\\ref{vary_Bz}, $B_2/B_0 \\sim 30$ for $\\Malf \\sim 80$ and $B_2/B_0 \\sim 400$ for $\\Malf \\sim 8000$. Qualitatively, a strong correlation between amplification and $\\Malf$ should not depend strongly on our approximations and may have important consequences. In young SNRs, the expansion of the ejecta material will drastically reduce any original, pre-SN circumstellar magnetic field. In fact, for any conceivable progenitor, the magnetic field inside of the reverse shock will drop to values too low to support the acceleration of electrons to radio emitting energies only a few years after the explosion \\citep[e.g.,][]{EDB2005}. Evidence for radio emission at reverse shocks in SNRs has been reported \\citep[see,][for example]{GotthelfEtal2001} and the strong amplification of low fields we see here, may make it possible for reverse shocks in young SNRs to accelerate electrons to \\rel\\ energies and produce radio \\syn\\ emission. If similar effects occur in \\rel\\ shocks, these large amplification factors will be critical for the internal shocks presumed to exist in \\gamray\\ bursts (GRBs). Even if large \\BFA\\ is confined to \\nonrel\\ shocks, which tend to be more efficient accelerators than \\rel\\ ones, amplification will be important for understanding GRB afterglows since the expanding fireball will slow as it moves through the interstellar medium and will always go through \\transrel\\ and \\nonrel\\ phases.  As expected, amplifying the magnetic field leads to a greater maximum particle momentum, $\\Pmax$, a given shock can produce. Quantifying $\\Pmax$ is one of the outstanding problems in shock physics because of the difficulty in obtaining parameters for typical SNRs that allow the production of cosmic rays to energies at and above the CR knee near $10^{15}$\\,eV. Assuming that acceleration is truncated by the size of the shock system, we determine $\\Pmax$ from a physical constraint: the relevant parameter is the distance to the free escape boundary in diffusion lengths. This means that the limit on acceleration feeds back on the shock structure and also mimics, in terms of the spectral shape, what happens in actual shocks where $f(x,p)$ must turn over smoothly at the highest energies (as in Figure~\\ref{vary_Bz_fp}). The spectral shape will be particularly important if the model is applied to the knee of the cosmic-ray spectrum or to nonthermal X-ray emission in SNRs.  Our results show that $\\Pmax$ does increase when field amplification is included, but the increase is considerably less than the amplification factor at the shock $B_2/B_0$ (compare the heavy dotted and heavy solid curves in Figure~\\ref{BL_noBL_fp}). The main reason for this is that high momentum particles have long diffusion lengths and the precursor magnetic field well upstream from the subshock strongly influences $\\Pmax$. We calculate the spatial structure of the amplified magnetic field (Figs.~\\ref{BL_noBL} and \\ref{vary_Bz}) and, as expected, field amplification is greatest near the subshock and $\\Beff$ merges into the ambient field far upstream. The diffusion length that determines $\\Pmax$ (i.e., $[\\DiffPmax/u(x)]_\\mathrm{ave}$) is some weighted average over the varying $u(x)$ and $\\Beff(x)$ and is considerably greater than that estimated from $B_2$ alone. If the shock size, in our case $\\Dfeb$, limits acceleration, $\\Pmax$ will be considerably less than crude estimates using a spatially independent $B_2$.  On the other hand, particles spend a large fraction of their time downstream from the shock where the field is high and collision times are short. If shock age limits acceleration rather than size, we expect the increase in $\\Pmax$ from the amplified field to be closer to the amplification factor, $B_2/B_0$. Thus, the spatial structure of the precursor field and $u(x)$, in addition to the overall amplification of $B$, will determine the relative time spent upstream versus downstream and will determine $\\Pmax$ for a given shock system. This will be even more critical for electrons than for ions since electrons experience \\syn\\ and \\IC\\ losses which will mainly occur downstream. Again, the qualitative nature of the above conclusions should not depend on the particular parameters and approximations we make here. We leave more detailed quantitative work for future study.  Finally, it is well known that DSA is inherently efficient. Field amplification reduces the fraction of shock ram kinetic energy that is placed in \\rel\\ particles but, at least for the limited examples we show here, the overall acceleration process remains extremely efficient. Even with large increases in $\\Beff(x)$, well over 50\\% of the shock energy can go into \\rel\\ particles (Figure~\\ref{inj_eff}).  As in all \\SC\\ calculations, the injection efficiency must adjust to conserve momentum and energy. In comparing shocks with and without field amplification, we find that field amplification lowers $\\Rtot$ and, therefore, individual energetic particles are, on average, accelerated less efficiently. In order to conserve momentum and energy, this means that more thermal particles must be injected when amplification occurs. The shock accomplishes this by establishing a strong subshock which not only injects a larger fraction of particles, but also more strongly heats the downstream plasma. This establishes a \\NL\\ connection between the field amplification, the production of cosmic rays, and the X-ray emission from the shocked heated plasma."
    },
    "0606/astro-ph0606355_arXiv.txt": {
        "abstract": "We discuss the colors of 99,088 galaxies selected from the Sloan Digital Sky Survey Data Release 1 ``main'' spectroscopic sample (a flux limited sample, $r_{Pet}<17.77$, for 1360 deg$^2$) in the rest-frame Str\\\"omgren system (\\U, \\V, \\B, \\Y). This narrow-band (\\about 200~\\AA) photometric system, first designed for the determination of effective temperature, metallicity and gravity of stars, measures the continuum spectral slope of galaxies in the rest-frame 3200--5800~\\AA\\ wavelength range. We synthesize rest-frame Str\\\"omgren magnitudes from SDSS spectra, and find that galaxies form a remarkably narrow locus (\\about 0.03~mag) in the resulting color-color diagram. Using the Bruzual~\\& Charlot population synthesis models, we demonstrate that the position of a galaxy along this locus is controlled by metallicity and age of the dominant stellar population. The distribution of galaxies along the locus is bimodal, with the local minimum corresponding to an \\about 1~Gyr old single stellar population. The position of a galaxy perpendicular to the locus is independent of metallicity and age, and reflects the galaxy's dust content, as implied by both the models and the statistics of IRAS detections.  Comparison of the galaxy locus in the rest-frame Str\\\"omgren color-color diagram with the galaxy locus in the $H_\\delta$-$D_n(4000)$ diagram, utilized by Kauffmann et al.~(2003) to estimate stellar masses, reveals a tight correlation, although the two analyzed spectral ranges barely overlap. Furthermore, the rest-frame $r-i$ color (5500--8500~\\AA\\ wavelength range) can be ``predicted'' with an rms of 0.05~mag using the rest-frame Str\\\"omgren colors. {\\it Overall, the galaxy spectral energy distribution in the entire UV to near-IR range can be described as a single-parameter family with an accuracy of 0.1~mag, or better.} This nearly one-dimensional distribution of galaxies in the multi-dimensional space of measured parameters strongly supports the conclusion of Yip et al. (2004), based on a principal component analysis, that SDSS galaxy spectra can be described by a small number of eigenspectra. Furthermore, the rest-frame Str\\\"omgren colors correlate tightly with the classification scheme proposed by Yip et al. (2004) based on the first three eigenspectra. Apparently, the contributions of stellar populations that dominate the optical emission from galaxies are combined in a simple and well-defined way. We also find a remarkably tight correlation between the rest-frame Str\\\"omgren colors of emission-line galaxies and their position in the Baldwin-Phillips-Terlevich diagram. These correlations between colors and various spectroscopic diagnostic parameters support earlier suggestions that rest-frame Str\\\"omgren photometry offers an efficient tool to study faint cluster galaxies and low surface brightness objects without performing time-consuming spectral observations. ",
        "introduction": "The studies of galaxies have been recently invigorated due to the advent of modern sensitive large-area surveys across a wide wavelength range. The Sloan Digital Sky Survey (SDSS, York et al. 2000, Stoughton et al. 2002, Abazajian et al. 2003) stands out among these surveys because it has already provided UV to near-IR five-color imaging data, and high-quality spectra for over 100,000 galaxies. The spectroscopic galaxy sample is defined by a simple flux limit (Strauss et al. 2002), and will eventually include close to 1,000,000 galaxies.  A number of detailed galaxy studies based on SDSS data have already been published. Strateva et al. (2001) and Shimasaku et al. (2001) demonstrated a tight correlation between the $u-r$ color, concentration of the galaxy's light profile as measured by the SDSS photometric pipeline {\\it photo} (Lupton et al. 2002), and morphology. Blanton et al. (2001) presented the SDSS galaxy luminosity function, and Kauffmann et al. (2003ab) determined and analyzed stellar masses and star-formation histories for 100,000 SDSS galaxies. A  large number of other, published and ongoing, studies are based on the rich dataset provided by SDSS.  SDSS galaxy spectra have high quality (R$\\sim$1800, 3800-9200~\\AA\\ wavelength range, spectrophotometric calibration better than 10\\%; for more details see Stoughton et al. 2002), and have been used in a number of detailed investigations, including studies of elliptical (Eisenstein et al. 2003, Bernardi et al. 2003abcd), star-forming (Hopkins et al. 2003), and active galaxies (Kauffmann et al. 2003c, Hao et al. 2005ab, Zakamska et al. 2003, 2004). Yip et al. (2004) analyzed the SDSS spectra of 170,000 galaxies using principal component analysis (Karhunen-Lo\\'{e}ve transform). They demonstrated that SDSS galaxy spectra can be described by a small number of eigenspectra: more than 99\\% of the  galaxies are found on a two-dimensional locus in the space spanned by the ratios of the first three eigencoefficients. Even the most extreme emission-line spectra can be described to within the measurement noise with only 8 eigenspectra. An efficient {\\it single-parameter} classification scheme is proposed, and it suggests three major galaxy classes. Another remarkable result of their analysis is a strong correlation between various spectral lines, encoded in, and readily discernible from the eigenspectra.  It is not yet known what the results of Yip et al. imply for the detailed distribution of galaxies in color space, but it is obvious that a high degree of structure should exist. Also, a strong correlation is expected between colors and various spectroscopic parameters, such as the strength of the $H_\\delta$ line and the 4000~\\AA\\ break, that were recently utilized by Kauffmann et al. (2003a) to estimate stellar masses and the the dust content for SDSS galaxies.  The main goal of this paper is to investigate whether the intrinsic simplicity of galaxy spectra implied by the results of Yip et al. can be reproduced in rest-frame color-color diagrams, and whether the position of a galaxy in these diagrams is related to parameters discussed by Kauffmann et al. (2003a). We present an analysis of correlations between colors and various spectroscopic diagnostic parameters, including model-dependent estimates of stellar masses and dust content. We utilize the ``rest-frame'' Str\\\"omgren photometric system that measures the continuum spectral slope of galaxies in the rest-frame 3200--5800~\\AA\\ wavelength range. In principle, any set of bandpasses could be used to synthesize the colors. An additional motivation for synthesizing galaxy colors in the rest-frame Str\\\"omgren photometric system is a suggestion that it is an efficient tool for studying faint cluster galaxies and low surface-brightness objects without performing time-consuming spectral observations\\footnote{Observations in the rest-frame Str\\\"omgren photometric system use filters adjusted (in hardware) to the known cluster redshift (say, from the brightest cluster galaxy).} (Fiala, Rakos \\& Stockton 1986). The unprecedented number of high-quality SDSS spectra can be used to analyze in detail the advantages and limitations of that method, and relate galaxy colors in the Str\\\"omgren system to parameters commonly measured using SDSS data.  We describe the synthesis of rest-frame Str\\\"omgren magnitudes from SDSS spectra in Section 2. In Section 3 we analyze the distribution of galaxies in the resulting color-color diagram, and interpret it using the Bruzual~\\& Charlot (1993, 2003) population synthesis models. In Section 4 we demonstrate and study in detail correlations between rest-frame Str\\\"omgren colors and various diagnostic spectroscopic parameters. In Section 5 we correlate the rest-frame Str\\\"omgren colors with the spectral eigencoefficients. In Section 6 we study the properties of the galaxies at sufficiently large redshifts ($\\gtrsim0.18$) to have reliable synthesized \\U~magnitudes. We discuss and summarize our results in Section 7.    \\begin{figure} \\centering \\includegraphics[bb=20 160 532 629, width=\\columnwidth]{smolcic.fig1.ps} \\caption{The top panel shows renormalized (see eq.~\\ref{phi}) filter transmission curves for the SDSS photometric system (ugriz), and for the Str\\\"omgren photometric system. The bottom panel emphasizes the 3800-7000~\\AA\\ region. The two spectra are typical for blue and red galaxies, and the four labeled emission lines are used to separate star-forming from AGN galaxies.  The horizontal bar in the top left corner marks the wavelength region used in the analysis by Kauffmann et al. (2003).} \\label{fltsplot} \\end{figure} ",
        "conclusions": "We analyzed the synthesized rest-frame Str\\\"omgren colors for \\about100,000 SDSS galaxies. This narrow-band photometric system is well suited for studying galaxy colors in the 4000--5800~\\AA\\ wavelength range. Galaxies form a narrow locus in the rest-frame Str\\\"omgren $bz-yz$ vs. $vz-yz$ color-color diagram, with a width of only 0.03~mag. This small width, in addition to demonstrating that the errors in synthesized colors are small, shows that the slope of galaxy spectra in the 3200--5800~\\AA\\ wavelength range is practically a one-parameter family. This conclusion is independent of the details of the employed photometric system.  Using the Bruzual~\\& Charlot population synthesis models, we find that the dependence on age and metallicity is {\\it fully absorbed} in the first principal color axis, $P1$. This color is essentially a measure of the Hubble type (in sense of early vs. late type galaxies, rather than detailed morphological classification). It is not possible to break the age-metallicity degeneracy with any of the observables discussed here. This may not be possible even when using standard methods such as Lick indices, at least in the case of elliptical galaxies (Eisenstein et al. 2003). As demonstrated by Eisenstein et al., the strengths of numerous metal lines, such as Mg b, can be accurately predicted from the strength of the $H_\\delta$ line (see their Fig. 5); hence, no additional information can be extracted from the measurements of those lines. On the other hand, Tremonti et al. (2004) report a very tight mass-metallicity relation (0.1 dex scatter in metallicity at a given mass) for \\about50,000 SDSS star-forming galaxies. Thus, it may be that age and metallicity are strongly correlated, too, which then would provide a trivial solution to the problem of breaking the age-metallicity degeneracy. Indeed, such a relation has been claimed for disk stars in our Galaxy (Twarog 1980; however, see also Feltzing, Holmberg \\& Hurley 2001).  The second principal color axis, $P2$, is independent of stellar age and metallicity, at least according to the Bruzual~\\& Charlot models. We stress that the definition of this color is purely {\\it observational}, and the models explain both its value and narrow distribution width. The $P2$ distribution width (0.03~mag) measured using SDSS data is smaller than typical observational errors in rest-frame colors using other data available to date ($\\ga$0.05~mag). The $P2$ color appears most sensitive to the galaxy dust content, as demonstrated by the analysis of model-dependent estimates of the dust content by Kauffmann et al. (2003a), and also by the redder $P2$ colors for galaxies detected by the IRAS survey. Of course, $P2$ could be {\\it indirectly} related to metallicity, if the dust content depends on metallicity (that is, the observed $P2$ color could be correlated with metallicity as a result of radiative transfer through interstellar dust, rather than due to the dependence of the source function on metallicity).  The principal colors are well correlated with numerous other spectral parameters determined from bluer and redder spectral ranges, as well as with the position of a galaxy in the BPT diagram constructed with emission-line strength ratios. These good correlations suggest that not only are the overall UV-IR SEDs a nearly one-parameter family, but that the same parameter also controls the properties of emission lines. Most notably, there is a good correspondence between the stellar mass-to-light ratio estimates by Kauffmann et al. (2003a) and the principal color $P1$.  The low dimensionality of galaxy spectra implied by the numerous correlations among the continuum and spectral parameters shown here is in good agreement with the conclusions by Yip et al. (2004), who applied the principal component analysis to SDSS galaxy spectra. The same conclusion was reached in a detailed study of elliptical galaxies by Eisenstein et al. (2003). These conclusions are valid for the overwhelming majority of all galaxies; however, we emphasize that the analysis presented here is based on low-order statistical moments (such as root-mean-square scatter) and cannot rule out a possibility that as many as 5-10\\% of galaxies behave differently.  Our results support the suggestion by Rakos and collaborators (Fiala, Rakos \\& Stockton 1986) that the rest-frame Str\\\"omgren photometry provides an efficient tool for studying faint cluster galaxies and low surface brightness objects, without the need for time-consuming spectroscopy. For example, the stellar mass-to-light ratio can be determined from $P1$ with an rms scatter of only 0.10. Since the brightest cluster galaxies can be selected from SDSS data to redshifts of \\about 0.8 (e.g. Ivezi\\'{c} et al. 2002), this method can be used to quantify the rise in star formation activity with cosmic epoch, and the Butcher-Oemler effect, with an unprecedented accuracy. We are currently undertaking such an observational program."
    },
    "0606/astro-ph0606163_arXiv.txt": {
        "abstract": "{} {Finding tracers of the innermost regions of prestellar cores is important for understanding their chemical and dynamical evolution before the onset of gravitational collapse. While classical molecular tracers, such as CO and CS, have been shown to be strongly depleted in cold, dense gas by condensation on grain mantles, it has been a subject of discussion to what extent nitrogen-bearing species, such as ammonia, are affected by this process. As deuterium fractionation is efficient in cold, dense gas, deuterated species are excellent tracers of prestellar cores. A comparison of the spatial distribution of neutral and ionized deuterated species with the dust continuum emission can thus provide important insights into the physical and chemical structure of such regions.} {We study the spatial distribution of the ground-state 335.5~GHz line of  ND$_2$H in LDN1689N, using APEX, and compare it with the distribution of the DCO$^+$(3--2) line, as well as the 350~$\\mu$m dust continuum emission observed with the SHARC~II bolometer camera at CSO. } {While the distribution of the ND$_2$H emission in LDN1689N is generally similar to that of the 350~$\\mu$m  dust continuum emission, the peak of the  ND$_2$H emission is offset by $\\sim$$10''$ to the East from the dust continuum and DCO$^+$ emission peak. ND$_2$H and ND$_3$ share the same spatial distribution. The observed offset between the ND$_2$H and DCO$^+$ emission is consistent with the hypothesis that the deuterium peak in LDN1689N is an interaction region between the outflow shock from IRAS16293--2422 and the dense ambient gas. We detect the $J = 4 \\rightarrow 3$ line of H$^{13}$CO$^+$ at 346.998 GHz in the image side band serendipitously. This line shows the same spatial distribution as DCO$^+$(3--2), and peaks close to the 350~$\\mu$m emission maximum  which provides further support for the shock interaction scenario.} {} ",
        "introduction": "Deuteration of nitrogen compounds, such as ammonia and N$_2$H$^+$, is spectacular in a number of environments including  dark clouds, such as L134N (Tin\\'e et al. 2000; Roueff et al. 2000; Roueff et al. 2005), low-mass star forming regions and prestellar cores such as LDN1689N and Barnard~1 (Loinard et al. 2001; Gerin et al. 2001; Lis et al. 2002a). The discussion of the respective contributions of grain and gas-phase processes in the deuteration is active, but no definite solution is yet available. The lack of detection of deuterated water in ices toward low-mass young stellar objects (YSOs) by Dartois et al. (2003) and Parise et al. (2003) suggests that ``another mechanism than pure solid state chemistry may be active to produce very high deuterium enrichment in the gas phase.''  Further constraints on the deuterium fractionation mechanisms are provided by the spatial distribution of deuterated species, as compared with the molecular gas distribution traced by the submillimeter dust continuum emission. Whereas maps of singly deuterated species have been published in many cores --- the relatively large line intensities of e.g. DCO$^+$, N$_2$D$^+$ allow easy mapping with state of the art detectors --- mapping multiply deuterated species has proven to be a challenge, given the relatively low line intensities. Ceccarelli et al. (2001) have shown that the D$_2$CO emission is extended around the class 0 protostar IRAS16293--2422. Roueff et al. (2005) present a limited map of the ND$_3$ ground-state transition at 309 GHz with the CSO in LDN1689N, which unfortunately suffers from rather poor pointing accuracy. Because the ground state ND$_2$H $1_{0,1} - 0_{0,0}$ lines, at 335.5 for the ortho species and 335.4 GHz for the para species (See Coudert and Roueff (2006) for NH$_3$ and its isotopologues line frequencies), are relatively strong (0.6 K in LDN1689N; Lis et al. 2006) and the atmospheric transmission is good at this frequency, we have carried out the first map of a doubly deuterated species in a dense core. ",
        "conclusions": "\\begin{enumerate} \\item We show that the ND$_2$H emission is extended in the \\object{LDN1689N} dense core. The abundance relative to H$_2$, derived from LTE, is $\\sim 1.8 \\times 10^{-9}$, which makes ND$_2$H a remarkably abundant molecule in this dense core. The [ND$_3$]/[ND$_2$H] abundance ratio is 0.01 -- 0.02. \\item The ND$_2$H emission is spatially offset from the dust continuum DCO$^+$(3--2) and H$^{13}$CO$^+$(4--3) emission by $\\sim 0.01$~pc, and appears to be blueshifted relative to the cloud envelope. Both findings are qualitatively consistent with the scenario of the formation of the LDN1689N dense core by the interaction of the blue lobe of the  \\object{IRAS16293--2422} molecular outflow with the ambient material. Detailed models could further test this scenario. \\end{enumerate}"
    },
    "0606/astro-ph0606480_arXiv.txt": {
        "abstract": "{}{ We present new radio continuum observations of ten  BIMA SONG galaxies, taken at 1.4 GHz with the Very Large Array. These observations allow us to extend the study of the relationships between the  radio continuum (RC) and CO emission to 22 CO luminous galaxies for which single dish CO images have been added to interferometric data. New \\emph{Spitzer} infrared (IR) images of six of these galaxies  have been released. The analysis of these high resolution images allowed us to probe the RC-IR-CO correlations down to linear scales of a few hundred pc.} {We compare the point-by-point RC, CO and mid-IR intensities across entire galaxy disks, producing radial profiles and spatially resolved images of the RC/CO and RC/mid-IR ratios.} {For the 22 galaxies analysed, the RC-CO correlation on scales from $\\sim10$ kpc down to $\\sim100$ pc is nearly linear and has a scatter of a factor of two, i.e. comparable to that of the global correlations. There is no evidence for any severe degradation of the scatter below the kpc scale. This also applies to the six galaxies for which high-resolution mid-IR data are available. In the case of \\object{\\object{NGC 5194}}, we find that the non-thermal radio spectral index is correlated with the RC/FIR ratio. } {The scatter of the point-by-point correlations does not increase significantly with spatial resolution. We thus conclude that we have not yet probed the physical scales at which the correlations break down. However, we observe local deviations from the correlations in regions with a  high star formation rate, such as the spiral arms, where  we observe a flat radio spectrum and a low RC/FIR ratio. In the intra-arm regions and in the peripheral regions of the  disk, the RC/FIR is generally higher and it is characterized by a steepening of the radio spectrum.} ",
        "introduction": "One of the major goals in studying spiral galaxies is to understand the relationship between the star formation process and the physical conditions in the interstellar medium. Since the discovery that stars form in molecular clouds, several efforts have been directed towards the study of the relation between the emission of the CO molecule, which traces the bulk of molecular gas, and the other star formation indicators.  Global studies  have revealed relationships between the radio continuum (RC), far-infrared (FIR) and CO emissions. The global correlation between the FIR and centimeter-wavelength RC emission is one of the strongest correlations in extragalactic research. Early works on this topic were reported by \\cite{helou85} and \\cite{dejong85}. \\cite{condon92} and more recently \\cite{yun01} found that these two emissions are linearly correlated over five orders of magnitude in luminosity, with an rms scatter of less than a factor of 2. On global scales the FIR emission is well correlated with the CO emission \\citep[e.g.][]{DevYou90}  and the CO emission is well correlated with the RC \\citep{rickard77,israel84,adler91}.  What is most extraordinary about the FIR-RC or CO-RC correlations is that they couple emissions arising from completely different processes. The usual explanation of the FIR-RC correlation invokes massive-star formation which accounts for the thermal radio emission via ionizing stars, for the non-thermal radio luminosity via supernova events, and for the FIR luminosity by means of massive stars heating the dust \\citep{wunderlich88}. However, this basic scenario has too many steps and too many parameters to explain the tightness of the observed correlations. More complex models have been put forward, see \\citet[ hereafter Paper I]{matteo05} for a brief summary, but a '{\\it{definitive}}' model has not yet emerged.  Although there is still no consensus  on the causes of the RC-FIR-CO correlations, nor any adequate explanation for their tightness, these correlations have been used to determine empirically   the star formation rate (SFR) \\citep{cram98} and the star formation efficiency (SFE)  in galaxies \\citep{gao01,gao03,murgia02,rownd99}. More generally it is found that these correlations  are well described by a ``Schmidt Law'' of the type SFR $ \\propto \\rho_{gas}^{N}$, where  $\\rho_{gas}$ is  the total gas density, and the exponent $N$ typically ranges from 1.3 to 1.5 \\citep{kennicutt98b}.    Detailed studies of individual objects at high spatial resolution are essential to understanding whether there is a common physical process behind the observed correlations and possibly synthesizing the various empirical descriptions.  An improved approach to the study of the spatial properties of the FIR-RC correlation was presented in \\cite{marsh95}. They studied the FIR-RC correlation on intermediate scales of $\\sim$ 1--3 kpc  and they found a small but systematic decrease in the FIR-RC ratio as a function of increasing radial distance from the center. In contrast, \\cite{murgia02} found that the ratio between the RC and CO emission is constant, to within a factor of 3, both inside the same galaxy and from galaxy to galaxy, down to kpc size scales. \\cite{leroy05} extended this result to a sample of dwarf galaxies, finding that this relationship between CO and RC is the same in all galaxies. In this context, dwarf galaxies appear to be simple low-mass versions of large spirals.   So far the low resolution of FIR observations has prevented  the comparison of the FIR and RC emission  on kpc scales for all but a few Local Group sources. \\cite{xu92}, studying the Large Magellanic Cloud found evidence for a breakdown in the FIR-RC correlation at a linear scale of  $\\sim$ 70 pc. In the Milky Way, \\cite{boulanger88} analysed the Orion region and showed that the FIR-RC correlation breaks down on the scale of a few hundred parsecs around a region of star formation.  The situation has recently changed with the advent of the satellite  {\\it Spitzer} that provides  the necessary resolution of IR observations to extend the study of the correlation to sub-kpc scales in many external galaxies. These observations can shed new light on the study of  the correlation on local scales. Recently, \\cite{murphy06} presented an initial look at the far-infrared-radio correlation within the star-forming disks of four nearby galaxies using {\\it Spitzer} images.     An ideal sample to investigate the spatially resolved FIR-RC-CO correlations is the BIMA SONG sample. The Berkeley-Illinois-Maryland Association Survey of Nearby Galaxies (BIMA SONG) is an imaging survey of  the CO (J=1-0) emission in 44 nearby spiral galaxies with a typical angular resolution of 6\\arcsec~\\citep{regan01, helfer03}. For half of the galaxies in the BIMA sample single-dish CO images at 55$''$ resolution obtained with the NRAO 12 m telescope\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} have been added to interferometric data. The combined images represent a significant improvement with respect to the previous CO extragalactic surveys, since in each galaxy  a large area can be mapped at high resolutions avoiding the missing flux problem.   Recently we  extended the study of CO-RC correlation down to  a spatial resolution of $\\sim$ 100 pc, by collecting 20 cm D and B-array data of a subsample of  nine bright BIMA SONG galaxies available in the VLA archive (Paper I). For these nine galaxies it has been possible to compare  CO and RC images at two different angular resolutions: a low resolution ($\\sim$ 55$''$) set obtained from the NRAO 12m and VLA D-array, and a high resolution ($\\sim$ 6$''$) set.  The point-by-point RC and CO brightnesses across the entire galaxy disks have been compared and the ratio of the CO and RC flux has been defined as a way to characterize the CO-RC correlation.  In this work we present the study of 13 more  BIMA SONG galaxies which allow us to complete the study, started with Paper I, of 22  CO-luminous BIMA galaxies.   In Section 2 we summarize the observations used in this paper and the data reduction procedure. Section 3 describes briefly how we performed the data analysis. The results are presented in Section 4 for the whole sample, and are discussed in Section 5. We analysed the  24 $\\mu$m-RC correlation for six galaxies of the sample, for which high resolution {\\it Spitzer} mid-IR images are available, at a spatial resolution ranging from $\\sim$400 pc to $\\sim$100 pc. We present the results of this analysis in Section 6. In Section 7 we compare the expectations of the leaky box diffusion model model, proposed in \\matteo, with the observed trend of spectral index versus RC-FIR ratio for \\object{NGC 5194}. In Section 8 we summarize our results.  \\begin{table*}[htbp] \\caption{Galaxy sample } \\begin{center} \\begin{tabular}{llllllll} \\\\ \\hline \\smallskip Source  & $\\alpha (J2000)$& $\\delta (J2000)$ &d &i&$P.A.$& RC3 Type, Nucleus\\\\ \\smallskip name &  (h m s)&($^{\\circ}$\\,$'$\\, $''$)& (Mpc)&$(^{\\circ})$&$(^{\\circ})$  &  & \\\\ \\hline \\\\ \\object{NGC 0628}&01 36 41.7 & +15 46 59 &7.3 &24 & 25 &SA(s)c\\\\ \\object{NGC 1068}&02 42 40.7 & - 00 00 48 &14.4& 33& 13&(R)SA(rs)b, Sy1.8\\\\ \\object{NGC 2903}&09 32 10.1 & +21 30 02 &6.3 &61 & 17&SAB(rs)bc,HII\\\\ \\object{NGC 3351}&10 43 31.1 & +11 42 14 &7.4 &40 & 13&SB(r)b,HII\\\\ \\object{NGC 3521}&11 05 48.0 &- 00 02 04 &7.2 &58 & 164&SAB(rs)bc,HII/L2\\\\ \\object{NGC 3627}&11 20 15.1 & +12 59 22 &11.1 &63 & 176& SAB(s)b, T2/Sy2\\\\ \\object{NGC 3938}&11 52 49.6 & +44 07 14 &17 &24 & 0& SA(s)c,HII\\\\ \\object{NGC 4303}& 12 21 54.9 & +04 28 25 & 15.2 & 27 & 0 & SAB(rs)bc, T2\\\\ \\object{NGC 4321}&12 22 54.8 & +15 49 20 &16.1 &32 & 154& SAB(s)bc,T2\\\\ \\object{NGC 4826}&12 56 43.6 & +21 40 59 &4.1 &54 & 111&(R)SA(rs)ab,T2\\\\ \\object{NGC 5248}&13 37 32.06 & +08 53 07 &22.7 &43 & 110& SAB(rs)bc, HII\\\\ \\object{NGC 5457}&14 03 12.48 & +54 20 55 &7.4 &27 & 40& SAB(rs)cd\\\\ \\object{NGC 7331}&22 37 04.09 &+34 24 56 &15.1&62 & 172&SA(s)b,T2\\\\ \\\\ \\hline \\\\ \\end{tabular} \\end{center} \\begin{list}{}{} \\item[] Coordinates, distances, inclinations, position angles, RC3 type and nuclear classifications from \\cite{helfer03}. Inclination is defined assuming i=0\\degr\\ for face-on galaxies. Position angles are measured in degrees from north  through east. For nuclear spectral classification: H II = H II nucleus, Sy = Seyfert nucleus, L = LINER, and T=transition object, which has [O I] strengths intermediate between those of H II nuclei and LINERs \\end{list} \\label{Gala} \\end{table*}  \\begin{table*}[htbp] \\caption{VLA B-array observation summary} \\begin{center} \\begin{tabular}{lllll} \\\\ \\hline \\smallskip Source   & Date & Duration & Beam & $\\sigma_{I}^{~a)}$\\\\ \\smallskip name &   & hours &HPBW($''$) & ($\\mathrm{\\mu Jy\\;beam^{-1}}$)\\\\ \\hline \\\\ \\object{NGC 0628} & 03-Nov.-2003 & 1.5&$5.7\\times 5.2$& $~40$\\\\ \\object{NGC 2903} & 04,06-Nov-2003 & 1.0&$5.9\\times5.3$& $~35$ \\\\ \\object{NGC 3351} & 28-Jul-1986$^{*}$&1.8&$4.5\\times4.1$&  $~46$\\\\ \\object{NGC 3521} & 04,06-Nov-2003 & 1.0&$4.9\\times4.8$& $~65$\\\\ \\object{NGC 3627} & 04,06-Nov-2003 & 1.0&$5.8\\times5.7$& $~35$\\\\ \\object{NGC 3938} & 04,06,10-Nov-2003 & 2.0 &$5.8\\times4.9$& $~50$\\\\ \\object{NGC 4303} & 13-Sep-1982$^{**}$ & 1.0 & $6.5\\times4.6$&$~100$\\\\ \\object{NGC 4321} & 04,06,10-Nov-2003 & 1.2 &$5.8\\times5.1$& $~30$\\\\ \\object{NGC 4826} & 06,10-Nov-2003 & 1.0&$5.5\\times5.2$& $~45$\\\\ \\object{NGC 5248} & 04,06,10-Nov-2003 & 1.8&$6.1\\times5.2$& $~50$ \\\\ \\object{NGC 5457} & 06-Aug-1990$^{***}$& 1.0&$5.8\\times5.3$& $~50$\\\\ \\object{NGC 7331} & 03-Nov.-2003 &1.0&$5.3\\times4.8$& $~50 $\\\\  \\\\ \\hline \\\\ \\end{tabular}\\\\ $^{*}$ Archival data: program A073;\\\\ $^{**}$ Archival data: program COND;\\\\ $^{***}$ Archival data: program AV180; \\end{center} \\begin{list}{}{} \\item[$a)$] Final noise of total intensity images. \\end{list} \\label{Obs} \\end{table*}  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.86\\textwidth]{5002fg1a_sc.ps} \\end{center} \\caption[]{$(a)$ Radio contours of galaxies overlaid on the optical images (grey scale), taken from the Digitized Sky Surveys (DSS) archive. See Table \\ref{Obs} for sensitivity and beam sizes of radio images. Contours start at 2~$\\sigma$, and increase by steps of a factor of 2.} \\label{figura1a} \\end{figure*} \\setcounter{figure}{0}   \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.86\\textwidth]{5002fg1b_sc.ps} \\end{center} \\caption[]{$(b)$ See Fig. \\ref{figura1a}.$a$. } \\end{figure*}  \\begin{figure*}[ht!] \\begin{center} \\includegraphics[width=16.5cm]{5002fg2a_sc.ps} \\end{center} \\caption[] { {\\bf (a) \\object{NGC 0628}.}  Radio continuum and CO comparison in BIMA SONG galaxies. {\\it (Upper left)} RC contours overlaid on CO emission for 55\\arcsec\\ images. {\\it (Lower left)} RC contours overlaid on CO emission for 6\\arcsec\\ images. Contours start at 5$\\sigma$ and scale by a factor of $\\sqrt{2}$. {\\it (Upper right)} \\i14\\ as a function of \\ico\\  for the 55\\arcsec\\ and 6\\arcsec\\ data sets.  The insets show the box grids used; the grid size corresponds to the resolution, so that each point is essentially independent. The dotted (dashed) lines show the mean (dispersion) of the \\i14\\ -- \\ico\\ relation  obtained for the  sample containing 22 galaxies (this work and \\matteo).  The solid blue line is a weighted fit to the points  which takes into account the errors in both coordinates. {\\it (Lower right)} Radial profiles of \\ico\\ and \\i14\\ {\\it (top)} and \\Qco\\ $\\equiv$ \\ico/\\i14 {\\it (bottom)} for the 55\\arcsec\\ and 6\\arcsec\\ images. The mean and dispersion of \\qco\\ $\\equiv \\log (Q_{\\rm CO/RC})$ is also given. } \\label{gal1} \\end{figure*} ",
        "conclusions": "\\begin{enumerate} \\item{We obtained new 1.4 GHz radio continuum images at 6\\arcsec\\ resolution for a sample of 13 nearby spiral galaxies belonging to the BIMA SONG survey.} \\item{We analysed the behaviour of the point-by-point correlation between the RC and CO intensities from kpc to sub-kpc scales in these objects, extending the previous study by \\cite{matteo05} to 22 galaxies.} \\item{For most galaxies we find that the RC and CO intensities are nearly linearly correlated (slope $\\simeq$ 1.1 $\\pm$ 0.3) over four orders of magnitude, with a scatter of a factor of 2 (0.3 dex).  However, for four galaxies, namely \\object{NGC 0628}, \\object{NGC 1068}, \\object{NGC 3938} and \\object{NGC 5457}, the correlation is not seen. \\object{NGC 1068} is a Seyfert galaxy whose RC emission is strongly contaminated by the AGN. \\object{NGC 0628}, \\object{NGC 3938} and \\object{NGC 5457} have low CO and RC surface brightness. The deviation from the RC-CO correlation is related to the fact that the RC emission is fainter than average at 6\\arcsec\\ resolution.}  \\item{In the inner 10 kpc of the disks of the 22 galaxies analysed in this work the \\qco\\ ratio is constant with radius. The few exceptions are again \\object{NGC 0628}, \\object{NGC 1068} and \\object{NGC 3938}.} \\item{In order to determine whether the goodness of the RC-CO correlation degrades going from kpc to sub-kpc scales, we studied the variation of the correlation scatter and the Pearson coefficient as a function of the spatial resolution of our images. We find that the median scatter increases from 0.15 to 0.25, going from $\\sim$10 kpc to 0.1 kpc. A similar trend characterizes the correlation coefficient from about 0.93 to 0.84. We conclude that the point-by-point RC-CO correlation deteriorates slightly when increasing the spatial resolution. However, we cannot identify any characteristic scale of decorrelation.} \\item{We present the point-by-point 24~$\\mu$m-RC and  24~$\\mu$m-CO correlations at sub-kpc scales for six galaxies of our sample. We probe the correlations on spatial scales ranging from $\\sim$\\,450 pc to $\\sim$\\,120 pc. The RC-CO and 24$\\mu$m-CO correlations are comparably tight,  with a scatter of $\\sim$\\,0.2 dex. The 24$\\mu$m-RC correlation has a scatter of only 0.1 dex and is somewhat flatter than the FIR-CO and RC-CO correlations. Moreover, we find that the scatter of the correlations does not increase with increasing spatial resolution. } \\item{For the galaxy \\object{NGC 5194}, using the {\\it Spitzer} 70 $\\mu$m image, it has been possible to test the leaky box diffusion model proposed in \\matteo. A point-by-point comparison between spectral index and RC-FIR ratio, measured at a resolution of $\\sim$\\,16\\arcsec, shows that regions with a  flat spectral index have a reduced RC-FIR ratio. } \\end{enumerate}"
    },
    "0606/astro-ph0606449_arXiv.txt": {
        "abstract": "We undertake a theoretical study of the near-infrared (NIR) lightcurves of Type~Ia supernovae (SNe~Ia).  In these bands, the lightcurves are distinguished by a secondary maximum occurring roughly 20 to 30 days after the initial one.  Using time-dependent multi-group radiative transfer calculations, we calculate the UBVRIJHK-band lightcurves of model SN~Ia ejecta structures.  Our synthetic NIR lightcurves show distinct secondary maxima, and provide favorable fits to observed SNe~Ia.  We offer a detailed explanation of the origin of the NIR secondary maximum, which is shown to relate directly to the ionization evolution of iron group elements in the ejecta.  This understanding provides immediate intuition into the dependence of the NIR lightcurves on the physical properties of the ejecta, and in particular explains why brighter supernovae have a later and more prominent secondary maximum.  We demonstrate the dependence of the NIR lightcurves on the mass of \\Nifs, the degree of \\Nifs\\ mixing, the mass of electron caputre elements, the progenitor metallicity, and the abundance of intermediate mass elements (especially calcium).  The secondary maximum is shown to be a valuable diagnostic of these important physical parameters.  The models further confirm that SNe~Ia should be excellent standard candles in the NIR, with a dispersion $\\la 0.2$~mag even when the physical properties of the ejecta are varied widely.  This study emphasizes the consummate value of NIR observations in probing the structure of SNe~Ia and in furthering their cosmological utility. ",
        "introduction": "The extensive monitoring of Type~Ia supernovae (SNe~Ia) has typically focused on optical band observations, with only occasional ventures into the infrared \\citep{Kirshner_1973,Elias_1981,Elias_1985}.  The situation is now beginning to change, and near-infrared (NIR) observations over the two or three months following explosion are becoming more common, at least for the more nearby events.  About a dozen SNe~Ia have published, well observed JHK-band lightcurves \\citep[e.g.,][]{Meikle_IR2000,Hernandez_98bu,Valentini_03E, Candia_00cx, Kris_2001AJ,Kris_01el,Kris_2004AJ2,Kris_2004AJ1}. NIR spectra have been obtained for several objects as well \\citep{Meikle_94D, Bowers_86G,Rudy_00cx, Hoeflich_99by, Marion_IR}. Future SN~Ia surveys promise to gather NIR lightcurves for a statistically interesting sample of events.  With the broadening of our wavelength horizons comes the potential to probe the nature of SN~Ia explosions in new ways.  Observational studies of SNe~Ia in the NIR have often emphasized their cosmological utility \\citep{Kris_IRHubble}.  Normal SNe~Ia are found to be excellent \\emph{standard} (as opposed to calibrated) candles in the NIR, with an intrinsic dispersion of less than 0.2~mag in the J,H, and K bands \\citep{Elias_1985,Meikle_IR2000}.  Moreover, dust extinction is smaller by a factor $\\sim 5$ in the NIR as compared to the V band, largely eliminating uncertainties in the reddening corrections. The primary challenge for the cosmology studies, naturally, is that SNe~Ia are much fainter and harder to observe in the NIR.  The NIR lightcurves of SNe~Ia possess a morphology distinct from those at optical wavelengths.  The I through K-band lightcurves are distinguished by the presence of a secondary maximum, occurring roughly 20 to 30 days after the initial one.  A corresponding ``shoulder'' can often be seen also in the R and V-band lightcurves. When the light from all bands is suitably integrated, the bolometric lightcurves sometimes show an inflection at these times as well \\citep{Contardo_bol}.  In the I-band, the properties of the secondary maximum are found to correlate with the lightcurve decline rate (and hence peak luminosity) of the supernova, being more prominent and occurring later in the broader/brighter SNe~Ia \\citep{Hamuy_96Templates, Nobili_Iband}.  Very subluminous objects may lack a secondary maximum entirely.  Similar trends appear in the J, H, and K band lightcurves, although there are interesting exceptions to the rule \\citep{Kris_2001AJ,Candia_00cx}.  Occasional theoretical studies have touched upon the NIR lightcurves of SNe~Ia.  \\cite{Hoeflich_DD} computed I-band and NIR lightcurves for several delayed-detonation models, many of which displayed secondary maxima.  The authors explained the double-peaked behavior as a ``temperature-radius'' effect, in which the expansion of the photosphere compensates for the declining temperature in the ejecta. \\cite{Wheeler_IR} and \\cite{Hoeflich_99by} have further demonstrated the value of NIR spectra in diagnosing the physical conditions in the ejecta. \\cite{Pinto-Eastman_II} described the NIR secondary maximum as the release of pre-existing, trapped radiative energy due to a sudden decrease in the flux mean opacity.  They placed special emphasis on the role played by stable iron group elements at the ejecta center in this process. Most significantly, \\cite{Pinto-Eastman_II} recognized the importance of the ionization state of the ejecta, in particular the increased NIR emissivity of singly as opposed to double ionized iron group elements. This last idea is confirmed and extended in the models studied here.  Despite the theoretical insights, the origin of the NIR secondary maximum and its dependence on the SN ejecta properties have remained in many ways, and to many people, obscure.  Moreover, theoretical models of SNe~Ia have often had difficulty in fitting the distinctive double-peaked I-band lightcurves \\citep[e.g.,][]{Hoeflich_Khokhlov_LC,Pinto-Eastman_II,Blinnikov_2004} whereas very few NIR lightcurve calculations have been attempted \\citep[c.f.][]{Hoeflich_DD}.  This paper focuses exclusively on theoretical models of the far-red and NIR lightcurves of SNe~Ia.  We calculate the synthetic lightcurves of parameterized ejecta configuration using the time-dependent multi-dimensional radiative transfer code \\Code\\ (\\S\\ref{Sec:Technique}).  Our model lightcurves show distinct secondary maxima and provide reasonable fits to the observed R, I, J and (to a lesser extent) H and K band lightcurves of SNe~Ia (\\S\\ref{Sec:Fiducial}).  Close examination of our model calculations provides a clear-cut explanation for the NIR secondary maximum in SNe~Ia (\\S\\ref{Sec:Explain}).  The double-peaked behavior is related directly to the ionization evolution of iron group elements in the SN ejecta. In particular, the NIR emissivity of iron/cobalt increases sharply at a temperature $T \\approx 7000$~K marking the transition between the singly and doubly ionized states.  The eventual cooling of the iron-rich layers of ejecta to this temperature is thus accompanied by a sudden increase in emission at far red and NIR wavelengths.  Understanding the origin of the secondary maximum provides immediate intuition into the dependence of the NIR lightcurves on the ejecta properties, and allows for theoretical explanations for the empirically established correlations (\\S\\ref{Sec:Depend}).  In particular we study the effect of variations in the mass of \\Nifs, the degree of mixing of \\Nifs, the mass of electron capture elements, the progenitor metallicity, and the abundance of intermediate mass elements (especially calcium).  All effects are confirmed with illustrative theoretical models.  The NIR secondary maximum is shown to be a valuable diagnostic of these important physical parameters. The models further demonstrate that SNe~Ia are indeed excellent standard candles in the NIR, even when the physical properties of the ejecta are varied widely (\\S\\ref{Sec:Dispersion}). ",
        "conclusions": "We have modeled the far-red and NIR lightcurves of SNe~Ia and given a detailed explanation of the characteristic secondary maximum.  Our synthetic model lightcurves were calculated using parameterized 1-D ejecta configurations and the time-dependent multi-group radiative transfer code \\Code.  The model lightcurves displayed distinct and conspicuous secondary maxima and provided favorable fits to the NIR observations of the normal Type~Ia SN~2001el.  By varying the mass of \\Nifs, the models also reproduced the observed trend that brighter SNe~Ia have later and more prominent secondary maxima.  We trace the origin of the secondary maximum directly to the ionization evolution of iron group elements in the ejecta. Specifically, the NIR emissivity of iron/cobalt gas peaks sharply at a temperature $\\Tto \\approx 7000$~K, marking the transition between the singly and doubly ionized states.  The recombination of iron-rich gas from \\rectwoone\\ thus leads to enhanced redistribution of radiation from blue to NIR.  Interestingly, as the supernova cools, the global ionization evolution takes the form a \\rectwoone\\ ``recombination wave'' gradually receding deeper into the ejecta.  The onset and propagation of this wave through the iron rich layers marks the rise and fall the secondary maximum.  Because the ejecta are transparent at longer wavelengths during these epochs, the NIR observations allow us to watch directly as the recombination wave scans progressively through deeper and deeper layers of ejecta.  While the models considered in this paper captured the essential aspects of SN~Ia NIR lightcurves, they also highlighted several outstanding issues for the radiative transfer calculations.  First, the use of a complete and accurate atomic linelist proved critical in modeling the NIR bands.  Further improvement of the currently available atomic data is likely needed to accurately model the H and K-band lightcurves.  Second, non-thermal ionization effects from radioactive gamma-rays become significant for $\\texp \\ga 70$~days when LTE predicts neutrality, and thus likely have a dramatic impact on the NIR lightcurves at these late epochs.  Third, a proper treatment of Ca~II IR triplet line source function is crucial in synthesizing accurate I-band model lightcurves, as the assumption of purely absorbing lines leads to unrealistic results.  Future advances to the \\Code\\ code will permit more detailed studies of these and other important effects.  However, one does not expect the technical developments to change the general NIR lightcurve trends and dependencies explained in this paper.  In this paper, we studied the dependence of the NIR lightcurves on a number of important physical parameters, which highlighted the many ways in which NIR lightcurves offer valuable diagnostics of the ejecta properties and powerful constraints on explosion models.  First, the double-peaked morphology of the NIR lightcurves can be taken as a direct consequence of the abundance stratification in SNe~Ia, in particular the concentration of iron group elements in the central regions.  This confirms previous inferences based upon post-maximum and nebular spectra \\citep[e.g.,][]{Branch_w7, Hoeflich_99by, Kozma_neb}.  Abundance stratification is generic to certain classes of explosion models, for example the delayed detonation models \\citep{Khokhlov_DD} and the detonation from failed detonation \\citep{Plewa_GCD}.  In contrast, models characterized by large-scale mixing, such as published 3-dimensional deflagration models \\citep{Gamezo_3D,Reinecke_3D} are likely inconsistent with the double peaked behavior in the NIR lightcurves.  Further NIR observations, coupled with the transfer models, should be useful in constraining the exact degree of \\Nifs\\ mixing in SNe~Ia, and thus in testing current and future explosion paradigms.  Second, the luminosity of the secondary maximum provides a measure of the amount of iron group elements (both stable and radioactive) synthesized in the explosion.  The observed correlation between B-band decline rate and the luminosity of the secondary maximum provides strong evidence that slower declining SNe~Ia have (on average) a larger production of iron group elements.  This conclusion coincides with several other inferences to the same \\citep{Contardo_bol, Mazzali_nebular, Stritzinger}.  Third, the NIR secondary maximum is sensitive to the amount of stable iron group elements produced in the explosion, and hence the progenitor metallicity.  \\cite{Timmes_metallicity} has suggested that metallicity variations may lead to 25\\% variations in \\Nifs.  We find that the NIR signature of this variation is distinct from that of varying \\Nifs\\ independently.  Our predicted correlation arising from metallicity variations is that earlier secondary maxima will be as bright or brighter than later ones.  This conflicts with the primary observed trend, and suggests that the metallicity is likely a sub-dominant cause of SN~Ia luminosity variations.  This signature further opens the possibility of using an observational sample of NIR lightcurves to constrain metallicity variations among different progenitor populations.   Fourth, the timing of the secondary maximum is a direct probe of the temperature and ionization evolution in the ejecta, and hence a useful tool in interpreting observations.  As it turns out, the ionization evolution of crucial importance to the optical lightcurve decline rate as well. The \\rectwoone\\ recombination of iron group elements contributes significantly to blanketing in the B and V-bands due to the development of blends of strong Fe~II and Co~II lines in the optical part of the spectrum. In this context, NIR observations may also be especially useful in diagnosing peculiar objects.  For example, the paradoxical SN~2002cx was $\\sim 2$ mag subluminous, but had a peculiar spectrum characteristic of ``hot'' iron-rich ejecta \\citep{Li_02cx}. NIR observations would compliment that spectroscopic analyses be providing independent measures of the temperature scale, ionization evolution, iron core size and mixing.  Finally, we reiterate the cosmological utility of the NIR lightcurves of SNe~Ia.  In our lightcurve models, the dispersion in NIR peak magnitudes is found to be small ($\\la 0.2$~mag) even when the physical properties of the ejecta are varied widely.  At certain epochs (e.g., the J-band local minimum occurring near day 15 after B-maximum) the predicted dispersion is even smaller.  Our models thus further testify to the robustness of SNe~Ia NIR lightcurves as standard candles.  Perhaps more importantly, the NIR observations provide deep and readily interpreted diagnostics of the physical conditions in the SN~Ia ejecta, and therefore offer an opportunity for controlling potential systematic errors in the cosmology studies.  We have shown that the physical parameters affecting the NIR secondary maximum are also those that may modulate the optical peak brightness and decline rate of SNe~Ia.  Thus the NIR lightcurves are one of the most promising places to search for empirical secondary parameters related to deviations from the standard width-luminosity relation.  In this context, the NIR observations strike a nice balance in their level of information content.  Given the double-peaked morphology, the NIR lightcurves carry considerably more information than do the single-peaked optical band observations; at the same time, they remain simple enough to submit to a much more ready statistical analysis than do the complex spectral time series.  For all these reasons, a sizable sample of observed NIR lightcurves would be a remarkably powerful tool in constraining the possible evolution or progenitor drift of SNe~Ia arising in differing stellar populations.  This potential should be exploited further with expanded observational programs and concerted theoretical efforts."
    },
    "0606/astro-ph0606213_arXiv.txt": {
        "abstract": "{\\it HST}/NICMOS PSF-subtracted coronagraphic observations of HD~181327 have revealed the presence of a ring-like disk of circumstellar debris seen in 1.1 $\\mu$m light scattered by the disk grains, surrounded by a diffuse outer region of lower surface brightness. The annular disk appears to be inclined by 31\\fdg7 $\\pm$ 1\\fdg6 from face on with the disk major axis PA at 107$\\degr$ $\\pm$ 2$\\degr$. The total 1.1 $\\mu$m flux density of the light scattered by the disk (at  1\\farcs2 $<$ {\\it r} $<$ 5\\farcs0) of  9.6 mJy  $\\pm$ 0.8 mJy  is  0.17\\% $\\pm$ 0.015\\% of the starlight. Seventy percent of the light from the scattering grains appears to be confined in a 36 AU wide annulus centered on the peak of the radial surface brightness (SB) profile 86.3 AU $\\pm$ 3.9 AU from the star, well beyond the characteristic radius of thermal emission estimated from {\\it IRAS} and {\\it Spitzer} flux densities assuming blackbody grains ($\\approx$ 22 AU). The 1.1 $\\mu$m  light scattered by the ring: (a) appears bilaterally symmetric,  (b) exhibits directionally preferential scattering well represented by a Henyey-Greenstein scattering phase function with g$_{HG}$ = 0.30 $\\pm$ 0.03, and (c) has a median SB (over all azimuth angles) at the 86.3 AU radius of peak SB of 1.00  $\\pm$ 0.07 mJy arcsec$^{-2}$.  No photocentric offset is seen in the ring relative to the position of the central star. A low surface brightness diffuse halo is seen in the NICMOS image to a distance of $\\sim$ 4$\\arcsec$. Deeper 0.6  $\\mu$m {\\it HST}/ACS PSF-subtracted coronagraphic observations reveal a faint  (V $\\approx$ 21.5 mag arcsec$^{-2}$) outer nebulosity from 4$\\arcsec < r < 9\\arcsec$, asymmetrically brighter to the North of the star.  We discuss models of the disk and properties of its grains, from which we infer a maximum vertical scale height of 4--8 AU at the 87.6 AU radius of maximum surface density, and a total maximum dust mass of collisionally replenished grains with minimum grain sizes of $\\approx 1~\\mu$m of $\\approx$ 4 M$_{Moon}$. ",
        "introduction": "Fifteen years elapsed between the acquisition of the prototypical scattered-light image of a debris disk around a young main-sequence star ($\\beta$ Pictoris; Smith \\& Terrille 1984), offering strong support to the conjecture of the ``Vega phenomenon'' (far-IR emission in excess of a stellar photosphere from orbiting thermally emissive dust \\citep{BAC93} as suggested from earlier {\\it IRAS} observations \\citep{aumann84}), and the second example of a dusty circumstellar disk -- in the form of a ring -- seen by starlight illuminating its constituent grains (HR 4796A; Schneider et al. 1999).  Advances in space- and ground- based high contrast imaging capabilities have given rise to increasing, though still very small, numbers of scattered-light images of circumstellar debris systems. The very recent discoveries of a high surface brightness $\\beta$ Pictoris-like nearly edge-on disk around the A-star HD 32297 (Schneider, Silverstone \\& Hines 2005 [=SSH05]), and the very low surface brightness 140 AU ring around an A-star much older than $\\beta$ Pictoris, i.e., Fomalhaut \\citep{KAL05} speak  to the diversity in disk morphologies and scattering properties. While ring-like structures are obscured in scattered-light images of nearly edge-on disks, their intrinsic morphologies may be inferred by inversion of their radial surface brightness profiles (egs.: \\cite{buit}, \\cite{pantin96}, \\cite {aug01} for $\\beta$ Pictoris; and models presented by \\cite{krist05} and \\cite{metch05} for AU Microscopii). We now report the discovery of a ring-like disk of circumstellar debris seen about the star HD~181327. This new addition to the growing menagerie of circumstellar debris disks is, in many respects, similar to the HR~4796A star:disk system.  The  recent imagery of a broader ring about the older solar analog, HD~107146 \\citep{ardila04}, suggests that such structures are a common (if not prevalent) outcome of the planet formation process.  HD 181327 (HIP 95270) is an F5/F6V star ({\\it T}$_{eff}$ $\\approx$ 6450 K; \\cite{nord04}, log($g$) = 4.5) at 50.6 pc $\\pm$ 2.1 pc \\citep{peri97} with a Vega-like far-IR excess noted by Backman \\& Paresce (1993).  HD 181327  was identified as a main sequence debris-disk candidate by Manning \\& Barlow (1998)  with an {\\it L}$_{ir}$/{\\it L}$_{*}$ = 0.2\\%.  Kalas et al (2002) favored the Vega phenomenon interpretation of its far-IR excess based upon its youth and its gas depleted environment. The star's youth was first suggested from its potential  membership in the $\\sim$ 40 Myr Tucanae association \\citep{ZUK00}, but later identified as a member of the younger ($\\sim$ 12 Myr) $\\beta$~Pictoris moving group (Zuckerman \\& Song 2004; Mamajek et al. 2004).  The conjecture of a stellar age of $\\sim$ 10$^{7}$ yr is further supported by two independent observations. First, HD~181327 appears to exhibit common proper motion  with the 7$\\arcmin$ (0.1 pc at 50 pc) distant HR 7329, also identified as a member of the $\\beta$ Pictoris moving group by \\cite{ZUK01}. Both stars have statistically identical Hipparcos parallaxes and (long baseline) Tycho 2 proper motions. HR 7329 is a ZAMS A-star possessing a close substellar companion\\footnote{The total mass of HD181327 + HR7329A/B is $\\sim$ 3.8 M$_{Sun}$, and stellar binaries with separations $\\sim$ 0.1 pc are not unheard of with similar systemic masses.  Hence, in light of recent discussion on brown dwarf formation and binary populations (egs., see  \\cite{close03} fig 15 and \\cite{burgasser05}) one might conjecturally consider HD181327 + HR7329A/B to be a triple system.} with an isochronal age of $\\sim$ 10--30 Myr \\citep{Lowrance2000}. Second, HD 181327's Hipparcos derived B-V color index of +0.480 is completely consistent with an unreddeded F5--F6 star\\footnote{In the Hipparcos catalog, the F5 and F6 stars within 75 pc with parallaxes of S/N $\\geq$ 8 have median B-V color indices of +0.464 and+0.496, respectively. Hence, there appears to be no evidence for reddening for HD 181327 (unsurprising given its distance and environment).}. This, then, is in agreement with HD181327's e(b-y) = 0.002 mag \\citep{nord04}, and the stellar luminosity  is then 3.1 $\\pm$ 0.3 L$_{sun}$ from which the stellar age estimated using evolutionary tracks by \\citet{dm97}, \\citet{ps01}, and \\citet{ba98} is consistent with $\\sim$ 20 Myr.  Assuming blackbody grains, \\cite{ZUK04} suggested a thermal equilibrium temperature of 65 K for the IR emissive dust at a characteristic orbital radius 35 AU from HD 181327. \\citet{NEU04}, using  {\\it NNT} + SHARP-I \\& SofI, and \\citet{CHA04}, with {\\it ESO} 3.6 m AO + ADONIS/SHARP-II, conducted searches for a companion candidate to HD 181327, but found none. ",
        "conclusions": "The inferred spatial distribution of dust detected about stars, from thermal infrared excesses (as measured from {\\it IRAS} and {\\it Spitzer}), often predicts central clearings or ``holes'' to explain the observed spectral energy distributions.  Scattered-light imagery of circumstellar debris directly provides  constraints on the spatial distribution of dust independent of the dust grain temperature in regions sufficiently far from the central star which can be instrumentally probed with high spatial resolution.  The central regions of dusty circumstellar disks that are nearly edge on to the line-of-sight (e.g., $\\beta$ Pictoris, AU Microscopii and HD~32297) remain hidden at optical and near-IR wavelengths due to the opacity of the dust.  The inner regions of debris disk systems which are more inclined to the line-of-sight, such as HD~181327, are revealed directly and give  support to the interpretation of inner holes from analyses of debris disk SEDs. While only a very small number of light-scattered images of circumstellar debris disks exist, the growing number allows inter-comparison of the disk characteristics with stellar type and age (when sufficiently constrained; see Table 3 and Figure 15).  With the notable exception of $\\alpha$ PsA, which thus far has only been observed in scattered light with {\\it HST}/ACS, all of the debris disks listed in Table 3 have been observed with scattered-light imagery in the optical and the near-infrared. Very recently, two newly discovered, scattered-light disks around older (0.3 -- 1 Gyr) stars, HD~53143 and HD~139664 \\citep{KAL06}, have been reported, adding importantly to this still very limited sample.  Of the debris disks imaged to date, the HR~4796A ring-like disk is, in many respects, morphologically most similar to HD~181327's. Both scattered-light disks appear as relatively narrow annuli, have similar near-IR scattering fractions and thermal IR excesses, and physically are of similar size (70 and 86 AU in radii, respectively).  Both stars, while not coeval (HR~4796A being a member of the TW Hya association, and HD~181327 in the $\\beta$ Pictoris moving group)  are of similar ages ($\\sim$ 10 Myr), and spectral types.  Directionally preferential scattering, symmetrical about the disk minor axis (i.e., with respect to the line-of-sight), is seen in the HD~181327 disk. Such non-isotropic scattering might be expected from circumstellar particles evolved from a primordial ISM-like population (comparable in size to, or larger than, the  1.1 $\\mu$m wavelength of the observation), and similar behavior is seen in the HR~4796A debris ring, and in $\\alpha$ PsA's and HD~107146's as well.  These four ring-like debris disks vary by factors of $\\sim$ 2 in radius (from the already noted smaller HR~4796A and HD~181327 to the larger HD~107146 (133 AU; \\cite{ardila04}) and $\\alpha$ PsA (145 AU; \\cite{KAL05}), and $\\sim$ 4 in ring width-to-radius (w:r) ratio (from $\\sim$ 0.17, for both HR~4796A \\citep{schgil} and  $\\alpha$ PsA, to 0.64 for the very broad HD~107146 ring, and HD~181327 in between at 0.38). The disks about HR~4796A and  $\\alpha$ PsA, both early A stars, are most disparate in radius but most similar in w:r ratios, suggesting a lack of simple scalability in ring geometries conflated (if not determined) instead by  dynamical and kinematic processes in play. For example, in the case of HR~4796A, the outer radius of its disk may be truncated by the dynamical influence of its M-star companion (e.g., \\cite{arty94}). The sharp inner and more shallow outer edges of the illumination-corrected scattered-light radial profiles of HD~181327 (Fig. 7) and  $\\alpha$ PsA (figure 3 of \\cite{KAL05}) are remarkably similar, suggestive of similar clearing and/or confinement mechanisms for the disk grains.  The photocenter of the $\\alpha$ PsA's dust ring is asymmetrically offset from the position of the central star. As suggested by \\cite{KAL05}, such an asymmetry may implicate the existence of one (or more) planetary perturbers. The HR~4796A debris ring exhibits an $\\sim$~20\\% SB asymmetry at the ring ansae at near-IR and optical wavelengths \\citep{sch04} (also seen in it's mid-IR thermal emission as reported by \\cite{telesco}), with no photocentric offset in the position of the star relative to the ring ansae detectable in the scattered-light observations. Such a non-axisymmetric SB asymmetry about the ring minor axis could also arise from a local enhancement (or depletion) in the dust density distribution due to perturbations from unseen, co-orbital, planetary-mass bodies (e.g., \\cite {ozer}).  Though, alternative mechanisms, dependent upon the amount of residual gas in the system, could account for the ring structure itself (e.g, \\cite{klar}; \\cite{tag}) without invoking planetary-mass companions.  The HD~181327 ring morphology is bi-axially symmetric, centered on the star, and devoid of any readily apparent azimuthal  SB anisotropies aside from those dependent upon scattering phase angle.  Such a ring structure could arise from gas drag in the absence of planets, though the gas-to-dust ratio in the HD~181327 disk is, currently, unknown."
    },
    "0606/astro-ph0606739_arXiv.txt": {
        "abstract": "Black hole masses predicted from the $M_\\bullet-\\sigma$ relationship conflict with those predicted from the $M_\\bullet-L$ relationship for the most luminous galaxies, such as brightest cluster galaxies (BCGs). This is because stellar velocity dispersion, $\\sigma,$ increases only weakly with luminosity for BCGs and other giant ellipticals. The $M_\\bullet-L$ relationship predicts that the most luminous BCGs may harbor black holes with $M_\\bullet$ approaching $10^{10}M_\\odot,$ while the $M_\\bullet-\\sigma$ relationship always predicts $M_\\bullet<3\\times10^9M_\\odot.$ Lacking direct determination of $M_\\bullet$ in a sample of the most luminous galaxies, we advance arguments that the $M_\\bullet-L$ relationship is a plausible or even preferred description for BCGs and other galaxies of similar luminosity. Under the hypothesis that cores in central stellar density are formed by binary black holes, the inner-core cusp radius, $r_\\gamma,$ may be an independent witness of $M_\\bullet.$ Using central structural parameters derived from a large sample of early-type galaxies observed by {\\it HST}, we argue that $L$ is superior to $\\sigma$ as an indicator of $r_\\gamma$ in luminous galaxies.  Further, the observed  $r_\\gamma-M_\\bullet$ relationship for 11 core galaxies with measured $M_\\bullet$ appears to be consistent with the $M_\\bullet-L$ relationship for BCGs. BCGs have large cores appropriate for their large luminosities that may be difficult to generate with the more modest black hole masses inferred from the $M_\\bullet-\\sigma$ relationship. $M_\\bullet\\sim L$ may be expected to hold for BCGs, if they were formed in dissipationless mergers, which should preserve ratio of black hole to stellar mass. This picture appears to be consistent with the slow increase in $\\sigma$ with $L$ and the more rapid increase in effective radii, $R_e,$ with $L$ seen in BCGs as compared to less luminous galaxies. If BCGs have large BHs commensurate with their high luminosities, then the local black hole mass function for $M_\\bullet>3\\times10^9M_\\odot$ may be nearly an order of magnitude richer than what would be inferred from the $M_\\bullet-\\sigma$ relationship. The volume density of the most luminous QSOs at earlier epochs may favor the predictions from the $M_\\bullet-L$ relationship. ",
        "introduction": "Nearly every elliptical galaxy and spiral bulge has a black hole at its center \\citep{mag}. The masses of the black holes, $M_\\bullet,$ are related to the $V$-band luminosity, $L,$ and average stellar velocity dispersion, $\\sigma,$ of their host galaxies \\citep{d89, k93, kr, mag, fm, g00, tr02, hr}. The $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relationships are powerful tools as they allow the prediction of black hole masses --- which are difficult to measure directly --- from readily available galaxy parameters.  The black hole population in the most massive galaxies has yet to be assayed, however, which means that estimates of $M_\\bullet$ in these objects are based on extrapolations of relationships defined by smaller galaxies. The current record for largest black hole mass measured directly is $M_\\bullet\\sim3\\times10^9M_\\odot$ in M87 \\citep{m87}, yet M87 is only the {\\it second}-ranked galaxy in a cluster of modest richness. Brightest cluster galaxies (BCGs) in nearby Abell clusters are typically $\\sim3\\times$ more luminous \\citep{pl} and may host proportionately more massive BHs.  Testing this hypothesis through measurements of stellar dynamics requires both high sensitivity and high spatial-resolution, given the low central surface brightnesses and relatively large distances of BCGs. Such observations were not possible with the {\\it Hubble Space Telescope} ({\\it HST}) even before the failure of the {\\it Space Telescope Imaging Spectrograph;} they are only now becoming feasible with the advent of adaptive optics spectroscopy on 10m class telescopes.  A number of arguments suggest that black holes with $M_\\bullet>3\\times10^9M_\\odot$ do exist, even if this conclusion is not universal (e.g. \\citealt{mclure}). \\citet{net} argues that some QSOs have $M_\\bullet>10^{10}M_\\odot$ based on an empirical relationship between $M_\\bullet,$ broad-line width and nuclear luminosity for AGN.  \\citet{bechtold} and \\citet{vester1} also argue that some QSOs have black holes approaching this mass. Of particular relevance for BCGs is the hypothesis that cluster cooling flows are inhibited by AGN heating from the central galaxy \\citep{bt, chur}.  Recent {\\it Chandra} observations support a picture in which episodic AGN outbursts in BCGs heat the intra-cluster medium \\citep{vd}; the energetics required to terminate cooling flows imply $M_\\bullet>10^{10}M_\\odot$ for many clusters \\citep{fab}.  Arguments for such massive black holes appear to be in conflict, however, with the expectations from the $M_\\bullet-\\sigma$ relationship applied to the local galaxy velocity-dispersion distribution function. \\citet{tr02} find \\begin{equation} \\log (M_\\bullet/M_\\odot)=(4.02\\pm0.32)\\log(\\sigma/200 {\\rm\\ km\\ s^{-1}})+8.19\\pm0.06, \\label{eqn:msig} \\end{equation} for $H_0=70~{\\rm km~s^{-1}Mpc^{-1}}$ (which we will use throughout this paper). The \\citet{sheth} local velocity dispersion function shows a strong cut-off at $\\sigma\\approx400{\\rm\\ km~s^{-1}},$ which implies that galaxies harboring black holes with $M_\\bullet>3\\times10^9M_\\odot$ would be extremely rare.  \\citet{bern3} have identified a handful of galaxies with $\\sigma>400{\\rm\\ km~s^{-1}},$ but their results do not alter this conclusion.  Extrapolation of the $M_\\bullet-\\sigma$ relationship to galaxies more massive than M87 assumes that $\\sigma$ (and not galaxy mass) is the fundamental parameter for determining $M_\\bullet.$ The uncertainty in such an extrapolation is underscored by \\citet{wyithe}, who argues that the $M_\\bullet-\\sigma$ relationship is curved rather than linear in log-log space, in the sense that, at the high-$\\sigma$ end, the ``log-quadratic'' relationship predicts higher $M_\\bullet$ than does equation (\\ref{eqn:msig}). The Wyithe $M_\\bullet-\\sigma$ relationship, implies that the space density of black holes with $M_\\bullet>5\\times10^9M_\\odot$ may be substantially higher than that implied by equation (\\ref{eqn:msig}) (although the exact difference is highly sensitive to both the details of the velocity dispersion distribution function, and the assumed level of cosmic scatter in the $M_\\bullet-\\sigma$ relationship).  In this paper we point out that the $M_\\bullet-L$ relationship applied to the most luminous galaxies predicts $M_\\bullet$ values that are significantly larger than those predicted by either the \\citet{tr02} or \\citet{wyithe} $M_\\bullet-\\sigma$ relationships. This difference arises because BCGs do not follow the \\citet{fj} relationship between $L$ and $\\sigma.$ The relationship between $L$ and $\\sigma$ ``plateaus'' at large $L$ in the sense that BCGs have relatively low $\\sigma$ for their high $L$ (\\citealt{ho}; see also \\citealt{bmq}, who have seen this effect in simulations.)  Resolution of which of the $M_\\bullet-L$ or $M_\\bullet-\\sigma$ relationships is most representative of the black hole population in the most massive galaxies will only be possible when black hole masses can be measured in such galaxies.  In advance of such work, however, we can advance a number of arguments that suggest that the $M_\\bullet-L$ is a plausible and perhaps even preferred description for such systems.  The first set of arguments are based on the central structure of BCGs and other luminous elliptical galaxies that have cores in their central brightness profiles \\citep{l95, laine}. A core is evident as a radius at which the steep envelope of the galaxy ``breaks'' and transitions to an inner cusp with a shallow slope in logarithmic coordinates. The favored theory for core formation posits that cores are formed when stars are ejected from the galaxy's center by the decay of a binary BH created in a merger \\citep{bbr, ebi, f97, qh, mnm}. The size of the core then reflects the total mass ejected, which should be a function of $M_\\bullet.$  The size of the core may thus be an independent witness of $M_\\bullet.$  In BCGs and other galaxies of similar luminosity, galaxy luminosity is more closely related to the physical scale of the cores than $\\sigma,$ and the observed core size $M_\\bullet$ relationship for galaxies with cores and directly measured black hole masses appears to be consistent with the $M_\\bullet-L$ relationship.  A second set of arguments come from considering the formation of BCGs.  If BCGs are formed in ``dry'' mergers, then the ratio of black hole to stellar mass should be preserved over mergers, leading to the observed $M_\\bullet-L$ relationship.  In contrast, $\\sigma$ may change little over such mergers, and no longer track black hole mass as well it does for the less luminous galaxies from which the $M_\\bullet-\\sigma$ has been determined.  Lastly, we consider the relative predictions of the $M_\\bullet-L$ and $M_\\bullet-\\sigma$ relationships for the volume mass distribution function of black holes and which we compare to the predictions from QSO luminosity functions.  A decisive discrimination between the two relationships is not possible without a better understanding of the cosmic scatter in both relationships, but the \\citet{tr02} version of the $M_\\bullet-L$ relationship probably predicts too few high mass black holes to support the QSO luminosity function. ",
        "conclusions": "The $M_\\bullet-\\sigma$ relationship has come to be the ``gold standard\" for predicting black hole masses from galaxy properties due to its small scatter and its implications for illuminating the co-formation of galaxies and their nuclear black holes.  At first sight, the $M_\\bullet-L$ relationship might be dismissed as a simple consequence of the Faber-Jackson relationship.  With $M_\\bullet\\sim\\sigma^4$ and $L\\sim\\sigma^4,$ one would expect something like $M_\\bullet\\sim L.$ The larger scatter in the $M_\\bullet-L$ relationship further suggests that the $M_\\bullet-\\sigma$ relationship is really more fundamental. But as galaxy luminosity increases, $\\sigma$ levels off and the basic Faber-Jackson relationship does not appear to hold. At BCG luminosities there are no direct measurements of $M_\\bullet$ and $M_\\bullet(\\sigma)$ versus $M_\\bullet(L)$ present contradictory extrapolations. The contradiction essentially begs the question, do these exceptionally luminous galaxies have exceptionally massive black holes? The $M_\\bullet-L$ relationship answers this in the affirmative, while for the $M_\\bullet-\\sigma$ relationship to be correct we must accept the puzzling result that the black holes in BCGs have relatively modest masses. But this question leads to a broader issue, namely. is is there a significant population of black holes with $M_\\bullet$ approaching $10^{10}M_\\odot$ in the local universe?  The best way to answer these questions is to attempt to weigh the BHs in BCGs.  With the advent of LASER-guided adaptive optics-fed spectrographs on 10m class telescopes, it is now possible to do this. This paper may therefore be premature.  However, given the high attention to the $M_\\bullet-\\sigma$ relationship and its implications for galaxy formation, we believe that advancing the implications of $M_\\bullet-L$ relationship offers an important alternative view that should not be overlooked. Lacking hard measures of $M_\\bullet$ in the most massive galaxies, we have marshalled a number of less-direct arguments that this hypothesis may be favored for the most luminous galaxies.  The first set of arguments is based on the hypothesis that cores in the most luminous galaxies are created in a ``core scouring'' process in which a binary BH created in the merger of two galaxies evacuates stars from the center of the newly-merged product.  There presently is little observational support for the creation of binary black holes in mergers, but abundant theoretical work shows that realistic cores can be created by binary black holes, and the prevalence of nuclear black holes in galaxies overall strongly argues that such binaries must be created as a natural consequence of mergers. If so, the physical scale of cores, which we have parameterized as $r_\\gamma$ may be an independent witness of $M_\\bullet,$ and thus use the large cores in BCGs as an indicator of their black hole masses.  Based on central structural parameters derived from {\\it HST} observations, we find that the large cores in BCGs are commensurate with their high luminosities, while $\\sigma$ is a poor predictor of $r_\\gamma$ for $r_\\gamma>300$ pc.  The scatter in the $r_\\gamma-L$ relationship is much smaller than that in the $r_\\gamma-\\sigma$ relationship, again implying that $L$ and core scale are more closely related. The observed $r_\\gamma-M_\\bullet$ relationship for 11 core galaxies with directly determined black hole masses has large scatter, but appears to be more consistent with the $M_\\bullet-L$ rather than the $M_\\bullet-\\sigma$ relationship.  Lastly. the core masses in BCGs are over an order of magnitude larger than the black hole masses estimated from $M_\\bullet(\\sigma),$ but are only a few times larger than those estimated from $M_\\bullet(L);$ making such large cores with the smaller $\\sigma$-based black hole masses would be a strong challenge for the core-scouring hypothesis of core formation.  The second set of arguments comes from considering theoretical arguments concerning whether or not $L$ rather than $\\sigma$ should predict $M_\\bullet$ in BCGs.  The favored origin of BCGs is that they are the remnants of dissipationless purely-stellar mergers of less-luminous elliptical galaxies, augmented by ongoing galactic cannibalism of elliptical galaxies in the rich environments at the center of galaxy clusters. The plateau in the $L-\\sigma$ relationship plus but the steeper $L-R_e$ relationship at high galaxy luminosity presented here strongly favor this formation scenario. The luminosity of a BCG is the sum of the luminosities of its progenitors. Similarly, setting aside the possible ejection of nuclear black holes in the final stages of a merger, the final nuclear black hole mass should be the sum of the progenitor black holes.  Stated more directly, the ratio $M_\\bullet/L$ should be largely invariant over dissipationless mergers, leading to $M_\\bullet\\sim L$ at the high end of the galaxy LF.  In contrast, $\\sigma$ appears to be nearly constant over these mergers.  In effect, even if a relationship between $M_\\bullet$ and $\\sigma$ were set up in the initial stages of galaxy formation, it might not survive in a dissipationless merging hierarchy.  A final argument comes from attempting to infer the $z=0$ space density of the remnant black holes associated with the most luminous QSOs seen at $z\\sim2.$  As noted in the Introduction, the properties of the broad-line regions in the most luminous QSOs argue that they are powered by black holes with $M_\\bullet\\sim10^{10}M_\\odot.$  The heating of the intra-cluster medium in the richest galaxy clusters may also demand that some black holes in BCGs approach this mass. The critical part of this analysis is understanding how to correct the QSO space density for QSO luminosity evolution.  The remnant black holes last forever, but the QSOs represent only those BHs made visible during an epoch of high mass-accretion, which presumably lasts only a small fraction of the age of the universe. We used the \\citet{hop06} simulations to estimate the QSO lifetimes. The resulting shape and space density of the bright end of the QSO LF falls between the higher space density of the most massive black holes implied by $M_\\bullet(L)$ and those implied by $M_\\bullet(\\sigma),$ while the simple ``lightbulb'' model of QSO duty cycles favors the $M_\\bullet(L)$ relation.  This treatment is sensitive to the assumed amount of cosmic scatter in both $M_\\bullet$ relationships; however, it appears difficult for the log-linear $M_\\bullet-\\sigma$ relationship to explain the the observed space densities of the most luminous QSOs without assuming that its cosmic scatter is larger than is likely to be the case."
    },
    "0606/astro-ph0606106_arXiv.txt": {
        "abstract": "Primordial stars are formed from a chemically pristine gas consisting of hydrogen and helium. They are believed to have been born at some early epoch in the history of the Universe and to have enriched the interstellar medium with synthesized heavy elements before the emergence of ordinary stellar populations.  We study the formation of the first generation of stars in the standard cold dark matter model. We follow the gravitational collapse and thermal evolution of primordial gas clouds within early cosmic structures using very high-resolution, cosmological hydrodynamic simulations. Our simulation achieves a dynamic range of $\\sim 10^{10}$ % in length scale. With accurate treatment of atomic and molecular physics, it allows us to study the chemo-thermal evolution of primordial gas clouds to densities up to $\\rho \\sim 2\\times 10^{-8} {\\rm g}\\;{\\rm cm}^{-3} (n_{\\rm H}\\sim 10^{16} {\\rm cm}^{-3})$ without assuming any {\\it a priori} equation of state; a six orders of magnitudes improvement over previous three-dimensional calculations. We implement an extensive chemistry network for hydrogen, helium and deuterium. All the relevant atomic and molecular cooling and heating processes, including cooling by collision-induced continuum emission, are implemented.  For calculating optically thick \\HH cooling at high densities, we use the Sobolev method (Sobolev 1960) and evaluate the molecular line opacities for a few hundred lines.  We validate the accuracy of the method by performing a spherical collapse test and comparing the results with those of accurate one-dimensional calculations that treat the line radiative transfer problem in a fully self-consistent manner.  We then perform a cosmological simulation adopting the standard $\\Lambda$CDM model.  Dense gas clumps are formed at the centers of low mass ($\\sim 10^{5-6} M_{\\odot}$) dark matter halos at redshifts $z\\sim 20$, and they collapse gravitationally when the cloud mass exceeds a few hundred solar masses.  To examine possible gas fragmentation owing to thermal instability, we compute explicitly the growth rate of isobaric perturbations. We show that the cloud core does not fragment in either the low-density ($n_{\\rm H} \\sim 10^{10} {\\rm cm}^{-3}$) or high-density $(\\sim 10^{15} {\\rm cm}^{-3}$) regimes, where gas cooling rate is increased owing to three-body molecule formation and collision-induced emission, respectively. The cloud core becomes marginally unstable against chemo-thermal instability in the low-density regime. However, since the core is already compact at that point and correspondingly the sound-crossing time as well as the free-fall time are short, or comparable to the perturbation growth timescale, it does not fragment. Run-away cooling simply leads to fast condensation of the core to form a single proto-stellar seed. We also show that the core remains stable against gravitational deformation and fragmentation throughout the evolution.  We trace in Lagrangian space the gas elements that end up at the center of the cloud, and study the evolution of the specific angular momentum. We show that, during the final dynamical collapse, small angular momentum material collapses faster than the rest of the gas and selectively sinks inwards.  Consequently, the central regions have little specific angular momentum, and rotation does not halt collapse. With the large physical dynamic range of our simulation, we, for the first time, obtain an accurate gas mass accretion rate within a 10$M_{\\odot}$ innermost region around the protostar. The protostar is accreting the surrounding hot ($T\\sim 2000$K) gas at a rate of $\\dot{M}> 10^{-2}-10^{-3} M_{\\odot}/{\\rm yr}$. From these findings we conclude that primordial stars formed in early cosmological halos are massive.  We carry out proto-stellar evolution calculations using the obtained accretion rate. The resulting mass of the first star when it reaches the zero-age main sequece is $M_{\\rm ZAMS}\\sim 100 M_{\\odot}$, and less ($\\ga 60 M_{\\odot}$) for substantially reduced accretion rates. ",
        "introduction": "The current standard theory of cosmic structure formation posits that the present-day clumpy appearance of the Universe developed through gravitational amplification of matter density fluctuations generated in its very early history. The energy content of the Universe and the basic statistics of the initial density field have been determined with great accuracy from recent observations of the cosmic microwave background (Spergel et al. 2006), large-scale structure (Tegmark et al. 2004; Cole et al. 2005), and distant supernovae (Riess et al. 2004; Astier et al. 2006). It has become possible to make accurate predictions for the formation and nonlinear growth of large-scale structure within this framework.  The standard model based on cold dark matter (CDM) predicts that the mass variance after the epoch of matter-radiation equality has progressively larger amplitudes on smaller mass scales, indicating that objects form hierarchically, with smaller ones originating first.  Although the fluctuation amplitudes on scales much smaller than galactic scales are not known directly from observations, most of the viable candidates for the dark matter predict nearly scale-invariant fluctuations down to stellar or even planetary masses.  Such CDM models generically argue that the first cosmological objects are formed in low mass ($\\sim 10^6 M_{\\odot}$) dark halos at high redshifts ($z \\ga 20$) when primordial gas condenses via molecular hydrogen cooling (Couchman \\& Rees 1986; Haiman, Thoul \\& Loeb 1996; Tegmark et al. 1997; Abel, Bryan \\& Norman 2002; Yoshida et al. 2003; see Bromm \\& Larson 2004 for a review). Hierarchical structure formation eventually leads to the emergence of a population of early generation stars/galaxies which terminate the Cosmic Dark Ages by emitting the first light (Madau 2000; Cen 2003; Miralda-Escud\\'e 2003; Sokasian et al. 2003, 2004; Reed et al. 2005).  The first stars are also believed to be the first source of heavy element production (those heavier than lithium).  Heavy elements must have been processed in massive stars and expelled into the interstellar medium by supernovae, to enable the formation of ordinary stellar populations.  Early metal-enrichment of the low density intergalactic medium is also suggested by observations of the Lyman-$\\alpha$ forest (Songaila 2001, 2005; Schaye et al. 2003). Interestingly, the recent discovery of extremely metal-poor stars (Christlieb et al. 2002; Frebel et al. 2005) suggests that relics from very early generation stars exist even today in our own Galaxy.  Stars with such low metallicities provide valuable information on the first stars that likely polluted the interstellar medium (Iwamoto et al. 2005; Tumlinson 2006).  The study of primordial star formation has a long history.  The formation of the first cosmological objects via gas condensation by molecular hydrogen cooling has been studied for many years (e.g. Saslaw \\& Zipoy 1967; Peebles \\& Dicke 1968; Matsuda, Sato, \\& Takeda 1969; Kashlinsky \\& Rees 1983). In the context of star formation, the evolution of a collapsing primordial gas cloud has also been studied extensively (Matsuda, Sato, \\& Takeda 1969; Yoneyama 1972; Carlberg 1981). Palla, Salpeter \\& Stahler (1983) used a one-zone model to follow gas collapse to very high densities. They identified a minimum Jeans mass scale of 0.1 solar mass from the thermal evolution of the cloud core, suggesting the formation of low-mass primordial stars. One-dimensional hydrodynamic simulations of spherical gas collapse were performed by a number of researchers (Villere \\& Bodenheimer 1987; Omukai \\& Nishi 1998; Ripamonti et al. 2002).  Omukai \\& Nishi (1998) included a detailed treatment of all the relevant chemistry and radiative processes and thus were able to provide accurate results on the thermal evolution of a collapsing primordial gas cloud up to stellar densities.  These authors found that, while the evolution of a spherical primordial gas cloud proceeds in a roughly self-similar manner, there are a number of differences in the thermal evolution from that of present-day, metal- and dust-enriched gas clouds.  Recently, three-dimensional hydrodynamic calculations were performed by several groups.  Abel et al. (2000; 2002) and Bromm et al. (1999; 2002) studied the formation and fragmentation of primordial gas clouds in CDM models.  From various analyses of the simulation results, these authors conclude that primordial stars (often called ``Population III'') that form in dark matter matter halos are rather massive. Yoshida et al. (2003) examined the statistical properties of primordial star-forming clouds from a large sample of early mini-halos located in their cosmological simulations.  The chemo-thermal evolution of the primordial gas clouds is rather complex, being coupled with the dynamical evolution of the assembly of dark matter halos (Yoshida et al. 2003; Gao et al. 2005), and simple analytic models based on crude assumptions almost always fail in predicting the evolution of primordial gas in CDM halos.  The three-dimensional calculations of Abel et al. (2002; hereafter, ABN02) were able to follow the gas evolution until a small ($\\sim 1M_{\\odot}$) fully molecular core formed.  However, since they did not include the effect of the opacity of the cloud core, the radiative cooling rate is over-estimated at the cloud center.  Thus their results for late stage evolution, including the important problem of fragmentation and mass accretion, are uncertain.  Bromm \\& Loeb (2004) made similar approximations and also continued their simulation by creating a sink particle to study the late time accretion phase. None of these simulations reproduce the correct density, temperature, and velocity profiles in the vicinity the protostar, which are among the most important quantities for proto-stellar evolution.  A critical technique we describe in the present paper is computation of molecular line opacities. With the implementation of optically thick line cooling, the gas evolution can be followed to much higher densities than in the previous studies. In principle, evaluating the net line cooling rate at such high densities involves costly line radiative transfer.  We circumvent this by using a local Sobolev method which employs local velocity information as well as knowledge of the density and temperature.  We show that the method works well in problems of collapsing gas clouds, in terms of computation of radiative cooling rates and resulting density and temperature structure.  We apply this technique to cosmological simulations.  An important, outstanding issue is to determine whether or not a gas cloud fragments into multiple objects during its collapse.  If it does fragment, it could yield multiple stars, suggesting the possibility of star-cluster formation including low-mass stars (Sabano \\& Yoshii 1977; Silk 1983).  To the contrary, if the gas cloud remains stable against fragmentation throughout its evolution, it will lead to the formation of a single proto-stellar seed, surrounded by a large amount of infalling gas - a typical condition for a proposed scenario of the formation of massive stars (Larson \\& Starrfield 1971; see the review by Larson 2003).  We show that the gas cloud formed at the center of a high-redshift halo remains stable against fragmentation, to form a single proto-stellar seed.  We also find that the rate of gas accretion is very high.  The mass accretion rate onto a protostar largely affects the final stellar mass (Omukai \\& Palla 2001; 2003). Stahler, Palla, \\& Salpeter (1986a,b) suggested that large accretion rates ($> 10^{-3} M_{\\odot}\\;{\\rm yr}^{-1}$) are expected for primordial protostars because of high gas temperatures in the infalling surrounding gas.  For a better evaluation of the mass accretion rate, Omukai \\& Nishi (1998) use a self-similar solution to derive a time-dependent rate, while Tan \\& McKee (2004) consider effects of rotation in accretion dynamics around proto-stellar disks. We address this issue using fully self-consistent cosmological simulations.  Finally, we carry out proto-stellar evolution calculations and determine the resulting stellar mass on the assumption that the true gas mass accretion rate and its time-dependence is close to the instantaneous accretion rate at the protostar formation epoch.  The rest of the paper is organized as follows.  In Section 2, we describe the chemistry network used for our simulations. The relevant cooling and heating processes are described in Section 3. There, we also provide details of numerical implementation. In Section 4, we present the results of a test problem of a spherically collapsing gas. We compare our results to previous work using one-dimensional codes. We present the results of our $\\Lambda$CDM cosmological simulations in Section 5.  Implications of the main conclusions are discussed in Section 6. ",
        "conclusions": "We have developed numerical techniques to implement the required chemical and radiative processes for following the thermal evolution of primordial gases to very high densities. Using a cosmological simulation with hydrodynamics and chemistry, we have studied the formation and evolution of pre-stellar gas clouds in a $\\Lambda$CDM universe. Our three-dimensional simulation for the first time resolves the inner $\\sim 1M_{\\odot}$ fully molecular core, allowing us to obtain correct density, temperature, and velocity structure around the primordial protostar. These quantities altogether yield an accurate gas mass accretion rate around the protostar that was uncertain in previous works. We have derived three important facts from our simulations: (1) the primordial gas cloud in a cosmological minihalo does not fragment, but yields a single proto-stellar seed, (2) the cloud core has a small angular momentum so that rotation does not halt the collapse, nor is a disk formed, and (3) the rate of accretion from the cloud envelope is large. From these facts, we conclude that the first stars are massive. By performing a proto-stellar evolution calculation, we derive that the protostar grows to $\\sim 100 M_{\\odot}$ for the particular case found in our simulation. We derive the stellar mass for the first time from a self-consistent calculation of proto-stellar evolution, coupled with a cosmological simulation.  While we simulated only a single halo in detail, the particular case shows a convincing and plausible case for the formation of massive stars in early mini-halos in the CDM model. The fate of such massive primordial stars (Population III), whether or not they are commonly born in various environments, are of considerable interest. They are likely responsible for various feedback effects in the early Universe (Haiman, Rees, \\& Loeb 1997; Bromm, Yoshida, \\& Hernquist 2003; Yoshida, Bromm, \\& Hernquist 2004; see Ciardi \\& Ferrara 2005 for a review). Stars with mass $140-260 M_{\\odot}$ are believed to trigger pair-instability supernovae and expel all the processed heavy elements in the explosion (Bond, Arnett, \\& Carr 1984; Heger \\& Woosley 2002). It has also been suggested that less massive stars die as hypernovae (Umeda \\& Nomoto 2002), producing a quite different metal yield from the pair-instability case. Interestingly, the observed abundance pattern of extremely metal poor stars appears to favor the latter scenario (Umeda \\& Nomoto 2003; Frebel et al. 2005). Energetic supernovae are also destructive, being effective in evacuating the halo gas (Bromm, Yoshida \\& Hernquist 2003; Wada \\& Venkatesan 2003; Kitayama \\& Yoshida 2005) and possibly quenching further star-formation in that region. A single supernova would suffice to pollute a large volume of the interstellar matter from which very metal-poor stars might have been formed. Metal-enrichment by the first supernovae may also lead to the formation of ordinary stellar populations including low-mass stars (e.g. Mackey et al. 2003) if metal-enrichment is relatively confined within a small region. Intriguingly, Jimenez \\& Haiman (2006) argue that metal-mixing could be inefficient and that primordial stars are formed in galaxies at $z<5$.  Massive stars emit a large number of ionizing photons and can ionize a large volume of the intergalactic medium around them (Kitayama et al. 2004; Whalen et al. 2004). Recently, Yoshida (2006) carried out radiation-hydrodynamic calculations of early \\HII regions and concluded that radiation from a massive Population III star can completely ionize the gas within the host halo, quenching further gas condensation and star-formation for a significant period of cosmic time. Both massive and very massive ($>300 M_{\\odot}$) stars likely leave a black hole remnant. It is plausible that the remnants become the seeds for supermassive black holes that power luminous quasars found at $z \\ga 6$ (Li et al. 2006, in preparation) and subsequently at lower redshifts (e.g. Di Matteo et al. 2005; Hopkins et al. 2005, 2006). Because bottom-up hierarchical structure formation is a generic prediction of the CDM model, it is natural to think that the early population stars affect and even set the initial conditions for primeval galaxy formation.  It is important to explore feedback from the first stars in more detail under a proper cosmological set-up.  The prospects for detecting the first generation of stars appear promising. Supernovae resulting from the collapse of massive stars can be observed if they trigger gamma-ray bursts (GRBs). There is growing evidence that long duration GRBs are associated with energetic supernovae (e.g. Hjorth et al. 2003). Gou et al. (2004) suggest that the afterglow of high-redshift GRBs can be detected in X-rays. There are also a broad range of observations through which we can, in principle, indirectly observe star-formation at high redshifts, through, for example, angular fluctuations in redshifted 21 cm emission from intergalactic gas as stars reionize the Universe (see, e.g. Zaldarriaga et al. 2004; Furlanetto et al. 2004; Zahn et al. 2006), or from secondary anisotropies imprinted from stellar radiation on the CMB (e.g. McQuinn et al. 2005; Zahn et al. 2005).  Finally, we remark that it is still too early to definitely conclude the exact mass or mass range of the first stars. The actual mass accretion could take place in a more complicated manner than we assume. Feedback from the protostar itself can affect the dynamics and thermal evolution of the accreting gas. Ultra-violet radiation from the protostar likely affects the gas infall rate. However, primordial gas has a much smaller opacity because of the absence of dust, and thus radiation pressure is generally not important (Omukai \\& Inutsuka 2002). The effect of radiation pressure can also be reduced if protostellar outflows provide optically thin cavities through which photons can escape (Krumholz, McKee \\& Klein 2005). Future calculations of proto-stellar evolution need to include these feedback effects onto the infalling gas. Also, accurate modeling of disk accretion may be necessary in general cases (Tan \\& McKee 2004). Although our simulation is currently resolution-limited (rather than physics-limited), with the maximum density being still five orders of magnitude smaller than stellar densities, we foresee that it will be possible to carry out an {\\it ab initio} calculation of the formation of primordial protostars in the near future.  \\bigskip  NY thanks Fumitaka Nakamura, George Field, Simon Glover, Gao Liang, and Brian O'Shea for insightful discussions. The simulations were performed at the Center for Parallel Astrophysical Computing at Harvard-Smithsonian Center for Astrophysics, at the Center for Computational Cosmology at Nagoya University, and at the Data-Reservoir at the University of Tokyo. We thank Mary Inaba and Kei Hiraki for providing the computing resource at U-Tokyo. The work is supported in part by the Grants-in-Aid for Young Scientists (A) 17684008 (N.Y.), and by Young Scientists (B) 18740117, Scientific Research on Priority Areas 18026008 (K.O.) by the Ministry of Education, Culture, Science and Technology of Japan. This work was supported in part by NSF grants AST 03-07433 and AST 03-07690, and NASA ATP grants NAG5-12140, NAG5-13292, and NAG5-13381."
    },
    "0606/astro-ph0606276_arXiv.txt": {
        "abstract": "We argue that the first stars may have spanned the conventional mass range rather than be identified with the Very Massive Objects ($\\sim 100-10^3 \\rm M_\\odot$) favoured by numerical simulations. Specifically, we find that magnetic field generation processes acting in the first protostellar systems suffice to produce fields that  exceed the threshold for MRI instability to operate and thereby allow the MRI dynamo to generate equipartition-amplitude magnetic fields on protostellar mass scales below $\\sim 50 \\rm M_\\odot$. Such fields allow primordial star formation to occur at essentially any metallicity by regulating angular momentum transfer, fragmentation, accretion and feedback in much the same way as occurs in conventional molecular clouds. ",
        "introduction": "The formation of the first stars should in principle be simpler to understand than present-day star formation.  The usual prediction that the mass of the first stars ${\\mathcal{O}}(100-1000)\\rm M_\\odot$ is attributable to dominance of H$_2$ cooling and consequent high temperature and accretion rate (Abel, Bryan and Norman 2002; Bromm, Coppi and Larson 2002).  It is appropriate to question this result, that only very massive stars formed, for at least four phenomenological reasons.  Firstly, one finds solar mass stars at $\\rm [Fe] \\simlt -4$.  These may be contaminated by companion stars that exploded as hypernovae, whose ejecta are depleted in Fe relative to N and C and this interpretation is consistent with the measured abundance ratios for a handful of extreme metal-poor stars. However not all metal poor halo stars display such anomalies, and a substantial fraction at the $10^{-4}$ solar level have abundance ratios that are consistent with conventional SNII precursors. There are at least two halo stars known at $\\rm [Fe] \\simlt -5$.  HE 0107-5240 is a giant with [Fe/H] =-5.3, but enhanced nitrogen, carbon and oxygen: [N/H]= -3.0, [C/H]= -1.3 and [O/H]=-2.9 (Christlieb et al. 2004).  HE 1327-2326 is a main sequence star (or subgiant) with an iron abundance about a factor of 2 lower than HE 0107-5240.  In this latter case, both nitrogen and carbon are enhanced relative to iron by about 4 dex, while there is only a comparable upper limit on oxygen (Frebel et al. 2005).  Appeal to preenrichment by a core collapse $\\sim 25 \\rm\\,M_\\odot$ supernova of Population III abundance with fallback fits the abundance patterns well (Umeda and Nomoto 2003), although oxygen may possibly be overproduced in the case of HE 0107-5240 (Bessell, Christlieb and Gustafsson 2004).  The example of HE 1327-2326 eliminates any internal enrichment source (e.g. convective dredge-up), nor is there evidence for a binary companion in the case of HE 0107-524, thereby also suggesting that mass transfer from an AGB star is an unlikely explanation. The high CNO abundances argue for a common explanation involving enrichment of the primordial cloud by Type II supernovae of primordial abundance. Indeed, the abundance patterns in extremely metal poor halo stars suggest that enrichment was produced by Population III stars in the mass range 20-130 $M_\\odot$ (Umeda and Nomoto 2005).   Secondly, the broad emission line regions of very high redshift quasars, reveal high elemental abundances that appear to have also been generated by conventional SNII precursors.  In particular, the nearly constant FeII/MgII emission line ratios over $0<z<5$ requires intense SNII activity at a redshift as high as 9 (Dietrich, Hamann, Appenzeller and Vestergaard 2003).  Thirdly, the chemical yield predictions from primordial very massive stars, when normalised to the inferred ionising photon output required to reionise the universe at high redshift, do not correspond to observed abundances in any primitive environments.  In particular, the pair-instability supernova nucleosynthetic signatures generated by stars with initial masses in the range 130 to 260 M$_\\odot$ are not seen either in the intergalactic medium, including both Lyman alpha forest and damped Lyman alpha absorption systems, or in extremely metal-poor halo stars (Daigne et al. 2005).   Finally, from the theoretical perspective with regard to cooling, the primordial metal abundance pattern has profound consequences for the thermal balance and chemical composition of the gas, and hence for the equation of state of the parental cloud.  Spaans and Silk (2005) find that the polytropic index is soft for low oxygen abundance enhancements, [O/H]$<-3$, as appropriate for Population III, but stiffens to a polytropic index $\\gamma$ larger than unity for [O/H]$>10^{-2}$ due to the large opacity in the CO and H$_2$O cooling lines.  Hence Pop III star formation is efficient, especially before hypernova enrichment and associated oxygen enhancement has occurred. There should be no obstacle to forming stars over a wide range of masses even in the absence of significant fine-structure cooling. On the contrary, once the polytropic index stiffens, at [O/H]$\\simgt -2$, the IMF should change and, specifically, flatten, as argued by Spaans and Silk (2000).  Star formation is subsequently less efficient, in part because of massive star feedback in addition to the stiffening of the equation of state, as is required in most discussions of galaxy formation in order to avoid premature consumption of the gas reservoir.  Of course, to do justice to the very massive primordial star hypothesis, the typical expectation of one very massive star per primordial cloud is unlikely to produce enough metals to raise the mean IGM metallicity to a value of $\\sim 10^{-4}-10^{-2.5}$ solar, as observed in the Lyman alpha forest (Schaye et al. 2003) and as also required to significantly modify the cooling and thereby enhance fragmentation.  This is especially plausible as the ejecta are mostly not retained in the shallow potential wells of the pregalactic clouds (Norman, O'Shea and Paschos 2004).  Rather, one has the impression that the first stars that produced metals in significant amounts, in particular as monitored in the abundance patterns of the Lyman alpha forest at modest overdensity and in halo stars at $\\rm [Fe] \\simlt -4,$ spanned the conventional stellar mass range.  A more fundamental issue may be that the elegant state of the art simulations with cosmological initial conditions have nothing directly to say about the primordial stellar initial mass function, nor for that matter about the physics underpinning the inferred transfer of angular momentum, but only about the mass function of gas clumps. The ultimate fate of these gas clumps has not yet been convincingly demonstrated.  Notice that recent simulations by Saigo, Matsumoto and Umemura (2004) show that Population III stars are prone to form in binary systems through the formation of a rotationally supported disk. For that to happen, the initial angular momentum of the parent cloud needs to be small, the centrifugal force being at most $10\\% - 30 \\%$ of the pressure force. However, the accretion efficiency seems not affected by the binary formation, and the binary Pop III stars are again predicted to be very massive.   Of course, all of this begs the question as to whether a softened EOS can lead to low or intermediate mass star formation in situations of very low metallicity, or at any rate, reduce the characteristic stellar mass to below the lower limit for pair instability SN. Certainly, the clouds from which stars with [Fe/H] $\\simlt -5$ formed seem to have had a soft equation of state.  The minimum fragment mass is lowered at [O/H]$\\simgt -3$ due to enhanced fine-structure cooling (Bromm and Loeb 2003), but it is likely that turbulent fragmentation (Klessen et al. 2004) enhanced by the softening of the equation of state is the key to low-mass fragmentation, for example via disk formation. According to this viewpoint, it is the equation of state in the initial collapsing cloud that controls the IMF.  An alternative view is that the final fragment mass remains high, as suggested by the numerical simulations, but that a second mode of fragmentation associated with the onset of dust cooling at very high density may lead to low mass star formation (Omukai et al. 2005).  The latter study finds that a metallicity of $Z\\simgt 10^{-5}Z_\\odot$ at a density in excess of $10^{10}\\rm cm^{-3}$ marks this transition. However Spaans and Silk (2005) argue that the gas properties, and in particular the equation of state, at much lower densities, typically the range associated with saturation of H$_2$ cooling at $\\sim 10^4\\rm cm^{-3}$, most likely determine the IMF.  The purpose of this paper is to demonstrate that a plausible alternative to dust cooling in primordial clouds involves recourse to a mechanism that is universally accepted as being an essential key to controlling the IMF in conventional molecular clouds at the onset of collapse, namely magnetically-regulated fragmentation and accretion. The essence of our argument will be to derive the magnetic field strength in primordial star formation.  An important conclusion is that one should continue to find low mass, and especially abundance signatures of intermediate mass, stars down to essentially zero metallicity. Note that the possible effects of magnetic fields on primordial star formation have also been investigated in a pioneering paper by Tan and Blackman (2004), and more recently by Machida et al. (2006). ",
        "conclusions": "Magnetic fields have probably played a role in Primordial Star Formation, even if starting from a medium free of magnetic fields. As we have argued, radiation drag or thermal pressure effects are able to generate magnetic seeds in the disk surrounding the central accreting stellar progenitor. Initially, those seeds are weak too much for being amplified by small scale turbulence, or even MRI effects on scales comparable to the disk scale. However, as matter accretion proceeds, the mass of the central object grows, and the gravitational potential it creates deepens, increasing the rotation velocity of the disk. Eventually, the rotation is fast enough so that long wavelength MRI modes become unstable, as the minimum magnetic field for MRI becomes smaller than the field generated in the disk. Depending on the properties of the disk, this happens once the mass of the central object reaches $4 - 28 M_\\odot$.  Subsequently, MRI dynamo will be at work and will amplify the magnetic field. The field can then rapidly reach dynamically important values, and magnetically-driven ejection contributes to lower the effective accretion efficiency. The actual model of magnetic winds is beyond the scope of this article.  but this suggests that feedback is likely to require magnetically-driven outflows, which could occur during Pop III formation already when, as we argued above, the mass of the stellar progenitor is rather small. Note that, in their model of protostellar disk in primordial star formation, Tan and Blackman (2004) estimated the power of magnetically driven outflows. In their study, the outflow feedback effects reduce the star formation efficiency once the protostar reaches roughly 100 $M_\\odot$. Incidentally, concurrent conclusions were reached by Machida et al.  (2006) who simulated the collapse of a magnetized primordial cloud in rigid rotation, and the subsequent formation of a Pop III star, within the ideal MHD approximation. Depending on the (high) initial value ($B_{\\rm init} \\gtrsim 10^{-9} G$ for an initial cloud density $n_{\\rm c}\\sim 10^3 {\\rm cm}^{-3}$) of the magnetic field, their simulations show that magnetically driven jets develop and effectively reduce the accretion rate. Further numerical studies with higher resolution seem necessary to address both the mass ejection rate and the amplification of magnetic fields in proto-stellar disks for initially weaker, even vanishing, magnetic seeds.  We have argued that the interplay between the two modes of star formation, primordial, massive and conventional, involving all masses, is controlled by $B$ and not by $Z$.  Effective feedback requires that of order 10 percent of the gas accretion rate is channelled into star formation, with an outflow rate that on the average must be of the order of the net star formation rate. Globally, for the star forming cloud, one expects that $\\dot M_{\\rm outflow} \\sim \\dot M_{\\ast}\\sim 0.1 \\dot M_{\\rm accretion} $, much as is found in nearby cases of star-forming clouds.  Moreover, feedback is likely to be responsible for the turbulent support in clouds that lowers the star-formation efficiency and helps to generate the conventional IMF.  Thus the onset of cloud fragmentation, due eventually but we have argued, not exclusively, to the enhanced role of cooling, would allow field amplification, angular momentum transfer, feedback and low mass star formation.  Alternatively, suppose we accept the hypothesis that the first stars were massive objects.  To avoid the empirical objections discussed above, one would have to argue that merging and coagulation of gas clumps resulted in formation of predominantly very massive objects, of characteristic mass $\\simgt 1000\\rm M_\\odot,$ whose fate is to form intermediate mass black holes with relatively low nucleosynthetic yields associated with their collapse.  In this case, the seed fields may come from jets and outflows driven by spin-up of turbulent accretion disk dynamos as a consequence of accretion onto these intermediate mass black holes.  One can argue that the first generation of primordial clouds, which cooled predominantly via Lyman alpha emission, preferentially formed IMBHs, since the associated minihalos of mass $\\sim 10^7\\rm M_\\odot$ are promising environments for forming IMBHs in view of the high core accretion rates (Zhao and Silk 2005). Accretion disks around the IMBHs provide promising sites for MRI dynamos.  Alternatively, the jets may generate magnetic fields via the Weibel instability, that although small, amounting to $\\sim m_e/m_p$ of the equipartition value, could still be useful as seeds (Wiersma and Achterberg 2004).  The IMBH outflows allow one to generate larger seed fields required for conventional star formation than inferred from the Biermann battery.  The associated nucleosynthetic implications are modest, because IMBHs have essentially zero yield. For any reasonable IMF, there are of course associated Population III stars.   Consider for example the inference from primordial star simulations that one star of mass $\\sim 10^2 \\rm M_\\odot$ forms per pristine pregalactic cloud of baryonic mass $\\sim 10^7 \\rm M_\\odot$, a typical galactic precursor not dissimilar in mass to dwarf satellite galaxies. The resultant local enrichment would amount to of order $10^{-5}$ the solar value, and would be diluted by a further factor that corresponded to the fractional number of primordial star-forming clouds that formed the eventual galaxy. However the associated population of IMBHs could amount to as much as $\\sim 10\\%$ of the current epoch stellar mass, if accretion onto these IMBHs also provides a possible explanation of the NIR background excess (Madau and Silk 2005).  If such IMBH-induced jets carry flux, produced by a circum-black hole accretion disk dynamo, or drive jets that generate magnetic fields vis the proton Weibel instability in the shock interaction zone, then one might easily imagine seeding the clouds with a magnetic field that more than sufficed to allow MRI to operate.  In summary, we have outlined two alternative pathways which dispense with the need for Population III to consist exclusively of very massive stars and for the transformation from first stars into Population II to be determined exclusively by the rise in gas and/or dust phase metallicity.  MRI-dynamo generated magnetic flux in the protostellar accretion disk is the key to these possible scenarios. The nature of the seed field is the major difference.  One option is that protogalactic Biermann battery-seeded MRI dynamos provide circumstellar disks with the necessary field strength to overcome high accretion rates characteristic of primordial environments. This allows, and indeed the mechanism itself requires, stars of mass below $10-100\\rm M_\\odot$ to form.  In this case, there is no Population III.  Another possibility is that the first objects (Population III is a misnomer since no nucleosynthetic tracers remain) were very massive ($\\simgt 1000\\rm M_\\odot$), and form IMBHs which expel magnetic-flux-loaded jets that have been generated by dynamo activity in the circum-IMBH disk. The MRI prostellar dynamos are seeded by jet-induced instabilities, and again, the first stars are expected to span a broad mass range. The IMF is a consequence of magnetic feedback. Of course, it may vary with epoch and/or environment, but the principal point we wish to emphasize is that it spans the conventional range of stellar masses. One could then make the transition from Primordial Star Formation to Current Epoch Star Formation while the mean overall metallicity remained extremely low."
    },
    "0606/astro-ph0606089_arXiv.txt": {
        "abstract": "{Results from kinematic solar dynamo models employing $\\alpha$-effect and turbulent pumping from local convection calculations are presented. We estimate the magnitude of these effects to be around 2--3\\,m\\,s$^{-1}$, having scaled the local quantities with the convective velocity at the bottom of the convection zone from a solar mixing-length model. Rotation profile of the Sun as obtained from helioseismology is applied in the models; we also investigate the effects of the observed surface shear layer on the dynamo solutions. With these choices of the small- and large-scale velocity fields, we obtain estimate of the ratio of the the two induction effects, $C_\\alpha/C_\\Omega \\approx 10^{-3}$, which we keep fixed in all models. We also include a one-cell meridional circulation pattern having a magnitude of 10--20\\,m\\,s$^{-1}$ near the surface and 1--2\\,m\\,s$^{-1}$ at the bottom of the convection zone. The model essentially represents a distributed turbulent dynamo, as the $\\alpha$-effect is nonzero throughout the convection zone, although it concentrates near the bottom of the convection zone obtaining a maximum around $30\\degr$ of latitude. Turbulent pumping of the mean fields is predominantly down- and equatorward. The anisotropies in the turbulent diffusivity are neglected apart from the fact that the diffusivity is significantly reduced in the overshoot region. We find that, when all these effects are included in the model, it is possible to correctly reproduce many features of the solar activity cycle, namely the correct equatorward migration at low latitudes and the polar branch at high latitudes, and the observed negative sign of $B_r B_{\\phi}$. Although the activity clearly shifts towards the equator in comparison to previous models due to the combined action of the $\\alpha$-effect peaking at midlatitudes, meridional circulation and latitudinal pumping, most of the activity still occurs at too high latitudes (between $5\\degr \\ldots 60\\degr$). Other problems include the relatively narrow parameter space within which the preferred solution is dipolar (A0), and the somewhat too short cycle lengths of the solar-type solutions. The role of the surface shear layer is found to be important only in the case where the $\\alpha$-effect has an appreciable magnitude near the surface.} ",
        "introduction": "The early solar dynamo models, initiated by Parker (\\cite{Parker1955}), relied on a scalar $\\alpha$-effect and a negative gradient of the angular velocity, $\\Omega$, in the convection zone (e.g. Steenbeck \\& Krause \\cite{SteenKrau1969}; Deinzer \\& Stix \\cite{DeinzerStix1971}; K\\\"ohler \\cite{Koehler1973}). These simple models were enormously succesful in reproducing many of the observed features of the solar cycle, most prominently the equatorward propagating activity belts.  During the recent years, however, helioseismology has revealed that the radial differential rotation shows a positive gradient near the equator, whilst a negative gradient occurs only at high latitudes. Moreover, the strongest gradients occur near the poles and at the bottom of the convection zone (e.g. Schou et al. \\cite{Schouea1998}; Thompson et al. \\cite{Thompsonea2003}). The sign of the $\\Omega$-gradient, if combined with positive (negative) $\\alpha$-effect in the Northern (Southern) hemisphere, leads to poleward migration of the activity belts near the equator and to equatorward migration in the polar regions according to mean-field dynamo theory (Parker \\cite{Parker1955}; Yoshimura \\cite{Yoshimura1975}). This is in contradiction with the observed migration of solar activity tracers, being equatorward at low latitudes, whilst a weaker polarward wave occurs at high latitudes. The observed strong differential rotation in the polar regions, together with the commonly adopted description of the $\\alpha$-effect being proportional to $\\cos{\\theta}$ therefore peaking at the poles, leads to dynamo solutions with high-latitude activity; virtually no sunspots, however, are observed above the latitude $\\pm 40\\degr$.  A suitable meridional circulation pattern is one possible remedy to the aforementioned problems. At the moment it is known that near the surface of the Sun, i.e. the topmost 10--20\\,Mm, there is a poleward flow of the order of 10--20\\,m\\,s$^{-1}$ (Zhao \\& Kosovichev \\cite{ZhaoKosov2004}; Komm et al. \\cite{Kommea2004}), but very little is known about the structure of the flow in deeper layers. Models of the solar differential rotation predict that there should be a return flow near the bottom of the convection zone with a magnitude of approximately $1$\\,m\\,s$^{-1}$ (e.g. Kitchatinov \\& R\\\"udiger \\cite{KitchRued1999}; Rempel \\cite{Rempel2005}). The meridional flow structure and a low value of the turbulent diffusivity, $\\eta_{\\rm t}$, are of critical importance for the so-called flux transport dynamos (e.g. Choudhuri, Sch\\\"ussler \\& Dikpati \\cite{Choudhuriea1995}; Dikpati \\& Charbonneau \\cite{DikChar1999}). These models are essentially $\\alpha \\Omega$-dynamos where the meridional flow works as a conveyor belt for the poloidal magnetic field, and the $\\alpha$-effect is only due to the decay of active regions near the solar surface, or due to instabilities in the magnetic layer below the convection zone (e.g. Schmitt, Sch\\\"ussler \\& Ferriz-Mas \\cite{Schmittea1996}; Brandenburg \\& Schmitt \\cite{BranSchm1998}; Dikpati \\& Gilman \\cite{DikGil2001}), whereas the shear in the tachocline is responsible for the generation of the toroidal field. However, in the present paper we do not exclude the possibility to have an $\\alpha$-effect due to cyclonic turbulent motions in the convection zone proper as is expected in the mean-field theory and suggested by the local convection calculations (e.g. Brandenburg et al. \\cite{Brandea1990}; Ossendrijver, Stix \\& Brandenburg \\cite{Osseea2001}, Ossendrijver et al. \\cite{Osseea2002}; K\\\"apyl\\\"a et al. \\cite{Kaepylaeea2006}, hereafter KKOS). In this case it is questionable whether the meridional flow alone is able to counter the missing activity and wrong migration of the activity belts near the equator.  Another possibility to solve the problems outlined above arises in the form of dynamo coefficients. Most of the solar models have up to date relied on a scalar $\\alpha$-effect, or they have taken into account a limited amount of the anisotropies, such as vertical pumping of the mean magnetic field (e.g. Brandenburg, Moss \\& Tuominen \\cite{BranMoTuo1992}), or the contributions to the $\\alpha$-effect due to the steeply decreasing turbulence intensity towards the bottom of the convection zone (R\\\"udiger \\& Brandenburg \\cite{RuedBran1995}). Recently, Ossendrijver et al. (\\cite{Osseea2002}) computed all of the coefficients $a_{ij}$ that appear in the expansion \\begin{eqnarray} \\mathcal{E}_i = a_{ij} B_{j} + b_{ijk}\\frac{\\pd B_j}{\\pd x_k} + \\cdots \\;, \\label{equ:emf} \\end{eqnarray} describing the small-scale effects on the large scales from local numerical 3D models of stellar convection. Taking only the first term on the rhs into account, $\\vec{\\mathcal{E}}$ can be written as \\begin{eqnarray} \\vec{\\mathcal{E}} = \\overline{\\vec{\\alpha}} \\vec{B} + \\vec{\\gamma} \\times \\vec{B}\\;, \\end{eqnarray} where $\\overline{\\vec{\\alpha}}$ and $\\vec{\\gamma}$ are defined via \\begin{eqnarray} \\alpha_{ij} & = & \\frac{1}{2}(a_{ij} + a_{ji})\\;, \\label{equ:alpha}\\\\ \\gamma_i & = & -\\frac{1}{2} \\varepsilon_{ijk}a_{jk}\\;, \\label{equ:gamma} \\end{eqnarray} and consist of the symmetric and anti-symmetric parts of $a_{ij}$. The diagonal components of $\\alpha_{ij}$ describe the generation of a large scale magnetic field through the $\\alpha$-effect. The vector $\\vec{\\gamma}$ describes turbulent pumping of the mean field with a velocity that is common for all field components. The off-diagonal components of $\\alpha_{ij}$ contribute to the field direction dependent part of the pumping effect (see e.g. Ossendrijver et al. \\cite{Osseea2002}; KKOS).  The main findings for the $\\alpha$-effect in the study of Ossendrijver et al. (\\cite{Osseea2002}) were that whereas it is highly anisotropic, the component corresponding to $\\alpha_{\\phi \\phi}$ is almost constant in magnitude as a function of latitude for Coriolis number of ${\\rm Co} = 2\\,\\Omega \\tau \\approx 2.5$, where $\\tau$ is an estimate of the convective turnover time. It was also established numerically that the turbulent pumping of the mean magnetic field depends upon the field component. The results also clearly indicate down- and equatorward pumping of the mean magnetic fields that can alleviate the problems facing solar dynamo models. Furthermore, recently KKOS extended the study of Ossendrijver et al. (\\cite{Osseea2002}) to the rapid rotation regime, i.e. ${\\rm Co} \\approx 10$, corresponding to the deep layers of the solar convection zone. The main result of this study is that $\\alpha_{\\phi \\phi}$, responsible for the generation of the poloical field from the toroidal one, no longer peaks at the poles but rather at around latitude $\\Theta = 30\\degr$, which could be of further help for the solar dynamo models employing the helioseismic rotation profile. The vertical pumping of the magnetic field is downward near the poles, but can be upward near the equator. The latitudinal pumping is always towards the equator, but it is concentrated in a rather narrow latitude range, i.e. $|\\Theta| < 30\\degr$ near the equator. The vertical pumping of the toroidal field was found to be very small or upward at latitudes $|\\Theta| < 45\\degr$.  In the present study we explore the implications of the local 3D results for the solar dynamo by means of axisymmetric mean-field dynamo models employing the observed internal rotation of the Sun. Starting from a simple $\\alpha$-profile ($\\cos \\theta$ in latitude and constant in radius within the convection zone) we demonstrate the problems facing solar dynamo models. Then we introduce an $\\alpha$-profile that captures the essentials of the magnitude and latitude dependence found in local convection calculations. Then we proceed by adding the pumping effects and finally a reasonable meridional flow constrained by the flow velocity observed on the solar surface. We also discuss the stability of quadrupolar and dipolar field configurations for the chosen $\\Omega$- and $\\alpha$-profiles with and without turbulent pumping and meridional flow. Finally, we discuss briefly the phase dilemma of the radial and toroidal magnetic fields which has been widely used as a constraint for solar dynamo models. The remainder of the paper is organised as follows: in Sect.~\\ref{sec:model} the mean-field model is described in detail, and in Sects.~\\ref{sec:results} and \\ref{sec:conclusions} we present the results and the conclusions of the study. ",
        "conclusions": "\\label{sec:conclusions} We have studied kinematic mean-field solar dynamo models using helioseismologically determined rotation profile and $\\alpha$-effect and turbulent pumping from the local convection calculations of K\\\"apyl\\\"a et al. (\\cite{Kaepylaeea2006}). We show that if a simple $\\alpha$-effect profile proportional to $\\cos \\theta$ in latitude and constant in radius is applied, the results are very sensitive to the presence of a surface shear layer at $r > 0.95R$. Omitting this feature, the migration of the activity belts is equatorward near the poles, weak at midlatitudes, and poleward at low latitudes. The inclusion of the shear layer results in equatorward propagation at virtually all latitudes.  The $\\alpha$-coefficients from local convection calculations tend to be much more concentrated near the bottom of the convection zone and also closer to the equator with peak values occuring at around latitude $30\\degr$. In this case, however, the critical dynamo numbers are significantly larger due to the more concentrated nature of the $\\alpha$-effect. The solution is dipolar when $C_\\alpha < 12$ and quadrupolar (S0) if $C_\\alpha$ is larger even if a purely dipolar initial condition is used. Adding a latitudinal pumping velocity of the order of 2\\,m\\,s$^{-1}$ at the bottom of the convection zone remedies this problem and helps to shift the activity belts closer to the equator, whereas the vertical pumping seems to have a much smaller overall signifigance for the stability and overall structure of the solution. The latitudinal pumping also helps to restore the equatorward migration of activity belts, lost by the implementation of the $\\alpha$-effect from convection calculations. Adding the pumping effects also makes the cycle period significantly longer. Additional upward pumping of the azimuthal field component through the component $\\alpha_{\\phi \\theta}$ at near equator regions has very little effect on the solution. In all the models investigated, the correlation $B_r B_{\\phi}$ is mostly negative, which is in agreement with the observations, but none of the models produces poleward migration at high latitudes (polar branch).  If a single cell meridional flow patterns where the flow is of the order of 10\\,m\\,s$^{-1}$ and poleward near the surface and significantly smaller and equatorward at larger depths helps to shift the activity belts further closer to the equator, although they still appear too high ($|\\Theta| \\approx 5\\degr \\ldots 60\\degr$) in comparison to the Sun. A region of poleward migration also appears at high latitudes.  If the upper shear layer is removed from this model, the butterfly diagram changes only little in contrast to the models where there was an appreciably large $\\alpha$-effect also near the surface.  The distributed dynamo model investigated in this paper correctly reproduces many of the general features of the solar cyclic activity, including realistic migration patterns and correct phase relation, but some problems persist. These include the excitation of the dipolar (A0) solutions only within a narrow range of parameters with the anisotropic $\\alpha$-effect combined with meridional circulation, activity at too high latitudes for all models investigated, and the somewhat too short dynamo cycles. With further tuning of the various parameters going into the model (magnitudes of $\\alpha$- and $\\gamma$-effects, strength and profile of the meridional circulation) it would most likely be possible to obtain a solution reproducing the missing features, as well. Instead of going further into the domain of tuning in the present study, we note that further work in the context of local convection modelling is needed in order to obtain the totally missing information on the anisotropies of turbulent diffusivity (work in progress). It is also necessary to consider the nonlinear saturation process of the $\\alpha$-effect in more detail, for instance in the form of a dynamical $\\alpha$-effect with helicity fluxes (e.g. Brandenburg \\& Subramanian \\cite{BranSubr2005}). Further improvements of the present mean-field model include the relaxation of the kinematic approach to include full dynamics, for which purpose information from local convection calculations also exist (e.g. K\\\"apyl\\\"a, Korpi \\& Tuominen \\cite{Kaepylaeea2004}). With the inclusion of hydro- and thermodynamics the nonlinear feedbacks of the kind investigated by e.g. Rempel (\\cite{Rempel2006}) could be included in a self-consistent way.  Despite of the shortcomings discussed above, we feel that the simple model presented in this paper demonstrates two important aspects concerning solar dynamo theory. Firstly, as recently discussed at length by Brandenburg (\\cite{Brandenburg2005}) the surface shear layer may be important in shaping the solar dynamo; according to our results this is indeed to be expected if the $\\alpha$-effect has an appreciable magnitude near the surface. Secondly, we are able to demonstrate that a distributed $\\alpha \\Omega$-dynamo model applying a realistic rotation profile and meridional flow can reproduce the correct phase relation and migration of the activity belts. Although meridional circulation is an essential ingredient also in our model, the magnitude of this effect is of the same order as the turbulent inductive effects and the turbulent diffusivity is an order of magnitude larger than in the advection dominated flux-transport dynamos."
    },
    "0606/astro-ph0606040_arXiv.txt": {
        "abstract": "We present the spectra of 14 quasars with a wide coverage of rest wavelengths from 1000 to 7300 \\AA. The redshift ranges from $z =$ 0.061 to 0.555 and the luminosity from $M_{B}$ = $-$22.69 to $-$26.32. These spectra of high quality result from combining {\\it Hubble Space Telescope} spectra with those taken from ground-based telescopes. We describe the procedure of generating the template spectrum of \\ion{Fe}{2} line emission from the spectrum of a narrow-line Seyfert 1 galaxy I Zw 1 that covers two wavelength regions of 2200$-$3500 \\AA\\ and 4200$-$5600 \\AA. Our template \\fe2\\ spectrum is semi-empirical in the sense that the synthetic spectrum calculated with the CLOUDY photoionization code is used to separate the \\fe2\\ emission from the \\mg2 $\\lambda$2798 line. The procedure of measuring the strengths of \\fe2\\ emission lines is twofold; (1) subtracting the continuum components by fitting models of the power-law and Balmer continua in the continuum windows which are relatively free from line emissions, and (2) fitting models of the \\fe2\\ emission based on the \\fe2\\ template to the continuum-subtracted spectra. From 14 quasars including I Zw 1, we obtained the \\fe2\\ fluxes in five wavelength bands ($U1$ (2200$-$2660 \\AA), $U2$ (2660$-$3000 \\AA), $U3$ (3000$-$3500 \\AA), $O1$ (4400$-$4700 \\AA), and $O2$ (5100$-$5600 \\AA)), the total flux of Balmer continuum, and the fluxes of \\mg2 $\\lambda$2798, H$\\alpha$, and other emission lines, together with the full width at half maxima (FWHMs) of these lines. Regression analysis was performed by assuming a linear relation between any two of these quantities. Eight correlations were found with a confidence level higher than 99\\%; (1) larger \\ion{Mg}{2} FWHM for larger H$\\alpha$ FWHM, (2) larger $\\Gamma$ for fainter $M_B$, (3) smaller \\ion{Mg}{2} FWHM for larger $\\Gamma$, (4) larger \\ion{Mg}{2} FWHM for smaller \\fe2($O1$)/\\ion{Mg}{2}, (5) larger $M_{BH}$ for smaller $\\Gamma$, (6) larger $M_{BH}$ for smaller \\fe2($O1$)/\\ion{Mg}{2}, (7) larger [\\ion{O}{3}]/H$\\beta$ for larger \\ion{Mg}{2} FWHM, and (8) larger \\ion{Fe}{2}($O1$)/\\ion{Mg}{2} for larger \\ion{Fe}{2}($O1$)/\\ion{Fe}{2}($U1$). The fact that six of these eight are related to FWHM or $M_{BH}$ ($\\propto$ FWHM$^2$) may imply that $M_{BH}$ is a fundamental quantity that controls $\\Gamma$ or the spectral energy distribution (SED) of the incident continuum, which in turn controls the \\fe2\\ emission. Furthermore, it is worthy of noting that \\ion{Fe}{2}($O1$)/\\ion{Fe}{2}($U1$) is found to tightly correlate with \\ion{Fe}{2}($O1$)/\\ion{Mg}{2}, but not with \\ion{Fe}{2}($U1$)/\\ion{Mg}{2}. ",
        "introduction": "There is a growing interest in observing prominent \\ion{Fe}{2} emission lines in the spectra of active galactic nuclei (AGNs) to understand the physics of clouds in the broad emission line regions (BELRs) and also to explore the \\ion{Fe}{2} abundance as a function of cosmic time (e.g., Elston, Thompson, \\& Hill 1994; Kawara et al. 1996; Dietrich et al. 2002, 2003; Iwamuro et al. 2002, 2004; Freudling, Corbin, \\& Korista 2003; Maiolino et al. 2003). According to explosive nucleosynthesis, much of the iron is produced by Type Ia supernovae (SNe Ia), while $\\alpha$ elements such as O and Mg come from Type II supernovae (SNe II). Because progenitors of SNe Ia are long-lived, accreting white dwarfs in binaries, while those of SNe II are short-lived massive stars, Fe enrichment delays relative to $\\alpha$ elements. A lifetime of progenitors of SNe Ia is generally considered to be $\\tau_{\\rm Ia}\\sim$ 1$-$2 Gyr \\citep{hf, ytn, ytk}. For a cosmology of $H_0 = 70$ \\kmsMpc, $\\Omega_M = 0.3$, and $\\Omega_\\lambda = 0.7$ \\citep{ben}, such lifetime corresponds to the cosmic time at $z \\sim$ 3.2$-$5.6. Therefore, hoping to detect a sudden break in the Fe/Mg abundance ratio at high redshift, various groups have measured the flux ratio of \\ion{Fe}{2} emission relative to \\ion{Mg}{2} $\\lambda 2798$, \\ion{Fe}{2}/\\ion{Mg}{2}, in high-redshift quasars.  However, no clear break has been seen in a plot of \\ion{Fe}{2}/\\ion{Mg}{2} as a function of lookback time up to $z \\sim 6$. If the Fe/Mg abundance ratio is reflected in the \\ion{Fe}{2}/\\ion{Mg}{2} flux ratio, no break in the plot implies that a lifetime of progenitors of SNe Ia should be much less than that generally considered. Accordingly, a significantly shorter lifetime of $\\tau_{\\rm Ia}\\sim$ 0.2$-$0.6 Gyr is recently suggested \\citep{ft, mr, gra04}, based on the binary prescription by \\citet{mg} where $\\tau_{\\rm Ia}$ was said to be constrained by observations of binaries in the solar neighborhood. However, it is well known that the observed break of [$\\alpha$/Fe] at [Fe/H] = $-$1 in the solar neighborhood is not explained by this prescription.  The expected break of Fe/Mg at a certain high redshift would be obscured in \\ion{Fe}{2}/\\ion{Mg}{2} by other causes. First of all, the excitation mechanism of \\ion{Fe}{2} emission lines has not been well understood. For example, although model simulations are often performed in the framework of photoionization \\citep{sp, v03, b04}, it is sometimes pointed out that pure photoionization models may not be sufficient to explain the spectrum of strong \\ion{Fe}{2} emitters and a non radiative heating mechanism such as a shock has been suggested especially for the optical \\ion{Fe}{2} emission (e.g. Collin \\& Joly 2000). In addition, such model simulations imply that the relative iron abundance is only one parameter and furthermore not the most important one at all. Such effects of non-abundance factors include the spectral energy distribution (SED), strength of the radiation field, and the gas density of BELR clouds. Recently, \\citet{v03} and \\citet{b04} pointed out that a large microturbulence with $V_{turb} \\geq$ 100 km s$^{-1}$ may be responsible for strong \\ion{Fe}{2} emission. It is therefore important to explore the \\femg\\ flux ratio in the parameter space of non-abundance factors and to find the dependence of \\femg\\ on these factors.  Another cause which might obscure the break of Fe/Mg would come from uncertainties in measuring the respective strengths of Fe II and Mg II. To accurately measure the \\ion{Fe}{2} strength especially in the UV region, it is necessary to determine the contribution from the continuum emission which dominates the UV and optical spectra of quasars. This continuum subtraction becomes more accurate when a wider wavelength coverage of the spectrum is available.  In this paper, we present the spectra of 14 quasars with a wide coverage of rest wavelengths from 1000 to 7300 \\AA, and measure the relative line strength \\femg\\ accurately. A search for non-abundance effects in the \\femg\\ will be made at the end of this paper. A cosmology of $H_0$ = 70 \\kmsMpc, $\\Omega_M = 0.3$, and $\\Omega_\\lambda = 0.7$ will be used throughout. ",
        "conclusions": "We have presented the spectra of 14 quasars with a wide wavelength coverage from 1000 to 7300 \\AA\\ in the rest frame. The redshift ranges from $z =$ 0.061 to 0.555 and the luminosity from $M_{B}$ = $-$22.69 to $-$26.32. These spectra are in high quality and resulted from combining {\\it HST} spectra with those taken on ground-based telescopes. We described the procedure of deriving the template spectrum of \\ion{Fe}{2} from the I Zw 1 spectrum, covering two wavelength regions of 2200$-$3500 \\AA\\ and 4200$-$5600 \\AA\\ where prominent \\ion{Fe}{2} emission is seen.  Our \\fe2\\ template spectrum is semi-empirical in the sense that the synthetic spectrum calculated with the CLOUDY photoionization code is used to separate the \\fe2\\ emission from the \\mg2 $\\lambda$2798 line.  Our procedure of measuring strengths of \\fe2\\ emission is twofold; (1) subtracting the continuum component by fitting models of the power-law and Balmer continua to the continuum windows which are relatively free from line emissions, and (2) fitting models of the \\fe2\\ emission made by the \\fe2\\ template to the continuum-subtracted spectra. From 14 quasars including I Zw 1, we obtained the \\fe2\\ fluxes in five wavelength bands from the UV to optical, the total flux of Balmer continuum, the flux and FWHM of \\mg2 $\\lambda$2798, H$\\alpha$, and other emission lines.  Regression analysis was performed to derive a linear relation between two variables ($Y = A + BX$), and eight correlations with a confidence level of 99\\% or higher were found. These are the correlations of (1) larger \\ion{Mg}{2} FWHM for larger H$\\alpha$ FWHM, (2) larger $\\Gamma$ for fainter $M_B$, (3) smaller \\ion{Mg}{2} FWHM for larger $\\Gamma$, (4) larger \\ion{Mg}{2} FWHM for smaller \\fe2($O1$)/\\ion{Mg}{2}, (5) larger $M_{BH}$ for smaller $\\Gamma$, (6) larger $M_{BH}$ for smaller \\fe2($O1$)/\\ion{Mg}{2}, (7) larger [\\ion{O}{3}]/H$\\beta$ for larger \\ion{Mg}{2} FWHM, and (8) larger \\ion{Fe}{2}($O1$)/\\ion{Mg}{2} for larger \\ion{Fe}{2}($O1$)/\\ion{Fe}{2}($U1$). Six of these eight correlations are related to FWHM or $M_{BH}$ ($\\propto$ FWHM$^2$). This fact may imply that $M_{BH}$ is a fundamental quantity that controls $\\Gamma$ or equivalently SED of the incident continuum that eventually controls \\fe2\\ emission.  Figure \\ref{ObservedSpectra} and the \\ion{Fe}{2} template are available from the web site:  \\newline {\\small http://www.ioa.s.u-tokyo.ac.jp/$^\\sim$kkawara/quasars/index.html}. \\newline"
    },
    "0606/astro-ph0606726_arXiv.txt": {
        "abstract": "We follow the contraction and evolution of a typical Jupiter-mass clump created by the disk instability mechanism, and compute the rate of planetesimal capture during this evolution. We show that such a clump has a slow contraction phase lasting $\\sim 3\\times10^5$ years. By following the trajectories of planetesimals as they pass through the envelope of the protoplanet, we compute the cross-section for planetesimal capture at all stages of the protoplanet's evolution.  We show that the protoplanet can capture a large fraction of the solid material in its feeding zone, which will lead to an enrichment of the protoplanet in heavy elements.  The exact amount of this enrichment depends upon, but is not very sensitive to the size and random speed of the planetesimals. ",
        "introduction": "There are nearly 200 known extrasolar planetary systems.  Many of these systems contain planets of several Jupiter masses.  For those few systems that transit the stellar disk, the radius can be measured, and most of these transiting planets appear to be gas giants (see, e.g. Henry et al. 2000; Udalski et al. 2002; Pont et al. 2004; Alonso et al. 2004). Currently, there are two main theories for forming such planets. The core accretion theory (Pollack et al. 1996) argues that a heavy element core is built up by the accretion of planetesimals. As the core grows, its ability to accrete gas from the surrounding disk increases. When the core is sufficiently massive, there is a rapid accretion of such gas and a giant planet is formed.  The big advantage of this model is that the same basic mechanism of planetesimal accretion will form terrestrial planets, Uranus and Neptune-sized intermediate planets, and giant planets.  The disadvantages are that the core mass required to form Jupiter is at the upper end of the mass estimated from interior models (Saumon and Guillot 2004), and that the time to reach a Jupiter mass is uncomfortably close to the upper limit estimated for lifetime of the gas disk (Haisch, Lada, \\& Lada 2001).  Both these objections can be ameliorated, if not entirely removed, if the accreted material does not sink entirely to the core (Pollack et al. 1996) and if the opacities are in fact lower than those estimated from interstellar grains (Podolak 2003; Hubickyj et al. 2005).  A competing mechanism for giant planet formation is a local disk instability (Boss 1998a). This model suggests that under the right conditions an instability can form in the protoplanetary disk. This instability can lead to the creation of a self-gravitating clump of gas and dust. Such clumps can contract to form giant gaseous protoplanets (Boss 1997; Boss 1998b). This model has several advantages, among them, that planets can be formed quickly, before the nebular gas dissipates. Calculations by Boss (2000) show that the unstable disk can break up into giant gaseous protoplanets in $\\sim10^3$ yrs. The instability occurs over dynamical timescales (some tens of orbital periods), and the formation of clumps (future giant gaseous protoplanets) can easily occur within the estimated lifetime of most circumstellar disks (Haisch, Lada, \\& Lada 2001).  One of the problems with the disk instability model is that the planets formed by this mechanism start with a solar abundance of elements. Observations as well as theoretical models, however, indicate that Jupiter's envelope (and Saturn's as well) is enriched with heavy elements (Young 2003; Saumon et al. 1995). The theoretical estimate for the mass of heavy elements in Jupiter is $\\sim 20 - 30 M_{\\oplus}$ (Saumon and Guillot 2004).  A solar composition planet of Jupiter's mass would be expected to have only $\\sim 6 M_{\\oplus}$ of heavy elements. This means that planets created by the disk instability mechanism must have accreted this additional material later, presumably as solids. In this paper we follow the evolution of an isolated clump and compute the planetesimal accretion rate as a function of time. The calculation is divided into three parts: the computation of the evolution of an isolated clump, the computation of the cross section for the two body interaction between the protoplanet and a planetesimal, and the actual calculation of the mass that is accreted. The details are presented in the following sections. ",
        "conclusions": "We present an evolutionary model of an isolated giant gaseous protoplanet such as might be formed from a clump caused by an instability in the solar nebula. Assuming that such a clump survives and becomes a protoplanet, we calculate the rate at which it would capture planetesimals. We considered three planetesimal sizes and two types of materials, ice + rock, and pure rock. The total mass of solids that is captured depends on the size distribution of bodies in the accretion zone and on their composition. 100 km-sized planetesimals can not be captured during the first $\\sim10^5$ years of the evolution. This is due to the fact that until that time the density and temperature throughout most of the volume of the protoplanet are too low. Bodies with sizes of 1 km or less can be accreted starting at times shortly after the clump's creation.  Icy bodies are captured more easily than pure rock planetesimals.  In the case of higher random velocities, the capture efficiency decreases and the time for capturing the same amount of solids is longer, but even with random velocities as high as 15\\% of the keplerian velocity, more than 16 $M_{\\oplus}$ of heavy elements can be captured. For a higher surface density (bigger amount of solids in Jupiter's feeding zone), the total mass of heavy elements that is captured increases proportionately. As can be seen from table 3 the accreted mass varies by a factor of two when the random velocity of the 100 km planetsimals is increased from 1 to 2 km s$^{-1}$.  Since the latter is already 20\\% of Jupiter's orbital speed, it is unlikely that the random velocities will be significantly higher than this, especially since most of the planetesimals are captured before they can be highly stirred.  Even in this case some 50\\% of the mass in the feeding zone is captured.  If, as is likely, there is a distribution of planetesimal sizes, with 100 km being at the high end, then the fraction of feeding zone mass that is captured will be even higher.  It has been suggested (e.g. Weissman 1986) that planetesimals are actually loosely bound \"rubble piles\".  If this is indeed the case, they would be held together even more weakly than our ice + rock planetesimals.  Such planetesimals would be captured more easily that those composed of ice + rock, and again, the values we have presented for the mass captured would be a lower limit.  Of course all this assumes that the planetesimals remain in the feeding zone of the protoplanet and are not scattered out (see, e.g., Hahn and Malhotra 1999). It should be noted, however, that simulations of planetesimal scattering treat the planets as point masses, and do not include the gas drag on the planetesimal as it goes through the protoplanet's extended envelope.  The energy loss which results from the gas drag will act to inhibit scattering.  In addition, the larger cross section for capture that an extended protoplanet presents, will also significantly increase the ratio of captured to ejected planetesimals.  Thus, in the simulations of Hahn and Malhotra, for example, significant mass ejection requires times much longer than the $3\\times 10^5$ years required for the protojupiter to accrete most of the mass in its feeding zone.  Therefore, if giant planets are created by the disk instability mechanism, there is a good chance that many of the planetsimals will be accreted before they are scattered out of the feeding zone.  As a consequence, these protoplanets can collect a substantial amount of solids during their evolution and end up with a significantly non-solar composition. This finding removes an important objection to the disk instability model as a mechanism for giant planet formation.   Our model assumes a non-rotating clump, in reality, the clump is likely to have some angular momentum. This might act to slow the contraction of the clump in its later stages, and lead to the formation of a subdisk.  Planetesimals in this subdisk could then provide the material for forming both regular and irregular satellites. It is also interesting to speculate on the possibility of observational verification of our results.  Some models of extrasolar planets whose densities are known seem to indicate non-solar compositions (e.g., Bodenheimer et al. 2001). The extent of the enhacement of the non-solar component may provide limits on the proposed accretion mechanism. However, both the fits to the observations and the models we have computed, have, as yet, too many free parameters, to allow useful constraints.   There are a number of issues which may affect the conclusions presented above. In the first place, the captured planetesimals will contribute high-Z material to the gaseous envelope, and change its composition.  This will also affect the opacity of the envelope. Since much of this additional material will end up in the convective region of the planet, this change in opacity should not have a large effect on its subsequent evolution.  In addition, there will be some small effect due to the additional mass of the planetesimals themselves.  Finally, the high-Z material which remains in the envelope in the form of grains will sediment towards the center of the protoplanet leading to the formation of a core, and this too should affect the final evolution of the body.  We hope to address these issues in more detail in future work."
    },
    "0606/astro-ph0606510_arXiv.txt": {
        "abstract": "We give a self-contained modern linear stability analysis of a system of $n$ equal mass bodies in circular orbit about a single more massive body. Starting with the mathematical description of the dynamics of the system, we form the linear approximation, compute all of the eigenvalues of the linear stability matrix, and finally derive inequalities that guarantee that none of these eigenvalues have positive real part.  In the end, we rederive the result that J.C. Maxwell found for large $n$ in his seminal paper on the nature and stability of Saturn's rings, which was published 150 years ago. In addition, we identify the exact matrix that defines the linearized system even when $n$ is not large.  This matrix is then investigated numerically (by computer) to find stability inequalities. Furthermore, using properties of circulant matrices, the eigenvalues of the large $4n \\times 4n$ matrix can be computed by solving $n$ quartic equations, which further facilitates the investigation of stability.  Finally, we have implemented an $n$-body simulator and we verify that the threshold mass ratios that we derived mathematically or numerically do indeed identify the threshold between stability and instability. Throughout the paper we consider only the planar $n$-body problem so that the analysis can be carried out purely in complex notation, which makes the equations and derivations more compact, more elegant and therefore, we hope, more transparent. The result is a fresh analysis that shows that these systems are always unstable for $2 \\le n \\le 6$ and for $n>6$ they are stable provided that the central mass is massive enough.  We give an explicit formula for this mass-ratio threshold. ",
        "introduction": "One hundred and fifty years ago, \\citet{Max59} was awarded the prestigious Adam's prize for a seminal paper on the stability of Saturn's rings.  At that time, neither the structure nor the composition of the rings was known.  Hence, Maxwell considered various scenarios such as the possibility that the rings were solid or liquid annuli or a myriad of small boulders.  As a key part of this last possibility, Maxwell studied the case of $n$ equal-mass bodies orbiting Saturn at a common radius and uniformly distributed about a circle of this radius.  He concluded that, for large $n$, the ring ought to be stable provided that the following inequality is satisfied: \\[ \\mbox{mass(Rings)} \\le 2.298 \\mbox{mass(Saturn)}/n^2 . \\] The mathematical analysis that leads to this result has been scrutenized, validated, and generalized by a number of mathematicians over the years.  We summarize briefly some of the key historical developments. \\citet{1889Tisserand} derived the same stability criterion using an analysis where he assumed that the ring has no effect on Saturn and that the highest vibration mode of the system controls stability.  More recently, \\citet{1986A&A...161..403W} used the theory of density waves to show that Maxwell's results are correct in the limit as $n$ goes to infinity. \\citet{1935RSPTA.234..145P} reformulated the stability problem so that it takes into account the effect of the rings on the central body.  He proved that, for $n \\le 6$, the system is unconditionally unstable.  Inspired by this work, \\citet{1988A&A...205..309S} studied coorbital formations of $n$ satellites for small values of $n$ where the satellites are not distributed uniformly around the central body.  They showed that there are some stable asymmetric formations (such as the well-known case of a pair of ring bodies in L4/L5 position relative to each other---i.e., one leading the other by $60\\deg$).  Finally, \\citet{1991CeMDA..51...83S} extended the analysis of Pendse to find the stability criterion as a function of the number of satellites when $n$ is small.   The resulting threshold depends on $n$ but for $n \\ge 7$, it deviates only a small amount from the asymptotically derived value.   In this paper, we give a self-contained modern linear stability analysis of a system of equal mass bodies in circular orbit about a single more massive body.  We start with the mathematical description of the dynamics of the system. We then form the linear approximation, compute all of the eigenvalues of the matrix defining the linear approximation, and finally we derive inequalities that guarantee that none of these eigenvalues have positive real part.  In the end, we get exactly the same result that Maxwell found for large $n$. But, in addition, we identify the exact matrix that defines the linearized system even when $n$ is not large.  This matrix can then be investigated numerically to find stability inequalities even in cases where $n$ is not large.  Furthermore, using properties of circulant matrices, the eigenvalues of the large $4n \\times 4n$ matrix can be computed by solving $n$ quartic equations, which further facilitates the investigation. Finally, we have implemented an $n$-body simulator based on a leap-frog integrator (see \\cite{ST94,HMM94}) and we verify that the threshold mass ratios that we derived mathematically or numerically do indeed identify the threshold between stability and instability.  Throughout the paper we consider only the planar $n$-body problem.  That is, we ignore any instabilities that might arise due to out-of-plane perturbations.  Maxwell claimed, and others have confirmed, that these out-of-plane perturbations are less destabilizing than in-plane ones and hence our analysis, while not fully general, does get to the right answer.  Our main reason for wishing to restrict to the planar case is that we can then work in the complex plane and our entire analysis can be carried out purely in complex notation, which makes the equations and derivations more compact, more elegant and therefore, we hope, more transparent. ",
        "conclusions": ""
    },
    "0606/astro-ph0606660_arXiv.txt": {
        "abstract": "{ This paper provides an alternate derivation of the equations used in the GRMHD code of De Villiers \\& Hawley using Stokes Theorem. This derivation places the published form of the discretized equations of GRMHD in a broader context, and suggests possible future directions for numerical work.} ",
        "introduction": "\\parskip=12pt  The general relativistic magnetohydrodynamic (GRMHD) code of De Villiers \\& Hawley (2003; DH) has been used to study the accretion of matter onto black holes. It is a highly stable $3+1$D code that has been used to simulate in great detail the evolution of accretion disks in the spacetime of Kerr black holes. This code has established the ubiquity of energetic axial jets in such environments (De Villiers {\\it et al.}, 2003, 2005), and has shed light on the magnetic environment produced accretion disks (Hirose {\\it et al.}, 2004; De Villiers, 2006).  The refinement of the code through testing and development is ongoing. One area of study that may help spur future improvements is the mathematical foundation of the GRMHD equations. There are many threads in the literature documenting the differential version of the equations. Misner, Thorne, \\& Wheeler (MTW; 1973) provide a clear derivation of the equations of hydrodynamics. Anile (1989) discusses the extension to magnetohydrodynamics (MHD) in a special relativistic background. Hawley, Smarr, \\& Wilson (1984; HSW) treat hydrodynamics in a fixed background (black hole spacetime), while DH completes the derivation of the complete set of GRMHD equations.  This paper presents a derivation of the GRMHD equations from the generalized Stokes Theorem (MTW; see also Page \\& Thorne, 1974), and was motivated by a derivation due to Clarke (1996) of the equations of non-relativistic MHD as discretized for the ZEUS code. Clarke provides an integral formulation of the equations of MHD using Stokes Theorem, and clearly shows the connection between the apparently simplistic finite-differencing techniques used in ZEUS (and also in the GRMHD code) to the notion of flux conservation embodied in Stokes Theorem. ",
        "conclusions": "This paper has outlined the development of discretized equations of ideal General Relativistic Magnetohydrodynamics (GRMHD) from the integral form of the conservation laws. Stokes Theorem transforms the conservation laws of GRMHD into simple flux integrals through the boundary of a closed four-volume, and many possible discretizations are possible for these integrals. The particular choices used in the GRMHD code for the equation of continuity and the induction equation represent but one implementation, arguably the simplest and most straightforward of these. Many other possible implementations are implied by equations (\\ref{continuity}) and (\\ref{induction}). The GRMHD code uses an alternate discretization of the energy-momentum conservation equation; the integral formulation discussed here, (\\ref{tmn}), offers another possible discretization in keeping with the treatment of the continuity and induction equations. A note of caution is in order: the integral relations derived here, dealing as they do with coupled sets of non-linear equations, may not yield stable numerical schemes. Since only limited stability analysis is possible in non-linear differential equations, theoretical work must be complemented by careful development and testing to validate any particular numerical treatment.  In addition, this discussion has implicitly assumed smooth flows; the issue of shock capturing has not been touched upon. The GRMHD code uses a relativistic treatment that is based on the original Von Neumann \\& Richtmyer (1950) artificial viscosity treatment, a relativistic extension of this technique by May \\& White (1967), and HSW. This remains the state of the art as of this writing.  As a final note, the zone-based, flux-conserving formulation discussed here may also be of benefit in treating boundary conditions. The boundary zones, or ghost zones, of the computational grid are used to impose physical boundary conditions on the computational volume and also to provide approximate data used by the interpolation routines to provide correct time-centering of zones on the computational grid adjacent to the physical boundaries. The main interest in the integral formulation of boundary data lies at the inner radial boundary, where frame dragging effects are important, and also at the difficult ``corner zones'' which sit just outside the pole caps of the black hole. Definitive statements regarding the flux of electromagnetic energy near the event horizon hinge on reliable boundary treatments, and it is hoped that a more sophisticated approach, based on the integral formulation discussed here, may help. This remains an open area of research.  \\noindent{\\bf Acknowledgements:}  \\noindent The GRMHD code was developed by the author while at the University of Virginia, under NSF grant AST-0070979 and NASA grant NAG5-9266 (J.F. Hawley and S. Balbus, Principal Investigators). A small portion of this work was also carried out while the author was at the University of Calgary, partially supported by the Natural Sciences and Engineering Research Council of Canada (NSERC), as well as the Alberta Ingenuity Fund (AIF)."
    },
    "0606/astro-ph0606383_arXiv.txt": {
        "abstract": "Several commercial telecommunication ventures together with a well funded US military program make it a likely possibility that an autonomous, high-altitude, light-than-air (LTA) vehicle which could maneuver and station-keep for weeks to many months will be a reality in a few years. Here I outline how this technology could be used to develop a high-altitude astronomical observing platform which could return high-resolution optical data rivaling those from space-based platforms but at a fraction of the cost. ",
        "introduction": "\\label{sect:intro}  %  There is currently considerable interest in the telecommunications industry and the US military in developing an autonomous, high-altitude, light-than-air (LTA) vehicle that could maneuver and station-keep for weeks to many months.  Several telecommunication companies located both in the US and abroad along with several branches of the US military are investing tens of million of dollars in the development of these so-called hybrid airships: LTA vehicles consisting of helium filled superpressure balloon(s) that are made maneuverable by means of specially designed high-altitude propellers and powered by solar panels during the day and solar charged batteries and/or power cells at night. If in the next few years such vehicles are successfully built and flown, then these technologies allow for the potential development of a high-altitude astronomical observing platform which could return optical and infrared data rivaling those from space-based platforms but at a fraction of the cost.   \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=10cm]{fig1_gimp.ps} \\end{tabular} \\end{center} \\caption[example] { \\label{haa} Artist concepts of the US Missile Defense Agency's high-altitude airship (upper left; Source: Lockheed Martin), Lindstrand's high-altitude airship design for ESA (upper right; Source: Lindstrand Balloons Ltd), Korea's HAP airship (lower left; Source: Korea Aerospace Research Institute), and a proposed spherical, high-altitude communications platform (lower right; Source: 21$^{\\rm st}$ Century Airships).} \\end{figure}  \\subsection{Commercial High-Altitude Telecommunication and Internet Platforms}  For more than a decade, wireless communications using a quasi-stationary, high-altitude platform station (HAPS) for broadcasting services has been considered by several telecommunication companies as a way to economically expand commercial high-bandwidth data services to consumers. Proposed HAPS vehicles have included high-altitude aircraft and LTA airships operating at altitudes typically between 17 and 22 km  (55 -- 72 kft) where stratospheric wind speeds are generally lowest. Companies and governments which have pursed or are currently pursuing this technology using airships include the US-based firms SkyStation, Platforms Wireless International, and Sanswire Networks, the UK firm Advanced Technology Group (ATG), SkyNet and the Wireless Innovation Systems Group of the Yokosuka Radio Communications Research Center in Japan, and the Korean Ministry of Commerce, Industry, and Energy. The advantages and technical problems of HAPS for wireless communications has been well laid out and discussed in several journal and magazine articles\\cite{Tozer01,Platt99,Djuknic97}.  While no full-scale autonomous vehicle has yet been successfully flown achieving commercial payload and station-keeping duration needs, several prototypes have been built and undergone successful initial test flights.  \\subsection{`Near-Space' Military Reconnaissance and Tactical Platforms }  Across the military services, there is increased demand for real-time communications and `over-the-horizon' surveillance capabilities. Recent advances in LTA vehicles that operate at high-altitudes and hence offer large surveillance areas and good air defense survivability factors have attracted considerable interest from several military organizations.  For example, an airship at an altitude of 70 kft would have a line-of-sight regional coverage some 325 miles in diameter, meaning that just one such vehicle could survey nearly all of Afghanistan.  Military interest also arises from the potentially much longer on-station times of high-altitude airship platforms compared to high-altitude UAVs (i.e., weeks to months vs.\\ Global Hawk's $\\sim$ 60 hours) along with a low probability of communication intercept due to stable, direct line-of-sight communications\\cite{Jamison05}.  In this way, an airship would function as surrogate satellite but offer shorter transmission distances in theater battlespace and shorter ranges and hence higher resolution for sensor surveillance of ground targets\\cite{Jamison05}.  Probably the most aggressive and well funded military effort to develop a high-altitude airship (HAA) has been pursued by the US Missile Defense Agency (MDA) through its Advance Concepts Technology Development (ACTD) program.  This program's goal is the development of an autonomous airship that would be stationed at 65 kft, have an endurance of at least one month (with one year desired), carry a payload of 4000 lbs and be capable of supplying that payload with 10 kW of power. The vehicle envisioned would have a top airspeed of over 30 knots to station-keep on a designated ground target and be some 450 ft long and 150 ft in diameter with over 5 million cubic feet in volume. Lockheed Martin's Naval Electronic and Surveillance Systems (Akron OH) was selected for the initial design study and awarded US\\$40M in September 2003. It was recently awarded US\\$149M in December 2005 to begin building a small prototype (top speed 25 kts, 500 lb payload with 3 kW of power) to fly before 2010. The general vehicle concept for the Lockheed Martin design is that of a conventional dirigible shape with photovoltaic cells across the upper surfaces (see Fig.\\ 1). Station-keeping target parameters are $<$ 2 km for 50\\% of the time and $<$150 km 95\\% of the time.  The US Air Force is also interested in an unmanned vehicle to loiter at 100 kft for several days carrying a 50 kg surveillance payload and awarded a contract to JP Aerospace through the Space Battlelab and Space Warfare Center in Colorado.  The European Space Agency (ESA) has also considered development of a solar-powered airship and awarded a design study to Lindstrand Balloons (see Fig.\\ 1). Their airship design is similar to the MDA airship but with a single 8-meter diameter propeller and 15 kW of power for a communications payload. ",
        "conclusions": ""
    },
    "0606/astro-ph0606456_arXiv.txt": {
        "abstract": "We are carrying out a survey to search for giant extrasolar planets around nearby, moderate-age stars in the mid-infrared L' and M bands (3.8 and 4.8 microns, respectively), using the Clio camera with the adaptive optics system on the MMT telescope.  To date we have observed 7 stars, of a total 50 planned, including GJ 450 (distance about 8.55pc, age about 1 billion years, no real companions detected), which we use as our example here.  We report the methods we use to obtain extremely high contrast imaging in L', and the performance we have obtained.  We find that the rotation of a celestial object over time with respect to a telescope tracking it with an altazimuth mount can be a powerful tool for subtracting telescope-related stellar halo artifacts and detecting planets near bright stars.  We have carried out a thorough Monte Carlo simulation demonstrating our ability to detect planets as small as 6 Jupiter masses around GJ 450.  The division of a science data set into two independent parts, with companions required to be detected on both in order to be recognized as real, played a crucial role in detecting companions in this simulation. We mention also our discovery of a previously unknown faint stellar companion to another of our survey targets, HD 133002.  Followup is needed to confirm this as a physical companion, and to determine its physical properties. ",
        "introduction": "\\label{sect:intro}  %  Slightly over a decade ago, the observational science of extrasolar planetary systems began with the discovery of the planet 51 Pegasi B by Michel Mayor and Didier Queloz.  Since that time nearly 200 additional extrasolar planets have been discovered, mostly by using precise spectroscopic radial velocity measurements to detect the reflex motion of the parent star in response to the orbiting planet's gravity -- the same method used in the 51 Pegasi detection. Direct imaging of extrasolar planets remains an elusive goal.  It is a very attractive goal, however.  A planet detected by direct imaging could be investigated with spectroscopy and multi-band photometry, and astrometric measurements over time would yield its exact orbit.  Far more could be learned about such a planet than about one that is only detected by the radial velocity method. Also, the radial velocity method is most sensitive to planets in fairly close-in orbits, unlike the giant planets in our own solar system. Direct imaging is needed to find planets further out, and thus systems more analagous to our own.  Immediately after their formation in the protoplanetary disk of a star, giant planets are extremely hot due to the gravitational potential energy released in their accretion, and therefore glow brightly at infrared wavelengths. Since they have no internal source of energy, they cool over time, and their infrared flux decreases and shifts toward longer wavelengths.  Theoretical models \\cite{Baraffe} show that the current generation of large astronomical telescopes with adaptive optics (AO) systems is capable of directly imaging giant extrasolar planets in the infrared. The cooling properties of extrasolar planets, and the distribution of stars in the sun's vicinity, suggest two good strategies.  First, a planet search may focus on very young stars using the near infrared H and K bands (1.6 and 2.2 microns, respectively).  Stars with ages in the range of 5-30 million years (Myr) are good targets for this method, because their giant planets will still be bright at short wavelengths. Such searches can detect planetary systems that are still in the process of formation -- a fascinating possibility.  However, stars that are sufficiently young are rare, and therefore tend to be distant from the sun, so only bright planets at relatively large physical seperations can be detected.  A second promising method is to using longer infrared wavelengths, such as the L' and M bands (3.8 and 4.8 microns, respectively) to observe nearer stars at more moderate ages.  At these wavelengths, giant planets can be seen around very nearby stars up to the 5 billion year (Gyr) age of our own solar system.  Younger stars are still preferred, but excellent sensitivity can be obtained for stars in the 100 Myr to 1 Gyr range, far older than the optimal range for H and K searches.  Stars at these ages are far more common than those younger than 100 Myr, and many examples can be found very close to the sun.  Very faint planets can be detected around such nearby stars, and planets at small physical seperations from their parent stars can be resolved.  We have begun an L' and M band survey of 50 nearby, moderate age stars for extrasolar planets, using the newly developed L' and M band camera Clio on the 6.5 meter MMT telescope of the University of Arizona, with its deformable secondary AO system.  We obtain slightly better sensitivity to low-mass planets in the L' band than in the M band, so we plan to do our initial observations in L' and use the M band for followup as needed. To date, we have observed 7 of our 50 survey targets.  We present here the methods we use to obtain extremely high contrast and sensitivity. In Section~\\ref{sect:obs} we describe our observations.  Section~\\ref{sect:data} presents our data analysis methods.  In Section \\ref{sect:sens} we give a detailed analysis of the sensitivity we have obtained for planets orbiting the star GJ 450, one of the 7 targets we have observed so far. No companions to GJ 450 have been detected.  In Section \\ref{sect:mstar} we announce the discovery of a probable low-mass star orbiting another of our survey targets, HD 133002.  Section \\ref{sect:concl} gives our conclusions. ",
        "conclusions": "\\label{sect:concl}  The combination of the Clio L' and M band camera and the MMT AO system is capable of sensitive, high contrast observations in the L' band.  Detection of planets in the L' band requires extensive data processing.  Our method of using sky rotation to remove super speckles is powerful, at least in the specific case of the MMT and its AO system.  Estimating the sensitivity of L' observations in extrasolar planet searches is tricky and very important in order for upper limits to be meaningful.  Simple sensitivity estimates based on the rms of image regions should never be trusted without confirmation by Monte Carlo simulation, preferably a carefully constructed blind simulation such as we have carried out. Our simulation demonstrated our ability to detect planets of 6 Jupiter masses orbiting a 1 Gyr old star, and if planets near this mass are sufficiently common in 10-40 AU orbits we will probably discover some of them in the course of our survey. Our method of splitting the data into two independent halves and requiring that any real source be detectable on both halves is a good one in order to sort through numerous spurious detections and find those that are real.  The acqusition of a small number of short, unsaturated exposures of our science targets also played an important role in our sensitivity calculation and our Monte Carlo simulation, as it allowed us to accurately know the true PSF a planet would be expected to have.  We have discovered a stellar companion to HD 133002.  If it is physically related to HD 133002, it is a low mass red dwarf between 0.1 and 0.21 solar masses.  It awaits proper motion confirmation and other followup studies."
    },
    "0606/astro-ph0606442_arXiv.txt": {
        "abstract": "We suggest a method which improves the precision of studies of the primary composition of ultra-high-energy cosmic rays. Two principal ingredients of the method are (1)~comparison of the observed and simulated parameters for {\\em individual} showers, without averaging over arrival directions and (2)~event-by-event selection of simulated showers by the physical observables and not by the reconstructed primary parameters. A detailed description of the algorithm is presented and illustrated by several examples. ",
        "introduction": "\\label{sec:intro} The determination of the primary composition of cosmic rays with energies higher than $\\sim 10^{17}$~eV is a real challenge. The lack of knowledge of the types of primary ultra-high-energy particles which induce extensive air showers makes it difficult to study their origin and in some cases even to determine their energy spectrum. More precise and less model-dependent determination of the cosmic-ray primary composition, especially in the highest-energy domain, is one of the most important tasks in contemporary astroparticle physics (for a review and discussions see e.g.\\ Ref.~\\cite{Watson}).  Ultra-high-energy cosmic rays are observed through air showers, and to extract information about the primary particle, selected observables of real air showers are compared to those of  simulated ones for different primaries. Because of large shower-to-shower fluctuations, one cannot determine the primary-particle type of an individual air shower. That is why traditional approaches to the composition studies are based on the determination of average characteristics of a large sample of cosmic-ray events. This approach has obvious advantages: for a homogeneous composition, averaging smoothens out the fluctuations, and large statistics results in higher precision. Furthermore, the computational time required to perform reliable simulations of an observable averaged over a large sample is much smaller than that necessary for detailed simulations of all events in the data set. However, for mixed composition, averaging might become a problem because fluctuations of the discriminating observables over their central values are larger when air showers from different arrival directions are combined in one sample. Different zenith angles correspond to different atmospheric depths, and air showers detected by surface arrays have different stages of development and hence may have quite different observable parameters even for the same energies and primaries.  Moreover, the geomagnetic field introduces the azimuthal dependence for photon~\\cite{gamma-gMF} and very inclined hadron~\\cite{inclined-gMF} showers. Even fluorescent detectors, which observe the shower development on its way through the atmosphere, respond to showers from different arrival directions in different ways, notably with different accuracies.  For gamma-ray primaries at $E\\gtrsim 5\\cdot10^{19}$~eV, the entire shower development is direction-dependent.  We suggest to improve considerably the precision of the composition studies by performing individual simulations for each observed high-energy event. In the case of a large number of events in a data sample, averaging in bins of arrival directions may be used. If one is not interested in the study of possible gamma-ray primaries with $E\\gtrsim 5\\cdot10^{19}$~eV, binning in zenith angle only is often sufficient.  This method, in its simplest form with one observable, has been successfully implemented to obtain a limit on the gamma-ray fraction in the primary flux at $E>10^{20}$~eV using the data of the AGASA and Yakutsk experiments~\\cite{gamma-limit}. Event-by-event simulations (without selection by reconstructed parameters) were previously used for the composition studies in Refs.~\\cite{Fly-s-Eye,Homola,Fly-s-Eye-Homola}.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:general}, we present the sketch of the event-by-event approach to the composition studies, discuss the choice of the air-shower observables (Sec.~\\ref{sec:EandC}), give a general idea of how to constrain the probable primary-particle type of an individual air shower (Sec.~\\ref{sec:general-event}) and how to use these shower-by-shower constraints to gain information about the chemical composition of the primary cosmic-ray flux using a sample of events (Sec.~\\ref{sec:general-ensemble}). Sec.~\\ref{sec:integrals} contains the detailed description of the procedure outlined in Sec.~\\ref{sec:general} and ready-to-use formulae implementing this procedure.  Sec.~\\ref{sec:example} presents several examples which illustrate the method by the analysis of small (and hence statistically insignificant) samples of events. There, we consider only one composition-related observable, the muon density, and use samples of highest-energy AGASA and Yakutsk events as examples. In Sec.~\\ref{sec:example-event}, we analyse in detail a single event (the highest energy air shower reported by AGASA). The procedure for estimating the limit on the fraction of particular primaries in a given energy range is illustrated in Sec.~\\ref{sec:example-gamma-limit} with the sample of six AGASA showers with reported energies higher than $10^{20}$~eV and known muon content, while the algorithm to determine the best-fit composition assuming two possible primaries is illustrated in Sec.~\\ref{sec:example-p-Fe} with a sample of four events with reported energies above $1.5\\cdot 10^{20}$~eV detected by the AGASA and Yakutsk experiments. Both samples are small and the analysis of Sec.~\\ref{sec:example} serves for illustrative purposes only. We briefly summarize and discuss novelties of our method in Sec.~\\ref{sec:concl}.  Some of our notations are summarized in Appendix~\\ref{app:notations}.  Appendices~\\ref{app:events} and \\ref{app:reconstruction} contain technical information related to the examples presented in Sec.~\\ref{sec:example}. ",
        "conclusions": "\\label{sec:concl} To summarise, we presented a simple method which allows to improve precision of the studies of the primary composition of ultra-high-energy cosmic rays. Each event is studied individually. The simulated showers are selected by the physical observables to be consistent with the real event.  Our approach exploits the fact that the uncertainty in discrimination of the primary-particle types by conventional methods is determined not only by intrinsic fluctuations of the shower development, but also by the spread related to the arrival direction. Averaging over arrival directions introduces artificial, easy-to-avoid, fluctuations in the determination of both the reconstructed energy (see Fig.~\\ref{Fig:discussion-gamma} for an illustration) and composition-related parameters (Fig.~\\ref{Fig:discussion-Fe}). \\begin{figure} \\centerline{ \\includegraphics[width=112mm]{E0-Erec-gamma.eps} } \\caption{ Direction dependence of the reconstructed energy for gamma-ray primaries. Plotted is the reconstructed energy (determined by the AGASA method from $S(600)$) versus the primary energy. Dark boxes: arrival direction of the event 1; crosses: arrival direction of the event 3; grey circles: arrival directions randomly distributed according to the AGASA acceptance ($0<\\theta<45^\\circ$). Both $E_0$ and $E_{\\rm rec}$ are measured in eV. \\label{Fig:discussion-gamma} } \\end{figure} \\begin{figure} \\centerline{ \\includegraphics[width=112mm]{rho_mu-E0_ferrum.eps} } \\caption{ Direction dependence of muon density for iron primaries. Plotted is $\\rho _\\mu (1000)$ versus the primary energy. Dark boxes: arrival direction of the event 7 ($\\theta \\approx 48^\\circ$); crosses: arrival direction of the event 8 ($\\theta\\approx 59^\\circ$); grey circles: arrival directions randomly distributed according to the Yakutsk acceptance ($0<\\theta<60^\\circ$). $E_0$ is measured in eV, $\\rho _\\mu $ is measured in m$^{-2}$. \\label{Fig:discussion-Fe} } \\end{figure}  While the direction-related systematical effects are controlled within our method, there are still many other sources of uncertainty, the most important among them are the low precision of the measurement of composition-related quantities in experiments and artificial fluctuations in simulated air showers due to use of the thinning approximation. Use of multiple c-observables, their careful choice and optimisation of simulations may help to improve further the precision of our method.   As for any other method based on air-shower simulations, the capabilities of our approach are limited (and severely limited in the case of distinguishing different hadronic primaries) because of persistent theoretical problems in modelling high-energy hadronic interactions.  As compared to the traditional methods, the event-by-event study requires more computational time, which makes it difficult to implement it for large samples in the form described here. However, almost the same level of precision may be reached with reasonable binning in zenith (and in some cases also azimuth) angle.  As we have seen in Sec.~\\ref{sec:example}, the method is really strong in studies of the small samples, notably for the highest-energy or somewhat special (e.g.\\ correlated with particular hypothetical sources) events, where statistics is insufficient to obtain confident conclusions with the help of traditional methods. In some cases, e.g.\\ to constrain the gamma-ray primaries, even a single event may be succesfully studied.    \\ack We are indebted to L.~Dedenko, M.~Kachelrie\\ss, M.~Pravdin, K.~Shinozaki, D.~Semikoz and M.~Teshima for numerous helpful discussions. This work was supported in part by the INTAS grant 03-51-5112, by the Russian Foundation of Basic Research grants 05-02-17363 (DG and GR) and 04-02-17448 (DG), by the grants of the President of the Russian Federation NS-7293.2006.2 (government contract 02.445.11.7370; DG, GR and ST) and MK-2974.2006.2 (DG), by the fellowships of the \"Dynasty\" foundation (awarded by the Scientific Council of ICFPM, DG and GR) and of the Russian Science Support Foundation (ST). Numerical part of the work was performed at the computer cluster of the Theoretical Division of INR RAS.   \\appendix"
    },
    "0606/astro-ph0606397_arXiv.txt": {
        "abstract": "The Telescope to Observe Planetary Systems (TOPS) is a proposed space mission to image in the visible (0.4-0.9 $\\mu$m) planetary systems of nearby stars simultaneously in 16 spectral bands (resolution R$\\approx$20). For the $\\approx$10 most favorable stars, it will have the sensitivity to discover $2 R_E$ rocky planets within habitable zones and characterize their surfaces or atmospheres through spectrophotometry. Many more massive planets and debris discs will be imaged and characterized for the first time. With a 1.2m visible telescope, the proposed mission achieves its power by exploiting the most efficient and robust coronagraphic and wavefront control techniques. The Phase-Induced Amplitude Apodization (PIAA) coronagraph used by TOPS allows planet detection at 2$\\lambda/d$ with nearly 100\\% throughput and preserves the telescope angular resolution. An efficient focal plane wavefront sensing scheme accurately measures wavefront aberrations which are fed back to the telescope active primary mirror. Fine wavefront control is also performed independently in each of 4 spectral channels, resulting in a system that is robust to wavefront chromaticity. ",
        "introduction": "\\label{sect:intro} There are now about 200 known exoplanets, most of them identified by indirect detection techniques. While detection of Earth-like planets is still out of reach for current instruments, it appears almost certain that such planets exist, and may even be in habitable zones of nearby stars \\cite{kast93}. To both improve planet detection capabilities and allow characterization of their atmospheres and surfaces, we propose in this paper a visible 1.2-m diameter imaging telescope: the Telescope to Observe Planetary Systems (TOPS). TOPS would, for the first time, be able to directly image planetary systems similar to ours. It would identify and characterize massive rocky planets, gas giants and exozodiacal clouds around nearby stars.  To overcome the enormous star/planet contrast challenge, TOPS utilizes the most efficient coronagraphic and wavefront control techniques recently developed. The coronagraph adopted is the highly efficient Phase-Induced Amplitude Apodization (PIAA) coronagraph, described in \\S\\ref{sec:sls}. Wavefront control, described in \\S\\ref{sec:wfc}, relies upon focal plane wavefront sensing and a 2-step wavefront correction (primary mirror and fine DM). The spacecraft, described in \\S\\ref{sec:spacecraft}, are designed to meet the challenging wavefront stability requirements of this mission. Expected performance for both detection and characterization of exoplanets is quantified in \\S\\ref{sec:performance}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606504_arXiv.txt": {
        "abstract": "This is a sequel to the previous paper, in which we investigated the structure and stability of the spherically symmetric accretion flows through the standing shock wave onto the proto-neutron star in the post-bounce phase of the collapse-driven supernova. Following the prescription in the previous paper, we assume that the accretion flow is in a steady state controlled by the neutrino luminosity and mass accretion rate that are kept constant. We obtain steady solutions for a wide range of neutrino luminosity and mass accretion rate. In so doing, as an extension to the previous models, we employ a realistic EOS and neutrino-heating rates. More importantly, we take into account the effect of convection phenomenologically. For each mass accretion rate, we find the critical neutrino luminosity, above which there exists no steady solution. These critical points are supposed to mark the onset of the shock revival. As the neutrino luminosity increases for a given mass accretion rate, there appears a convectively unstable region at some point before the critical value is reached. We introduce a phenomenological energy flux by convection so that the negative entropy gradient should be canceled out. We find that the convection lowers the critical neutrino luminosity substantially, which is in accord with the results of multi-dimensional numerical simulations done over the years. We also consider the effect of the self-gravity, which was neglected in the previous paper. It is found that the self-gravity is important only when the neutrino luminosity is high. The critical luminosity, however, is little affected if the energy transport by convection is taken into account. ",
        "introduction": "Core-collapse supernovae are triggered by the gravitational collapse of massive stars. The mechanism of explosion is still not well understood. Many researchers think that the so-called neutrino heating mechanism is the most promising at present. In this scenario, after the shock is launched by the core bounce, it stalls inside the core due to the energy loss by photo-dissociations of nuclei as well as neutrino-cooling, and then it is revived by the irradiation of neutrinos diffusing out of the proto-neutron star \\citep{wil85}. Although many researchers have studied this scenario, numerical simulations have shown that the stalled shocks do not revive as long as the collapse is spherically symmetric \\citep{lie01,bur03,tho03,lie05,sum05}.  After the shock is stagnated due to the energy losses, the accretion through the standing shock wave onto the proto-neutron star is approximately steady. \\citet{bur93} and \\citet{yam05} attempted to approximate this phase of evolution by solutions of the time-independent Euler equations, assuming constant mass-accretion rates and neutrino luminosities. Adopting such a method, they demonstrated clearly what is the cause of the failure of explosions found in the numerical simulations. Varying these controlling parameters, the mass-accretion rate and neutrino luminosity, they found that for a given mass-accretion rate there is a critical neutrino luminosity, above which there exists no steady solution. Based on this fact, they argued that the revival of stalled shock occurs when the neutrino luminosity exceeds this critical value.  \\citet{yam05} found further that there are in general two types of solutions for a given mass accretion rate and a neutrino luminosity. The shock radii in these two solutions differ. They referred to the solution with a smaller shock radius as the inner solution and to the other solution as the outer solution. As the neutrino luminosity is increased with the mass accretion rate fixed, the shock radius for the inner solution becomes larger, whereas that for the outer solution becomes smaller. Two solutions coincide with each other when the luminosity reaches the critical value. For the luminosity over the critical value, there is no solution. They also studied the linear stability of the system against radial perturbations and found that the inner solution is stable while the outer solution is unstable and they are merged at the critical luminosity and become marginally stable. In that paper, the equation of state and the reaction rates for neutrino processes were simplified, and the self-gravity of the accreting matter is neglected. We improve these aspects in this paper and discuss their effects on the results.  The other concern in this paper is the convection in the accreting flow. The convection in the supernova core and its implication for the explosion has been studied by many researchers using multi-dimensional numerical simulations over the years \\citep{her94,bur95,jan96,mez98}. They have demonstrated that the convection tends to assist the revival of shock owing to the non-radial motions and the enhanced heating of the accreting matter although it seems that the convection alone can not give a successful explosion. Our main concern here is whether we can understand in the frame work by \\citet{yam05} the results obtained by the large scale simulations . In fact, we have recognized that some solutions in the paper \\citep{yam05} have a region that has a negative entropy-gradient and is unstable against convective motions. In this paper, we discuss the effect of convection by incorporating a phenomenological energy flux in our time-independent, spherically symmetric models. The convective energy flux is introduced in such a way that the negative entropy-gradient should be canceled out. This assumption will correspond to the maximal mixing by convection and, hence, we can hopefully estimate the upper limit for the effect of convection on the critical neutrino luminosity. ",
        "conclusions": "In this paper, we have obtained spherically symmetric steady accretion flows through the standing shock wave onto the proto-neutron star. We have improved the treatment of microphysics such as the equation of state and neutrino reaction rates. We have found that our previous findings are essentially unchanged after these modifications. For a given mass accretion rate, there are in general two solutions, the inner and outer solutions. There is a critical neutrino luminosity for a given mass accretion rate, above which no steady solution exists. The outer solution is unstable against radial perturbations and the two branches of solution coincide with each other for the critical luminosity. There is a quantitative difference, though. The critical luminosity, in particular, becomes much higher when the realistic microphysics are implemented.  We have also investigated the effect of convection, employing a phenomenological description of convection. We have introduced the convective fluxes of energy and proton number, which then have been determined in such a way that the negative entropy gradient is canceled out by them. This is roughly corresponding to the maximally efficient mixing by convection. We have found that the critical luminosity is lowered by a factor of $\\sim2$ by the convection, consistent with results from the large-scale numerical simulations done so far. The latest results of detailed numerical simulations suggest that the convection is not sufficient for the shock revival \\citep{bur06}. Comparing the values of the critical luminosity in our models with the actual luminosities obtained by detailed simulations \\citep{sum05}, we have found that the former is larger. Thus, our results agree with the numerical simulations also in this respect. It is noted that our models are supposed to correspond to the maximal convection. Furthermore, we have found that the profiles of various quantities, e.g. the density profile with $1/r^3$, as well as the shock radius obtained here are in good agreement with the results of detailed simulations for relatively low luminosities of $\\lesssim 6 \\times 10^{52}{\\rm ergs}$ ${\\rm s}^{-1}$. For higher luminosities, however, our results differ from those of the simulations, suggesting that the latter are far from the critical lines for shock revival.  In this paper, we have employed the classical criterion for the convective stability, which is, rigorously speaking, applicable only to the configuration without advection. Recently, \\citet{fog05} gave the modified criterion that takes into account the advection. They introduced the parameter $\\chi$ defined as \\begin{equation} \\chi\\equiv\\int_{r_{\\rm gain}}^{r_{\\rm sh}} \\left|\\frac{N}{u_r}\\right|dr, \\label{eq26} \\end{equation} where $N$ is the Brunt-V$\\ddot{\\rm a}$is$\\ddot{\\rm a}$l$\\ddot{\\rm a}$ frequency given as \\begin{equation} N^2 =\\left\\{\\frac{1}{p}\\left(\\frac{\\partial p}{\\partial S}\\right)_{\\rho,Y_e} \\frac{dS}{dr}+\\frac{1}{p}\\left(\\frac{\\partial p}{\\partial Y_e}\\right)_{\\rho,S} \\frac{dY_e}{dr} \\right\\}\\frac{g}{\\Gamma_1}. \\label{eq27} \\end{equation} In the above expression, $\\Gamma_1$ and $g$ are given by \\begin{equation} \\Gamma_1 =\\left(\\frac{\\partial\\ln p}{\\partial\\ln \\rho}\\right)_{S,Y_e}, \\label{eq28} \\end{equation} and \\begin{equation} g=\\frac{GM}{r^2}. \\label{eq29} \\end{equation} They claimed that the convection occurs in the presence of advection if the entropy-gradient is negative and $\\chi$ is larger than $3$. In order to see the consequence of this modification to our models, we have calculated the values of $\\chi$ for the solutions without convection taken into account. The values of $\\chi$ for our models with self-gravity are shown in figure \\ref{fig8}. We can see that $\\chi$ is always smaller than $3$ for the mass accretion rates larger than $1.0 {M}_{\\odot}{\\rm s}^{-1}$ and that even for the mass accretion rate of $0.1 {M}_{\\odot}{\\rm s}^{-1}$, $\\chi$ is smaller than $3$ except for a narrow range of neutrino luminosity. Thus, taken at face values, the criterion by \\citet{fog05} predicts that the convection does not occur for most of our models. This seems to be at odds with the results of numerical simulations. It is, we think, worth notice that the steady solutions we have obtained here are qualitatively different from the unperturbed states assumed in \\citet{fog05}. \\citet{fog05} assumed simplified formulae for the neutrino heating and equation of state while we employed more realistic ones and, as a result, the post-shock flows are much complicated. Hence we suspect that the critical value of $\\chi$ may be lowered in this case. Indeed, the detailed linear analysis using our steady solutions suggests that the critical value of $\\chi$ may be smaller than $3$ \\citep{yam06}. However, the present models employing the classical criterion will give the upper limit for the effect of convection.  There are a couple of things that have not been considered in this paper but will be important. For example, we have introduced the phenomenological energy- and proton-number fluxes in the convection region, but have neglected the kinetic energy of turbulence. This may not be ignored for the large velocity fluctuations as observed in the numerical simulations. Taken into account, it might increase the critical luminosity because the energy of down-flows will be sucked. It is also pointed out that the non-spherical deformation of the shock front by convection may be an important ingredient that cannot be considered in the present frame work. Thus far we have paid attention only to the condition for the shock revival. This is, however, a minimum requirement. The explosion energy and the mass of neutron star left behind are equally important issues. We are currently attempting to estimate these values based on the models presented here. Our model is also being extended to accommodate rotation, magnetic field, anisotropic neutrino irradiation as well as being utilized for the linear analysis of the non-local instability \\citep{fog02,blo03,ohn06}, whose results will be published in the near future."
    },
    "0606/astro-ph0606732_arXiv.txt": {
        "abstract": "We present a brief review of {\\sl Chandra} X-ray Observatory observations of neutron stars. The outstanding spatial and spectral resolution of this great observatory have allowed for observations of unprecedented clarity and accuracy. Many of these observations have provided new insights into neutron star physics. We present an admittedly biased and overly brief overview of these observations, highlighting some new discoveries made possible by the Observatory's unique capabilities. We also include our analysis of recent multiwavelength observations of the putative pulsar and its pulsar-wind nebula in the IC 443 SNR. ",
        "introduction": "\\label{intro} For many users, the view of the {\\sl Chandra} X-Ray observatory is as pictured in Figure~\\ref{f:fig1}. We therefore provide Figure~\\ref{f:fig2} to remind us that the Observatory, with Inertial Upper Stage (IUS) attached, was the largest and heaviest payload ever launched by NASA's Space Transportation system. The Observatory was launched on 1999 July 23 using the ill-fated orbiter Columbia. The IUS, a two-stage solid rocket booster, was subsequently fired and separated, sending {\\sl Chandra} toward a high elliptical orbit. On August 7, the fifth burn of {\\sl Chandra}'s integral propulsion system placed it into a 63.5-hour (80,800-km semi-major axis) orbit. On August 12, the Observatory's forward door opened, exposing the focal plane to celestial X-rays.  The Observatory was designed for 3 years of operation with a goal of five years. We are therefore extremely pleased that in July of this year the Observatory will have completed its seventh year of operation. Certain difficulties (slow degradation of the thermal shielding and contamination buildup on the Advanced CCD Imaging Spectrometer [ACIS] instrument's cold filters) not withstanding, the Observatory continues to operate successfully. The gas supply, used to point the Observatory, has a lifetime of much more than 15 years. The orbit will be stable for an even longer time period, and we anticipate many more years of usage.  Even the first observations with {\\sl Chandra} provided numerous startling insights concerning astrophysical systems in general, and neutron stars in particular. Two of the early images, those of the supernova remnant Cas A (Fig.\\ \\ref{f:casa}) and that of the Crab Nebula and its pulsar (Fig.\\ \\ref{f:crab}) have become two of the most spectacular and interesting observations made with {\\sl Chandra}.  \\begin{figure} \\centering \\includegraphics[width=8.0cm]{chandra_observation_search_sm.ps} \\caption{{\\sl Chandra} as perceived by an observer today.} \\label{f:fig1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8.0cm]{chandra_cargobay_sm.ps} \\caption{{\\sl Chandra} with IUS attached in Columbia's cargo bay.} \\label{f:fig2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8.0cm]{casa_sm.ps} \\caption{{\\sl Chandra} image of Cas A with the newly discovered point source at the center. Courtesy {\\sl Chandra} X-Ray Center (CXC). } \\label{f:casa} \\end{figure} ",
        "conclusions": ""
    },
    "0606/hep-th0606133_arXiv.txt": {
        "abstract": "\\large The three-year data from WMAP are in stunning agreement with the simplest possible quadratic potential for chaotic inflation, as well as with new or symmetry-breaking inflation. We investigate the possibilities for incorporating these potentials within supergravity, particularly of the no-scale type that is motivated by string theory. Models with inflation driven by the matter sector may be constructed in no-scale supergravity, if the moduli are assumed to be stabilised by some higher-scale dynamics and at the expense of some fine-tuning. We discuss specific scenarios for stabilising the moduli {\\it via} either D- or F-terms in the effective potential, and survey possible inflationary models in the presence of D-term stabilisation. ",
        "introduction": "The dominant paradigm for explaining the great size, age and flatness of the Universe, as well as the absence of unwanted primordial relics such as magnetic monopoles or gravitinos, is cosmological inflation. According to this hypothesis, very early in its history, the Universe underwent an epoch of near-exponential expansion, dubbed cosmological inflation. During this period, the energy density of the Universe is hypothesized to have been dominated by a near-constant vacuum energy density. Within the general context of quantum field theory, the favoured origin of this vacuum energy density is the potential energy $V$ of some elementary scalar field $\\phi$, called the inflaton. Consistency with the magnitude of primordial density perturbations inferred from the pioneering observations by the COBE satellite, as well as subsequent experiments, suggests that this potential energy density $V \\ll m_P^4$.  In order for the inflationary epoch to last long enough, and for the density perturbations to remain small, the inflationary potential must obey certain slow-roll conditions: \\begin{equation} \\epsilon \\; \\equiv \\; \\frac{m_P^2}{2}\\left( \\frac{V^\\prime}{V}\\right)^2, \\; \\eta \\; \\equiv \\; m_P^2 \\left(\\frac{V^{\\prime\\prime}}{V}\\right) \\; \\ll \\; 1. \\label{slowroll} \\end{equation} In principle, both the slow-roll parameters $\\epsilon$ and $\\eta$ may be extracted from the spectra of density perturbations, in particular {\\it via} the scalar spectral index $n_s = 1 - 6\\epsilon + 2\\eta$ and the amplitude ratio of tensor and scalar perturbations $r = 16\\epsilon$. Also observable in principle is the running of the scalar spectral index, $d n_s / d {\\rm ln} k$, but this is expected to be very small in slow-rolling inflationary models.  Important constraints on the slow-roll parameters have been obtained in the past, using the first-year data from WMAP and other observational inputs. Qualitatively new insight has recently been provided by the new three-year WMAP data (WMAP3), which are strikingly consistent with simple models of inflation based on an elementary scalar inflaton field whose potential is just a low-order polynomial~\\cite{Spergel:2006hy} {}\\footnote{Note that this statement is also supported by the WMAP3 data analysis given in~\\cite{Kinney:2006qm}, which yields a larger allowed region.}. In particular, the data are consistent with the simplest model for chaotic inflation with a quadratic potential $V(\\phi)=\\frac{1}{2}m^2\\phi^2$, as seen in Fig.~\\ref{pro:fig}. They also constrain severely possible modifications, e.g., of the form: \\beq \\label{pro:pot1} V(\\phi) = \\frac{1}{2}m^2\\phi^2\\left(1+\\frac{1}{2}\\kappa_c^2\\phi^2\\right) . \\eeq As seen in the right panel of Fig.~\\ref{pro:fig}, the supplementary parameter $\\kappa_c$ cannot exceed a few per mille, with the exact number depending of the number of e-foldings during the inflationary epoch. As seen in Fig.~\\ref{pro:fig}, the WMAP3 data~\\cite{Spergel:2006hy} are also consistent with another very simple potential: \\beq \\label{pro:pot2} V(\\phi) = \\lambda m_P^4 \\left( 1-\\kappa_s^2\\frac{\\phi^2}{m_P^2}\\right)^2 , \\eeq with the requirement, seen also in the right panel of Fig.~\\ref{pro:fig}, that $\\kappa_s^2$ cannot exceed one per cent if inflation is to start with the small initial value of the inflaton field~\\footnote{Note, however, that this potential also allows inflation with a large value of the inflaton field.}, i.e., on the branch $|\\phi|\\leq m_P/\\kappa_s$. This scenario is known as new or symmetry-breaking inflation. These and other models make characteristic predictions for the spectral index of scalar density perturbations, $n_s$, and for the tensor-to-scalar ratio $r$. The left and central panels of Fig.~\\ref{pro:fig} display the constraints on $n_s$ and $r$ inferred from a combination of WMAP3~\\cite{Spergel:2006hy} with other data~\\cite{Seljak:2006bg}. These are confronted with the different predictions of the potentials given in (\\ref{pro:pot1}) and (\\ref{pro:pot2}), assuming negligible running of the scalar spectral index, as predicted in these models.  \\begin{figure} \\hspace{-1.0cm} \\includegraphics*[height=7cm]{inflation1.ps} \\hspace{0.5cm} \\includegraphics*[height=7cm]{inflation2.ps} \\hspace{0.5cm} \\includegraphics*[height=7cm]{inflation3.ps} \\caption{\\it Confrontation of low-order polynomial inflationary potentials with the three-year data from WMAP~\\protect\\cite{Spergel:2006hy}, in combination with other data~\\protect\\cite{Seljak:2006bg}. The deviation from scale invariance of the scalar spectral index, $1-n_s$, is shown as a function of the tensor-to-scalar ratio $r$ for chaotic inflation (left) and new inflation (centre). The lines correspond to $N=47,50,53,56,59$ e-foldings, as indicated. The green (light gray) lines denoted by $\\mathsf{a}-\\mathsf{d}$ correspond to $\\kappa_c^2=10^{-4},10^{-3},10^{-2},10^{-1}$ in (\\ref{pro:pot1}), and the green (light gray) lines denoted by $\\mathsf{e},\\mathsf{f}$ correspond to $\\kappa_s^2=10^{-2},10^{-3}$ in (\\ref{pro:pot2}).  The dashed curves show the regions of the parameters allowed by WMAP3 and other data~\\protect\\cite{Seljak:2006bg} at the 95\\% CL. Right panel displays the corresponding 95\\% upper limits on $\\kappa_c^2$ and $\\kappa_s^2$ as functions of the numbers of inflationary e-foldings in the chaotic and new inflationary models, respectively. \\label{pro:fig}} \\end{figure}  It has long been a challenge to embed inflation in realistic models of particle physics. There are two aspects to this problem: identifying a suitable candidate for the inflaton, and ensuring that it has a suitable effective potential. We take the point of view that realistic particle models should incorporate supersymmetry~\\cite{ENOT}, which incidentally provides many scalar fields that one might hope to exploit as an inflaton. None of the MSSM fields fits the bill, but the scalar partner of one of the heavy singlet neutrinos in a seesaw model of neutrino masses might be suitable~\\cite{MSYY}. In the framework of global supersymmetry, such a sneutrino would naturally possess a simple quadratic potential without significant perturbative corrections, and its mass would fit naturally with the estimate of $m \\sim 2 \\times 10^{13}$~GeV required for the simplest chaotic inflation model (\\ref{pro:pot1}).  However, a more appropriate framework for describing inflation is presumably that offered by local supersymmetry. The low-energy limit of any supersymmetric theory of quantum gravity is necessarily some such supergravity theory. Moreover, since inflation is a fundamental property of the evolution of space-time, it must involve the gravitational sector of the microscopic theory, in which case supergravity is likely to play an essential role. The support of the data for the simplest inflationary models is then highly non-trivial information, since the effective potentials yielded by supergravity theories generically have forms that are more complicated than (\\ref{pro:pot1}) or (\\ref{pro:pot2}). Indeed, many supergravity potentials have minima with vacuum energy densities that are negative and ${\\cal O}(m_P^4)$.  The general term for this troublesome feature of supergravity is the $\\eta$ problem, namely the statement that supergravity with simple forms for the K\\\"ahler potential for scalar superfields, such as minimal supergravity, naturally predict large mass terms for all the scalars, destroying the flatness of any candidate inflationary direction in the effective potential. One possible way out of this problem is to postulate a K\\\"ahler potential whose prefactor $e^K$ dependens on the scalar fields in a manner softer than an exponential. The non-minimal form of the K\\\"ahler potential which is particularly well motivated from the point of view of string theory is the no-scale structure that appears naturally in compactifications of higher-dimensional superstring models. The great advantage of the no-scale structure is that it leads to a semi-positive definite scalar potential, very similar to a globally supersymmetric one~\\cite{Ellis:1983sf}.  Many superstring theories contain moduli fields $T_i$ whose effective potentials are flat classically. These no-scale models are characterized by effective low-energy supergravity theories containing terms in the K\\\"ahler potential $K$ of the characteristic logarithmic form: $K \\ni - \\log(T_i + \\bar{T}_i)$. Since the effective potentials of these moduli fields vanish classically, one might expect them to remain small even when corrections are calculated, so that they might be candidates for the inflaton field. Efforts to implement modular inflation continue and provide interesting and encouraging results. However, here we prefer to explore the other possibilities offered by the matter sectors of string theories, which are typically rather rich in the structures of their superpotentials and K\\\"ahler potentials, potentially offering ample opportunities to generate inflationary behaviour.   Specifically, in this paper we examine in some detail conditions necessary for matter-driven inflation to take place in a locally supersymmetric model. For simplicity, we consider here a no-scale model with a single modulus $T$ and a matter field $\\Phi$, with a no-scale K\\\"ahler potential of the form: \\begin{equation} K \\; = \\; - 3 \\log (T + \\bar{T} - 2 | \\Phi |^2) \\label{simpleK} \\end{equation} and a matter superpotential to be discussed below. We find that the matter sector in such a simplified no-scale supergravity may indeed lead to acceptable inflation, provided that the $T$ modulus is stabilised appropriately. One must also apply specific fine-tuning conditions on the inflationary superpotential, which depend on the precise form of the constraint on the modulus $T$. In at least some such models, the naturalness of these fine-tuning requirements is not obvious from the point of view of the effective low-energy theory.  We then go on to consider explicit models for the stabilization of the modulus $T$. One is D-type stabilisation made possible by supersymmetric D-terms, and the other is F-type stabilisation, where the superpotential depends on a $T$ modulus. As we show below, model building is much easier in the case of D-stabilisation, as the scales in the sector of the matter-type inflaton $\\Phi$ and in the $T$-modulus sector can decouple from each other. ",
        "conclusions": "We have investigated in this paper several possible embeddings of simple inflationary potentials in supergravity theories, in manners consistent with the recent WMAP3 data. We have concentrated on models with the no-scale structure K\\\"ahler potential in the inflaton sector, so as to avoid the $\\eta$ problem and to take the advantage of the semipositivity of the potential which holds in the exact no-scale structure limit of the models that  we have considered. We have found that the no-scale matter sector may lead to acceptable inflation, however  the modulus $T$ must be stabilized by some additional mechanism which does not spoil the no-scale structure too violently, and some apparent fine-tuning is necessary.  We have considered  D-type stabilisation which in the leading approximation keeps intact  the semi-positivity of the potential, and F-type stabilisation, which breaks this feature. With D-stabilisation, the model building becomes unexpectedly  simple, as the scales in the $\\Phi$-sector (by $\\Phi$ we denote the matter-type inflaton) and in the $T$-modulus sector decouple. Depending on the assignments of charges and the specific form of the K\\\"ahler potential for charged fields, both chaotic and symmetry-breaking inflation are possible in this case. In particular, slow-roll inflation starting from the saddle points in the domain of sub- or near-Planckian field strengths With F--stabilisation, the post--inflationary minimum is naturally of the AdS type but, after suitable uplifting of the vacuum energy, the structure necessary for  an inflationary epoch may appear. However, with a brane-antibrane-type potential for the modulus, the requirement that the uplifted vacuum energy is zero curves the potential too strongly to allow for inflation, at least if one stays with a reasonable level of fine tuning between expectation values of the modulus and the matter scalar, say up to one part per mille. In fact, in all the cases when $\\Phi$ inflation takes place, an interesting fine-tuning between the parameters of the superpotential for the $\\Phi$ field and the parameters of the sector which stabilises the $T$ modulus is necessary.  In summary, we have demonstrated that inflation from matter fields is a nontrivial feature of supergravity models, not only because of the $\\eta$ problem, but also because  it is subtly intertwined with the issue of moduli stabilisation. Nevertheless, models with inflation driven by the matter sector may be constructed in no-scale supergravity, if the moduli are stabilised by some higher-scale dynamics and at the expense of a certain level of fine-tuning.  \\bigskip  \\centerline{\\Large \\bf Acknowledgements}  \\vspace*{0.5cm}  The authors thank D. Langlois and V. Mukhanov for helpful discussions. The visit of D.L. and V.M. to ITPUW was supported by the TOK Project  MTKD-CT-2005-029466 ``Particle Physics and Cosmology: the Interface''. Z.L. thanks the CERN Theory Division for hospitality. This work was partially supported by the EC 6th Framework Programme MRTN-CT-2004-503369, by the Polish State Committee for Scientific Research grant KBN 1 P03D 014 26 and by POLONIUM 2005. K.T.\\ acknowledges support from the Foundation for Polish Science (Programme START).  \\vspace*{.5cm}"
    },
    "0606/astro-ph0606054_arXiv.txt": {
        "abstract": "% {The metal-rich bulge globular cluster NGC~6388 shows a distinct blue horizontal-branch tail in its colour-magnitude diagram (Rich et al. 1997) and is thus a strong case of the well-known {\\it 2$^{nd}$ Parameter Problem}. In addition, its horizontal branch (HB) shows an upward tilt toward bluer colours, which cannot be explained by canonical evolutionary models. Several noncanonical scenarios have been proposed to explain these puzzling observations. In order to test the predictions of these scenarios, we have obtained medium resolution spectra to determine the atmospheric parameters of a sample of the blue HB stars in NGC~6388. Using the medium resolution spectra, we determine effective temperatures, surface gravities and helium abundances by fitting the observed Balmer and helium lines with appropriate theoretical stellar spectra. As we know the distance to the cluster, we can verify our results by determining masses for the stars. During the data reduction we took special care in subtracting the background, which is dominated by the overlapping spectra of cool stars. The cool blue tail stars in our sample with effective temperatures \\teff\\ $\\approx$ 10,000~K have lower than canonical surface gravities, suggesting that these stars are, on average, ${\\approx}$0\\fm4 brighter than canonical HB stars in agreement with the observed upward slope of the HB in NGC~6388. Moreover, the mean mass of these stars agrees well with theoretical predictions.  In contrast, the hot blue tail stars in our sample with \\teff\\ $\\ge$ 12,000~K show significantly lower surface gravities than predicted by any scenario, which can reproduce the photometric observations.  Their masses are also too low by about a factor of 2 compared to theoretical predictions. The physical parameters of the blue HB stars near 10,000~K support the helium pollution scenario. The low gravities and masses of the hot blue tail stars, however, are probably caused by problems with the data reduction, most likely due to remaining background light in the spectra, which would affect the fainter hot blue tail stars much more strongly than the brighter cool blue tail stars.  Our study of the hot blue tail stars in NGC~6388 illustrates the obstacles which are encountered when attempting to determine the atmospheric parameters of hot HB stars in very crowded fields using ground-based observations.  We discuss these obstacles and offer possible solutions for future projects. ",
        "introduction": "\\label{sec:intro} Ever since its discovery over 30 years ago (Sandage \\& Wildey \\cite{sawi67}; van den Bergh \\cite{vdbe67}), the 2$^{\\rm nd}$ parameter effect has stood as one of the major unsolved challenges in the study of the Galactic globular clusters.  While it was recognized quite early that the horizontal branch ({\\bf HB}) becomes redder on average with increasing metallicity, many pairs of globular clusters are known with identical metallicities but markedly different HB morphologies, e.g., M~3 versus M~13.  Thus some parameter(s) besides metallicity (the 1$^{\\rm st}$ parameter) must affect the evolution of the HB stars in these globular clusters. Possible 2$^{\\rm nd}$ parameter candidates include the globular cluster age, mass loss along the red-giant branch ({\\bf RGB}), helium abundance $Y$, $\\alpha$-element abundance, cluster dynamics, stellar rotation, deep mixing, etc.  Hubble Space Telescope observations by Rich et al. (\\cite{riso97}) have found an unexpected population of hot HB stars in the metal-rich globular clusters NGC~6388 and NGC~6441 ([Fe/H] $\\approx-0.5$), making these clusters the most metal-rich clusters to show the 2$^{\\rm nd}$ parameter effect.  Ordinarily metal-rich globular clusters have only a red HB clump.  However, NGC~6388 and NGC~6441 possess extended blue HB tails containing ${\\approx} 15$\\% of the total HB population.  Quite remarkably, the HBs in both clusters slope upward with decreasing $B - V$ with the stars at the top of the blue tail being nearly 0\\fm5 brighter in $V$ than the well-populated red HB clump. Moreover, the RR Lyrae variables in these clusters have unusually long periods for the cluster metallicity, leading Pritzl et al. (\\cite{prsm00}) to suggest that NGC~6388 and NGC~6441 may represent a new Oosterhoff group.  For all of these reasons the HBs of NGC~6388 and NGC~6441 are truly exceptional.  In the next section we will discuss the implications of NGC~6388 and NGC~6441 for the 2$^{\\rm nd}$ parameter effect and will review a number of scenarios for explaining the HB morphology of these clusters. Sect.~\\ref{sec:obs} describes the observations and reduction of the medium resolution spectra that we have obtained to test these scenarios, while Sect.~\\ref{sec:analysis} gives the atmospheric parameters derived from these spectra.  These results are compared with the theoretical HB tracks in Sect.~\\ref{sec:theory}. Sect.~\\ref{sec:conclusions} summarizes our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Our results for the cool blue tail stars in NGC~6388 are consistent with both the upward slope of the HB in NGC~6388 and the predictions of the helium pollution scenario.  Moreover, they resolve the puzzling conundrum discussed by Moehler et al. (\\cite{mosw99}), who found higher than canonical surface gravities for a sample of blue tail stars in NGC~6388 and NGC~6441 in contradiction to all scenarios for explaining the upward sloping HBs in these globular clusters.  Most likely, these large gravities can be attributed to problems with the background subtraction from the spectra.  Our results for the hot blue tail stars remain problematical.  The low surface gravities and masses found for these stars might be due to problems with the background subtraction or possibly to inadequacies in the stellar atmospheres used in the analysis.  Further work is needed to clarify this point.  Unfortunately the atmospheric parameters for the hot blue tail stars do not permit a stringent test of the helium pollution scenario.  In particular, the present data cannot test the prediction for a greater gravity offset from the canonical ZAHB in the hot blue tail stars compared to the cool blue tail stars, as one would expect if the helium abundance increased along the blue tail.  The present study illustrates the difficulties encountered when doing spectroscopy in crowded fields.  From our experience we conclude that the problem of the hot stars in NGC~6388 cannot be solved by means of current ground-based observations.  One would need either optical spectra at a {\\em much} better spatial resolution (around 0\\farcs1) or multi-colour photometry including the near and far UV.  Spectroscopy at such high spatial resolution is currently available only at near-IR wavelengths, which would worsen the contrast between the hot and cool stars and, in addition, would allow only the use of hydrogen lines, for which the model atmospheres have not been tested\\footnote{IR observations of hot stars are usually done for massive stars, not for evolved low-mass stars}. Multi-colour photometry might provide an estimate of the effective temperatures and luminosities of these stars, but as they are hot, the satellite UV is definitely required."
    },
    "0606/astro-ph0606438_arXiv.txt": {
        "abstract": "{We combine bright {\\xmm} data  with the \\chandra Deep Field South observations in order to explore the behavior of the intrinsic AGN absorption, as a function  of redshift and luminosity. Our sample consists of 359 sources selected in the hard 2-8 keV band, spanning the flux range $6\\times 10^{-16}$-$3\\times10^{-13}$ \\funits with a high rate of spectroscopic or photometric redshift completeness (100 and 85 per cent respectively for the {\\chandra} and {\\xmm} data). We derive the column density values using X-ray spectral fits. We find that the fraction of obscured AGN falls with increasing luminosity in agreement with previous findings. The fraction of obscured AGN shows an apparent increase at high redshifts ($z>2$). Simulations show that this effect can be most probably attributed to the fact that at high redshifts the column densities are overestimated. { ",
        "introduction": "Ultra-deep surveys with {\\chandra} have resolved the bulk of the X-ray background at hard energies (2-10\\,keV), shedding light on the nature of the AGN population at faint fluxes  ($f_X(2-8~\\rm keV) \\sim 1.4\\times10^{-16}$ \\funits; Mushotzky et al. 2000, Giacconi et al. 2002, Alexander et al. 2003). These surveys have revealed luminous unobscured QSOs but also more importantly, large numbers of obscured AGN ($N_H>10^{22}$ \\cunits). The fraction of absorbed sources rises steeply with decreasing flux (e.g. Alexander et al. 2003) dominating the X-ray population at the flux limit of the 2\\,Ms \\chandra Deep Field North (CDF-N) survey (e.g. Perola et al. 2004). To the first approximation, this trend is in qualitative agreement with the predictions of the X-ray background (XRB) population synthesis models (e.g. Comastri et al. 1995; Gilli et al. 2001). A more detailed quantitative comparison however, reveals a number of inconsistencies. For example, the models above also predict large numbers of luminous ($L_X > 10^{44} \\rm erg \\, s^{-1}$) heavily obscured ($N_H>10^{22} \\, \\rm cm^{-2}$) type-II QSOs, although only a handful of such objects have been identified todate, despite painstaking efforts (e.g. Barger et al. 2003; Fiore et al. 2003; Szokoly et al. 2004). Parallel to the \\chandra deep fields, surveys with the {\\xmm}, probing on average brighter fluxes, also provide a wealth of complementary information on the nature and the  evolution of the AGN population. These brighter surveys also show a clear scarcity of obscured AGN relative to the model expectations (e.g. Piconcelli et al. 2002, Georgantopoulos et al. 2004, Perola et al. 2004).  Although the source of the inconsistency between observations and model predictions remains unclear, two main scenarios are put forward to reconcile the problem. The first  argues that observational biases are affecting our conclusions, while the second proposes major revision of the basic assumptions of the XRB models.  In the former case, it is proposed that obscured AGN are present in current surveys but lie in poorly explored regions of the parameter space. For example a large fraction of X-ray sources, particularly in the \\chandra deep fields (about 30\\%; Alexander et al. 2003) are optically faint ($R > 24.5$\\,mag) and therefore, hard to study in detail. Treister et al. (2004) argue that it is precisely this poorly explored population that comprises a large fraction of obscured AGN at high redshift. This is because although the energetic X-ray photons of these systems can penetrate the large obscuring columns of gas and dust, the rest-frame UV/optical light is heavily extincted resulting in faint observed optical magnitudes. Consequently the type-II AGN are difficult spectroscopic targets, even with 10-m class telescopes, and their intrinsic properties (including their $N_H$) remain ill constrained.  In the case of the model revision scenario, one of the key input parameters to the XRB models is the intrinsic column density ($N_H$) distribution of AGN. This is assumed to be that of Seyfert galaxies (i.e. low-luminosity AGN) measured from the local Universe (Risaliti et al. 1999). It is possible that the luminous AGN identified in deeper surveys do not follow the same column density distribution, leading to the failure of the population synthesis models. Indeed, Ueda et al. (2003) by combining sources with redshift information from the CDF-N and  brighter {\\asca} surveys find that the fraction of obscured AGN diminishes with increasing luminosity. The physical interpretation of this model could be that the highly luminous AGN blow away the obscuring screen or they photoionize the gas  around them.  More recently, La Franca et al. (2005) estimated the fraction of absorbed AGN as a function of luminosity combining data from the CDF-N and CDF-S as well as  bright {\\xmm} data. They also suggest a decrease of the obscured AGN fraction (hereafter F) with luminosity as well as a possible increase with redshift, in contrast to the standard model assumption that F is independent of $L_X$ or $z$. Such studies are obviously important for furthering our understanding of the AGN unification models and  moreover, they are crucial for the construction of X-ray background population synthesis models.  In this paper we attempt to further investigate the intrinsic absorption in AGN, using a large X-ray sample from {\\xmm} and {\\chandra} all with X-ray spectroscopic information. In particular, we combine the  {\\chandra} Deep Field South, CDF-S observations (162 sources, {\\it all}  with redshift information, from  Szokoly et al. 2004 and Zheng et al. 2004), with {\\xmm} data at brighter fluxes ($>3\\times10^{-14}$ \\funits), from  the HELLAS2XMM survey of Perola et al. (2004) (44 sources) and our own {\\xmm} survey which overlaps  with the Sloan Digital Sky Survey, SDSS (153 sources). The advantages of our approach are the following:   $\\bullet$ We have derived X-ray spectra for all the X-ray sources in our sample, trying to avoid the uncertainties that can be introduced by a simple X-ray colour.  $\\bullet$ We estimate the fraction of obscured AGN using the $1/V_{m}$ method. In this way we properly take into account the 'accessible' volume of each source, i.e. the fact that many obscured sources are not detected at a given flux limit.  $\\bullet$ We use the CDF-S, instead of the CDF-N, as the former has photometric or spectroscopic redshifts for all the sources. This means that there are no missing type-II AGN at high redshifts which would introduce a spurious correlation between obscured fraction and luminosity.  Throughout this paper we adopt H$_{\\circ}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$=0.3 and $\\Lambda$=0.7 ",
        "conclusions": "In this paper we use the largest {\\xmm} sample with X-ray spectroscopic information available in order to investigate the  behavior of intrinsic absorption in AGN as a function of redshift and luminosity. We also take  advantage of the complete optical coverage (photometric or spectroscopic) of the CDF-S observations in order to extend our results at fainter fluxes (down to $\\sim6\\times 10^{-16}$ \\funits).  In Fig. \\ref{vmaxl} we  plot the fraction of obscured objects as a function of luminosity separately for the {\\xmm} (upper panel), the {\\chandra} data (middle panel) and their combination (lower panel). Despite the weak indications for a decline of this fraction in the individual samples, the only way to unambiguously verify this correlation is to increase the coverage of the L-z plane by combining the two samples. It is likely that the rapid decline appearing in the combined sample could not be easily observed in the {\\xmm} and the {\\chandra} individual datasets due to the limited available volume of each independent, flux limited survey. For example {\\chandra} does not cover a large enough volume to sample a large number of luminous sources. However, we caution that the combination of these different subsamples may introduce some bias in favor of an $F-L_X$ correlation. Indeed, as we move toward higher luminosities we sample more {\\xmm} sources and less {\\chandra} sources in each luminosity bin according to Fig. \\ref{lz}. But the {\\xmm} sources are found in much brighter fluxes in comparison with the {\\chandra} ones. Thus, as we move progressively to higher luminosities we sample higher fluxes and hence less obscured sources according to the well known flux-absorption correlation. This effect is summarized in Table \\ref{medianf2} were we list the median 2-8 keV flux for each luminosity bin presented in Fig. \\ref{vmaxl}.  The decrease  of the fraction of obscured sources as a function of luminosity is consistent with the results of La Franca et al. (2005). The physical interpretation of this model could be that the radiation pressure flattens the torus in luminous objects (K\\\"{o}nigl \\& Kartje 1994) or increases the degree of photoionization of the gas around them. Another possible scenario is that of the 'receding torus' which has been proposed by Lawrence (1991) and has been recently updated by Simpson (2005), where because of the effects of dust sublimation the inner radius of the torus increases with luminosity.  One way to explore whether the relation between the absorbed fraction and luminosity ($F-L_X$ relation) is real (or is induced up to some degree by the strong flux-absorption correlation) is to model the number of absorbed sources as a function of flux. In Fig. \\ref{xrbmodel} we plot the fraction of obscured AGN, F, as a function of the flux. We plot separately the {\\xmm} points, the CDF-S points as well those from the {\\chandra} survey in the Extended Groth Strip (Georgakakis et al. 2006b) and compare with various model predictions. In the above models we use the Ueda et al. (2003) luminosity function. The solid line gives the predictions of our model and uses the $F-L_X$ relation derived here. The long dashed line corresponds to a model with a ratio of obscured to unobscured AGN, R=4 (or F=0.8) with no dependence on luminosity. This is the ratio derived in the local Universe by Risaliti et al. (1999) and Maiolino \\& Rieke (1995). The short dashed line corresponds to the R=1 case with no dependence on luminosity. The R=1 model nicely represents the {\\xmm} data while the R=4 model follows better the faint CDF-S data. However, only a model which includes a decrease in the $F-L_X$  relation can explain the abrupt increase of the fraction of absorbed sources with decreasing flux.  Previous estimates are in agreement with our results. Piconceli et al. (2003), analyzing hard X-ray {\\xmm} data, claimed that the observed fraction of obscured sources at bright fluxes ($>10^{-13}$ \\funits) is about 30 per cent, much lower than that predicted by the (R=4) XRB model.  La Franca et al. (2005), combining data from different X-ray samples, studied the behavior of obscuration in a much wider flux range. Their findings for the fraction of obscured sources are very similar to ours in both faint and bright fluxes. Their proposed $F-L_X$ relation is in agreement with our model (see Fig. \\ref{vmaxl}). Note that here we have used a power-law best fit model instead of a linear one.  Fig. \\ref{red_model} shows the dependence of $F$ on redshift. The fraction shows an apparent increase with redshift, with a more abrupt increase at high redshifts, $z>2$. Ueda et al. (2003) find no dependence of the fraction $F$ on redshift. Ballantyne et al. (2006) propose that the obscuration may be related to star-formation within the host galaxy and thus there may be some increase of the obscured AGN fraction  with redshift. They test their models by comparing with the observed type-I AGN in the CDF-N (Barger et al. 2003). These authors find an obscured  fraction F evolving as $(1+z)^{0.3}$, together with a dependence of F on luminosity. La Franca et al. (2005) find a dependence on redshift very similar with ours (see their Fig. 6 right panel): the fraction $F$ increases from $\\sim0.2$ at low redshift to $F\\sim0.6$ at $z>2$. In our case however, our simulations show that the lowest energy bins fluctuations can introduce a significant artificial correlation of the fraction F with redshift. This effect becomes particularly important at high redshifts due to the K-correction effect. This suggests that the observed fraction of obscured sources at high $z$ is erroneously  enhanced. Consequently there is no significant evidence for an increase in the fraction F with redshift.   \\begin{figure} \\centering \\includegraphics[width=8.5cm]{XRB_model.ps} \\caption{The fraction of the sources with $N_H>10^{22}$ cm$^{-2}$ for the {\\xmm} sample (open circles) and the {\\chandra} sample (filled circles). The triangles represent the EGS \\chandra survey of Georgakakis et al. (2006b). The solid line gives the  predicted fraction using the $F-L_X$ relation derived here. The long and short dashed lines give the models with R=4 and R=1 ratios of obscured to unobscured AGN respectively with no dependence on luminosity.} \\label{xrbmodel} \\end{figure}"
    },
    "0606/astro-ph0606112_arXiv.txt": {
        "abstract": "{}{Galactic globular cluster (GC) stars exhibit abundance patterns which are not shared by their field counterparts, e.g.  the well documented O-Na and  Mg-Al anticorrelations. A widely held  hypothesis is that the gas was ``polluted\" by stars more massive than the presently observed low-mass stars. In the framework of this ``self-enrichment\" scenario for GCs, we present a new method to derive the Initial Mass Function (IMF) of the polluters, by using the O/Na abundance distribution.} {We focus on NGC~2808, a GC for which the largest sample of O and Na abundance determinations  is presently available. We use the abundance distribution of [O/Na] to derive the amount of polluted material with respect to that of original composition. We explore in details two scenarios for the self-enrichment of the cluster, which differ by the assumptions made on the composition of the polluter ejecta. In each case we consider two classes of possible ``culprits\" : massive Asymptotic Giant Branch (AGB) stars (4-9~\\ms)  and winds of massive stars (WMS) in the mass range 10-100~\\ms.} {We obtain upper limits for the slope of the IMF (assumed to be given by a power-law) of the stars initially more massive than the present turnoff mass. We also derive lower limits for the amount of stellar residues in NGC~2808.} {We find that the polluter IMF had to be much flatter than presently observed IMFs in  stellar clusters; this is  in agreement with the results of two other methods for GC IMF determination, which we also discuss. Additionaly, we find that the present mass of the GC should be totally dominated by stellar remnants if the polluters were AGB stars; this is not the case if the culprits are WMS. We critically analyse the advantages and shortcomings of each potential polluter class, and we find the WMS scenario more attractive.} ",
        "introduction": " ",
        "conclusions": "Today, there is not a unique, neither a well-defined, nor even a fully self-consistent scenario for the self-enrichment of globular clusters. The most commonly invoked one assumes that the H-processed ejecta of massive AGB stars of a first stellar generation (of unspecified, but presumably flatter than normal, IMF) form a second stellar generation (of equally unspecified IMF) and of appropriate chemical composition (i.e. in agreement with observed  O/Fe vs Na/Fe ratios in GC stars). Most of the theoretical effort up to now has been devoted in trying to reproduce the observed O vs Na or Mg vs Al relations from the nucleosynthesis point of view (see references in Charbonnel 2005).  In this section, we discuss several aspects of the two self-enrichment Scenarios (I and II) which are sketched in \\S~3 and which involve not only AGBs but also winds of massive stars as alternative sources of the chemically polluted material.  \\subsection{The mass and composition of the polluting ejecta}  We start by noting that any scenario of that kind necessarily involves assumptions about three quantities:  a) the IMF (mass limits and shape) of the first stellar generation and, when it exists, of the second one also (there are two generations in our Scenario I but only one in our Scenario II);  b) the amount and composition of the stellar ejecta that contribute to the formation of the contaminated stars, as well as the mass range of the polluting stars;  c) the amount of unprocessed gas, which is mixed to various degrees with the processed ejecta (it is equal to zero in our Scenario I).  From those quantities, only the second can be evaluated, in principle, from stellar nucleosynthesis calculations; the other two can only be guessed at, or constrained, by a combination of theory and observations.  Since we lack, at present, the required input for item (b)\\footnote{As mentioned in \\S~3, the yields of massive AGB stars have been computed recently by several groups in the context of the self-enrichment hypothesis. Nevertheless, the predictions remain highly uncertain (mainly because they depend on the modelling of complex physical mechanisms such as HBB, third dredge-up, convection, mass loss, rotation, ...) and none of the present models does properly account for the O-Na and Mg-Al anticorrelations yet. On the other hand, models of rotating massive stars with the appropriate initial composition will soon be available (Decressin et al. in preparation).} we had to make here some assumptions. We assumed then that the {\\it amount of H-processed ejecta} is the maximum possible from both potential sources, AGBs or massive stars (Eqs. 12, 13 and Fig.~\\ref{residueejecta}). Obviously, any other assumption reducing the amount of the processed ejecta would result in IMFs even flatter than those found in \\S~4 (all other things being kept the same). As discussed already in \\S~4.2, a smaller amount of H-processed ejecta than assumed here is much more plausible physically, for both AGB stars and  massive stars.  We also made two different assumptions about the {\\it composition of the ejecta}, namely that they are processed to various degrees (Scenario I) or to an extreme degree (Scenario II). Obviously, the latter assumption is utterly implausible, and it was used only to place absolute  lower limits to the slope of the polluter IMF. Indeed, it is hard to conceive that the {\\it whole envelope of all the polluters} (AGBs or massive stars) is processed to such an extreme degree ([O/Na]=-1.4); this may happen at some point of the evolution of some of the polluters, but it cannot characterize to totality of the envelope evolution, nor the whole mass range of the polluters. Thus, the lower horizontal line of Scenario II (middle panel of Fig.~\\ref{IMFgeneral}) is an extreme and unrealistic lower limit; the upper line of Scenario I is more realistic in that respect. Still, observations suggest that it cannot be true, either.  Indeed, Pasquini et al. (2005) report the presence of Li  in 9 turn-off stars of the intermediate metallicity globular cluster NGC~6752 ([Fe/H]=--1.4), a cluster which displays the anticorrelation of O vs. Na. The Li abundance in NGC~6752 correlates with the one of O and anticorrelates with those of Na and N. The maximum Li abundance corresponds to stars with [O/Fe]$\\sim$0.4, i.e. to stars formed from unprocessed material with composition similar to the one of field stars. It is no surprise then that the abundance of Li in those stars (log(Li/H)$\\sim$--9.6) is at the level of the so-called ``Spite plateau\" observed in field halo stars (e.g., Charbonnel \\& Primas 2005). On the other hand, the lowest detected Li values are about 3 times smaller and are found in stars with very low [O/Na]$\\sim$--1.1 in NGC~6752.  The presence of so much Li (only within a factor of three below the plateau level)  in stars of modified composition in NGC~6752 clearly is at odds with our Scenario I: if such stars are made exclusively from the H-processed ejecta of some class of polluters (as assumed in our Scenario I), then no Li should be present there, in view of the nuclear fragility of that element. And even if one assumes that some fresh Li can be produced in H-burning (and ejected before subsequent destruction) through the ``Cameron-Fowler (1971) mechanism\", it is utterely implausible that the abundance of this new Li is almost as high as the ``Spite plateau\". The most natural way to explain the presence of Li in the O-depleted stars of NGC~6752 is by assuming that those stars are formed from a mixture of nuclearly processed material  with material of normal  composition, as in our Scenario II (Prantzos \\& Charbonnel, in preparation).  Finally, we note that in both Scenarios I and II we assumed that, while the original IMF contained stars in the full mass range of 0.1 to 100 \\ms, the ejecta of the polluters formed only stars in the 0.1-0.8 \\ms \\ range, i.e. still alive  today. This assumption has no physical justification and was only made in order to maximize the ratio of ejecta to live stars, in the same spirit with the other assumptions discussed in this section. Assuming that the ejecta formed stars in the full 0.1-100 \\ms \\ range would obviously lead to flatter IMFs in order to satisfy the observational requirements\\footnote{Assuming that the second (``polluted\") generation of stars contained its own massive and fast evolving ``polluters\" (the ejecta of which would give rise to a third generation, and so on) would make the whole problem much more complex, and it would invalidate the analysis made here.}.  We conclude then that both our Scenarios I and II are grossly underestimating ``reality\", in the framework of the self-enrichment hypothesis for globular clusters. A realistic scenario should invoke both ejecta with varying composition (as in Scenario I, more realistic on nucleosynthesis grounds) {\\it and} mixture of those ejecta with pristine gas (as in Scenario II, to account for Li observations); this implies that the horizontal ``observational\" lines in Fig.~\\ref{IMFgeneral} (middle panel) should be replaced by a line lying above the one of Scenario I.  Also, a realistic scenario should consider an amount of ejecta considerably lower than  the (deliberate) overestimate we made in Eqs. (12) and (13); consequently, the curves of Fig.~\\ref{IMFgeneral}  (middle panel) should shift downwards, as discussed in \\S~4.3. Overall, in a realistic scenario the slope $X$ of the polluter IMF should be much smaller than  derived in \\S~4.3 and certainly much smaller than 0 in case the polluters are AGB stars. In other terms, an extremely flat IMF should be necessary to satisfy the constraints imposed by the abundance distribution of Fig.~\\ref{histogramme}.  In such a case, by extrapolating the trends obtained in the lower panel of Fig.~\\ref{IMFgeneral} (towards the left) one sees  that:  - if the polluters are AGB stars, the mass of the cluster today should be completely dominated by residues (by a factor of several);  - on the contrary, if the polluters are massive stars, the total mass of the residues today would be much smaller than the mass of live stars in the cluster.  With those implications in mind, we now take a closer look at the self-enrichment scenario, for each one of the two invoked classes of polluters.  \\subsection{Polluters: AGBs vs. WMS}  \\subsubsection{AGBs?}  According to a commonly held idea, the peculiar chemical composition of GC stars results from pollution by the ejecta of AGB stars, massive enough (4-9 \\ms) as to develop HBB in their envelopes. Such stars evolve in time scales of $\\sim$ 20 to 150 Myr and release their ejecta at velocities $v<30$ km/s (e.g., Loup et al. 1993), lower than the escape velocity from the cluster; the ejecta are trapped in the cluster and form new stars, either directly (our Scenario I) or indirectly, after mixing to various degrees with pristine material (our Scenario II, which accounts for the Li observations, as argued in Sec. 5.1).  In our opinion, this idea suffers from several drawbacks (in addition to the problematic nucleosynthesis predictions of the current models discussed in \\S~3). In particular, it gives no satisfactory answer as to the roles of stars more massive (M$>$9 \\ms) and less massive (M$<4$ \\ms) than the presumed polluters.  1) Stars more massive than 9 \\ms \\ evolve on time scales lower than 20 Myr and are quite numerous (since a flat IMF is required in any case to account for the observed abundance distribution, as argued in this work). They release material processed through various nuclear burning stages, both through their winds (slow in the Main Sequence, quite fast in the Wolf-Rayet stage) and through the final stellar explosion as supernovae. A commonly given answer as to the fate of their ejecta is that the fast moving supernova material ($v>$400 km/s even for the innermost SN regions) escapes the system and leaves no trace behind. On its way, it sweeps up the previously released wind ejecta, but also any ambient gas of pristine composition. Consequently, no original Li rich gas is left behind to get subsequently mixed  with the AGB ejecta. The second (polluted) stellar generation should then be formed exclusively with AGB ejecta, which implies that  AGBs should produce their own Li, at a level close to the Spite plateau. This is an extremely strong constraint on the nucleosynthesis  of AGB stars.  Alternatively, one may argue that an insignificant amount of nuclearly processed ejecta escapes the supernovae, the vast majority being trapped within a black hole; this is an utterly implausible scenario, since neutron stars and pulsars are common in the Milky Way and there is no reason why massive stars in GCs would produce no such residues.  Finally, one might also invoke a peculiar IMF of the first stellar generation in GCs, lacking massive stars altogether. Such an ad hoc assumption, however, is not supported by any theoretical arguments or numerical simulation results.  2) Stars in the mass range 2-4 \\ms \\ evolve on timescales of 150 to 1000 Myr and release almost as much mass as those of 4-9 \\ms (depending on the assumed IMF: for X=0.15 we obtain $E$(2-4)=0.35 $E$(4-9)). Even if they do not develop HBB in their envelopes, they go through the third dredge-up phase; this brings He-burning products (in particular C) and s-elements to the surface, which are slowly released to the ISM. If new stars are formed from these ejecta, C and s-elements excesses should be observed in GC stars, which is not the case. One then has to assume that the ejecta of 2 to 4 \\ms \\ form no new stars, why those of their slightly more massive sisters manage to do it very efficiently. Perhaps, on time scales $>$100 Myr tidal interactions of the GCs with the Galactic potential are able to remove most of the cluster gas (as well as part of its stars); still some ejecta should be left behind to form stars with C and s-elements excess. Until a consistent dynamical scenario is developed to account  convincingly for that point, we consider it to be one of the drawbacks of the AGB scenario for the pollution of GCs.  \\subsubsection{WMS?}  The idea that  WMS may be at the origin of some anomalies in the composition of GCs has been recently suggested by Norris (2004) and Maeder and Meynet (2006), in order to explain the blue main sequence of the cluster $\\omega$ Centauri: the high helium content of the stars of that sequence could originate from WMS, producing a large helium/metal ratio.  Omega Cen is, however, a peculiar case,  showing a dispersion in its heavy metal content, unlike any other GC. If the polluters of GCs are indeed the winds of massive stars, the scenario could run along the following lines:  The winds of those stars, rich with H-burning products (especially if the stars are rotating), are slowly released in the ambient ISM and they push it, opening circumstellar cavities; they are finally mixed with pristine gas, still left over in the system after the formation of the first stellar generation. By reaching the walls of those cavities, the shock waves of the subsequent SN explosions induce the formation of new stars from this gas of mixed composition. The hot SN ejecta inside the shock waves are expanding mainly along the cavities opened  up previously by the massive star winds, and they finally find their way out of the cluster, escaping altogether the system. Thus, the stars we now observe in the cluster are either of original composition, or of composition resulting from the mixture of pristine gas with WMS. On longer timescales ($>$20 Myr) the wind ejecta of lower mass stars (M$<$10 \\ms) that do not explode, are slowly released in the cluster. However, they never form new stars, since there is no external agent to trigger condensation of that gas and to induce star formation. As discussed before, tidal interactions ultimately strip that gas off the cluster.  The advantages of the WMS scenario are that: i) it may ``naturally'' account for the formation of the second stellar generation with mixed composition, which is triggered by the SN shock waves, and ii) it may also explain why no trace of the He-burning products of the third dredge-up (which certainly occurs in low-mass AGBs) is seen today in GC stars: these products, as well as those of HBB in more massive AGBs, are slowly released in a rarefied medium (since almost all of the gas of the system has been swept up or condensed previously, after the passage of the shock waves); as a result, they never condense to form new stars (for lack of triggering agent) and they are ultimately removed from the cluster through tidal interactions.  Of course, this novel scenario has also its potential drawbacks. It assumes that WMS mix thoroughly with pristine gas and that the subsequent shock wave triggers an efficient star formation in the mixture, while the SN ejecta escape the system. One may, in fact, imagine the opposite situation, where the shock waves disperse the gas, or the SN ejecta mix with that gas (at least partially). At the current stage of our knowledge, both alternatives are possible.  The WMS scenario sketched here is, admittedly, at a very qualitative level (but no more than the scenario involving AGBs). It should be substantiated at least by nucleosynthesis calculations demonstrating that WMS can indeed provide the chemical composition required to explain the observations of abundances in GCs; such calculations are already performed and they will be presented elsewhere (Decressin et al. in preparation). We feel, however, that the WMS scenario opens very interesting perspectives for the understanding of the self-enrichment of GCs."
    },
    "0606/astro-ph0606324_arXiv.txt": {
        "abstract": "Opher et al. \\cite{opher06} showed that an interstellar magnetic field parallel to the plane defined by the deflection of interstellar hydrogen atoms can produce a north/south asymmetry in the distortion of the solar wind termination shock. This distortion is consistent with Voyager 1 and Voyager 2 observations of the direction of field-aligned streaming of the termination shock particles upstream the shock. The model also indicates that such a distortion will result in a significant north/south asymmetry in the distance to the shock and the thickness of heliosheath. The two Voyager spacecraft should reveal the nature and degree of the asymmetry in the termination shock and heliosheath. ",
        "introduction": "After 27 years of anticipation, Voyager 1 crossed the inward moving termination shock (TS) at $94~AU$ and is now exploring the heliosheath \\cite{burlaga,decker,stone05}. The twin Voyager spacecraft are probing the northern and southern hemispheres of the heliosphere.  As Voyager 1 crossed the TS and began exploring the heliosheath it has become increasingly clear that this previously unexplored region is full of surprises. Using in situ spacecraft data to constrain the shape of the heliosphere is challenging because they are single point observations. For a quantitative global understanding of the three dimensional structure of the heliosphere, it is necessary to use modeling in conjunction with observations to probe this region.  In mid 2002, Voyager 1 began observing strong beams of energetic termination shock particles (TSPs) streaming outward along the spiral magnetic field. The strong upstream TSP beams were observed much of the time until Voyager 1 crossed the shock at $94~AU$ in December 2004. Jokipii et al. \\cite{jokipii} and Stone et al. \\cite{stone05} suggested that the upstream beaming resulted from a non-spherical shock. For a spherical shock, Voyager 1 would observe upstream TSPs streaming inward along the magnetic field. With a non-spherical shock, Voyager 1 could be connected to the TS along magnetic field lines that crossed the TS (the source of TSPs) and then crossed back into the supersonic solar wind.  In order to know Voyager 1 measurements in 2002 it is important to understand the shape of the heliosphere. The interaction of the solar system with the local interstellar medium requires an intensive modeling effort because this problem is inherently three dimensional and involves solar and interstellar magnetic fields, ionized and neutral atoms, and cosmic rays\\citep{suess,zank}. The size and shape of the heliosphere will depend on the properties of both the solar wind and the local interstellar medium.  Based on the polarization of light from nearby stars, Frisch \\cite{frisch,frisch1} suggested that the magnetic field direction is parallel to the galactic plane (and directed toward $l \\sim 70^{\\circ}$). Voyager 3kHz radio emission data also show preferred source locations in a plane parallel to the galactic plane \\cite{kur}. On the other hand, Lallement et al. \\cite{lall}, mapping the solar Lyman-$\\alpha$ radiation that is resonantly backscattered by interstellar hydrogen atoms, found that the neutral hydrogen flow direction differs from the helium flow direction by $4^{\\circ}$. The plane of the H deflection (HDP) is tilted from the ecliptic plane by $\\sim 60^{\\circ}$ and is consistent with an interstellar magnetic field parallel to the HDP plane\\cite{izmo}.  However, Pogorolev et al.\\cite{pog06} show that the H deflection does not uniquely determine the plane of the magnetic field.  If the interstellar magnetic field is inclined to the interstellar velocity, it can produce a lateral or north-south asymmetry in the heliospheric shape \\cite{pog,ratk}. Opher et al. \\cite{opher06} showed that that an interstellar magnetic field parallel to the H-deflection plane (HDP) can produce a north/south asymmetry in solar wind termination shock that is consistent with Voyager 1 and Voyager 2 observations of the direction of streaming of the particles. We review these results and comment on the Heliosphere global asymmetry. ",
        "conclusions": "The recent results by Opher et al.\\cite{opher06} indicate that an interstellar magnetic field $(B\\sim 2\\mu G)$ in the H-deflection plane can distort the heliosphere in a manner consistent with the streaming observed by Voyager 1 and Voyager 2. The model also indicates that such a distortion will result in a significant north/south asymmetry in the distance to the shock and the thickness of heliosheath, and it is reasonable to expect that Voyager 2 will encounter the shock in the next two years. When combined with a more realistic heliospheric model with a tilted current sheet, this should provide additional information about the strength and inclination of the local interstellar magnetic field.  \\begin{theacknowledgments} The authors would like to thanks the use of Columbia cluster NASA Ames. Part of this work is the result of research performed at the Jet Propulsion Laboratory of the California Institute of Technology under a contract with the National Aeronautics and Space Administration. \\end{theacknowledgments}"
    },
    "0606/astro-ph0606655_arXiv.txt": {
        "abstract": "Periodic outbursts are observed in many AGNs and usually explained with a supermassive black hole binary (SMBHB) scenario. However, multiple periods are observed in some AGNs and cannot be explained with it. Here we analyze the periodicity of the radio light curves of AO 0235+164 at multi-frequencies and report the discovery of six QPOs in integer ratio 1:2:3:4:5:6 of QPO frequencies, of which the second with period $P_2 = (5.46 \\pm 0.47) \\, {\\rm yr}$ is the strongest.  We fit the radio light curves and show that the initial phases of six QPOs have zero or $\\pi$ differences relative to each other. We suggest a harmonic relationship of QPOs. The centroid frequency, relative strength, harmonic relationship and relative initial phases of QPOs are independent of radio frequency. The harmonic QPOs are likely due to the quasi-periodic injection of plasma from an oscillating accretion disk into the jet. We estimate the supermassive black hole mass $M_{\\rm BH} \\simeq (4.72\\pm 2.04) \\times 10^8 M_\\odot$ and the accretion rate $\\dot{m}\\simeq 0.007$.   With the knowledge of accretion disk, it implies that the inner region of accretion disk of AO 0235+164 is a radiatively inefficient accretion flow. The oscillation accretion is due to the p-mode oscillation of the thick disk probably excited by a SMBHB. The theoretical predications of fundamental oscillation frequency and the harmonics are well consistent with the observations. Harmonic QPOs would be absent when the thick disk becomes geometrically thin due to the increase of accretion rate. We discuss the observations of AO 0235+164 basing on the SMBHB-thick disk oscillation scenario. ",
        "introduction": "Blazars, consisting of BL Lac objects and flat-spectrum radio quasars, are a subclass of AGNs whose relativistic jet is nearly aligned to the line of sight, and show extreme variabilities at from radio through optical and X-ray to high-energy $\\gamma$-ray wavelengths \\citep{antonucci93,urry95,wagner95,ulrich97}. Long term multi-wavelength monitorings show that the variabilities of blazars are very complex with large amplitude and on time-scale ranging from tens of minutes, hours and days to months and years at all wavelengths. The outbursts with time-scale of order of years or longer in some blazars are  periodic \\citep{sill88,liu95,liu97,rai01,qian04} and of particular interest as it might connect to the presence of supermassive black hole binary (SMBHB) at center \\citep{sill88,liu02,qian04, osto04}, whose ultimate coalescence would generate an outburst of gravitational wave radiation \\citep{thorne76} and be the main target of a future space gravitational wave detector, Laser Interferometer Space Antenna (LISA) \\citep{merritt05}.  SMBHBs in galactic nuclei are expected by the hierarchical galaxy formation model in the cold dark matter (CDM) cosmology \\citep{kauffmann00} in which the present galaxies are the products of frequent galaxy minor mergers. Active SMBHB and final coalescence have been suggested to be the physical origin of the peculiar radio morphologies in some AGNs \\citep[for a recent review, see][]{komossa03}, e.g. the helical morphology of radio jets \\citep{begelman80}, the interruption and recurrence of activity in double-double radio galaxies \\citep{liu03}, the X-shaped feature of winged radio sources \\citep{merritt02,liu04}, the orbital motion of radio core \\citep[e.g.][]{sudou03}. Since they were discovered in BL Lac object OJ287 \\citep{sill88}, periodic optical outbursts of blazars have been ascribed to the orbital motion of SMBHB \\citep{sill88,abraham99,valtaoja00,liu02, rieger04,lobanov05}. However, it is unclear how SMBHBs in AGNs trigger the observed periodic outbursts if present. The proposed scenarios in literature include (1) direct impact of the secondary black hole against a standard thin accretion disk \\citep[e.g.][]{lehto96}) or an inefficient accretion flow, e.g. advection dominated accretion flow \\citep[ADAF;][]{liu02}, (2) rotating helical jets due to the orbital motion of SMBHB \\citep{katz97,villata98,osto04,rieger04}, (3) periodic change of jet orientation introduced by the Lens-Thirring precessing of a warped disk \\citep{abraham99,lobanov05}. Although the detailed physical mechanisms in these models are different, all of them suggest a common feature of single period in the light curves.  One more complication to the investigation of the periodicity in blazars is from the observations of low frequency X-ray quasi-periodic oscillations (QPOs) in black hole X-ray  binaries \\citep{has86,klis89,stroh96} and possibly in Sgr A$^*$ of our Galactic center \\citep{torok05}, which are believed to be due to the disk oscillations. Black hole X-ray binaries, so called micro-quasars, are physical analogue to AGNs and are powered by accretion of stellar black holes of mass about $10 {\\rm M}_\\odot$. As the variation time-scale of accretion disk around black holes is proportional to the mass of black hole and the formation of jet is coupled with accretion disk \\citep{gallo03,fend04,fend04a}, the low frequency QPO with typical frequency $\\sim 1 {\\rm Hz}$ in micro-quasars corresponds to a period of order of years or longer in a blazar system of black hole mass larger than $10^8 M_\\odot$. Disk oscillations may cause a periodic change of accretion and the QPOs in X-ray black hole binary system \\citep{rezzolla03}, which could lead to periodic jet emission due to the jet-disk coupling.  To understand the physical origin of periodic outbursts, it is essential to investigate in details the structure of periodic outbursts in both light curves and power density spectrum, and the dependence of periodicity on the observational wavelengths. For example, the discovery of double peaks in the periodic optical outbursts of the BL Lac object OJ287 \\citep{sill96} leads to the second version of the SMBHB model \\citep{lehto96}, while the investigation of the relation of the optical outburst structure and the radio variations leads to the third version of the SMBHB scenario \\citep{valtaoja00,liu02}. It is generally believed that the short-term variabilities of blazars are non-periodic \\citep[e.g.][]{wagner95,ulrich97}, but the short-term periodic brightness flickering may have been found to superpose on the long-term periodic variations in some blazars, e.g. OJ287 \\citep{wu06} and 3C66A \\citep{lain99}. In order to give more constraints on the models for periodic outbursts in AGNs, recently we started a program to investigate the relationships of periodic major outbursts and minor events and the long- and short-term periodicities of a large sample of AGNs \\citep{zhao05}. Here we report the results of multiple quasi-periodic oscillations (QPOs) and the harmonic resonant relationship 1:2:3:$\\dots$ of QPO frequencies in BL Lac object AO 0235+164.  BL Lac object AO 0235+164 with redshift $z=0.94$ is one of the most violently variable object whose periodicity of outbursts in optical and radio bands has been investigated and reported in literature \\citep[e.g.][]{web88,web00,roy00,rai01,rai05}. \\citet{roy00} analyzed the data of the University of Michigan Radio Astronomy Observatory (UMRAO) at frequencies 4.8, 8.0 and 14.5 GHz from 1975 to 1999 and discovered a period $\\sim 5.8 \\, {\\rm yrs}$ at 8.0 and 14.5 GHz, which is consistent with the observations of a period $5.7 \\, {\\rm yrs}$ in the optical \\citep{rai01,rai05}. \\citet{osto04} interpreted the long term periodicity of major outbursts with a helical-jet model and the minor events with rotation of the non-periodic inhomogeneity. However, the radio outbursts are double-peaked and the burst peaks do not appear as regularly as the period predicted \\citep{rai01,rai05}, implying that the substructure of the outburst are complex. On the other hand, the multiple short periods though at much lower significant level of signal-to-noise and the poor identification of central frequency  have been claimed in literature. \\citet{web88} first identified the three peaks with three periods $\\sim 2.79 \\, {\\rm yrs}$, $\\sim 1.58\\, {\\rm yrs}$, and $\\sim 1.29 \\, {\\rm yrs}$ in the Fourier power spectrum of the optical light curves. Subsequent observations marginally confirmed the periods with 2.7 and 1.2 years in optical \\citep{web00} and with 1.8, 2.8 and 3.7 years in radio wave-bands with the Discrete Correlation Function (DCF) analysis method \\citep{rai01}.  Our work is based on the three radio databases of the University of Michigan Radio Astronomy Observatory (UMRAO), the National Radio Astronomy Observatory (NRAO), and the Mets\\\"ahovi Observatory. In Sec.~\\ref{database}, we start the investigations from analyzing the consistency of UMRAO's and NRAO's databases, by comparing the best-sampled observations at 8.0 GHz at UMRAO and those at 8.2 GHz at NRAO and then merging the observational data at these two frequencies. The periodic analysis results for the combined light curves at 8 GHz and identification of  harmonic QPOs are given in Sec.~\\ref{result8}. In Sec.~\\ref{resmult}, we investigate the dependence of QPO properties on radio frequency. After estimating the mass of central supermassive black hole (SMBH) and the accretion rate in Sec.~\\ref{sec:bhmass}, we analytically discuss the p-mode oscillation of thick disk in Sec.~\\ref{sec:diskosc}.  A SMBHB-disk oscillation scenario for multiple harmonic QPOs is suggested and discussed in Sec.~\\ref{sec:pertbinary}.  Our discussions and conclusions are given in Sec.~\\ref{concl}. We assume a flat cosmology with $H_0 = 75 \\, {\\rm Km \\; s^{-1} Mpc^{-1}}$, $\\Lambda = 0$, and $q_0 = 0.5$ throughout the paper, leading to a source distance $d \\simeq 4.4 Gpc$. ",
        "conclusions": "\\label{concl}  We have investigated the radio variabilities of AO 0235+164 and analyzed the periods of radio light curves at 4.8, 8, 14.5, 22, and 37 GHz, basing on the databases of UMRAO, NRAO and Mets\\\"{a}hovi Observatory. We first study the consistence of the UMRAO's and NRAO's data at 8 GHz for AO 0235+164 and then construct a combined radio light curve with the observational data from the two databases. The periodic analyses with three kinds of classical periodic analysis methods show with unprecedentedly high signal-to-noise ratio that the variations of the combined radio light curves are composed of six periodic variations plus random fluctuations. In the power spectrum, the peaks are Gaussian and identified with QPOs of periods $P_1 = 12.02 \\, {\\rm yr}$, $P_2 = 5.45 \\, {\\rm yr}$, $P_3 = 3.62 \\, {\\rm yr}$, $P_4 = 2.83 \\, {\\rm yr}$, $P_5 = 2.15 \\, {\\rm yr}$, and $P_6 = 1.81 \\, {\\rm yr}$. The coherent quality factor $Q$ of the six QPOs are limited by the monitoring program time. We discover a harmonic relationship for the six QPOs, based on the ratio of QPO frequencies in a sequence of small integers 1:2:3:4:5:6 and on the zero or $\\pi$ difference of initial phases of QPOs. The harmonic relationship suggests that the six QPOs should have the same physical origin. Among the six harmonic QPOs, the second one with period $P_{\\rm 2} \\simeq 5.46 \\pm 0.47 \\, {\\rm yrs}$ is the strongest and is identified with the fundamental period. The five strong harmonics reported in some literature are confirmed here with an unprecedentedly high signal-to-noise and accuracy, while the weakest QPO of period $P_5 = 2.15 \\, {\\rm yrs}$ is reported first time in this paper.  We investigate the dependence of harmonic relationship, coherent quality factor, the relative {\\it rms}, the differences of the initial phase of QPOs on radio frequency. The results show that all the QPO properties are independent of radio frequencies. The six harmonic QPOs give a good fit to the outburst structures of the radio light curves at 4.8, 8, 14.5, 22 and 37 GHz, and predict major outbursts and outburst structures significantly different from that given with a single period of $P \\simeq 5.46 \\, {\\rm yr}$ \\citep{rai01}. Both the single-period and the multiple harmonic QPO (for the low frequencies) scenarios predict a major outburst in the early half of 2004, while the combinations of six harmonic QPOs suggest only minor fluctuations during that period but a major burst in the late half of 2004 or the early half 2005 at 37 GHz, as it is shown in Fig.~\\ref{fit8}. In our periodic analyses of radio light curves at 22 GHz and 37 GHz, we have used the data after 2001 and the harmonic QPO scenario predict a major outburst in the early 2005. At lower frequencies we cannot get the observational data in recent years and make the analyses basing on the data before 2001. Therefore, the predictions are slightly different for different wave-bands and the predicted major bursts may be in the early 2005. However, the observations of a WEBT (the Whole Earth Blazar Telescope) campaign started from 2003 show that AO 0235+164 is quiet in recent years and the predicted major outbursts in the spring of 2004 did not come \\citep{rai05}. Is the predicted major outburst just delayed significantly or is the activity of the object interrupted due to some reasons? When will the object re-activate with a different active cycle? To answer the questions, more observations for another several years are needed.  Either a significant delay of outbursts or the interruption of activity for some time would exclude the single period scenario and the simple version of a helical jet model but may be consistent with the SMBHB-thick disk oscillation model. In a SMBHB-thick disk oscillation scenario, the multiple QPOs in the radio light curves of AO 0235+164 are due to quasi-periodic injection of plasma from an oscillation thick disk of a finite radial extent into relativistic jet. A torus or a thick accretion disk of finite radial extent oscillates and accretes matter quasi-periodically with small integer ratio of frequencies with 1:2:3:$\\dots$, if it is perturbed by a global perturbation or a local strong periodic agent \\citep{zanotti03,rezzolla03,rubio05b}. The fundamental oscillation frequency of a thick disk is determined by the radial extent and sound speed. We estimated the central black hole mass $M_{\\rm BH} = (4.72\\pm 2.04) \\times 10^8 M_\\odot$ and the dimensionless accretion rate $\\dot{m} \\simeq 0.007$ of AO 0235+164, based on the observations of emission line properties. The estimated dimensionless accretion rate with the current knowledge of accretion around black hole suggests that the inner region of accretion disk in AO 0235+164 is a geometrically thick disk (i.e. ADAF or RIAF) with radial extent about $1.2 \\times 10^3 r_{\\rm G}$ and the outer region is a truncated cool thin disk. With the self-similar solution of ADAF, we analytically estimate the fundamental oscillation frequency $f_{\\rm obs} \\simeq 0.0938 \\, {\\rm yr}^{-1}$ or period $P_{\\rm obs} = 10.7 \\,{\\rm yr}$ in the observer's frame, which together with the predicted overtones are consistent with the observed multiple QPOs with lowest frequency $P_1 \\simeq 12.02 \\pm 5.79 \\, {\\rm yr}$ (or $2P_2 = 10.9 \\, {\\rm yrs}$ at higher accuracy) and harmonics. As there is no natural global or strong local perturbation source in standard SMBHB-accretion disk system to excite the global oscillations of a thick disk of a large radial extent $R_{\\rm d} \\sim 10^3 r_{\\rm G}$, we suggest a SMBHB rotating within the thick disk to excite the p-mode oscillations of RIAF in AO 0235+164. The rotation of the secondary triggers acoustic waves in an ADAF (or RIAF), which propagate inwards and outwards in the thick disk and are reflected backwards by the inner and outer boundaries. The p-mode oscillations with frequencies of about the fundamental oscillation frequency and the overtones are trapped by the cavity. The orbital motion of a SMBHB could lead to the formation of a helical morphology of relativistic jet and to the rapid changes of VLBI jet position angles. Therefore, in the SMBHB-disk oscillation model, the relativistic jet can be also helical. The orbital motion may correspond to the 6th QPO with $P_6 = 1.81 \\, {\\rm yrs}$, the 2nd QPO with $P_2 = 5.46 \\, {\\rm yrs}$, or have not been observed. Long term VLBI monitoring of rapid changes of jet positional angles would help to answer the question. One of the differences with the simple helical jet model is that in the SMBH-thick disk oscillation scenario the minor burst events in the radio light curves is no longer random but together with the major outbursts are the combinations of harmonic six QPOs which are the p-mode oscillations of the RIAF and the binary orbital motion. This is a first time report of multiple harmonic QPOs and the possible detection of perturbed disk p-mode oscillations in AGNs.  The second difference of predictions of the helical jet model and the SMBHB-thick disk oscillation scenario is that periodic outbursts in the former model are persistent and regular while in the later scenario are temporary with moderate coherence quality factor $Q$ and depend on the dimensionless accretion rate. In the SMBHB-thick oscillation scenario, resonant harmonic QPOs would disappear from the radio light curves or the periodic activities of the radio source are interrupted, when the accretion rate increases only about two times aand becomes larger than the critical accretion rate $\\dot{m}_{\\rm cr}$ and the inner region of accretion disk becomes a geometrically thin cool disk. Therefore, the quiescence and the absence of outbursts from the radio light curves of AO 0235+164 may be because of the transition of accretion mode from a RIAF to a standard thin cool accretion disk due to the moderate increase of accretion rate. The X-ray spectral observations with the Chandra satellite in August 2000 and with the XMM-Newton satellite in August 2004 show tentative detections of the red-shifted Fe K$\\alpha$ fluorescent emission line in AO 0235+164 \\citep{rai06}, which suggests that the accretion disk at the immediate vicinity of the SMBH is very likely a standard thin accretion disk \\citep{fabian00}. Therefore, the absence of outbursts from the radio light curves and the interruption of the harmonic QPOs can be understood in the SMBHB-thick disk oscillation scenario as the transition of accretion mode due the increase of accretion rate. In a jet-disk coupling scenario, an object becomes radio quiet when the accretion disk becomes a standard thin disk \\citep{fend04}.  We notice that in the helical jet model, the separation of SMBHB is much smaller than the transitional radius $R_{\\rm tr}$ of inner thick RIAF and outer thin disk. Therefore, the SMBHB in the helical jet model for AO 0235+164 should also excite the p-mode oscillations of an RIAF, leading to quasi-periodic accretion and quasi-periodic injection of plasma into the relativistic jet. Therefore, multiple QPOs with harmonic relationship are also expected by the helical jet model for AO 0235+164, if we consider the RIAF p-mode oscillations.  The embedding secondary black hole spherically accretes matter from the RIAF and radiates energy in X-ray. The accretion radius of the secondary black hole is $r_{\\rm acc} \\approx {Gm \\over c_{\\rm s}^2 + (\\Delta{v})^2}$, where $\\Delta v$ is the differential velocity of the secondary black hole and the plasma in the thick disk. For a minor merger with $q \\ll 1$, we obtain $r_{\\rm acc} \\ll R$ and $(\\Delta{v})^2 \\ll c_{\\rm s}^2$. From the self-similar solution $c_{\\rm s}^2 = (2/5) V_{\\rm K}^2$ for $\\gamma = 5/3$ \\citep{narayan94,narayan00}, we have $r_{\\rm acc} \\approx (5/2) q (1+q) R$. The accretion rate of the secondary black hole is \\begin{equation} \\dot{M}_{\\rm s} \\sim 4 \\pi r_{\\rm acc}^2 c_{\\rm s} \\rho , \\label{eq:secacc} \\end{equation} where the mass density $\\rho$ is given with the self-similar solution and $\\rho = {5 \\sqrt{5} \\over 12 \\sqrt{2} \\pi} \\dot{M}_{\\rm Edd} \\dot{m} \\alpha^{-1} R^{-2} V_{\\rm K}^{-1}$. From Eq.~(\\ref{eq:secacc}), we have $\\dot{M}_{\\rm s} \\sim {25 \\sqrt{5} \\over 6 \\sqrt{2}} q^2 (1+q)^2 \\dot{M}_{\\rm Edd} \\dot{m} \\alpha^{-1}$. Defining the Eddington accretion rate for the secondary black hole $\\dot{M}_{\\rm Edd}^s = L_{\\rm Edd}^s / 0.1 c^2$ with $ L_{\\rm Edd}^s = 1.26 \\times 10^{38} (m / M_\\odot) \\, {\\rm ergs\\; s^{-1}}$, we have the dimensionless accretion rate of the secondary black hole, $\\dot{m}_{\\rm s} \\equiv \\dot{M}_{\\rm s} / \\dot{M}_{\\rm Edd}^s \\sim 6.6 (1+q)^2 (q/\\alpha) \\dot{m}$. For the typical parameters $q = 0.01$ and $\\alpha = 0.13$, the relative accretion rate of the secondary is $\\dot{m}_{\\rm s} \\sim 0.51 \\dot{m} \\sim 5.1 \\times 10^{-3}$, while $\\dot{m}_{\\rm s} \\sim 5.1 \\times 10^{-2}$ for $q=0.1$. The X-ray emission from the accretion of the secondary black hole is small and may not be observable in AO 0235+164, because of the strong X-ray radiations both from the relativistic jet and from accretion disk around the primary. A different $\\gamma$ value does not change the conclusion.  Our work implies that if a thick accretion disk with finite radial extent is perturbed globally or locally by, e.g. SMBHB, the p-mode oscillations can be observed as multiple harmonic QPOs, when the relativistic jet is along the line of sight. Periods have been reported and SMBHBs have been suggested in literature for many AGNs. So, multiple QPOs with a harmonic relationship of frequencies are expected by the SMBHB-disk oscillation model in moderately low luminosity AGNs with SMBHB at center, which may have accretion configuration like the one in the BL Lac object AO 0235+164. We will report our periodic analysis results on a large sample of AGNs in a coming paper \\citep{zhao05}."
    },
    "0606/astro-ph0606149_arXiv.txt": {
        "abstract": "Spectroscopic monitoring of 141 southern field B type stars, 114 of them known to exhibit the Be phenomenon, allowed the estimation of their projected rotational velocities, effective temperatures and superficial gravities from both line and equivalent width fitting procedures. Stellar ages, masses and bolometric luminosities were derived from internal structure models. Without taking into account for the effects of gravity darkening, we notice the occurrence of the Be phenomenon in later stages of main sequence phase. ",
        "introduction": "\\indent Despite the large amount of works on the subject of B and Be stars published in the past years, both theoretical and observational, there are still many remaining unsolved problems concerning their outstanding peculiarities. Among them the origin of their high rotation velocities, dependence of Be frequency counts with evolutionary stages in the main sequence, presence and origin of magnetic fields and mass loss (Vinicius et al. 2006, McSwain \\& Gies 2005, Levenhagen \\& Leister 2004, Levenhagen et al 2003, Porter \\& Rivinius 2003).  \\indent Concerning particularly the case of Be stars, the most outstanding observable characteristic in the visible domain is the existence of emission in Balmer lines, being often accompanied by emission of singly ionized metals and changes of the spectrophotometric characteristics (Moujtahid et al. 1999; Levenhagen 2004). In these stars, emissions and energy distributions are variable and the most important among these variations are ``phase transitions'', which most likely reflect changes in the structure of the circumstellar envelope (CE) and processes related with its formation.  \\indent One of the fundamental questions related to the Be phenomenon concerns the origin of the fast rotation rates of the central star. Rotation has quite different effects on stellar evolution, as changes of internal hydrostatic equilibrium, transport of chemical species by meridional circulation and effects on mass loss rates (Maeder 1999).  These days it is a consense that fast surface rotation plays a crucial role in triggering the Be phenomenon, where the onset of new observational evidences indicate that these objects are in fact near critical rotators, with $\\omega \\simeq 0.88$ (Zorec, Fr\\'emat \\& Cidale 2005, Fr\\'emat et al. 2005, Townsend, Owocki \\& Howarth 2004).  There are actually three main possibilities for that: (i) it would be a permanent innate property; (ii) it would be due to spin up by binary mass transfer; (iii) or it would be acquired somehow during the main sequence evolution phase (Crampin \\& Hoyle 1960, Schild \\& Romanishin 1976). Early attempts to describe the Be phenomenon suggested the occurrence of Be stars during the secondary contraction phase (Schmidt-Kaler 1964). A few years later, it was seen the appearance of a significant fraction of Be stars close to the ZAMS (Schild \\& Romanishin 1976). In the 80's it was generally accepted that those objects occupy the whole main sequence band at different evolutionary phases (Mermilliod 1982; Slettebak 1985).  \\indent Maeder, Grebel \\& Mermilliod (1999) explained the origin of high rotation rates by two possible scenarios of low metallicity effects in star formation regions. In the first scenario a weak coupling of magnetic fields with star forming matter leads to less angular momentum losses and a consequent spin-up. The second possibility is based on a low opacity effect in the pre main sequence phase which in turn determines the external convective zones to fade away and thus leading to less angular momentum losses.  Today some works present the Be phenomenon as a result of an evolutionary effect instead (Fabregat \\& Torrej\\'on 2000). In a recent work on Be stars in open clusters, McSwain \\& Gies (2005) pointed out that the Be phenomenon is not strongly dependent on metallicity or cluster density, instead it is influenced by the evolutionary stage. One problem with those studies is that they rely on photometric measurements only, which tends, in the case of Be stars, to systematically overestimate the magnitudes due to the presence of a circumstellar environment. On the other hand, Zorec et al (2006) determined photospheric parameters of a sample of 97 field Be stars near the Sun based on spectroscopic measurements considering both rotating and non-rotating model atmosphere scenarios. In this work they found that, without taking into account for the effects of fast rotation in the models, Be stars usually appear near the TAMS rather than the ZAMS in the HR diagram, i.e. the Be phenomenon would arise due to an apparent evolutionary effect. Later on, considering the effects of gravity darkening, the stellar distribution in the HR diagram becomes more homogeneous, and Be stars are allowed to occur at any stages within the main sequence phase. In order to elucidate this problem further investigations are still needed.  \\indent In this paper we present the results of spectroscopic analyses carried out with 141 southern B stars (where 114 are field Be stars), some of them located away from the solar neighbourhood, in order to provide physical parameters estimations of \\vsini, \\teff~ and \\logg, and interpolations in the tabulated evolutionary tracks by Schaller et al. (1992) necessary for further investigations on these objects. These parameters have been estimated from NLTE model atmospheres without taking into account for gravitational darkening (von Zeipel 1924a,b) and their positions in the HR diagram were established. We also estimated the errors committed in the physical parameters (\\teff,\\logg) by neglecting the rotational effects. ",
        "conclusions": ""
    },
    "0606/astro-ph0606180_arXiv.txt": {
        "abstract": "In this paper we investigate the possible reasons why \\HI~21-cm absorption in damped Lyman-$\\alpha$ systems (DLAs) has only been detected at low redshift: To date, no 21-cm absorption has yet been detected at $z_{\\rm abs}>2.3$ and at redshifts less than this, there is a mix of detections and non-detections in the DLAs searched. This has been attributed to the morphologies of the galaxies hosting the DLAs, where at low redshift the DLAs comprise of both large and compact galaxies, which are believed to have low and high spin temperatures, respectively. Likewise, at high redshift the DLA population is believed to consist exclusively of compact galaxies of high spin temperature \\citep{ck00,kc01a,kc02}. However, in a previous paper \\citep{cmp+03} we found that by not assuming or assigning an, often uncertain, value for the coverage of the radio continuum source by the 21-cm absorbing gas, that there is generally no difference in the spin temperature/covering factor ratio between the 21-cm detections and non-detections or between the low and high redshift samples. Furthermore, only one of the 18 non-detections has a known host morphology, thus making any link between morphology and 21-cm detectability highly speculative.   We suggest that the lack of 21-cm absorption detections at high redshift arises from the fact that these DLAs are at similar angular diameter distances to the background quasars (i.e. the distance ratios are always close to unity): Above $z_{\\rm abs}\\sim1.6$ the covering factor becomes largely independent of the DLA--QSO distance, making the high redshift absorbers much less effective at covering the background continuum emission. At low redshift, small distance ratios are strongly favoured by the 21-cm detections, whereas large ratios are favoured by the non-detections. This mix of distance ratios gives the observed mix of detections and non-detections at $z_{\\rm abs}\\lapp1.6$. In addition to the predominance of large distance ratios and non-detections at high redshift, this strongly suggests that the observed distribution of 21-cm absorption in DLAs is dominated by geometric effects. ",
        "introduction": "Redshifted absorption systems lying along the sight-lines to distant quasars are important probes of the early to present day Universe.  In particular, damped Lyman-$\\alpha$ absorption systems (DLAs), where $N_{\\rm HI}\\ge2\\times10^{20}$ \\scm, are useful since they account for at least 80\\% of $\\Omega_{\\rm neutral}$ in the Universe \\citep{phw05}. Since the Lyman-$\\alpha$ transition occurs in the ultra-violet band, direct ground based observations of neutral hydrogen are restricted to redshifts of $z\\gapp1.7$. However, observations of the \\HI ~spin-flip transitions at $\\lambda_{\\rm rest}=21$ cm can probe from $z=0$, thereby providing a useful complement to the high redshift optical data.  Provided the 21-cm and Lyman-$\\alpha$ absorption arise in the same cloud complexes, the column density $N_{\\rm HI}$ [\\scm] of the absorbing gas in a homogeneous cloud is related to the velocity integrated optical depth of the 21-cm line via \\begin{equation} N_{\\rm HI}=1.823\\times10^{18}\\,T_{\\rm spin}\\int\\!\\tau\\,dv\\,. \\label{enew} \\end{equation} Here $T_{\\rm spin}$ [K] is the spin temperature of the gas and so in principle, armed with the neutral hydrogen column density from an observation of the Lyman-$\\alpha$ line, this quantity may be derived from the optical depth of the 21-cm absorption.  However, the observed optical depth of the 21-cm line also depends upon on how effectively the background radio continuum is covered by the absorber via, $\\tau\\equiv-\\ln\\left(1-\\frac{\\sigma}{f\\,S}\\right)$, where $\\sigma/S$ is the depth of the line relative to the flux density and $f$ is the covering factor of the flux by the absorber.  In the optically thin regime ($\\sigma/f.S\\lapp0.3$), which applies to all but one of the known 21-cm absorbing DLAs\\footnote{0235+164 \\citep{rbb+76}.}, Equation~\\ref{enew} reduces to $N_{\\rm HI}=1.823\\times10^{18}\\frac{T_{\\rm spin}}{f}\\int\\!\\frac {\\sigma}{S}\\,dv\\,,$ thus giving a direct measure of the spin temperature of the gas for a known column density (from the Lyman-$\\alpha$ line) and covering factor.  However, in the absence of any direct measurement of the size of the radio absorbing region, this latter value is often assumed or at best estimated from the size of the background emission region.  From the literature, 16 of the 35 DLAs searched have been found to exhibit 21-cm absorption (see Table \\ref{t1})\\footnote{Since radio frequency interference (RFI) prevented reliable observations of 0432--440 and 1228--113 \\citep{cmp+03}, these are not included in the analysis.}, all of which occur at redshifts below $z_{\\rm abs}\\leq2.04$, although there are a near equal number of non-detections also below this redshift. \\citet{ck00,kc01a,kc02} therefore advocate a scenario where low redshift DLAs have a mix of low (21-cm detections) and high (21-cm non-detections) spin temperatures, with the high redshift absorbers having exclusively high spin temperatures.  However, in a previous paper \\citep{cmp+03}, we find evidence that the importance of the covering factor is underestimated and the common practice of setting $f=1$ could possibly have the effect of assigning artificially high spin temperatures to DLAs, particularly those not detected in 21-cm. Furthermore, we found no statistical difference in the spin temperature/covering factor ratio between the low and high redshift samples, although the larger absorbing galaxies (spirals) group together at low values of ${T_{\\rm spin}}/{f}$ and $z_{\\rm abs}$. Since the ratio of spin temperature/covering factor is ${T_{\\rm spin}}/{f}\\propto N_{\\rm HI}/\\int\\!\\frac{\\sigma}{S}\\,dv$ (Equation~\\ref{enew}), we have taken into account the total \\HI\\ column density of the absorber and integrated optical depth of the 21-cm absorption (incorporating the radio flux). This suggests that the difference between the detections and non-detections is due to some other effect, a possibility which we investigate in this paper. ",
        "conclusions": "Regarding the detectability of 21-cm absorption in DLAs, we have found: \\begin{itemize} \\item In general, the non-detections of 21-cm absorption in DLAs have been searched as deeply as the detections, meaning that the ratio of spin temperature/covering factor does not differ significantly between the two samples.  \\item There is an apparent bias for 21-cm absorption to be detected in DLAs originally discovered through the Mg{\\sc \\,ii} doublet rather than the Lyman-\\AL\\ line. This, however, is superficial and merely reflects the true bias introduced by the redshift distribution of the DLAs: \\end{itemize} At $z_{\\rm abs}\\gapp1.6$, the absorbers are effectively at the same (angular) distances as the background quasars. After ruling out the possibility that the closer proximity of the undetected DLAs to the background quasars significantly raises the spin temperatures above those of the detections, we believe that the non-detections are due to low covering factors, the result of the flattening of the angular diameter distance at $z\\gapp1$.   Since DLAs are not detected in 21-cm absorption at $z_{\\rm abs}\\geq2.04$\\footnote{The highest redshift of a confirmed detection is now $z_{\\rm abs}=2.347$ (Table \\ref{t1}).}, \\citet{ck00,kc01a,kc02} suggested that the non-detections are due to the high redshift ($z_{\\rm abs}\\gapp2$) DLAs having exclusively high spin temperatures and the presence of both 21-cm detections and non-detections at low redshift is attributed to a mix of spin temperatures. However, the distribution of 21-cm detections and non-detections closely follows that of low and high angular diameter distance ratios, respectively, so that high redshift DLAs have exclusively high ratios, whereas $z_{\\rm abs}\\lapp2$ systems exhibit a mix of ratios. This geometric effect means that absorbers at high redshift will always cover the background quasar much less effectively than at low redshift and the degeneracy between spin temperature and covering factor may only ever be resolved by targetted searches for 21-cm absorption in high redshift DLAs towards very compact radio sources."
    },
    "0606/astro-ph0606463_arXiv.txt": {
        "abstract": "We have conducted wide-field \\HI\\ mapping of a $\\sim 5.5\\degr \\times 5.5\\degr$ region surrounding the NGC 3783 galaxy group, to an \\HI\\ mass limit of $\\sim 4\\times 10^8$\\Msun.  The observations were made using the multibeam system on the Parkes 64-m radiotelescope, as part of the Galaxy Evolution Multiwavelength Study (GEMS).  We find twelve \\HI\\ detections in our Parkes data, four more than catalogued in HIPASS. We find two new group members, and discover an isolated region of \\HI\\ gas with an \\HI\\ mass of $\\sim 4 \\times 10^8$\\Msun, without a visible corresponding optical counterpart. We discuss the likelihood of this \\HI\\ region being a low surface brightness galaxy, primordial gas, or a remnant of tidal debris. For the NGC 3783 group we derive a mean recession velocity of 2903$\\pm$26 \\kms, and a velocity dispersion of 190$\\pm24$ \\kms. The galaxy NGC 3783 is the nearest galaxy to the luminosity weighted centre of the group, and is at the group mean velocity. From the X-ray and dynamical state of this galaxy group, this group appears to be in the early stages of its evolution. ",
        "introduction": "The physical processes operating in galaxy groups are likely to be somewhat different to those operating in galaxy clusters. For example, the intragroup medium is significantly less dense than the intracluster medium and galaxy motions slower -- suggesting that ram pressure stripping is less effective in groups than in clusters (Mulchaey \\& Zabludoff 1998). On the other hand, as relative galaxy motions are smaller, gravitational interactions will be more effective and outright mergers more common in groups (Zabludoff \\& Mulchaey 1998).  Neutral hydrogen (\\HI) observations are a particularly useful method of probing the relative importance of ram pressure stripping and tidal interactions. The \\HI\\ is typically more extended than the optical disk of a galaxy, and thus is more sensitive to tidal interactions than the stars (e.g. Haynes, Giovanelli \\& Chincarini 1984), and ram pressure stripping affects the gas rather than the stars in a galaxy. In addition, \\HI\\ surveys can detect new, optically faint but gas-rich galaxies, previously unidentified in optical surveys (e.g. Kilborn et al. 2005; Banks et al. 1999). However, systematic, optically unbiased \\HI\\ surveys of loose galaxy groups are few. These \\HI\\ surveys of groups have found galaxies with extended HI envelopes (Haynes 1981), some galaxies with significant \\HI\\ deficiencies (Kilborn et al. 2005; Omar \\& Dwarakanath 2005) and that there are no intergalactic \\HI\\ clouds in groups similar to the Local Group (Pisano et al. 2004; Zwaan 2001).  Tidal interactions between gas-rich galaxies (e.g. Barnes \\& Hernquist 1986; Toomre \\& Toomre 1972) are known to produce \\HI\\ bridges (e.g. Koribalski \\& Manthey 2005; Koribalski \\& Dickey 2004; Yun, Lo \\& Ho 1994; Li \\& Seaquist 1994), tails (e.g. Hibbard et al. 2001b; Hibbard \\& Yun 1999) or even `clouds' that are spatially distinct from the host galaxy (e.g. Bekki, Koribalski \\& Kilborn 2005a; Bekki et al. 2005b; Ryder et al. 2001; Schneider 1985). Such \\HI\\ clouds, with no obvious optical counterpart, have been detected in clusters (Oosterloo \\& van Gorkom 2005; Davies et al. 2004; Giovanelli \\& Haynes 1989) but their origin is not always immediately obvious (Bekki et al. 2005a; Minchin et al. 2005; Djorgovski 1990). Determining the origin of \\HI\\ `clouds' is vital. For example, cosmological n-body simulations of the Local Group predict many more dark matter halos, than galaxies are optically observed in the Local Group (Moore et al. 1999; Klypin et al. 1999). Could these halos be detected in an \\HI\\ survey?  Despite deep \\HI\\ surveys of groups (Pisano et al. 2004; Zwaan 2001), and the HIPASS southern sky \\HI\\ survey (Meyer et al. 2004), we are yet to identify any \\HI\\ sources that are primordial, i.e. are not associated with a star forming galaxy.    This paper describes \\HI\\ observations of the NGC 3783 group, and in particular, the discovery of several new group members, and a region of HI emission that has no optical counterpart. This work is part of a larger investigation into the properties of galaxies in groups, the Group Evolution Multiwavelength Study (GEMS). GEMS is a large survey of about 60 groups that were selected to have previous ROSAT PSPC X-ray observations (see Forbes et al. 2006). Osmond \\& Ponman (2004; hereafter OP04) detail the selection criteria and group properties for the 60 groups. We have made wide-field \\HI\\ observations of 16 GEMS groups with the Parkes telescope. These observations will allow us to characterise the \\HI\\ content of loose groups, and the relationship between \\HI\\ content and X-ray emission. Kilborn et al. (2005, 2006) describe the \\HI\\ sub-sample and characteristics.  The NGC 3783 group was first identified by Giuricin et al. (2000), who found three galaxies including the well-studied Seyfert galaxy NGC 3783. The group lies at a distance of 36 Mpc (OP04); at this distance, 1\\arcsec = 174.5 pc.  OP04 find extended X-ray emission surrounding NGC 3783, out to a radius of 69 kpc (see Section~\\ref{xray} for details).    We briefly discuss our \\HI\\ observations and data reduction in Section~2, and the results from the wide-field \\HI\\ imaging in Section~3. High resolution observations of NGC 3783 itself, and a nearby dwarf galaxy are presented in Section~4, and we discuss the discovery of a region of intra-group \\HI\\ gas in Section~5.  We finish with discussions in Section~6, and conclusions in Section~7. Throughout we use a heliocentric velocity in the optical convention. ",
        "conclusions": "We have made a wide-field \\HI\\ survey of the NGC~3783 galaxy group, using the Parkes radiotelescope. We found 12 \\HI\\ detections in the region, of which one is a region of extended \\HI\\ emission that we cannot match with corresponding optical emission. We believe the origin of this \\HI\\ region is likely to be tidal debris rather than ram-pressure stripped, or primordial \\HI. We found two previously uncatalogued dwarf galaxies -- one from our Parkes observations (GEMS\\_N3783\\_8), and one from high resolution ATCA observations (ATCA\\_1134-37).  We used the 6dFGS DR2 and NED databases to find previously catalogued galaxies in the region, and determine parameters for the NGC~3783 group. We calculate a mean velocity for the group of 2903$\\pm$26 \\kms, and velocity dispersion of 190$\\pm$24 \\kms. The NGC~3783 group is a rare example of a group at the early stage of its evolution. Comparison with other GEMS groups will help in determining the evolutionary path of galaxies within groups in general."
    },
    "0606/astro-ph0606239_arXiv.txt": {
        "abstract": "{One possible way to investigate the nature of the primordial power spectrum fluctuations is by investigating the statistical properties of the local maximum in the density fluctuation fields.} {In this work we present a study of the mean correlation function, $\\xi_r$, and the correlation function for high amplitude fluctuations (peak-peak correlation) in a slighlty non-Gaussian context.} {From the definition of the correlation excess, we compute the Gaussian two-point correlation function and, using an expansion in Generalized Hermite polynomials, we estimate the correlation of high density peaks in a non-Gaussian field with generic distribution and power spectrum. We also apply the results obtained to a scale-mixed distribution model, which correspond to a nearly Gaussian model.} {The results reveal that, even for a small deviation from Gaussianity, we can expect high density peaks to be much more correlated than in a Gaussian field with the same power spectrum. In addition, the calculations reveal how the amplitude of the peaks in the fluctuations field is related to the existing correlations.} {Our results may be used as an additional tool to investigate the behavior of the N-point correlation function, to understand how non-Gaussian correlations affect the peak-peak statistics and extract more information about the statistics of the density field.}  \\keywords {fluctuations fields -- random variable -- correlation function} ",
        "introduction": "Investigation of the statistical properties of cosmological density fluctuations is a very useful tool to understanding the origin of the cosmic structure. Roughly, cosmological models to describe primordial fluctuations can be divided in two classes: Gaussian and non-Gaussian. The most accepted model for structure formation assumes initial\tquantum fluctuations created during inflation and amplified by gravitational effects. The standard inflationary models predicts an uncorrelated random field, with a scale-invariant power spectrum, which follows a nearly-Gaussian distribution \\cite{b9,b10}. However, non-Gaussian fluctuations are also allowed in a wide class of alternative models, such as: the multiple interactive fields \\cite{b11,b12}, the cosmic defects models \\cite{b13} and the hybrid models \\cite{b14,b15}. By discriminating between different classes of models, the statistical properties of the fluctuations field can be used to investigate the nature of cosmic structure. However, non-Gaussian models comprehend an infinite range of possible statistics. As a consequence, performing statistical tests of this kind are not a straightforward task, since there is no adequate general test for every kind of model. To attack this problem, any effort to better understand how the statistical properties of the density fluctuation field affect the observed Universe is welcome, since it may bring extra pieces of information to the investigation of cosmic structure.  Due to the great importance of characterizing non-Gaussian signatures, many statistical approaches have been used to study the distribution of fluctuations in the cosmic microwave background radiation (CMBR) \\cite{b18,b46,b27,b31,b32} and the large scale structure (LSS) \\cite{b28,b19,b29,b20}. One possible way to investigate the nature of the primordial power spectrum fluctuations is by investigating the statistical properties of the local maximum in the density fluctuation fields. Since some of the peak properties, such as number, frequency, correlation, height and extrema, are highly dependent upon the statistics of the fluctuation field, we can gather information about the statistical distribution function by studying the morphological properties of the fluctuations fields \\cite{b8,b22,b35}. Other useful statistical estimator applied to investigate density fluctuations are the wavelets tools \\cite{b24,b25,b38}, the phase correlations \\cite{b17,b26} and the most widely used estimator: the N-point correlations in phase \\cite{b23,b6,b21} or density spaces \\cite{b5,b30,b33,b34,b37,b39}.  The extensive use of the two-point correlation function to characterize the statistical properties of the fluctuations field is justified by the mathematical simplicity in the Gaussian condition, since the mean correlation function can be obtained analytically and it completely specifies a Gaussian distribution (as well as its power spectrum). However, this assumption is not true for a non-Gaussian case, where higher order correlations may give a significant contribution, despite the great effort demanded to compute a wide range of correlations. However, it is possible to detect primordial non-Gaussianity with a non-zero measure of the N-point correlation $ (n \\ensuremath {\\ge} 3)$. In order to achieve a good description of non-Gaussian signatures in cosmic structure, many works have been done to estimate the three-point correlation function and the related bi or tri-spectrum for a few classes of non-Gaussian models \\cite{b6,b32,b27,b21,b41,b42,b43,b44}. It is believed that, with the advent of the CMBR experiments and the high quality surveys in cosmology, the N-point correlation function will be the main statistical descriptor for the cosmic structure.  In this work we present a technique to extract non-Gaussian components from a two-point correlation function. In presence of a non-Gaussian component, we point out that even between two points it is possible to estimate the influence of higher order correlations for models with different statistical description. By performing the calculation of the high order correction terms for the peak-peak correlation function, we show how the amplitude of peaks is dependent of the correlations involved. We point out that for a non-Gaussian statistics the higher order correlations impose some restrictions to the amplitude of high density peaks.  This paper is divided as follows: in Section 2 we give a general description of a random variable field and estimate the peak-peak correlation function for a Gaussian field. In Section 3 we describe the general treatment to obtain the peak-peak non-Gaussian correlation function. In Section 4 we apply the calculations for a slightly non-Gaussian model in two steps: first we consider an approximated solution for a non-Gaussian model with null high order correlations and finally we estimate the complete solution for a slightly non-Gaussian field. In section 5, we summarise our results and discuss the possibilities to use the peak-peak correlation function as a statistical descriptor of the density fluctuation field. ",
        "conclusions": "In this work, we have estimated the two-point mean correlation function and the peak-peak correlation function for the density field. In the Gaussian case, the calculations are simplified, since the Fourier transform of the power spectrum completely describes the random variable. However, for a non-Gaussian field, the calculations are much more complicated, since high order correlations between these two-points may not vanish and strongly contribute to the final function, even in a small deviation from the Gaussian case.  In this work, we showed that, when considering a mixed model, both the mean correlation and the peak-peak correlation functions are much more intense in small scales than in the Gaussian case. This result can be particularly relevant since it is generally accepted that galaxies form in high density regions. In addition, we conclude that the peak-peak correlation function is quite sensitive to the PDF of the fluctuations field, especially for a mixed model. This result suggests that it is possible to use the peak-peak correlation function as a test for the nature of cosmic structures. Nevertheless, we have to be careful when approximating terms for high order correlations, since the peak-peak correlation function is very sensitive to the correction terms. We also point out that correlations of order \\ensuremath{> 2} can be a very important tool to characterize non-Gaussian fields and definitely deserve a deeper investigation. \tEstimating high order correlations allows us to investigate the behaviour of the N-point correlation function, as well as gather more information about the amplitude of the expected high density fluctuation field. As seen in Figure 3, correlations may restrict the amplitude of the density peaks. Furthermore, from the amplitude of the peaks found in the density field (in CMB or LSS fluctuations) we are able to extract more information about the statistics of the density field. It is also good to remember that the presence of correlations of order $ s \\ge 3$ lead to formation of structures in earlier times than would be expected for a model with the same power spectrum but with weaker spatial correlations. The resuls presented in this paper may be used to set new constraints in structure formation models.  One possible application of this method on investigations of primordial non-Gaussianity could be implemented in the search for maximum amplitude fluctuations of the full sky CMB temperature maps, such those derived from WMAP \\cite {b49} and the future Planck mission \\cite {b50}. However, variations in the number of density peaks and their correlation may as well be related to non-Gaussian Galactic foregrounds or other contaminants. Once detected such non-Gaussian trace, we have to be very careful before assigning it to a primordial origin. To avoid misunderstanding, the investigation of the peaks statistics in CMB datasets should be realized over several datasets and different frequencies. Indeed, one possible manner to minimize the foregrounds effects is to analyse the most sensitive and cleaned map, such the WMAP three-year di-biased internal linear combination, WMAP-DILC map, or the WMAP coadded map, with the combination of Q+V+W frequencies. This analysis will be the next test applied to the non-Gaussian mixed model.    \\begin{acknowledgement}  APAA thanks the financial support of FAPESB and CNPq, under grant 1431030005400. ALBR thanks the financial support of CNPq, under grants 470185/2003-1 and 306843/2004-8. CAW was partially supported by CNPq grant 307433/2004-8.  \\end{acknowledgement}"
    },
    "0606/astro-ph0606713_arXiv.txt": {
        "abstract": "A model for the infrared emission of the high redshift ultraluminous infrared galaxy FSC 10214+4724 is presented. The model assumes three components of emission: a dusty torus viewed edge-on, clouds that are associated with the narrow-line region and a highly obscured starburst. It is demonstrated that the presence of clouds in the narrow-line region, with a covering factor of 17\\%, can explain why the mid-infrared spectrum of FSC 10214+4724 shows a silicate feature in emission despite the fact that its torus is viewed edge-on. It is also shown that the same model, but with the torus viewed face-on, predicts a spectrum with silicate emission features that is characteristic of the spectra of quasars recently observed with $Spitzer$. ",
        "introduction": "\\label{}  The z=2.285 galaxy FSC10214+4724 was discovered 15 years ago by Rowan-Robinson et al. (1991). At the time of its discovery it was thought to be one of the most luminous objects in the Universe. It was later found to be gravitationally lensed by a factor 10-100 (Broadhurst \\& Lehar 1995, Graham \\& Liu 1995, Serjeant et al. 1995, Eisenhardt et al. 1996). It is still the best studied ultraluminous infrared galaxy at that redshift. Rest frame optical-UV spectroscopy and spectropolarimetry showed that FSC 10214+4724 exhibits similar characteristics to Seyfert 2 galaxies (Elston et al. 1994, Lawrence et al. 1993, Goodrich et al. 1996). A large quantity of molecular gas is suggested by CO observations (Brown \\& Vanden Bout 1991, Solomon et al. 1992, Scoville et al. 1995). The object is also bright in the submillimetre. It is therefore likely that FSC 10214+4724 is similar to some local ultraluminous infrared galaxies which harbour simultaneously a luminous starburst and an AGN (e.g. Genzel et al. 1998).  Radiative transfer models of emission from dust in a starburst or an AGN were used in an attempt to understand the origin of the infrared emission of FSC10214+4724. Rowan-Robinson et al. (1993) showed that good fits to the overall spectral energy distribution (SED) could be obtained with both a starburst and an AGN model. The dust density in the latter falls off as $r^{-1}$ and varies exponentially with azimuthal angle (Efstathiou \\& Rowan-Robinson 1995). Green \\& Rowan-Robinson (1996) concluded that starburst or AGN activity alone can not explain the SED of FSC 10214+4724 but instead a combination of the two is needed. Granato, Danese \\& Franceschini (1996) found that a good fit to the SED of FSC 10214+4724 and other hyperluminous infrared galaxies could be obtained with pure dust-enshrouded AGN models. The limited wavelength coverage in the infrared is the main reason why modeling attempts in the mid-90s remained inconclusive.  Recently Teplitz et al. (2006) presented a mid-infrared spectrum for FSC 10214+4724 which was obtained with the Infrared Spectrograph onboard $Spitzer$ (Houck et al. 2004). This spectrum places strong constraints on models of the infrared emission of this object. What is interesting about the spectrum presented by Teplitz et al. is that the silicate feature appears in emission. This is surprising given the type 2 nature of the AGN in FSC 10214+4724. Local examples of type 2 AGN (e.g. NGC1068) show silicate features in absorption in agreement with radiative transfer models of dusty tori viewed almost edge-on (Pier \\& Krolik 1992, Granato \\& Danese 1994, Efstathiou \\& Rowan-Robinson 1995). The spetrum presented by Teplitz et al. shows no signs of PAH features except possibly a weak 6.2$\\mu m$ feature. This implies that the starburst if present is weak or highly obscured.   In this letter radiative transfer models of dusty AGN and dusty starbursts are used to fit the now detailed SED of FSC 10214+4724. Section 2 describes briefly the radiative transfer models to be used. In section 3 the results are presented and briefly discussed. A flat Universe is assumed with $\\Lambda =0.73$ and H$_0$=71Km/s/Mpc. ",
        "conclusions": "Detailed modeling of the emission of NLR clouds that are illuminated by a central source is beyond the scope of this letter. The emission of such clouds is approximated  with that of spherical clouds of dust at a constant temperature. It is assumed that the density in the cloud is uniform and the optical depth in the V band from the center to the surface is 10. The same grain mixture used for the dusty starburst radiative transfer calculations is assumed. A good fit to the $Spitzer$ spectrum can be obtained by assuming that the NLR dust is concentrated in clouds at a temperature of 200 and 610K with the 200K cloud(s) being about an order of magnitude more luminous. This is broadly in agreement with the findings of Teplitz et al. (2006). The starburst model assumed is the $\\tau_V = 200$ model shown in Fig. 1.  The model assumed for the torus is the most geometrically thin of the models considered by Efstathiou \\& Rowan-Robinson ($\\Theta_1=30^o$). This implies a torus half opening angle of 60$^o$ which, however, is very poorly constrained. The model also assumes a ratio of inner to outer disc radius of 0.05, a ratio of height to outer radius of 0.5 and an equatorial optical depth at 1000\\AA~ of 1000. It is assumed that the tapered dusty disc in FSC 10214+4724 is viewed edge-on. This means that its mid-infrared radiation is heavily suppressed. The sum of the emission of the NLR dust and the disc is shown in Fig. 2 and 3 whereas the individual spectra are shown in Fig. 4. Note that because the starburst is very optically thick it does not make a significant contribution to the mid-infrared flux. This can explain why there is no evidence of PAH features in the $Spitzer$ spectrum published by Teplitz et al. except for a weak 6.2$\\mu m$ feature that is reproduced by the model.  As can be seen from the plot of the edge-on and face-on spectra emitted by the torus in Fig. 4, the AGN luminosity is highly anisotropic. To obtain the total luminosity one needs to multiply the `apparent' torus luminosity with the anisotropy correction factor $A$ which is defined to be  $$  A(i) = {{ \\int_0^{\\pi/2} S(i')\\ \\sin i' \\ di'} \\over {S(i)}}  $$ where $S(i)$ is the bolometric emission at angle $i$. For the particular tapered disc model and inclination, $A$ is found to be 18.1. The emission by the NLR dust is assumed to be isotropic. With this correction the AGN luminosity is predicted to be $6.9 \\times 10^{14} L_\\odot$. The luminosity of the 200K NLR dust is predicted to be $1.0 \\times 10^{14} L_\\odot$ or 15\\% of the total AGN luminosity. As mentioned above the luminosity of the 610K dust is about an order of magnitude lower or 2\\% of the total. As the NLR clouds are assumed to be optically thick, the fraction of total AGN luminosity emitted by the 200 and 610K dust is also the covering factor of the NLR dust. The starburst luminosity is predicted to be $7.6 \\times 10^{13} L_\\odot$. Note that all of the quoted luminosities need to be reduced by a factor  10-100 due to gravitational lensing. Lens models suggest that the optical emission may be more highly magnified than the infrared emission that dominates the bolometric luminosity. Eisenhardt et al. (1996) estimate a magnification factor for the infrared emission of 30 which implies that the AGN luminosity is of the order of $2.3 \\times 10^{13} L_\\odot$. Note, however, that since the emission by the NLR dust and the torus does not arise from the same region it is possible that their magnification will be different.   \\begin{figure} \\epsfig{file=figure2.ps, angle=0, width=8cm} \\caption{ Fit to the spectral energy distribution of FSC 10214+4724. The dotted and dashed lines show the contributions of the AGN and starburst respectively to the total emission (solid line). Data from Teplitz et al (2006), IRAS, Rowan-Robinson et al. (1993). } \\end{figure}  \\begin{figure} \\epsfig{file=figure3.ps, angle=0, width=8cm} \\caption{ As Fig. 2 but showing the fit to the mid-infrared spectrum in more detail. } \\end{figure}   \\begin{figure} \\epsfig{file=figure4.ps, angle=0, width=8cm} \\caption{ Comparison of the edge-on and face-on views of FSC 10214+4724 according to the model presented here. The thick (red) dotted line gives the spectrum emitted by the edge-on torus and the dashed line (magenta) gives the emission of the dust in the cones. The thick (red) solid line gives the total emission in the edge-on case. The thin (blue) dotted line gives the predicted spectrum for a face-on view of the torus and the thin (blue) solid line gives the sum of that emission to the emission of the dust in the cones which is assumed to radiate isotropically. The vertical scale is arbitrary. } \\end{figure}    As it is demonstrated in Figs. 2, 3 and 4  the presence of dust in the narrow-line region of an AGN nucleus, with a covering factor of 17\\%, can explain the fact that silicate features appear in emission both in luminous type 2 and type 1 AGN. The presence of these features in type 2 AGN is entirely consistent with the unified scheme (Antonucci 1993) as a dust torus that makes an insignificant contribution in the mid-infrared may still be present. The anisotropy of the emission from such a torus is so high that when the same system is viewed face-on the torus emission dominates and the effect of the presence of the NLR dust is to make the silicate emission features more prominent. The face-on spectrum shown in Fig. 4 is similar to the spectra of quasars recently observed with $Spitzer$.  A question that naturally arises is why lower luminosity type 2 AGN do not show the same effect. There are probably two reasons for this: firstly in lower luminosity type 2s the ratio of the luminosity of the acompanying starburst to that of the AGN is higher so what we see in the mid-infrared is the complex of PAH features and silicate absorption features emitted by the starburst. Secondly the torus opening angle in lower luminosity AGN may be smaller and therefore the fraction of the source luminosity that is intercepted by NLR dust is smaller. This was already suggested by Lawrence (1991) who estimated the statistics of type 1 and type 2 AGN from different surveys. As proposed by Lawrence this dependence on luminosity may arise if the thickness of the torus is independent of luminosity and its inner radius is controlled by dust sublimation."
    },
    "0606/astro-ph0606525_arXiv.txt": {
        "abstract": "We present B, V, R, and H$\\alpha$ photometry of 8 clusters in the Small Magellanic Cloud, 5 in the Large Magellanic Cloud, and 3 Galactic clusters, and use 2 color diagrams (2-CDs) to identify candidate Be star populations in these clusters.  We find evidence that the Be phenomenon is enhanced in low metallicity environments, based on the observed fractional early-type candidate Be star content of clusters of age 10-25 Myr.  Numerous candidate Be stars of spectral types B0 to B5 were identified in clusters of age 5-8 Myr, challenging the suggestion of \\citet{fab00} that classical Be stars should only be found in clusters at least 10 Myr old.  These results suggest that a significant number of B-type stars must emerge onto the zero-age-main-sequence as rapid rotators.  We also detect an enhancement in the fractional content of early-type candidate Be stars in clusters of age 10-25 Myr, suggesting that the Be phenomenon does become more prevalent with evolutionary age.  We briefly discuss the mechanisms which might contribute to such an evolutionary effect.  A discussion of the limitations of utilizing the 2-CD technique to investigate the role evolutionary age and/or metallicity play in the development of the Be phenomenon is offered, and we provide evidence that other B-type objects of very different nature, such as candidate Herbig Ae/Be stars may contaminate the claimed detections of ``Be stars'' via 2-CDs. ",
        "introduction": "Classical Be stars are rapidly rotating, near main sequence B-type stars which show or have shown hydrogen Balmer emission \\citep{jas81}.  It is widely accepted that Be stars have gaseous, geometrically thin circumstellar disks \\citep[for a review, see][]{por03}. As summarized in \\citet{por03}, numerous mechanisms have been proposed to explain how circumstellar disks may form around classical Be stars, including the wind-compressed disk (WCD) model \\citep{bjo93}, the magnetically torqued wind-compressed disk model \\citep{cas02}, and processes related to non-radial pulsations of the stellar photosphere \\citep{riv01}, yet none have been completely successful in explaining the observed phenomena.  Beginning with the work of \\citet{str31}, it has often been speculated that the fundamental source of the Be phenomenon is tied to the rapid rotation of these stars: the canonical assertion is that classical Be stars rotate at v$_{eq}$/v$_{crit}$ $\\sim$ 70-80\\% of their critical velocity \\citep{por96, por03}.  However, recent theoretical work incorporating the effects of equatorial gravity darkening into studies of rotation rates \\citep{tow04} suggests there is a degeneracy in the measurement of rotation rates such that increasing v$_{eq}$ from 80\\% to 100\\% of the critical velocity has no effect on observed line widths.  Thus classical Be stars may actually be rotating at or near their critical breakup velocity.  Alternatively, \\citet{cra05} suggest that a subset of classical Be stars might be rotating at sub-critical rates as low as 0.4-0.6 v$_{crit}$.  Studying stars from the Bright Star Catalog \\citep{hof82} and its supplement \\citep{hof83}, \\citet{zor97} found that the mean frequency of Be stars with respect to normal B stars was $\\sim$17\\% in the Galaxy.  \\citet{zor97} also found that the peak frequency occurred for spectral type B1, where 34\\% of B stars were Be stars, and noted no difference in frequencies between giant, dwarf, and subgiant type B stars.  It should be noted however that the reliability of Bright Star Catalog spectral types is often questionable.  Early studies also concluded that Galactic Be stars were present from the zero-age-main-sequence to the terminal-age main sequence \\citep{mer82,sle85}.  Due in part to the advent of CCDs, systematic studies of the Be populations of star clusters and associations have expanded past our Galaxy to include the Magellanic Clouds \\citep{fea72,gre92,gre97, die98,kel99,gre00,kel00,ols01}.  These extragalactic survey programs have typically used 2-color diagram (2-CDs) photometric techniques to identify the frequency of B-type stars exhibiting excess H$\\alpha$ emission (``Be stars'') relative to normal B-type stars in stellar associations and clusters.  General trends in the fractional Be content of a cluster as a function of its age and/or metallicity have been observed, leading to suggestions that secondary mechanisms, besides rapid rotation, might influence the development of Be circumstellar disks.  \\citet{mer82,gre97,fab00} and \\citet{kel04} have found that the frequency of the Be phenomenon seems to peak in clusters with a main sequence turn-off of B1-B2, leading to the suggestion that the Be phenomenon is related to a star's evolutionary age.   \\citet{fab00} suggested the Be phenomenon will start to develop only in the second half of a B star's main sequence lifetime, owing to structural changes in the star.  Specifically, \\citet{fab00} noted that Be star-disk systems should start to appear in clusters 10 Myr old, corresponding to the mid-point main sequence lifetime of B0 stars \\citep{zor97}, and their frequency should peak in clusters 13-25 Myr old, corresponding to the mid-point main sequence lifetime of B1-B2 stars \\citep{zor97}.  Curiously, \\citet{kel99} claimed to detect numerous Be stars in the Small Magellanic Cloud (SMC) cluster NGC 346, which has an age between 2.6 Myr \\citep{kud89} and 5 Myr \\citep{mas89}; this result seems to conflict with Fabregat \\& Torrejon's assertion.  \\citet{hil93} also suggested that numerous classical Be stars might be present in the young cluster NGC 6611; however, it is possible that these objects are pre-main-sequence Herbig Ae/Be stars \\citep{fab00}.  Numerous mechanisms have been proposed to explain this apparent evolutionary effect, including structural changes within the central star \\citep{fab00}, an evolutionary spin-up of the central star \\citep{mey00, kel04}, and spin-up due to mass-transfer in binary systems \\citep{mc05b}.   Several photometric studies \\citep{gre92,maz96,gre97,mae99,kel04} have suggested that the Be phenomenon may be more prevalent in low metallicity environments, based on comparisons of the apparent fractional Be populations of Galactic, Large Magellanic Cloud (LMC), and SMC clusters.  IUE, HST, and FUSE observations \\citep{gar85,bia96,ful00} show that stars tend to have lower wind velocities in metal poor environments. The WCD model \\citep{bjo93} suggests that reducing terminal wind velocities would increase the likelihood that a B star experiencing mass-loss would be able to retain this matter in the form of a circumstellar disk, although \\citet{owo96} claim non-radial line forces may inhibit such a scenario.  \\citet{ma99a} suggests that the increase in the frequency of Be stars in low metallicity environments is related to the presence of a greater number of rapidly rotating stars in these locations, owing to a reduced coupling of the magnetic field and the young stellar object (\\citealt{mae99}; see also \\citealt{pen04}).  The accuracy of these initial interpretations is hampered by numerous factors.  \\citet{fab00} note that the photometric surveys used by many to identify Be stars only detect Be stars with high levels of emission, and likely miss weak emitters. Furthermore, since Be stars are known variables which can experience active and quiescent stages \\citep{tel00,bjo02}, it is likely that some Be stars in a cluster might be in a quiescent stage and hence not identified as Be stars by these techniques.  This possibility has been demonstrated by various followup investigations \\citep{hum99,kel99}, which do not detect 100\\% of previously identified Be stars in certain clusters; additionally, a significant number of previously unidentified Be stars are uncovered in these studies.  \\citet{kel00} note some spurious detections in photometrically identified Be stars located in crowded cluster cores. Even some spectroscopically identified Be stars \\citep{maz96} have been shown to be spurious detections caused by the presence of diffuse background H$\\alpha$ emission \\citep{kel98}. The number of clusters with accurately determined Be frequencies is still limited, especially regarding SMC clusters, thus small number statistics become important when attempting to draw broad conclusions from the literature data set.  In this paper, we provide a significant increase in the number of extragalactic clusters whose candidate Be populations have been identified photometrically.  In section 2, we outline our observations and data reduction techniques.  In section 3, we detail our identification of candidate Be stars in 8 SMC, 5 LMC, and 3 Galactic clusters.  We offer a detailed discussion of how these results may shed insight into the role age and/or metallicity play in the development of the Be phenomenon in Section 4. ",
        "conclusions": "In order to better probe the role of evolutionary age and/or metallicity may play in the development of the Be phenomenon, we approximated rough spectral types for our candidate Be stars on the basis of rough magnitude bins associated with spectral ranges.  We first transformed the apparent magnitudes of our candidates to the absolute scale, using the standard equation m$_{V}$ - M$_{V}$ = 5 log d - 5 + A$_{V}$, where A$_{V}$ = R$_{V}$ * E$_{(B - V)}$.  For LMC clusters we used a distance modulus of 18.5 \\citep{wes90} and R$_{V}$ = 3.41 \\citep{gor03}, while for SMC clusters we adopted a distance modulus of 18.9 \\citep{wes90} and R$_{V}$ = 2.74 \\citep{gor03}.  For our Galactic clusters we assumed R$_{V}$ = 3.1 and used a distance modulus of 11.31 for NGC 2186 \\citep{mof75} and 13.24 for NGC 2439 \\citep{whi75}.  E$_{(B - V)}$ values for most clusters are listed in Table \\ref{allsum}.  Due to the lack of available published reddening values for Bruck 60, Bruck 107, and HW 43, we assigned these clusters the mean reddening value for the SMC, 0.037, as determined by \\citet{sch98}.  Following the work of \\citet{gre97}, we used the calibration for main sequence stars given in Table 5 of \\citet{zor97} to transform the absolute magnitude of each of our candidate Be stars into a spectral type.  As discussed by \\citet{gre97}, a large uncertainty is associated with such a transformation technique; hence, our estimated spectroscopic classifications should be considered only crude approximations.  We summarize these rough spectroscopic classifications for our LMC and SMC clusters in Table \\ref{photspectypesum}. In column 2 of Table \\ref{photspectypesum}, we assigned an age label to each cluster given by: 5 Myr $<$ Age $<$ 8 Myr = very young (vy), 10 Myr $<$ Age $<$ 25 Myr = young (y), and 32 Myr $<$ Age $<$ 158 Myr = old (o).    Note that we classified LH 72 as ``very young'' (vy) although a range of ages have been suggested for this cluster (see Table \\ref{allsum}). Given the suggested multiple epochs of star formation in NGC 1850 (see Section 3.3.2), we designated this cluster as ``old?'' (o?).  We have included all photometrically identified candidate Be stars in Table \\ref{photspectypesum}, including those classified as ``possible detections'' in Table \\ref{photsum}.   We have combined the results of Table \\ref{photspectypesum} into two bins: ``early-type'' candidate Be stars, corresponding to rough spectral types B0-B3, and ``later-type'' candidate Be stars, corresponding to rough spectral types B4-B5.  The results of this binning are given in Table \\ref{earlylate}, and we suggest that such averages should mask the general level of uncertainty associated with the spectral types we have assigned.  Clusters with fewer than 20 B-type main sequence objects in either of these spectral bins were excluded from consideration to lessen the effects of small number statistics.  The standard deviation errors quoted in Table \\ref{earlylate} were calculated assuming the data followed a binomial distribution.  We have also averaged the fractional Be content, binned to early- and later-type stars, of clusters having like age and metallicity properties, as summarized in Table \\ref{agezstats}.  The uncertainties quoted in this table merely represent the propagation of the previously calculated standard deviations, i.e. $\\sigma$ = N$^{-0.5}$ ($\\sigma_{i}^{2}$ + $\\sigma_{i+1}^{2}$ + ...)$^{0.5}$.   \\subsection{Metallicity Effects}  We first examined our dataset to identify any trends in the fractional candidate Be content as a function of metallicity.  Following the work of \\citet{mae99}, we first considered the fractional early-type (B0-B3) Be content of ``young'' clusters having ages of 7.0 $<$ log (t) $<$ 7.4.  Data matching these criteria in Table \\ref{earlylate} are plotted in Figure \\ref{midagez}, where open circles correspond to literature data, filled circles correspond to data presented in this study, and crosses represent the average of all clusters' fractional Be content within a metallicity bin.  It is possible to find a wide range of reported metallicities for many of the individual clusters presented in this survey, as noted by \\citet{mae99}; hence, following their example, we have chosen to represent all SMC clusters with one average metallicity and all LMC clusters with another average metallicity value.  We also follow the practice of \\citet{mae99} of assigning single, average metallicity values to both Galactic clusters located exterior and interior to the Solar location (see Figure \\ref{midagez}).  We recognize that this is an oversimplification, but given the large scatter in the current available values for individual cluster metallicities, it seems the best approach at present.  From Figure \\ref{midagez}, it is clear that there exists a wide range of Be / (B + Be) ratios within each metallicity bin.  The systematic errors inherent in the use of the 2-CD technique undoubtedly contribute to this scatter; differences in the environmental properties of specific individual clusters also might influence the observed scatter of fractional Be content.  However, it is clear that a trend in the fractional Be content with metallicity exists.  This trend is better seen in Figure \\ref{averagedz}, where the filled circles represent the average fractional Be content of ``young'' clusters.  The average fractional Be content of SMC metallicity clusters is $>$ 2$\\sigma$ higher than the average Galactic (interior or exterior) value, as documented in Table \\ref{agezstats}.  While Figure \\ref{averagedz} shows the same general trend depicted by \\citet{mae99}, we believe that the improved statistics provided by the present study significantly strengthens the claim of a trend with metallicity.  The trend claimed by \\citet{mae99} was based upon the observation of one cluster in the SMC metallicity bin of z=0.002.  The present study's use of 4 SMC clusters reduces the average fractional Be content within this metallicity bin from 39\\% \\citep{mae99} to 32\\%.  We must note that our results are not inconsistent with the general conclusion reached by \\citet{mc05b}, who found no evidence of a metallicity trend in their study of Galactic clusters.  As seen in Figure \\ref{averagedz}, we also observe little convincing evidence of a trend over the smaller metallicity range probed by Galactic clusters; however, evidence of a metallicity trend develops when one extends surveys to include broader regions of metallicity space, such as the LMC and SMC.  If one assumes that the fractional Be content of clusters depends on metallicity, then such a trend also should be apparent when examining clusters in age denominations other than ``young'' clusters.  We have plotted the average fractional early-type Be content of ``very young'' clusters, depicted as open triangles, and ``old'' clusters, depicted as open squares, in Figure \\ref{averagedz}.  Although the quality of available data for these age groups is severely limited, and no data exist for Galactic clusters, we suggest that the ``very young'' and ``old'' cluster data plotted in Figure \\ref{averagedz} are not inconsistent with the Be phenomenon being more prevalent in low metallicity environments.  We do recognize, however, that there is still some controversy over the metallicity effect.  For example, the results of \\citet{mar06} find a Be fraction in the LMC cluster NGC 2004 which is about the same as the standard Galactic value of 17\\%. To resolve these lingering uncertainties, we strongly suggest that additional observations of clusters in these extreme age ranges must be made, spanning a wide range of metallicities, in order to provide the statistics necessary to more definitively identify any trend present.    \\subsection{Evolutionary Age Effects}  Inspection of Tables \\ref{photspectypesum}, \\ref{earlylate}, and \\ref{agezstats} clearly reveal that we have found a substantial number of candidate Be stars in clusters ranging in age from 5 Myr to 8 Myr.  From Table \\ref{photspectypesum}, it is clear that these young clusters not only contain B0 type candidate Be stars, but they also have a significant number of later-type candidate Be objects.  If the Be phenomenon develops in the second half of a B star's main sequence lifetime, we would not expect to find B0 type stars, let alone B4-B5 type stars, in these very young clusters.  Furthermore, the average fractional candidate Be content in these very young clusters, detailed in Table \\ref{agezstats}, is similar to the nominal frequency of the Be phenomenon in the Galaxy of 17\\% \\citep{zor97}.  Clearly our present understanding of how the Be phenomenon is related to a star's evolutionary age would be altered if we could establish that these very young \\textit{candidate} Be stars were not remnant pre-main-sequence star-disk systems.  The similarities between the fractional candidate Be population found in these very young clusters with respect to the nominal frequency of the Be phenomenon observed in the Galaxy is interesting.  Classical Be stars of such a young age clearly would not have spent enough time on the main sequence to significantly spin-up via the mechanism proposed by \\citet{mey00}, or via mass transfer in a binary system \\citep{mc05b}.  Since it is widely accepted that the Be phenomenon is inherently tied to rapid rotation, our data seem to suggest that stars emerging from their pre-main-sequence phase must possess a wide range of rotational velocities; namely, a significant number of objects must be rotating near their critical breakup velocities at the zero-age-main-sequence (ZAMS).  We will explore the true nature of many of these extremely young candidate Be stars in a later paper.  From Table \\ref{agezstats}, it is clear that the fractional early-type Be content of ``young'' clusters is significantly higher than the nominal frequency of the Be phenomenon (17\\%); furthermore, the fractional later-type Be content of these clusters is significantly lower than the early-type content. Examining the fractional early- and later-type Be content of ``old'' clusters in Tables \\ref{earlylate} and \\ref{agezstats}, we also find evidence that the Be phenomenon in these clusters is more prevalent than the nominal value of 17\\%.  It is unclear from our dataset whether there is any difference between the frequency of early-type and later-type Be stars in these old clusters.  Our detection of early-type objects in ``old'' clusters is very curious, and we can not find a simple explanation in the literature which explains why such objects would be detected.  The slight enhancement of the fractional later-type Be content in our old clusters and the much larger enhancement of the fractional early-type Be content in our young clusters supports previous suggestions that the Be phenomenon is more prevalent in the later stages of a B star's main sequence lifetime \\citep{mer82,gre97,fab00,kel04,mc05b}.  Recall however, that the results from our very young clusters indicate that such an evolutionary effect \\textit{may not} be the sole mechanism responsible for the development of the Be phenomenon.  We now consider the mechanisms which might be responsible for the observed evolutionary effect in the Be phenomenon.  The detection of a sizable fraction of candidate Be stars in clusters younger than 10 Myr suggest that stars emerging onto the ZAMS might display a wide range of rotational rates, including stars which are rotating near their critical breakup velocity, hence exhibiting the Be phenomenon.  The theoretical models of \\citet{mey00} and \\citet{mae01} illustrate that the ratio of angular velocity to critical angular velocity ($\\Omega$ / $\\Omega_{critical}$) steadily increases throughout the main sequence lifetime of early-type B stars, which these authors suggest might explain why the Be phenomenon is more prevalent in the later part of a B star's main sequence lifetime.  While we agree that the general notion of an evolutionary spin-up provides a reasonable explanation for most of the observational properties documented here and in the literature, we also consider some of other implications of Meynet's and Maeder's models.  Table \\ref{lifetimes} lists the main sequence lifetimes predicted by the models of \\citet{mey00} and \\citet{mae01} for rotating and non-rotating stars.  Interestingly, stars starting with an enhanced rotation rate on the ZAMS will experience a longer main sequence lifetime than stars of similar mass which start with a small or null rotation rate.  Recall that the fundamental underlying mechanism believed to differentiate classical Be stars from ``normal'' B-type objects is rapid rotation.  We speculate that the different main sequence lifetimes experienced by rapidly rotating versus slowly rotating B-type stars might influence the observed enhancement of the Be phenomenon in the later stages of a star's main sequence lifetime. Consider a SMC cluster of age 25 Myr, which Table \\ref{lifetimes} suggests will only have B3 or later-type ``normal B stars'' remaining on the main sequence, assuming they started on the ZAMS as slow rotators.  In addition to B3 or later-type rapidly rotating classical Be stars, B2 type classical Be stars will also still reside on the main sequence when the cluster is 25 Myr old.  Hence, the fractional Be content derived for this cluster via the 2-color diagram technique will not be comparing Be to normal B stars of a similar range of spectral types.  The additional presence of B2 type Be stars in this cluster might inflate the observed fractional Be content.  Since clusters younger than 10 Myr old should still have their entire B0-B9 sequence present on the main sequence, regardless of whether one considers rapidly or non-rapidly rotating stars, we expect no enhancement of the Be phenomenon due to this effect.  Conversely, we would expect that the expansion of main sequence lifetimes due to rotation will greatly affect clusters in the 16-26 Myr age range.  In these clusters, ``normal'' B1-B2 stars which were slowly rotating when they reached the ZAMS would have evolved off of the main sequence. B1-B2 type classical Be stars, which were rotating much more rapidly when they reached the ZAMS, will still be observable on the main sequence, hence B1-B2 and B3-later type Be stars will be detected via 2-CDs.  Because the Be phenomenon is most prevalent amongst B1-B2 type stars, one might expect a very large enhancement of the Be phenomenon to be detected in these clusters.  When one considers the fractional Be content of older clusters, the Be phenomenon would likely still be enhanced by this effect, but to a lesser degree than clusters 16-26 Myr old, as the prevalence of the Be phenomenon is lower in later spectral types.  Our data are consistent with this speculative scenario.  No strong enhancement in the Be phenomenon is seen in our very young clusters, while moderate and strong enhancements in the Be phenomenon are seen in our old and young clusters respectively.  From Table \\ref{agezstats}, one might question why we observe early-type objects in our ``old'' clusters.  We would expect that most slowly rotating early-type stars should have evolved off of the main sequence in these clusters.  A limited number of early-type classical Be stars might still reside on the main sequence, owing to the extended main sequence lifetime afforded to rapid rotators.  It is a little more difficult to explain the nature of early-type non-Be stars detected in these clusters; however, it is possible that these objects are blue stragglers.  Note that the binary scenario for the formation of Be stars predicts that some should appear as blue stragglers \\citep{pol91}; hence, if ``normal'' B-type blue stragglers populate these clusters, some of the clusters' Be population may have their origin linked to binarity, not the suggested extension of main sequence lifetimes.  It is clear that additional observations are needed to investigate the nature of both the early-type candidate Be stars and normal B stars in such ``old'' clusters.  Differentiating between an enhancement in the Be phenomenon which is due to the spin-up of stars as they evolve along the main sequence versus an extension of the main sequence lifetime in rapidly rotating stars will likely be a difficult task.  Of course, it is entirely possible that both mechanisms could play a role in producing the observed enhancements.  One could examine the statistical distribution of rotational velocities of early-type versus later-type B stars in clusters spanning a range of ages to look for evidence of a systematic spin-up of the overall distribution with age.  Conversely, if one could find numerous classical Be stars in a cluster having earlier spectral types than all of the normal main sequence stars present, this would provide evidence that the extended main sequence lifetimes of rapid rotators might contribute to the observed enhancements of the Be phenomenon.  In practice, measuring the spectral types of classical Be stars to high accuracy is difficult, so this later type of observational undertaking would likely be challenging.  \\subsection{Galactic Clusters}  We also used the calibration of \\citet{zor97} to assign crude spectral types to the candidate Be stars we identified in our Galactic clusters (see Table \\ref{galacticspectypes}).  Due to the large uncertainty in the distance to NGC 2383, we have not attempted to convert the observed apparent magnitudes of this cluster's candidate Be stars into spectral types.  Recall that our NGC 2186 photometry includes systematic errors which make our targets appear $\\sim$1 magnitude too faint; to correct for these systematics, we have applied a 1 magnitude offset to our data to compute the spectral types listed in Table \\ref{galacticspectypes}.  As discussed in Section 3.3.3, we detected 4 of 5 previously identified Be stars in NGC 2439 via our 2-CD and 1 previously undetected candidate Be star.  Spectral types for two of the known Be stars are given in \\citet{sle85}: star White $\\# 6$ \\citep{whi75} is a B2 V, which generally agrees with our rough B0 classification of this object (NGC 2439:WBBe 1 in Table \\ref{photdetails}), and star White $\\# 81$ \\citep{whi75} is a B1 V, which generally agrees with our rough B0 classification of this star (NGC 2439:WBBe 2 in Table \\ref{photdetails}).  \\subsection{ELHC Fields}  As previously described, we have detected ELHC \\# 1,3,4,6,7,8,12,13, and 19 as candidate emission-line objects via 2-CDs and failed to detect ELHC 5, 11, and 20. This 75\\% detection rate offers supporting evidence that ELHC stars might be bonafide pre-main-sequence objects.  Recent low-resolution spectroscopic and JHK photometric observations of some of these ELHCs \\citep{dew05} have revealed that some objects might have nearby, previously unresolved neighbors; furthermore, the lack of a strong near IR excess and lack of forbidden emission lines has led these authors to question whether many of these ELHCs are truly Herbig Ae/Be stars or whether they are classical Be stars.  With respect to our results, we also had noted that ELHC 11 seemed to be comprised of 2 components.  The low-resolution spectroscopy of ELHC 11 and 20, which were both not detected as emission-line objects in our study, revealed these sources to have H$\\alpha$ absorption lines slightly filled in with emission \\citep{dew05}.  Because the 2-CD technique fails to detect weak H$\\alpha$ emitters, it is not surprising that we failed to detect these objects.  Based on these results, we would predict that when spectroscopic observations of ELHC 5 are made, they will also reveal a partially filled in absorption line at H$\\alpha$.  In a future publication, we will present the results of moderate resolution spectroscopic and near-IR photometric observations of many of these ELHC stars which were obtained to further probe the true nature of these objects.  \\subsection{Limitations of the 2-CD technique}  We briefly discussed some of the known limitations of the 2-CD technique in the Introduction, and now consider this situation in more detail.  The 2-CD technique typically assumes that all stars within an appropriate range of colors that have a (R-H$\\alpha$) magnitude in excess of some threshold are ``Be stars''. However, as many Galactic O, B, and A type supergiants are known to show frequent H$\\alpha$ emission, one expects that the detection of such objects via the 2-CD technique would artificially inflate the number of detected ``Be stars''.  For example, \\citet{ols01} obtained follow-up spectroscopy of star S132 in the LMC cluster LH 72, which they identified as a Be star from a 2-CD, and classified this object as a B8 Ia star.  The luminosity class of S132 excludes it, by definition, from being a classical Be star.  Similarly, \\citet{kel99} identified NGC 346:KWBBe 13 as a ``Be star'' via a 2-color diagram. Unfortunately this object is the well known HD 5980 Wolf-Rayet/Luminous Blue Variable system, thus it is abundantly clear that it has been mis-classified by the simple 2-CD method as a (classical) ``Be star''.  In principle, one would expect that the location of these objects on a CMD would identify them as ``contaminants''; however, this is clearly not always the case.  It is possible that these stars are background objects with respect to the clusters with which they have been associated.  It is also possible that a significant amount of their radiation is attenuated by a localized region of dust.  These two examples demonstrate that it is not unreasonable to expect OBA supergiants to ``contaminate'' 2-CD detections. We suggest the frequency of contamination by such objects is likely to be greater when one uses the 2-CD technique to investigate field star populations \\citep{kel99, kel01}, compared to cluster populations, as the evolutionary age of a cluster population should restrict the range of spectral types of OBA supergiants present.  Recent follow-up low-resolution spectroscopic and near IR photometric observations of some ELHCs \\citep{dew05} have raised issues with their classification as Herbig Ae/Be objects. Even assuming the most pessimistic interpretation of the results of \\citet{dew05}, i.e. that ELHC 7 is the only true Herbig Ae/Be star of these candidates as it is the only one which exhibits a large near IR excess, the implications of these results regarding the proper interpretation of 2-CDs can not be ignored.  We detected ELHC 7 on a 2-CD as being an excess H$\\alpha$ emitter; therefore, we have solid evidence that Herbig Ae/Be stars may be detected via this technique.  Clearly, it is not sufficient to merely label all excess H$\\alpha$ emitters detected via 2-CDs as ``Be stars'' as is nearly universally done in the literature. Given a cluster of sufficient age, and assuming coeval development, one would not expect to find any remaining pre-main-sequence stars; hence, this type of contamination may not be an issue.  However, when examining the background field Be population surrounding a cluster using the 2-CD technique \\citep[see][]{kel99}, we assert that one must be wary of detecting Herbig Ae/Be stars, which would serve to artificially inflate the ratio of detected classical Be stars.  \\citet{kel00}, \\citet{ols01}, Olsen (2003, private communication), and the present study all found evidence that diffuse background contamination likely inflates the number of true classical Be stars detected via the 2-CD technique.  A typical form of these spurious detections we encountered were random, quasi point-source-like pockets of emission which were misconstrued as stars, especially when the \\textit{daofind} image roundness parameters were slightly relaxed.  One also might consider how diffuse sky nebulosity affects the measurement of real sources.  The IRAF photometry methods we used sampled the sky background within an annulus of a given radius around sources and used the median pixel value as an estimate of the background.  In extremely variable backgrounds, for example containing H II filamentary structures, it is not always clear that this technique fully accounts for the true background present. While we do not propose a better technique to estimate background fluxes in regions of complex nebulosity, we feel it is important to remind the reader of some of the difficulties associated with observations of embedded objects when considering the results of these and similar 2-CD investigations.  The addition of other independent techniques, such as IR photometry and broadband or spectropolarimetry, could provide much needed cross-checks on such identifications."
    },
    "0606/astro-ph0606536.txt": {
        "abstract": "We present the evolution of the volume averaged properties of the rest-frame optically luminous galaxy population to $z\\sim3$, determined from four disjoint deep fields with optical to near-infrared wavelength coverage.  Our use of independent lines of sight substantially mitigates field-to-field variations.\u00a0 We select galaxies above a fixed rest-frame \\mv-band luminosity of $3\\times 10^{10}$\\lumsol\\ and characterize their rest-frame ultraviolet through optical properties via the mean spectral energy distribution (SED).  To measure evolution we apply the selection criteria to a sample of galaxies from the Sloan Digital Sky Survey (SDSS) and COMBO-17 survey.  The mean rest-frame 2200\\ang\\ through \\mv-band SED becomes steadily bluer with increasing redshift but at all redshifts $z<3$ the mean SED falls within the range defined by ``normal'' galaxies in the nearby Universe.  Rest-frame 4000\\AA/Balmer breaks are present in the volume averaged SED at all redshifts, indicating significant light and mass contributions from evolved populations, even at $z\\sim 2.8$.  We measure stellar mass-to-light ratios (\\mlstar) by fitting models to the rest-frame UV-optical SEDs. Multiplying our volume averaged \\mlstarlam{V}\\ estimates at each redshift by the measured mean \\mv-band luminosity density, we derive the stellar mass densities \\rhostar.  The stellar mass density in galaxies selected at a fixed luminosity has increased by a factor $3.5-7.9$ from $z=3$ to $z=0.1$, where the range includes the uncertainty due to field-to-field variance within our own data. If we use our observed \\mlstarlam{V} evolution to select galaxies at a fixed mass, the stellar mass density evolves by a factor of $5.3- 16.7$.  After correcting to total, the measured mass densities at $z<2$ lie below the integral of the star formation rate (SFR) density as a function of redshift as derived from UV selected samples after a standard correction for extinction.  This may indicate a systematic error in the total \\rhostar\\ or SFR($z$) estimates.  We find large discrepancies between recent model predictions for the evolution of \\rhostar\\ and our results, even when our observational selection is applied to the models.  Finally we determine that Distant Red Galaxies (selected to have $J_s - K_s>2.3$) in our \\l{V}\\ selected samples contribute $30\\%$ and 64\\% of the stellar mass budget at $z\\sim2$ and $z\\sim 2.8$ respectively.  These galaxies are largely absent from UV surveys and this result highlights the need for mass selection of high redshift galaxies. ",
        "introduction": "\\label{Intro}  Our knowledge of the high redshift Universe has increased rapidly over the past decade and is becoming ever more comprehensive.  The initial large advances in this field were enabled by the efficient selection of $z\\sim3$ starforming galaxies based on their optical light, e.g. \\citet{Stei96} and \\citet{Stei99}, and by the development of efficient multi-object spectrographs on large 8-10m class telescopes.  Objects selected by these techniques allowed us for the first time to study statistically significant samples of galaxies at $z>1.5$.  Resultant follow-up work on these optically (rest-frame ultraviolet; UV) selected objects at $z\\sim3$ (Lyman Break Galaxies; LBGs) and subsequent samples at $z\\sim 2$ (BM/BX objects; Steidel et al. 2004; Adelberger et al. 2004) have demonstrated that these galaxies have generally low extinctions $E(B-V)\\sim 0.15$ and modest ages ($<1$ Gyr) and stellar masses (\\mstar$\\sim 10^{10}$\\msol) with a tail to high end values existing at bright \\mk\\ magnitudes, e.g. \\citet{Saw97}, Papovich, Dickinson, \\& Ferguson (2001; hereafter P01), \\citet{Shap01}, \\citet{Shap05}, \\citet{Reddy05}.  Although efficient in telescope time for observing objects with $R\\lesssim25$, optical surveys will miss objects that have high extinctions or those with a lack of current vigorous star formation. In fact, most galaxies with \\mstar$\\gtrsim10^{11}$\\msol\\ have optical magnitudes too faint for optical spectroscopic followup \\citep{Dokkum06}.  A comprehensive view of the high redshift Universe requires that UV selection be complemented with selection from light at least as red as the NIR, which corresponds to the rest-frame optical out to $z\\sim 3$.  Deep NIR observations allow us to detect highly extincted (e.g. SCUBA galaxies; Smail et al. 1997) or passive galaxies at $z>1.5$ (e.g. Daddi et al. 2005), which have very little rest-frame UV emission.  However, getting deep NIR data is observationally expensive and even over small fields requires substantial investments of 8-10m telescope time to reach depths fainter than \\Ks$=20-21$.  Until the recent development of megapixel NIR arrays progress was measurable but slow, e.g. \\citet{Hogg97} and \\citet{Ber98}.  A large step forward was taken with the Faint Infrared Extragalactic Survey (FIRES; Franx et al. 2000; Rudnick et al. 2001; Labb\\'e et al. 2003; F\\\"orster Schreiber et al. 2005) and with the NICMOS observations of the Hubble Deep Field North (HDF-N; P01; Dickinson et al. 2003). FIRES is an ESO large program which imaged in \\Js, \\mh, and \\Ks\\ the HDF-S and a field centered around the x-ray luminous cluster MS1054-03 at $z=0.83$.  At the time these fields had a unique combination of area and deep HST optical imaging and the ESO data extended this coverage to the \\Ks-band.  The FIRES data were particularly valuable because the associated deep \\Ks-band data provide access to the rest-frame \\mv-band all the way out to $z\\sim 3$.  The very deep FIRES data allowed us to probe rest-frame $V$-band luminosities \\l{V}\\ down to $\\sim 50\\%$ of the local \\lstarlam{V}\\ value.  The powerful combination of very deep NIR and optical data now make it possible to assemble representative samples of the Universe at $z>1.5$.  For example, \\citet{Franx03} and \\citet{Dokkum03} have shown that galaxies selected to have \\jk$>2.3$ are luminous in the rest-frame optical, likely have high \\mstar, but would be almost completely missed by rest-frame UV selected surveys.  Subsequent studies of these Distant Red Galaxies (DRGs) have shown that they have systematically older ages, higher extinctions, higher \\mstar, and comparable or even higher star formation rates than the optically selected LBGs at comparable redshifts and \\mk\\ magnitudes \\citep{Dokkum04, Forster04, Labbe05, Pap05, Dokkum06} and may contribute significantly to the stellar mass density at $z>1.5$ (Rudnick et al. 2003 -- hereafter R03).  Other author have also found galaxy populations that would be missed by rest-frame UV selection, highlighting the need for a comprehensive selection of galaxies (e.g. Daddi et al. 2004).  % Using the new sets of multi-wavelength data various authors have % achieved new insights into the nature of the high redshift galaxy % population.  For example, studies have measured the evolution of the % stellar mass density (Dickinson et al. 2003 -- hereafter D03; Fontana % et al. 2003 -- hereafter F03; R03), the rest-frame optical % morphologies of galaxies \\citep{Trujillo04, Pap05}, and their % spectral properties \\citep{Mello04, Shap04}.  While individual properties of galaxies are indeed important, for some quantities the modeling of the volume averaged population may yield more robust results.  For example R03 demonstrated that bursty SFHs cause a smaller systematic error in the mean stellar mass-to-light ratio \\mlstar\\ of the galaxy population - and hence estimates of \\rhostar\\ - than when it is determined from averages of \\mlstar\\ estimates of individual galaxies.  The utility of using cosmically averaged quantities has also been demonstrated by other authors across a range of topics.  Using the luminosity density determinations at many wavelengths \\citet{Mad98} tried to constrain the shape of the cosmic SFH by fitting the different bands simultaneously.  \\citet{Pei99} derived independent constraints on the cosmic SFH by modeling the cosmic gas content, metallicity, and extragalactic background light.  Using the FIRES data on the HDF-S R03 modeled the rest-frame optical mean colors and luminosity densities to derive the evolution in the stellar mass density at $z<3$.  At low redshifts \\citet{Baldry02} and \\citet{Gla03} modeled the mean galaxy spectrum from 2dF and the Sloan Digital Sky Survey (SDSS; York et al. 2000) to constrain the cosmic SFH.  Using the deepest existing \\Ks-band data R03 showed for the HDF-S that the evolution in the volume averaged rest-frame optical colors, e.g. \\ubr, \\bvr, and \\uvr, at $z<3$ was monotonic, with bluer colors at higher redshifts.  The colors also lay close to the locus of colors occupied by individual normal galaxies in the local Universe \\citep{Lar78, Jansen00a}.  They also found that the evolution in the rest-frame colors could be approximated by smooth SFHs with a constant metallicity and dust.  Models with these SFHs were then used to convert the \\uvr\\ to \\mlstarlam{V}\\ and hence to a mass density.  A primary goal of this paper is to extend the analysis of R03.  One major area of improvement is that this paper will model the full rest-frame UV through optical volume averaged SED of luminous galaxies in the Universe, allowing more freedom in the choice of dust attenuation, SFH, and age.  The full rest-frame UV through optical information also gives us more insight as to the nature of the mean stellar population of luminous galaxies at high redshift.  Perhaps the most important improvement with respect to R03 is that we include data from four disjoint fields, covering a total area of $\\approx 98.8$ arcminutes.  This is crucial since the field-to-field variations in the number densities of galaxies are expected to be large in\u00a0small fields.  In addition, multiple fields are important to characterize the mean contribution of different galaxy populations, e.g. DRGs, to the volume averaged SED and \\rhostar.  For example the field-to-field variations in the number densities of DRGs is known to be high with large differences between the HDF-N and HDF-S (Labb\\'e et al. 2003; hereafter L03).  In this paper we combine FIRES data from the HDF-S and MS1054-03 fields with those from the HDF-N catalog of P01 and D03 and a catalog from the Great Observatories Origins Deep Survey (GOODS) images of the Chandra Deep Field South (CDF-S; Wuyts et al. in preparation). Using photometric redshifts from these four catalogs we derive the rest-frame luminosities, luminosity densities and colors.  We examine the trends in color and redshift and compute the full rest-frame 2200\\ang\\ through \\mv-band volume averaged SED.  We then apply the same analysis to a galaxy sample from SDSS and from the COMBO-17 survey and use it to derive the evolution of mean SED properties in a consistent manner from $z\\sim 3$ down to $z=0$.  Finally, we fit the full rest-frame mean SED with a set of simple models and use these to derive the mean \\mlstarlam{V}\\ and \\rhostar\\ in each redshift bin.  In \\S~\\ref{data_sec} we describe the data.  In \\S~\\ref{methods_sec} we review the methods that we use to calculate photometric redshifts, rest-frame luminosities, luminosity densities, and global colors.  In \\S~\\ref{results_sec} we present the results including the evolution of the rest-frame luminosity density, color, and volume averaged SED. In this section we also present the model fits to the mean SED and the derived \\mlstar\\ and \\rhostar\\ evolution.  We discuss these results in \\S~\\ref{discuss_sec} and present our conclusions and summary in \\S~\\ref{sum_sec}.  Throughout this paper we assume $\\Omega_\\mathrm{M}=0.3,~\\Omega_{\\Lambda}=0.7,~\\mathrm{and~H_o}=70~{\\rm h_{70}~km~s^{-1} Mpc^{-1}}$ unless explicitly stated otherwise. ",
        "conclusions": "\\label{sum_sec}  In this paper we have presented the evolution in the properties of the volume averaged stellar population from $z=0$ to $z=3$.  We have used deep NIR and optical imaging over $\\sim 100$ square arcminutes in four disjoint fields to mitigate field-to-field variations.  The deep NIR data allow us to select and analyze galaxies at rest-frame optical wavelengths, where dust extinction and stellar mass-to-light ratio variations are far less than in the rest-frame UV.  We derived photometric redshifts for all objects lacking a spectroscopic redshift and measure rest-frame UV through optical luminosities of all galaxies at $z<3.2$.  We create a rest-frame \\mv-band luminosity selected sample with \\l{V}$>3\\times 10^{10}$\\lumsol\\ above which we are complete at all redshifts.  From these individual measurements we construct volume averaged properties such as the luminosity density, \\jrest, and color, e.g. \\ubr\\ and \\bvr.  The redshift evolution in luminosity density and mean color is very different.  As found in R03 for a much smaller sample, \\jrestlam{V} is roughly constant within a factor of $\\approx 2.2$ from $z\\sim3$ to $z=0$.  This modest evolution is broadly consistent with the results of D03 and \\citet{Giallongo05}.  The evolution in the rest-frame colors, however, is dramatic and the colors become systematically redder with decreasing redshift.  The colors of our SDSS comparison sample at $z=0.1$ are \\bvr$=0.72$ and \\ubr$=0.07$.  In contrast the average color of luminous galaxies at $z\\sim3$ is \\bvr$=0.40$ and \\ubr$=-0.38$.  By averaging over the population as a whole we are also averaging over the individual SFHs of the galaxies and the resultant volume averaged SFH will be smoother than that of the individual galaxies. We expect the mean colors to be more easily modeled than those of individual galaxies, which presumably have bursty SFHs.  We find that the mean colors at all redshifts lie close to the locus of morphologically normal local galaxies in \\bvr\\ vs. \\ubr\\ space but compared to the local sample are slightly bluer in \\ubr\\ at a fixed \\bvr.  Some of this difference could result from photometric redshift errors but could also be indicative of a residual signature of bursts in the mean colors.  We derive the luminosity weighted volume averaged UV-optical rest-frame SED for all galaxies with \\l{V}$>3\\times 10^{10}$\\lumsol. At all wavelengths the mean SED becomes gradually more red with decreasing redshift and at all redshifts is within the range of SEDs spanned by normal galaxies in the local Universe.  Even at our highest redshift bin the volume average SED does not look like that of a starburst from \\citet{Kinn96} with little extinction but rather, is redder than an irregular template.  We also compute the mean SED from DRGs chosen to have \\jk$>2.3$.  This mean DRG SED redder than average in the rest-frame optical and in the rest-frame UV.  With simple models we transform the mean SED into an evolution in the global \\mlstarlam{V}.  Because the volume averaged SFH should be smoother than that of the individual galaxies it is therefore more appropriate to use smooth models to fit the mean SED than to fit the SEDs of individual galaxies.  Nonetheless, the colors cannot be fit well at all redshifts with a single model having an exponential SFH and a constant dust obscuration and metallicity.  We can, however, fit the average SED at each redshift reasonably well with some combination of the parameters, as long as they can change as a function of redshift.  We use these model fits to derive luminosity weighted stellar mass-to-light ratios for all galaxies with \\l{V}$>3\\times 10^{10}$\\lumsol.  \\mlstarlam{V} also declines smoothly from $z=0.1$ to $z\\sim3$ with \\mlstarlam{V}$=4.9$ and $0.5$ at $z=0.1$ and $z\\sim3$ respectively.  There are strong variations in \\mlstarlam{V} across the population.  The DRGs have an \\mlstarlam{V} a factor of $\\approx 1.1$ and 3.6 higher than the non-DRGs in the $z=2.01$ and $z=2.8$ redshift bins.  This is consistent with the results of \\citet{Labbe05} and \\citet{Forster04} who modeled the SEDs of individual DRGs.  Multiplying the \\jrestlam{V} and mean \\mlstarlam{V} measurements we derive the evolution in \\rhostar\\ out to $z\\sim3$.  \\rhostar\\ in luminous galaxies declines by a factor of $4.8$ out to $z=3$.  At $z\\sim 2$ and 2.8, the DRGs contribute 30\\% and 64\\% of \\rhostar, showing that UV selection techniques miss much of the stellar mass at $1.6<z<3$ in rest-frame optically luminous galaxies.  We also compare our \\rhostar\\ measurements to a set of recent model predictions subjected to our same rest-frame luminosity cut and find a large disagreement between the models themselves and between the models and observations.  The different models can match the observed evolution in \\rhostar\\ over a limited redshift range, but both models predict an evolution that is too shallow to match the observations. This therefore implies that the masses of luminous galaxies in the models, or their abundance, are not consistent with the observations.  To better understand how different types of galaxies contribute to the stellar mass budget at high redshifts it will be necessary to reduce our reliance on photometric redshifts by performing NIR spectroscopy on a representative sample of DRGs.  This will confirm the redshifts of the optically faint population that contributes over 50\\% of the stellar mass at high redshift.  Re-performing this experiment with IRAC data will reduce the uncertainty due to dust obscuration and obtaining very deep NIR/MIR imaging will allow us to constrain the low mass end of the galaxy stellar mass function."
    },
    "0606/astro-ph0606243_arXiv.txt": {
        "abstract": "{From a 2$^\\circ\\times$ 2$^\\circ$ survey of NGC 6822  we have previously established that this Local Group dwarf irregular galaxy  possesses a huge spheroid having more than one degree in length. This spheroid is in rotation but its rotation curve is known only within $\\sim 15'$ from the center. It is therefore critical to identify bright stars belonging to the spheroid to characterize, as far as possible, its outer kinematics.} {We use the new wide field near infrared imager CPAPIR, operated by the SMARTS consortium,  to acquire $J$, $K_s$ images of two 34.8$'\\times 34.8'$ areas in the outer spheroid to search for C stars. } {The colour diagram of the fields allows the identification of  192 C stars candidates but a study of the FWHM of the images permits the rejection of numerous non-stellar objects with colours similar to C stars. } { We are left with 75 new C stars, their mean K$_s$ magnitude and mean colour are similar to the bulk of known NGC 6822 C stars. } {This outer spheroid survey confirms that the intermediate-age AGB stars are a major contributor to the stellar populations of the spheroid. The discovery of some 50 C stars well beyond the limit of the previously known rotation curve calls for a promising spectroscopic follow-up to a major axis distance of 40$'$. } ",
        "introduction": "NGC 6822, a Local Group dwarf irregular galaxy, has been the first extragalactic object identified by Hubble (1925) who determined its distance to be 214,000 pc or (m--M) = 21.65.  NGC 6822 is located in the constellation Sagittarius at a low Galactic latitude ($\\ell$ = 25$^\\circ$, b = --18$^\\circ$). It is thus seen behind a relatively heavy stellar foreground and its reddening is far from negligible. Today, we adopt for NGC 6822 a distance of (m -- M)$_0$ = 23.35, based on the RR Lyrae observations of Clementini et al. (2003), the Cepheids observations of Pietrzy\\'nski et al. (2004) and the near infrared observation of the tip of the red giant branch (TRGB) by Cioni \\& Habing (2005). At 470 kpc, it is the nearest Magellanic-type galaxy, after the Magellanic Clouds.  Contrasting strongly with its discovery description by Barnard (1885) who depicted it as: ``a 2$'$ nebula, diffuse and difficult to see with his 5-inch instrument'', NGC 6822 has been found to be much more extended. Its global structure was first studied by Hodge (1977) who described the bar  as a 8$'$ long structure with a position angle of 10$^\\circ$. In a more detailed investigation,  by Hodge et al. (1991), the galaxy could be traced to 10$'$. They adopt the point of view that the galaxy is circular rather than following the shape of the bar. We realize today that the low-density of the periphery of the galaxy can easily be masked by the substantial foreground density seen along the line of sight.  The first use of a wide-field imager, to search for C stars, by Letarte et al. (2002) has revealed that NGC 6822 is surrounded by a huge stellar structure extending to at least 20$'$  from its center and containing a substantial intermediate-age population. This huge spheroid was also seen by Lee \\& Hwang (2005).  Young stars were identified far from the center, along the HI disk by Komiyama et al. (2003). Recently, Battinelli et al. (2006) have mapped the elliptical spheroid of NGC 6822 to a semi-major axis distance of 36$'$ . From radial velocities of numerous C stars, Demers et al. (2006) have shown that the spheroid is dynamically decoupled from the HI disk and rotates at right angles to it. Since C stars used as kinematical probes were then known only to a distance of $\\sim 15'$ , it is of great interest to identify other C stars further out in the spheroid to extend its rotation curve  to compare it to the HI rotation curve, observed to $\\sim30'$ (Weldrake et al. 2003).  Few near-infrared investigators have targeted NGC 6822. Elias \\& Frogel (1985) have obtained $J$,$K$ magnitudes of 18 mostly spectroscopically confirmed red supergiants, with a $\\langle K\\rangle$ = 13. These stars are much brighter than C stars. Davidge (2003) has observed three tiny areas near the NGC 6822 bar to determine the metallicity from the slope of the RGB. He concluded that the bulk of the NGC 6822 red giants are of intermediate-age rather than $\\sim 12$ Gyr. The wide-field survey of a 20$'$ $\\times$ 20$'$ central region of NGC 6822 by Cioni \\& Habing (2005) is a first attempt to present a global view of the AGB of NGC 6822 in the near infrared. More recently, Kang et al. (2006) have obtained $J,H,K$ photometry of small areas near the center, where the space density of C stars is quite high.  We then present results of a near-infrared survey, which targets the outer spheroid and takes advantage of the new wide-field imager CPAPIR. It is essentially complementary to Cioni \\& Habing (2005) observations. We survey two regions in the outer spheroid of NGC 6822 with a small overlap with their central field. ",
        "conclusions": "The main goal of the present paper was to find bright kinematic probes, namely C stars, to trace the rotation curve of the NGC 6822 spheroid up to its detected limit ($\\approx 36'$). We have obtained accurate $J$,$K_s$ photometry for stars in two 34$'\\times 34'$ fields located in the periphery of the  spheroid along the major axis. On the basis of their $(J- K_s)$ colours we find 142 {\\it bona fide} C stars, half of them are newly discovered objects. Maps of these C stars are shown in Figure 10, for the NE field, and in Figure 11 for the SW field. Solid dots represent known C stars while the new ones are indicated by open circles for those with SHARP $<$ 0.20, (75\\% of them) and open triangles for those with 0.20 $<$ SHARP $<$ 0.30. As expected a few of the high SHARP objects match known C stars. However, one notes a unmistakable surplus of triangles in the areas corresponding to the outer spheroid. It is expected that a number of high SHARP candidates are non-stellar. Nevertheless, future observers can prioritize their candidates from the SHARP given. The inner square traces the border of the CPAPIR field. These new C stars will permit to more than double the length of the rotation curve of the spheroid and reach even further than the HI survey.   \\begin{figure*} \\centering \\includegraphics[width=6cm]{NEnewplot.ps}  \\caption{Map of the NE CPAPIR field, C stars found in the CFH12k survey are shown by solid dots while the C stars identified here are represented by circles, for those with SHARP $<$ 0.20 and open triangles for stars with 0.20 $<$ SHARP $<$ 0.30 } \\label{NE field} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=6cm]{SWnewplot.ps}  \\caption{Map of the SW CPAPIR field, C stars found in the CFH12k survey are shown by solid dots while the C stars identified here are represented by circles, for those with SHARP $<$ 0.20 and open triangles for stars with 0.20 $<$ SHARP $<$ 0.30 } \\label{SW field} \\end{figure*}"
    },
    "0606/astro-ph0606596_arXiv.txt": {
        "abstract": "We report observations with the {\\sl Chandra X-ray Observatory} of a field in the $\\gamma$-Cygni supernova remnant (SNR78.2+2.1) centered on the cataloged location of the unidentified, bright \\gray\\ source \\EG. In this search for an X-ray counterpart to the \\gray\\ source, we detected 30 X-ray sources. Of these, we found 17 strong-candidate counterparts in optical (visible through near-infrared) cataloged and an additional 3 through our optical observations. Based upon colors and (for several objects) optical spectra, nearly all the optically identified objects appear to be reddened main-sequence stars: None of the X-ray sources with an optical counterpart is a plausible X-ray counterpart to \\EG---if that \\gray\\ source is a spin-powered pulsar. Many of the 10 X-ray sources lacking optical counterparts are likely (extragalactic) active galactic nuclei, based upon the sky density of such sources. Although one of the 10 optically unidentified X-ray sources could be the \\gray\\ source, there is no auxiliary evidence supporting such an identification. ",
        "introduction": "}  \\EG\\ is the brightest of the many unidentified sources detected by the Energetic Gamma-Ray Experiment Telescope (EGRET). The third and final official EGRET catalog 3EG \\citep{har99} lists 271 sources, of which only a quarter have firm identifications. The paucity of identifications results primarily from large uncertainties in \\gray-source positions, due to the angular resolution of EGRET's multilevel-spark-chamber detector and to poor photon statistics. Typical positional uncertainties are a few degrees \\citep{mat96}, improving to $\\approx\\!0.5\\arcdeg$ (FWHM) for bright sources with adequate statistics above several GeV, for which determination of the photon's direction is more precise. Systematic uncertainties---due to non-uniform background or the presence of adjacent sources---also can affect positional determinations \\citep{har99}. This is most severe for sources near the Galactic plane, where the diffuse background is bright and highly structured and the concentration of \\gray\\ sources is high. Consequently, most EGRET sources identified through positional correlations are  extragalactic---e.g., blazars---at high Galactic latitude. Even so, about half the high-confidence blazar identifications from the 3EG catalog lie outside the stated 68\\%-confidence contours \\citep{har99}.  Clearly, identifying counterparts to \\gray\\ sources at low Galactic latitude requires more than positional correlations. For example, the identification of seven EGRET sources as spin-powered pulsars relied upon detection of periodic modulation of the \\gray\\ flux. Based upon pulsar population models \\citep{mcl00,zha00}, pulsar emission models \\citep{rom95,gon04}, the association of EGRET sources with young objects, and the fact that the only definite identifications of persistent EGRET Milky-Way sources are bright pulsars, we expect many more EGRET sources to be pulsars. However, further searches for pulsations in EGRET data---either blind \\citep{cha01} or folding on known pulsar periods \\citep{nel96}---have been unsuccessful.  Analyses of the sample distribution of \\gray\\ spectrum, flux, variability, and location has produced some progress in distinguishing classes of EGRET sources. A population of hard, steady, Galactic-plane sources seem similar to confirmed \\gray\\ pulsars; another of soft, variable, high-Galactic-latitude sources are likely unidentified blazars \\citep{sow04}. In addition, there are a group of weaker, steady sources, spatially correlated with the Gould belt \\citep{geh00}; a group of luminous, soft, highly variable (and perhaps older) sources that form a halo about the Galactic center \\citep{gre01,gre04}; and a subgroup of variable Galactic-plane objects \\citep{nol03}. While the nature of most of these \\gray\\ sources is unknown, spatial correlations with tracers of recent star formation---including supernova remnants \\citep[SNRs;][]{stu95}, OB associations \\citep{kaa96}, and massive stars \\citep{rom99}---support an association with young objects.  Proposed \\gray-emission mechanisms are consistent with young source populations. Shocks in massive-star supersonic winds \\citep{cas80}, high-mass X-ray binaries \\citep{tav97,kir99}, microquasar jets \\citep{geo02}, pulsar magnetospheres, and pulsar-wind nebulae \\citep{rob05} are all environments conducive to producing high-energy \\grays---either directly through electron bremsstrahlung, Compton scattering, or synchrotron emission; or indirectly through hadronic collisions with surrounding material followed by $\\pi^o$-decay emission. Such processes also work on larger spatial scales: Cosmic rays produced in supernova blast waves can interact with the interstellar medium to produce \\grays\\ \\citep{gai98}. These same mechanisms may produce the diffuse \\gray\\ background \\citep{ber93,hun97}.  \\EG\\ is in the Galactic plane \\citep[$l,b=78.05,2.08$;][]{har99}, as are all the known \\gray\\ pulsars. It exhibits a steady \\citep{nol03}, hard \\citep[photon index $\\Gamma=2.08\\pm0.04$;][]{esp96}, power-law \\gray\\ spectrum breaking at about 4\\,GeV \\citep{mer96,buc98}, typical of spin-powered pulsars \\citep{tho99}. Hence, \\EG\\ is one of the best candidates for a \\gray\\ pulsar \\citep{mer96}. However, efforts to detect pulsations in EGRET data \\citep{bra96,koh95,cha01} or to discover radio pulsars in the region \\citep{bec04,nic97} have been unsuccessful. Pulsations are difficult to detect in EGRET data, due to low signal-to-noise ratios and to the propensity of \\gray\\ pulsars to spin-down rapidly and to exhibit glitches \\citep[e.g.,][]{cha01}. In fact, the successful detection of \\gray\\ pulsations in most EGRET pulsars resulted from epoch-folding the \\gray\\ data with the known period of the pulsar \\citep{nol93,kan94,ram95}.  The absence of a radio source---to a sensitivity of about 40\\,$\\mu$Jy \\citep{bec04}---within the \\EG\\ error circle, is not unprecedented amongst \\gray\\ pulsars: The nearby Geminga pulsar is radio quiet. The discovery of pulsations in a coincident \\ROSAT\\ X-ray source \\citep{hal92}, followed by verification in the EGRET data \\citep{ber92}, are the basis for Geminga's classification as a pulsar. With Geminga as an archetype, several authors \\citep[e.g.,][]{hel94,der94,tho94,rom95} have suggested that radio-quiet \\gray\\ pulsars may account for most of the unidentified EGRET sources. Likewise, the recently discovered \\citep{obr05} transient radio pulsars---Repeating Radio Transients (RRATs) and intermittent (short-burst-mode) radio pulsars---may also contribute to this population.  If \\EG\\ is a radio-quiet pulsar similar to Geminga, then we can estimate its X-ray emission, under certain assumptions. A 0.5-MK blackbody component dominates Geminga's soft-X-ray flux \\citep{kar05,hal97}, producing a ratio of energy flux $S_{x}/S_{\\gamma}=1.8\\!\\times\\!10^{-3}$ for the 0.1--2.4-keV and 0.1-5-GeV bands \\citep{tho96}. From \\ROSAT\\ observations, we \\citep{bec04} previously set an upper limit $S_{x}<1.8\\!\\times\\!10^{-13}$~erg/(cm$^{2}$ s) to any soft, point-like X-ray counterpart to \\EG: This corresponds to $S_{x}/S_{\\gamma}\\leq 2.3\\!\\times\\!10^{-4}$, well below expectations based on Geminga. The analog of the harder, nonthermal, X-ray emission from Geminga is better matched to the effective bandpass of the \\Chandra\\ ACIS I array. \\citet{jac02} measure a 0.7--5.0-keV flux $2.6\\!\\times\\!10^{-13}$~erg/(cm$^{2}$ s) from Geminga. Again, if we assume a similar $S_{x}/S_{\\gamma}$ ratio, we would expect a flux $<10^{-13}$~erg/(cm$^{2}$ s) from \\EG. Our \\Chandra\\ observation easily detects X-ray sources at this level; however, there are dozens of faint X-ray sources within the error contour for \\EG.  The direction of \\EG\\ lies within the 1\\arcdeg-diameter \\citep{hig77}, young supernova remnant SNR78.2+2.1 \\citep{gre74}, often called the ``$\\gamma$-Cygni SNR'' after the unrelated bright near-by star. The distance to the $\\gamma$-Cyg SNR is approximately 1.6~kpc \\citep{hig77,lan80}, corresponding to a (true) distance modulus $\\Delta m \\approx 11.0$.  Having noticed the proximity of the \\gray-source and SNR directions, \\citet{stu95} suggested that the \\gray\\ emission from \\EG\\ could result from protons accelerated in the interaction of the supernova blastwave with surrounding gas. They emphasized that the SNR's young age, radio brightness, and flat radio spectrum all favored this cosmic-ray acceleration mechanism. Following \\citet{dru94}, they also remarked that the higher density medium of the molecular cloud associated with the SNR should enhance the \\gray\\ flux.  Using similar arguments, \\citet{esp96} concluded the $\\gamma$-Cyg SNR is one of the best candidate cosmic-ray origin sites. However, they found that the \\gray\\ spectrum of \\EG\\ is much flatter than expected from $\\pi^{\\rm o}$ decay. \\citet{gai98} showed that larger contributions from electron bremsstrahlung and Compton scattering would make the model more consistent with observations. However, the model is still inconsistent with observed upper limits to the TeV \\gray\\ emission from \\EG\\ \\citep{buc98}, thus questioning its applicability to this source. Furthermore, the EGRET data are more consistent with a point source than with the extended SNR emission.  X-ray searches for cosmic-ray acceleration sites in the $\\gamma$-Cyg SNR have been unsuccessful. While \\ROSAT\\ \\citep{bra96,bec04}, \\ASCA\\ \\citep{uch02}, and {\\sl INTEGRAL} \\citep{byk04} observations all detect X-ray emission within the remnant, none originates near the cataloged position of \\EG. The soft X-ray emission could be obscured by intervening material. However, contrary to expectations for the cosmic-ray acceleration model \\citep[e.g.,][]{byk00}, there is no radio emission from the location of \\EG.  \\citet{bra96} analyzed six \\ROSAT\\ PSPC observations of the $\\gamma$-Cyg SNR region, pointed within 40\\arcmin\\ of the reported position of \\CG. They found a point source---\\RX---within the 95\\%-confidence contour of the 2EG position, which they suggested as the X-ray counterpart to the \\gray\\ source. Further, \\citet{bra96} and \\citet{car00} found a possible optical counterpart for \\RX. Follow-up optical observations identified a $14.5^{\\rm mag}$ K0-V star within the $\\approx\\!6\\arcsec$-radius \\ROSAT\\ error circle, which would not account for the \\gray\\ emission. In that the X-ray--to--optical flux ratio is only marginally consistent with that of late-type stars \\citep{sto91,fle95,bra96} one could not totally exclude the association of \\RX\\ with the \\gray\\ source.  The 3EG catalog \\citep{har99} gave an improved  position for \\EG---$20^{\\rm h} 21^{\\rm m} 1\\fs0 +40\\arcdeg 17\\arcmin 48\\arcsec$---displaced by about 10\\arcmin\\ from the 2EG position used by Brazier \\et\\ (1996). With this improved position, the proposed counterpart \\RX\\ no longer lies within the 95\\%-confidence contour of \\EG, although it is within the 99\\%-confidence contour. Previously \\citep{bec04}, we reported follow-up studies of \\RX\\ with \\Chandra\\ and the Green Bank Radio Telescope. The \\Chandra\\ observations determined the position and spectrum of \\RX\\ with high precision, showing that \\RX\\ is indeed the X-ray counterpart to the K0-V star, with a soft spectrum indicative of stellar coronal emission. The Green Bank observations set upper limits to the presence of radio pulsations.  Here we report results of a \\Chandra\\ observation centered on the 3EG position for \\EG. The \\Chandra\\ sensitivity is so deep that many X-ray sources lie within the formal error contours for the position of \\EG. Consequently, we seek a distinguishing feature amongst the detected X-ray sources, in order to establish one (or more) of them as a probable counterpart to \\EG. Section~\\ref{s:xray_obs} describes the results of our search for point-like X-ray sources. Section~\\ref{s:oi} presents the results of our searches for candidate optical counterparts to the detected X-ray sources through cross correlating with visible and near-infrared catalogs, and through new optical imaging and spectroscopic observations. Finally, Section~\\ref{s:summary} discusses and summarizes the results. ",
        "conclusions": "}  Using the \\Chandra\\ X-ray Observatory, we continued our search \\citep{bec04} for possible X-ray counterparts to the intriguing \\gray\\ source \\EG\\ (\\CG). We found 30 X-ray sources in a field centered on the 3EG-cataloged position, located within the $\\gamma$-Cyg SNR. One of these---S10---lies $2\\farcs3\\!\\pm\\!1\\farcs0$ from a previously detected source---S211 of \\citet{bec04}---and could be the same source. Most \\Chandra\\ sources in the field are naturally near the completeness limit  ($\\approx\\!10$ counts), slightly above the detection limit ($\\approx\\!7$ counts). The strongest source---S10---is about 10 times the detection limit. For the 14.3-ks observing time, detected sources have (\\S\\ref{s:analysis_flux} and Table~\\ref{t:data_x}) photon spectral fluxes $K_{E} \\ga 1.6\\!\\times\\!10^{-6}$~photons/(cm$^2$ s keV) at 1~keV, X-ray energy fluxes $S \\ga 9\\!\\times\\!10^{-15}$~erg/(cm$^2$ s) in the 0.5--8-keV band, and energy spectral fluxes (flux densities) $S_{\\nu} \\ga 1.1$~nJy at 242~PHz (1~keV). For sources at distances comparable to that of the $\\gamma$-Cyg SNR (about 1.6~kpc), the X-ray luminosity would be $L_{x} > 3\\!\\times\\!10^{30}$~erg/s in the 0.5--8-keV band.  \\citet{zav04} show that the effective optical--X-ray energy index $\\alpha_{ox} \\approx 0.35\\pm 0.1$---corresponding to $\\log(L_{x}/L_{o}) \\approx +2$---for spin-powered pulsars detected in both these bands. Hence, if one of the X-ray sources in the field is a (similar) spin-powered pulsar, its flux density in an optical (or near-IR) band would  be no brighter than about $0.1\\, \\mu$Jy ($\\approx\\!25^{\\rm mag}$)---i.e., undetected in either the 2MASS or the USNO-B catalog or by our visible or near-infrared observations. Consequently, those 20 X-ray sources with strong-candidate counterparts (Tables~\\ref{t:data_oi} and \\ref{t:wirc_data}) are {\\em not} good candidates for spin-powered pulsars. Our optical spectroscopy (\\S\\ref{s:oi_obs} and Table~\\ref{t:obs_data}) shows that most strong-candidate counterparts are reddened main-sequence stars, with $\\alpha_{ix} \\approx 2$---corresponding to $\\log(L_{x}/L_{i}) \\approx -3$. Furthermore, our X-ray observations of variability (\\S\\ref{s:analysis_var}) in S10 and S25 and of apparently highly-absorbed spectra (\\S\\ref{s:analysis_spec}) in S03, S12, and S22 would suggest that none of these sources is a spin-powered pulsar.  The available data do not disqualify the remaining 10 X-ray sources--- S07, S09, S15, S18, S19, S21, S23, S27, S29, and S30---as possible candidates for a spin-powered pulsar. However, there is nothing distinguishing any of these X-ray sources from the others: They could represent several classes of objects. In fact, based upon deep surveys, we would expect approximately 10 extragalactic X-ray sources in this field. As we (Becker \\et\\ 2004) reported previously, there is {\\em no} radio evidence for a pulsar in the $\\gamma$-Cyg SNR field, down to a limiting sensitivity of $L_{\\nu} = 0.09$ mJy kpc$^{2}$ at $\\nu_{r}=$ 820~MHz (for an assumed 4\\% pulse duty cycle). Thus, if one of the X-ray sources is a spin-powered pulsar, it would be rather radio-quiet, as is Geminga. Finally, in that accurate determination of a \\gray\\ source position against a strong gradient in the diffuse background is difficult, we must also admit the possibility that the EG3 position for \\EG\\ is incorrect. Finally we note that without high-precision X-ray spectra of each of the candidate X-ray sources, and detailed follow up in other wavelength bands, there is essentially no satisfactory way in which to eliminate some of the candidates from consideration. This is especially true for neutron stars as the the corresponding infrared-visible fluxes are very weak making such observatiobs especially difficult. In such cases, the principal and important {\\sl Chandra} contribution is to provide target lists as we have done with accurate positions as a basis for future studies."
    },
    "0606/astro-ph0606075_arXiv.txt": {
        "abstract": "We report the discovery of XMMXCS~J2215.9-1738, a massive galaxy cluster at $z =1.45$, which was found in the {\\em XMM Cluster Survey}. The cluster candidate was initially identified as an extended X-ray source in archival XMM data.  Optical spectroscopy shows that 6 galaxies within a $\\sim$60 arcsec diameter region lie at $z = 1.45\\pm 0.01$.  Model fits to the X-ray spectra of the extended emission yield $kT = 7.4^{+2.7}_{-1.8}$ keV (90\\% confidence); if there is an undetected central X-ray point source then $kT = 6.5^{+2.6}_{-1.8}$ keV.  The bolometric X-ray luminosity is $L_x = 4.4^{+0.8}_{-0.6} \\times 10^{44}$ ergs s$^{-1}$ over a 2 Mpc radial region.  The measured $T_x$, which is the highest known for a cluster at $z > 1$, suggests that this cluster is relatively massive for such a high redshift. The redshift of XMMXCS~J2215.9-1738 is the highest currently known for a spectroscopically-confirmed cluster of galaxies. ",
        "introduction": "High redshift galaxy clusters provide important laboratories for the study of structure formation and galaxy evolution. They also can be used to constrain cosmological parameters independent of the Cosmic Microwave Background and supernova methods. The number of known clusters at $z>1$ is currently small ($\\simeq 10$ with X-ray confirmation: e.g. \\citealt{lynx,0910,1252,mullis}), but growing rapidly thanks to new surveys being carried out in the X-ray \\citep*[e.g.][]{xcs,mullis,xmmlss,champ}, optical \\citep*[e.g.][]{RCS}, and infrared \\citep*[e.g.][]{bootes}. In the near future, Sunyaev-Zel'dovich surveys should also provide large numbers of high redshift clusters \\citep{sz}.  We announce here the discovery of a high redshift cluster in the XMM Cluster Survey \\citep*[XCS:][]{xcs} and the measurement of its X-ray temperature. At $z=1.45$, XMMXCS~J2215.9-1738 is the most distant cluster with spectroscopic redshift confirmation to date. Unless otherwise noted, we assume $\\Omega_m = 0.27$, $\\Lambda = 0.73$, and $H_o = 70$~km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Analyses of our combined X-ray imaging spectroscopy, optical spectroscopy, and optical-IR imaging provide strong evidence that XMMXCS~J2215.9-1738 is a massive galaxy cluster at $z=1.45$. At least 10 galaxy clusters with $z>1.4$ and $T_{x}>6$ keV are expected to be present in our current XCS area of 168.2 deg$^{2}$, if the universe is assumed to be flat with $\\Omega_{m}=0.238$ and $\\sigma_{8}=0.744$ (WMAP - Spergel et al. 2006).  This estimate requires the relation between cluster mass and X-ray temperature to be normalized so as to reproduce the local abundance of X-ray clusters (HIFLUGCS - Reiprich \\& B\\\"{o}hringer 2002). However, it should be noted that these assumptions require that a local 5 keV cluster would need to have a mass, within an overdensity of 500 with respect to the critical density, of around $2\\times10^{14}\\,h^{-1}\\,M_{\\sun}$ \\citep{pier,viana}, about 2/3 smaller than is implied by the most recent observational data \\citep{app,av05}.  With the combination of the serendipitous searches for clusters in the {\\it XMM} archival data \\citep{xcs,mullis}, and the on-going wide-area surveys such as UKIDSS (Lawrence et al.\\ in preparation) and the RCS \\citep{RCS} the time finally is ripe for the identification of large samples of $z > 1$ clusters.  The construction of such samples will pave the way towards a better understanding of the origin of galaxy clusters and their constituent galaxy populations."
    },
    "0606/gr-qc0606032_arXiv.txt": {
        "abstract": "Non-perturbative quantum geometric effects in Loop Quantum Cosmology predict a $\\rho^2$ modification to the Friedmann equation at high energies. The quadratic term is negative definite and can lead to generic bounces when the matter energy density becomes equal to a critical value of the order of the Planck density. The non-singular bounce is achieved for arbitrary matter without violation of positive energy conditions. By performing a qualitative analysis we explore the nature of the bounce for inflationary and Cyclic model potentials. For the former we show that inflationary trajectories are attractors of the dynamics after the bounce implying that inflation can be harmoniously embedded in LQC. For the latter difficulties associated with singularities in cyclic models can be overcome. We show that non-singular cyclic models can be constructed with a small variation in the original Cyclic model potential by making it slightly positive in the regime where scalar field is negative. ",
        "introduction": "One of the cornerstone problems in cosmology is that of the initial conditions of our universe. In standard big bang cosmology the weight of this problem is shifted to quantum gravity which is expected to generate ideal conditions for the genesis of our universe. Generally, it is envisioned that our universe started in a hot expanding phase from the big bang and entered a stage of inflation before becoming radiation and matter dominated. However the idea that our expanding universe was preceded by a contracting phase of a classical universe has had many avatars. This idea, though promising to eliminate some of the fundamental problems related with initial conditions in standard cosmology and providing an alternative to inflation, has remained unsuccessful due to absence of  generic non-singular evolution from the contracting to the expanding branch. For a flat $k = 0$ cosmological model, which is of interest of the present work, based on classical general relativity (GR) such an evolution is forbidden by  powerful singularity theorems, unless one assumes an exotic form of matter which violates positive energy conditions. Though quantum corrections have been introduced to achieve a non-singular evolution between contracting and expanding phases in special conditions, it has been realized that a generic non-singular evolution is difficult to achieve unless one incorporates non-perturbative effects of quantum gravity. In absence of any non-perturbative quantum gravitational modifications to dynamics, innovative ideas like pre big bang string cosmology \\cite{pbb} and Ekpyrotic/Cyclic models \\cite{ekp-cyclic,StTur} have so far had limited viability \\cite{cyc-nonsing}.  Loop quantum gravity (LQG) is a leading non-perturbative background independent approach to quantize gravity \\cite{lqg_review}. The underlying geometry in LQG is discrete and the continuum spacetime is obtained from quantum geometry in a large eigenvalue limit. The application of LQG techniques to homogeneous spacetimes results in loop quantum cosmology (LQC) \\cite{mblr} which has led to important insights on the resolution of singularities in various situations \\cite{bigbang,abl,sing}. Using extensive analytical and numerical methods the analysis of the evolution of semi-classical states for the flat $k=0$ model with a massless scalar field has shown that a contracting semi-classical universe passes through the quantum regime bouncing to an expanding semi-classical universe \\cite{aps,aps1,aps2}. Thus within the context of the model considered the idea of the non-singular bounce is realized in a natural fashion.    The underlying dynamics in LQC is governed by a discrete quantum difference equation in quantum geometry. However, using semi-classical states we can construct an effective Hamiltonian description on a {\\em continuum} spacetime which has been shown to very well approximate the quantum dynamics \\cite{sv,aps1,aps2}. In particular, we can obtain effective equations for the modified Friedmann dynamics which can be used to investigate the role of non-perturbative quantum corrections. An important feature of the dynamics is that the classical Friedmann equation is modified to include a $\\rho^2$ term which is relevant in the high energy regime. The modified term in the Friedmann equation is negative definite implying a bounce (or more generally a turn-around) when the energy density reaches a critical value on the order of the Planck density in accordance with the results from the quantum evolution.  An important question which gains immediate attention is the viability of the bounce picture for more general matter sources and compatibility of inflationary and Cyclic model ideas with LQC. In this regard various questions arise: Does the bounce occur for various inflationary potentials? Do inflationary solutions form the attractors of the dynamics after the bounce?  Since the modifications which lead to a non-singular bounce are obtained from non-perturbative quantum gravity effects is it possible that these can lead to construction of non-singular cyclic models using Cyclic/Bicyclic potentials? What insights are obtained to alleviate the problems of Cyclic models?  This work seeks to  answer  these questions using the effective theory in LQC. The plan of this paper is as follows. In the next section we briefly review the effective theory of LQC and classify the differences between quantum bounce generated by loop quantum dynamics and the classical bounce originating from general relativity. In Sec. III we perform a qualitative analysis using phase space diagrams to study the dynamics for inflationary potentials ($\\phi^2$ and $\\phi^4$) and show that inflationary trajectories are attractors of the dynamics after the bounce. In Sec. IV we apply the qualitative analysis to negative potentials of the Cyclic/Bicyclic models. We show that the Bicyclic potential \\cite{bike} can successfully lead to non-singular cycles, whereas the original Cyclic potential can exhibit cycling of the scale factor but not the scalar field. A truly cyclic model can be easily constructed if the Cyclic potential becomes slightly positive in the regime where the scalar field is negative, as the Bicyclic potential does. The non-singular cyclic behavior obtained for negative potentials is unavoidable and occurs for generic values of parameters. We conclude the paper with a discussion and open issues in Sec V. ",
        "conclusions": "The non-perturbative quantum geometric effects lead to a $\\rho^2$ modification with a negative sign in the Friedmann equation in LQC. Since the correction term is negative definite it can lead to a quantum bounce in the high energy regime when loop quantum modifications are dominant. We have investigated the qualitative details of this bounce for inflationary and negative potentials. The existence of the outer boundary given by (\\ref{outerb}) in the phase portraits guarantees non-singular behavior of all the solutions with a scalar field in LQC with any kind of potential. For negative potentials the inner boundary also appears corresponding to the classical recollapse. The presence of the two boundaries for negative potentials leads to the possibility of solutions having cyclic behavior. We have also briefly reported on the nature of quantum turn-arounds for a phantom field.  The massless scalar field gives a good example of one feature of cosmological solutions within LQC, namely the non-singular bounce. In effect, at least solutions with $\\phi^2$ and $\\phi^4$ potentials can be well approximated near the bounce by the corresponding exact solutions for a  massless scalar field, because most of them approach the outer boundary of phase space with kinetic energy  much larger than the potential energy, thus the potential can be neglected. For these potentials we showed that inflationary trajectories are attractors of dynamics after the bounce. In the same way a massive scalar field with negative potential represents a good example of cycling between expansion and contraction accompanied with inflation, exhibiting the main features of the more complicated potentials. Our analysis shows that there is little qualitative difference between the massive scalar field with negative constant and Bicyclic potentials, although their functional form is very different. In both cases cyclic solutions can be found with an infinite number of inflationary stages.  Though the problem of the big crunch can be overcome with potential of the Cyclic model, a limitation remains in the lack of a  turn-around of $\\phi$ from the negative side of the potential. In the Cyclic model this instant is envisioned as the collision of two branes. If the effective potential between branes can become positive prior to collision (arising from higher order perturbative string effects/or non-perturbative corrections), then a turn-around of the field can occur and a viable cyclic model can be constructed. On the other hand, the Bicyclic potential provides the possibility for the field to turn around and may be used for further development of cyclic models.    The qualitative analysis reported in this work has shown that both inflationary paradigm and the Cyclic/Bi-cyclic scenarios may be incorporated in LQC. Our analysis has not touched on the aspects which can quantify the viable parameter space of these models with respect to observational data. Further, we have not explored the way LQC may shed insights on the problem of evolution of perturbations through the bounce. Work on including perturbations in the quantum theory and deriving effective equations is in progress, which will eventually address the viability of these ideas with respect to observations."
    },
    "0606/astro-ph0606305_arXiv.txt": {
        "abstract": "We present optical integral field spectroscopy of the H$\\alpha$-luminous ($>10^{42}$\\ergps) central cluster galaxies in the cores of the cooling flow clusters A1664, A1835, A2204 and Zw8193. From the [NII]+H$\\alpha$ complex in these moderate resolution (70--150\\kmps) spectra we derive two-dimensional views of the distribution and kinematics of the emission line gas, and further diagnostics from the [SII] and [OI] lines.  The H$\\alpha$ emission shows a variety of disturbed morphologies, ranging from smooth but distorted to clumpy and filamentary, with velocity gradients and splittings of several hundred \\kmps~on spatial scales of 20\\kpc~or more. Despite the small sample size, there are some generic features. The most disturbed H$\\alpha$ emission appears to be associated with secondary galaxies within 10--20\\kpc~(projected) of the central galaxy and close in velocity to the H$\\alpha$. The global H$\\alpha$ kinematics match those of the CO(1-0) emission in single-dish data. The [NII]/H$\\alpha$, [SII]/H$\\alpha$ and [OI]/H$\\alpha$ ratios vary little with position, local H$\\alpha$ surface brightness or between clusters.  We propose that the H$\\alpha$ and CO emission arise in molecular clouds heated by a starburst, and that the latter has been triggered by interaction with a secondary galaxy. Such CO emission is known to trace massive ($>10^{10}$\\Msun) compact ($<20$\\kpc) reservoirs of cool molecular gas, and it is plausible that an infalling galaxy would disturb this gas, distorting the H$\\alpha$ morphology and initiating widespread star formation. We also examine the role of cloud-cloud collisions in the undisturbed molecular gas reservoir, and suggest that they might be an important source of excitation for the emission line gas in the cores of lower H$\\alpha$ luminosity clusters with less intense star formation. ",
        "introduction": "In recent years our understanding of the X-ray cooling flow phenomenon in galaxy cluster cores has been revolutionised. Throughout most of the 1980s and 1990s, X-ray observations suggested that gas in the central 100\\kpc~is cooling out at rates of up to several hundred solar masses per year, but the lack of evidence for a reservoir of cooled gas led to heated debate (summarised by Fabian~1994) over this interpretation of the X-ray data. Results from {\\em XMM-Newton} and {\\em Chandra} have since led to a sharp downward revision in X-ray cooling rates (e.g. Schmidt, Allen \\& Fabian 2001) and also reveal a strong deficit of line emission from gas cooling below $T_{\\rm{virial}}/3$ (Peterson et al.~2003). The implication is that X-ray cooling is quenched, for which numerous mechanisms have been proposed, including: rapid mixing of hot and cold phases, inhomogeneously distributed metals in the intracluster medium (Fabian et al.~2001,2002); active galactic nucleus (AGN) heating by jets (Br\\\"{u}ggen \\& Kaiser 2003) and sound waves (Fabian et al. 2003); thermal conduction of heat from the hotter outer parts of the cluster into the cooler core (Voigt et al. 2002); a significant relativistic cosmic ray component frozen into the thermal gas (Cen 2005); the release of gravitational energy from blobs of gas which detach from the bulk flow and fall directly into the core (Fabian 2003).  Concurrently, significant progress has been made in identifying cool gas and dust in cluster cores. Edge~(2001) detected CO emission in the centres of 16 cooling flows, consistent with $10^{9}-10^{11.5}$\\Msun~of \\H2~at 20--40\\K~for a standard CO:\\H2~conversion (see also Salom\\'e \\& Combes 2003). These are roughly the masses expected, given the revised cooling rates and likely ages. Interferometry shows further that the CO emission is localised within the central few arcsec of the cluster (Edge \\& Frayer 2003; Salom\\'e \\& Combes 2004). The frequent occurrence of smaller masses ($\\sim 10^{5}-10^{6}$\\Msun) of hot \\H2~has also been established (e.g. Edge et al.~2002; Jaffe, Bremer \\& van der Werf~2001), and excitation analysis suggests that this hot \\H2~is a high pressure, transiently-heated component (Wilman et al.~2002). Both CO and \\H2~emissions correlate well with the strength of the H$\\alpha$ emission from ionized gas at $10^{4}$\\K, whose prevalence in these environments, often in the form of spectacular filaments, has long been known (e.g. Hu et al.~1983; Crawford et al.~1999). Despite the clear association between optical line emission and short central X-ray cooling times (Peres et al.~1998; Bauer et al.~2005), their physical relationship is ill-understood. Photoionisation by the observed excess population of hot massive stars can energetically account for the H$\\alpha$ luminosities in the most luminous systems (Allen~1995; Crawford et al.~1999). {\\em Spitzer} MIPS photometry of 11 CCGs by Egami et al.~(2006) also shows that the most H$\\alpha$-luminous in their sample (A1835, A2390 and Zw3146) have prominent far-infrared thermal dust emission plausibly powered by star formation, two of them with $L_{\\rm{IR}} > 10^{11}$\\Lsun. At lower H$\\alpha$ luminosities the picture is less clear: the tapping of energy from the intracluster medium (ICM) through turbulence (Crawford \\& Fabian~1992) and heat (Sparks et al. 2004) are just two mechanisms which have been invoked to explain the optical nebulosity in such systems.  In this paper we present integral field spectroscopy of the ionized gas in the cores of four such clusters, A1664, A1835, A2204 and Zw8193. The principal aim is to obtain a full two dimensional view of the distribution and kinematics of the gas through high resolution spectroscopy of the H$\\alpha$+[NII] emission line, with additional ionization information being gleaned from the [SII]$\\lambda \\lambda$6717,6731 and [OI]$\\lambda \\lambda$6300,6363 lines where possible. These four central cluster galaxies (CCGs) all have H$\\alpha$ luminosities exceeding $10^{42}$\\ergps, making them 4 of the top 6 most-H$\\alpha$ luminous systems in the extensive CCG spectroscopic survey by Crawford et al.~(1999). In this regime of H$\\alpha$ luminosity, photoionisation by a young stellar population can account energetically for the luminosity of the H$\\alpha$ nebulosity (Allen~1995; Crawford et al.~1999). In addition to an analysis of the CCGs, we also present spectroscopy of other sources within the IFU field of view, including other cluster galaxies and (in the case of A2204) a serendipitous gravitationally-lensed background galaxy. We first present results for the clusters individually and then summarise and interpret their generic features. Throughout the paper we assume a cosmology with $H_{\\rm{0}}=70$\\kmpspMpc, $\\Omega_{\\rm{M}}=0.3$ and $\\Omega_{\\rm{\\Lambda}}=0.7$ and all physical quantities quoted from other papers have been converted accordingly. ",
        "conclusions": "The results presented in this paper have shed new light on the kinematic and morphological complexity of the emission line gas in the cores of these H$\\alpha$-luminous clusters. Our principal findings are that: $\\bullet$ In many cases the H$\\alpha$ emission appears to have been disturbed by interaction with a small companion galaxy to the central galaxy. $\\bullet$ The CO and H$\\alpha$ emission share the same kinematics and by implication trace each other closely. $\\bullet$ The emission line ratios of the gas are uniform, showing no strong variations in [NII]/H$\\alpha$, [SII]/H$\\alpha$ or [OI]/H$\\alpha$, as a function of position or H$\\alpha$ surface brightness, suggesting that its excitation state is independent of the processes which disturb it on the observed kiloparsec scales.  Our interpretation is that the H$\\alpha$ and CO emission are produced in a population of molecular clouds warmed by a starburst, which has itself been triggered by the passage of a secondary galaxy through the gas-rich cluster core. In these most luminous H$\\alpha$ CCGs, this galaxy-wide starburst leads to the saturation of the optical emission line ratios at a uniform level characteristic of massive star formation. In the absence of such triggering, the cooled gas reservoir will be in a more quiescent state with a much lower level of H$\\alpha$ and CO emission; optical emission line ratios will be governed by other processes, e.g. excitation through cloud--cloud collisions (as described above) or mixing layers (Crawford \\& Fabian~1992), to mention just two possibilities. Starting from this low H$\\alpha$ luminosity quiescent state, adding some star-forming clumps will increase the overall H$\\alpha$ and CO luminosity and give rise to spatial variations in the optical line ratios. At the highest H$\\alpha$ luminosities, star-formation will be widespread, with little spatial variation in optical emission line ratios.  To test this interpretation, future IFU observations should be carried out on CCGs covering the [OII]$\\lambda 3727$ and [OIII]+H$\\beta$ emission lines and the sub-4000\\AA~continuum in order to probe current star formation, and with near-infrared IFUs to study the H$_{\\rm{2}}$ and Pa$\\alpha$ emission. Observations should also be performed on lower H$\\alpha$ luminosity CCGs in order to examine whether spatial variations in optical emission line ratios are indeed larger, due to less widespread star formation. On a longer timescale, mm-interferometers such as ALMA will reveal the spatial distribution of the CO emission (and dust) for a full comparison with the H$\\alpha$. Our findings also highlight the need for numerical modelling of these dense CO reservoirs, with and without external perturbations."
    },
    "0606/astro-ph0606133_arXiv.txt": {
        "abstract": "An effective quantum theory of gravitation in which gravity weakens at energies higher than $\\sim 10^{-3}~\\rm{eV}$ is one way to accommodate the apparent smallness of the cosmological constant. Such a theory predicts departures from the Newtonian inverse-square force law on distances below $\\sim 0.05$ mm. However, it is shown that this modification also leads to changes in the long-range behavior of gravity and is inconsistent with observed gravitational lenses. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606419_arXiv.txt": {
        "abstract": "We present deep integrated-light spectroscopy of nine M33 globular clusters taken with the Hectospec instrument at the MMT Observatory. Based on our spectroscopy and previous deep color-magnitude diagrams obtained with HST/WFPC2, we present evidence for the presence of a genuine intermediate-age globular cluster in M33. The analysis of Lick line indices indicates that all globular clusters are metal-poor ([Z/H]~$\\la-1.0$) and that cluster M33-C38 is $\\sim\\!5\\!-\\!8$ Gyr younger than the rest of the sample M33 star clusters. We find no evidence for a population of blue horizontal branch stars in the CMD of M33-C38, which rules out the possibility of an artificially young spectroscopic age due to the presence of hot stars. We infer a total mass of $5-9\\times10^{4}~M_{\\odot}$ for M33-C38, which implies that M33-C38 has survived $\\sim2-3$ times longer than some dynamical evolution model predictions for star clusters in M33, although it is not yet clear to which dynamical component of M33 -- thin disk, thick disk, halo -- the cluster is associated. ",
        "introduction": "Star clusters which have survived for several ($\\sim2-8$) billion years (intermediate-age clusters) can provide important clues concerning both the evolution history of their parent galaxies and for understanding the destruction processes which erode star cluster systems \\citep[e.g.,][and references therein]{goudfrooij04}. In the Galaxy, clusters of this age all fall in the category of ``old open clusters'' \\citep[see compilation in][]{friel95} -- they all reside in the Galactic thin disk, and have lower masses ($\\sim$few $10^3$~$M_{\\odot}$) than typical ancient globular clusters ($\\ga10^{4.5}~M_{\\odot}$). The oldest of these old open clusters help to pinpoint the age of the thin disk, while the ages and masses of the entire population help constrain the survival rates of clusters in the Galactic disk. Intermediate-age clusters are also known to exist in the Large and Small Magellanic Clouds (LMC and SMC). The LMC contains a population of clusters which are $1-3$~Gyr, while the SMC formed numerous clusters $4-8$~Gyr ago.  In more distant galaxies, it becomes increasingly difficult to establish the presence of intermediate-age clusters, primarily because the techniques available to study compact clusters become more limited. For example, the integrated optical colors of intermediate-age clusters are degenerate with those of ancient ($\\geq8$~Gyr) clusters. Although a number of clusters have absorption line index strengths (measured from integrated-light spectroscopy) which suggest that these objects are of intermediate age, their absolute ages cannot be definitively established due to the possible presence of hot stars in the core helium burning stage (hot ``horizontal branch'' or HB stars). These hot HB stars could potentially boost the Balmer absorption line strengths sufficiently to mimic younger ages for old globular clusters \\citep[e.g.,][]{lee05, maraston00}. Ideally, cluster ages should be determined from deep color-magnitude diagrams (CMDs) which reveal the main sequence turnoff; to date however, this has only been accomplished for a single cluster beyond the Magellanic Clouds \\citep[M31-SKHB 312;][]{brown04}, using 129 orbits of {\\it Hubble Space Telescope} (HST) time.  In this Letter, we present integrated-light (optical) spectroscopy for nine ancient star cluster candidates in the nearby spiral galaxy M33. These clusters have deep CMDs available from HST/WFPC2 observations, which clearly reveal the HB morphology. The combination of HB morphology and absorption line measurements result in robust relative age estimates, and are ideal to search for intermediate-age clusters in M33. ",
        "conclusions": "\\subsection{Evidence for an Intermediate-Age Cluster} In Figure~\\ref{ps:am}, we show the measurements for the metal-sensitive absorption line index [MgFe]\\arcmin\\ versus the age-sensitive Balmer line indices H$\\delta_{\\rm A}$, H$\\gamma_{\\rm A}$, and H$\\beta$, as is typically done to estimate the ages and chemical compositions of stellar populations from integrated-light spectroscopy \\citep[e.g.,][]{puzia05}. For comparison, we show simple stellar population (SSP) model predictions from \\cite{thomas04} for six different ages (solid lines) and six different values of [Z/H] (dashed lines). We also compared the M33 cluster measurements with the solar-scaled SSP model predictions of \\cite{bc03} at 4.7~\\AA\\ resolution (not shown), which yield results consistent with the \\cite{thomas04} models.~All panels in Figure~\\ref{ps:am} show a ``cloud'' of points located towards low metallicities and old ages, although we cannot rule out an age spread of the order of $\\Delta t/t\\la 0.3$, which translates into a $\\sim\\!4$~Gyr spread at 12 Gyr absolute age. Cluster C38 is clearly offset from the ``cloud\" in every panel, and clearly appears to have a younger age than the other clusters -- recall that our measurements have not been absolutely calibrated to the models, but should be robust in a relative sense. Therefore, Figure~\\ref{ps:am} implies that C38 is $\\sim\\!5-8$ Gyr younger than the other M33 clusters plotted here. Clusters with H$\\alpha$ seen in emission (R14 and U77) are plotted with different symbols in Figure~\\ref{ps:am}. The Balmer absorption-line strengths for these clusters thus represent lower limits.  \\begin{figure*} \\centering \\includegraphics[bb= 0 0 400 400, width=5.5cm]{f2a.eps} \\includegraphics[bb= 0 0 400 400, width=5.5cm]{f2b.eps} \\includegraphics[bb= 0 0 400 400, width=5.5cm]{f2c.eps} \\caption{Age-metallicity diagnostic plots for the sample M33 globular clusters, constructed from three different Balmer line indices: H$\\delta_{\\rm A}$ ({\\it left panel}), H$\\gamma_{\\rm A}$ ({\\it middle panel}), and H$\\beta$ ({\\it right panel}) vs.~[MgFe]\\arcmin\\ \\citep[see also Fig.~9 of][for details]{puzia05}. Over-plotted are SSP models of \\cite{thomas04} for [$\\alpha$/Fe]~$=0.3$, metallicities [Z/H]=$-2.25, -1.35, -0.33, 0.00$, 0.35, and 0.67 dex ({\\it dashed lines}), and ages 15, 10, 5, 3, 2, 1 ({\\it solid lines}). The thick solid line is the 5 Gyr isochrone, and is used to split between old and intermediate age globular clusters. Note that some clusters are not plotted in the individual panels due to contamination of their index passband definitions by bad pixels. Clusters with H$\\alpha$ observed in emission are indicated by open circles. All clusters are labeled as in Fig.~\\ref{fig1}. We point out that the Lick index measurements are not calibrated to the Lick system so that only relative measures should be derived from the shown diagrams.} \\label{ps:am} \\end{figure*}  Could the stronger Balmer line indices for C38, which suggest a younger age, be due to the presence of blue/hot HB stars? In Figure~\\ref{fig:cmd}, we show the resolved CMD of C38 based on deep HST/WFPC2 observations for this cluster, and also for M9 (one of the older clusters) for comparison. C38 shows no blue HB stars, while M9 does. Note that the HST observations have sufficient depth and resolution to detect blue HBs if they are present, but they clearly do not exist in C38. Therefore, the presence of hot HB stars, which has been suggested as a plausible explanation for apparent intermediate ages of clusters in different galaxies \\citep[such as in M31;][]{puzia05b}, cannot be the correct explanation in this case. The line indices and the knowledge of the HB morphology taken together, indicate that it is the {\\it age\\/} of the cluster which drives the observed Balmer line strengths.  There are two features from the resolved CMD of C38 shown in Figure~\\ref{fig:cmd} which add weight to this conclusion.  First, S00 found that the absolute I band magnitude of the red clump for this cluster is more luminous than all of the others in their sample, and that this behavior is reasonably well fit by the $\\sim2$~Gyr theoretical HB models of \\cite[see also Figure~26 in S00]{girardi99}. Secondly, when the CMD for C38 is cleaned to exclude the most crowded central regions, there appears to be a hint of a main sequence turnoff (MSTO) in the CMD, as illustrated in the left panel of Figure~\\ref{fig:cmd}. In contrast, M9 does not show evidence for a MSTO.~Also in this figure, we show a Z = 0.004 isochrone at an age of 2~Gyr from Girardi et al. (2000) for comparison. Taken together with the spectroscopic results and the abnormally bright red clump, the presence of a MSTO in the CMD of C38 provides strong evidence that C38 is a genuine intermediate-age ($\\sim2-5$~Gyr) star cluster in M33.  \\begin{figure}[!t] \\centering \\includegraphics[bb=46 291 585 752,width=8.5cm]{f3.eps} \\caption{$V-I$ versus $I$ CMD for M33 clusters C38 and M9. The CMDs have been restricted to include stars in an annulus between 1 and 2\\arcsec\\ away from the center of each cluster, in order to minimize blending and contamination. A Girardi isochrone for an age of 2 Gyr and $Z=0.004$ is overplotted onto the C38 CMD, which shows the red clump stars and a hint of a main sequence turnoff. A fiducial HB for the Galactic globular cluster M3 \\citep{johnson98} is overplotted onto the M33-M9 CMD, which clearly reveals the blue horizontal branch stars in this cluser.} \\label{fig:cmd} \\end{figure}   \\subsection{Properties of M33-C38}  Here we compile known properties of M33-C38. This cluster lies at a projected distance of $\\sim\\!4.5$ kpc from the center of M33, has a total $V$-band magnitude of 18.01 and $B-V=$ 0.73 and $V-I=0.89$ colors \\citep{christian88, chandar99, chandar02}. S00 estimated that it has an [Fe/H] of $-0.65\\pm0.16$ from the slope of the red giant branch, which is consistent with our spectroscopy given the uncertainties ([Fe/H]$=-1.1\\pm0.6$).  \\cite{chandar02} measure a radial velocity of $-145\\pm30$ km/s for this cluster. The local H\\,{\\sc i} (disk) velocity at the location of C38 is found to be $\\sim-100$~km/s from the radio maps of \\cite{warner73}. Although the cluster velocity measurement deviates from the local disk motion in Figure~\\ref{fig:cmd} of \\cite{chandar02}, (showing a comparison of cluster velocities versus the local disk motions relative to cluster age, C38 falls in a region which is ambiguous as to which M33 component C38 belongs, clusters with the most deviant velocities are almost certainly in a halo/thick disk component). Therefore, this object could belong to either the thin disk or a halo/thick disk component. Implications for these possibilities are discussed in the next section.  Unfortunately M33-C38 does not have any published $M/L$ ratios, which would provide a direct estimate of its total mass. There are, however, published $M/L_V$ values for three of the clusters presented in this work \\citep[H38, M9, and R12;][]{larsen02}. They have an average $M/L_V=1.53\\pm0.18$, which is consistent with $M/L_V$ measurements for Galactic GCs \\citep{mclaughlin00}. The integrated spectra presented here corroborate that these objects are ancient, even though their exact ages are not known. Given the fact that $M/L_V$ {\\it decreases} with decreasing age, as an upper limit for C38 we take the highest value found empirically by \\cite{larsen02} of $M/L_V\\!=\\!1.87$. As a lower limit, we assume the $M/L_V$ value predicted by the \\cite{bc03} models at an age of 2~Gyr and for a low metallicity system ($M/L_V\\!=\\!1.125$). Assuming a foreground reddening $E_{B-V}=0.1$ mag towards M33 and a distance modulus of 24.64 mag, we estimate that C38 has a mass in the range $5-9\\times10^{4}\\, M_{\\odot}$. This is about an order of magnitude higher than the masses for old Galactic open clusters, and within the range of ancient globular clusters.  \\subsection{Implications of Intermediate-Age Clusters in M33}  The confirmed presence of an intermediate-age cluster in M33 has important implications for star/cluster formation processes in spiral galaxies, regardless of which structural component C38 belongs to. Below we explore the implications of the presence of intermediate-age clusters in both a thick disk versus halo component, and also in the thin disk.  If C38 belongs to either a thick disk or halo component, it is likely that at least a few of the other clusters in our sample belong to the same structural component, so that a relatively large age spread among M33 clusters would be present in that component. In the Galaxy, halo GCs are known to be old and have an age spread of only $\\sim\\!3$~Gyr, whereas the GCs associated with the thick disk are essentially coeval \\citep{deangeli05}. We suggest that a large age spread for M33 clusters would favor a halo over a thick disk origin, for the following reasons. The thick disk in the Milky Way is relatively metal-rich, and believed to have formed during a single event, likely the accretion of a relatively massive satellite galaxy which puffed up the disk \\citep[e.g.][]{morrison93, abadi03, martin04}. Such an origin for a thick disk in M33 would imply that clusters in M33 formed for many Gyr relatively undisturbed in the disk, until several Gyr ago when an accretion event took place, resulting in a thick disk with a cluster population spanning a large range of ages. Given that to date there is no evidence for such an event in M33, and that this type of scenario would require that no significant merging occurred before or since, we suggest that it is unlikely that clusters residing in a thick disk would have an age spread of several Gyr.  Therefore, the presence of C38 seems to favor a halo component over a thick disk component. This would point to a longer timescale for the buildup of the M33 halo relative to that found for the Milky Way, as suggested previously by S00 and \\cite{chandar02}.  If instead of the halo/thick disk, C38 resides in the thin disk of M33 (the velocity measurements are ambiguous in discriminating between these two possibilities), the intermediate age of this object also has important implications. First, this would confirm the presence of relatively massive, intermediate-age disk clusters, which are not widely observed in our own Galaxy (but might be hidden behind large columns of dust). Second, the ability of massive clusters $(\\sim\\!10^4~M_{\\odot}$) to survive in the disk of M33 has been suggested to be limited to $<1$~Gyr. \\cite{lamers05} used cluster samples in M33 and their theory of cluster dissolution \\citep[e.g.][]{boutloukos03} to estimate directly from the cluster age and mass distributions the ability of clusters to survive in different environments. For M33, they find that $10^4~M_{\\odot}$ clusters will disrupt in $\\log({\\rm age})=8.8$ years ($\\sim\\!600$~Myr), and that a $5\\times10^{4}\\, M_{\\odot}$ cluster will disrupt in $\\sim1.5$~Gyr.~Mass estimates for C38 put it at $5-9\\times10^{4}\\, M_\\odot$. Although we do not know its precise age, it appears that C38 has likely survived $2-3$ times longer than this prediction. Regardless, if C38 is part of M33's thin disk, this points to the formation of a star cluster which is more massive than known counterparts of a similar age in the Local Group, and also to the ability of such an object to survive disruption forces in the environment of M33 for several billion years.  \\vspace{-0.5cm}"
    },
    "0606/astro-ph0606091.txt": {
        "abstract": "The origin of the soft X-ray excess emission observed in many type-1 AGN has been an unresolved problem in X-ray astronomy for over two decades. We develop the model proposed Gierli\\'nski \\& Done (2004), which models the soft excess with heavily smeared, ionized, absorption, by including the emission that {\\it must} be associated with this absorption. We show that, rather than hindering the ionized absorption model, the addition of the emission actually helps this model reproduce the soft excess. The emission fills in some of the absorption trough, while preserving the sharp rise at $\\sim$1 keV, allowing the total model to reproduce the soft excess curvature from a considerably wider range of model parameters. We demonstrate that this model is capable of reproducing even the strongest soft X-ray excesses by fitting it to the \\xmm~EPIC PN spectrum of PG1211+143, with good results. The addition of the emission reduces the column density required to fit these data by a factor $\\sim$2 and reduces the smearing velocity from $\\sim$0.28c to $\\sim$0.2c. Gierli\\'nski \\& Done suggested a tentative origin for the absorption in the innermost, accelerating, region of an accretion disk wind, and we highlight the advantages of this interpretation in comparison to accretion disk reflection models of the soft excess. Associating this material with a wind off the accretion disk results in several separate problems however, namely, the radial nature, and the massive implied mass-loss rate, of the wind. We propose an origin in a `failed wind', where the central X-ray source is strong enough to over-ionize the wind, removing the acceleration through line absorption before the material reaches escape velocity, allowing the material to fall back to the disk at larger radii. ",
        "introduction": "\\label{1}  The origin of the soft X-ray ($\\leq$2 keV) excess emission observed in many type-1 AGN has remained a long standing problem in X-ray astronomy despite considerable attention over the past two decades ({\\it e.g.} Arnaud \\etal 1985, Saxton \\etal 1993, Mineshige \\etal 2000). Several different ideas have been proposed for its origin including the thermal tail from the accretion disk ({\\it e.g.} Pounds \\etal 1986) and a separate cool Comptonizing region ({\\it e.g.} Zdziarski \\etal 1996) in either an optically thick transition region between the disk and a hot inner flow (Magdziarz \\etal 1998) or in an ionized skin on the accretion disk (Janiuk, Czerny \\& Madejski 2001). These models encounter considerable problems; specifically, that the best-fit temperature for such thermal components is remarkably constant across a wide range of objects that exhibit a considerable range of black hole masses, luminosities and accretion rates ({\\it e.g.} Czerny \\etal 2003; Gierli\\'nski \\& Done 2004: hereafter GD04; Crummy \\etal 2005).  Detailed modelling of the \\xmm~EPIC PN spectrum of the Narrow-line Seyfert 1 galaxy 1H 0707-495 suggested that Compton reflection from a partially ionized skin on the accretion disk can reproduce the strong soft excess (Fabian \\etal 2002). This model provides a natural explanation for the `constant temperature' of the soft excess emission since the `excess' is now primarily associated with the abrupt change in opacity at $\\sim$0.7 keV. Partially ionized material is much more reflective below this energy, and the rising reflected continuum is enhanced by associated CNO lines and radiative recombination (RRC) emission features. Together, these can produce the soft excess, with the advantage that the characteristic energy of the emission is fixed by atomic physics, so is largely independent of the properties of the AGN (Zycki \\& Czerny 1994). The narrow atomic features associated with ionized reflection are not evident in the spectra of the soft excess suggesting that large (relativistic) velocity smearing is present in this material (Fabian \\etal 2002). In addition, the largest soft excesses need a reflected fraction considerably greater than unity, arguing for an unusual geometry such as {\\it i}) a highly clumpy/corrugated disk surface so that the observed spectrum is dominated by multiple reflections of the mostly hidden intrinsic hard X-ray source (Fabian \\etal 2004; Ross \\etal 2002), or {\\it ii}) a model in which strong light-bending towards the black hole increases the X-ray flux incident on the disk ({\\it e.g.} Fabian \\etal 2005). The first of these models implies a complex physical configuration, perhaps as the result of disk instabilities that may be produced in high mass accretion rate sources such as Narrow-line Seyfert 1s (NLS1s), while the latter naturally produces the required large velocity smearing from the inner regions of the accretion disk.  Gierli\\'nski \\& Done (2004) proposed an alternative origin for the soft excess originating from similar atomic processes in partially ionized material, but with a different geometry. They suggested that the partially ionized material is seen in absorption, rather than in reflection. Again, a large velocity dispersion is required to hide the atomic line and edge features, and results in a broad, smooth absorption trough that adds considerable curvature to the X-ray spectrum between $\\sim$0.3-5 keV. The advantage of this model over reflection is that a larger range soft excess sizes can be produced without invoking geometric changes. However, while the reflection models include the self-consistent emission from the partially ionized material, the absorption model presented in GD04 does not. This is equivalent to assuming that the material subtends a small solid angle, in contrast with the large fraction of type-1 AGN that show soft excess emission. Chevallier \\etal (2005) include this emission as well as reflection and suggest that all three components work together to produce the soft X-ray excess.  GD04 suggest a tentative origin for the absorbing material in a UV line-driven accretion disk wind; that is material launched from the disk that is subsequently accelerated to high velocities by the radiation pressure on UV absorption lines, in similar fashion to the winds from hot stars ({\\it e.g.} Lamers \\& Cassinelli 1999), reasoning that a large velocity dispersion is implicit to material launched from the inner disk and hence strong velocity smearing is to be expected. The physical properties of such line-driven accretion disk winds have been investigated in considerable detail in other work ({\\it e.g.} Arav, Li \\& Begelman 1994, Murray \\& Chiang 1995, 1997 \\& 1998, Murray \\etal 1995, Proga, Stone \\& Kallmann 2000 - hereafter PSK00, Proga \\& Kallman 2004 - hereafter PK04, etc). The general properties implied for accretion disk winds match well to the blue-shifted ($\\leq$10$^{3}$ kms$^{-1}$) optical-UV-X-ray absorption lines seen in the Broad Absorption Line (BAL) quasars (Proga \\& Kallman 2004), and may also be responsible for the broad emission lines seen in all Quasars (Murray \\& Chiang 1997; Elvis 2000). Furthermore, such a wind provides an attractive and simple way of dynamically linking the otherwise distinct regions within the standard AGN unification scheme.  In this paper we extend the absorption model proposed in GD04 to include the emission that {\\it must} be associated with this absorption. In Section~\\ref{2} we briefly present the details of the photoionization modelling and the velocity smearing. In Section~\\ref{3} we examine the properties of the absorption and emission in more detail, focusing on the impact of the emission on the resulting spectral shape, and demonstrate that the combination of emission and absorption removes many of the fine-tuning issues present in the pure absorption model of GD04.  We test the impact of the addition of the emission on fits to real data by fitting the \\xmm~spectrum of the bright, strong soft excess Quasar PG1211+143 with our models of the absorption and emission from a disk wind, together with ionized reflection and the narrow absorption systems known to be present in this source (Section~\\ref{4}). Sections 5 and 6 discuss the physical origin of the absorbing material and its properties. ",
        "conclusions": "\\label{6}  GD04 proposed a model for the soft excess, observed in the X-ray spectra of many type-1 AGN, based on heavily smeared, ionized, absorption. They suggested a tentative origin for this absorption in the innermost, accelerating, region of an accretion disk wind. We expand on this model by including the emission that {\\it must} be associated with this absorption, focusing on the impact that the emission has on the total spectrum. Initially, we examine the properties of the emission, the absorption and the total spectrum as we vary the free parameters in the model, before testing the model against data for PG1211+143, a bright nearby quasar known to have one of the strongest soft X-ray excesses.  The main conclusion from this work is that the emission associated with heavily smeared, ionized, absorption actually helps this model reproduce the soft excess, by filling-in some of the absorption trough. This removes some of the fine tuning inherent in the original pure absorption model and, rather than reducing the curvature imposed on the spectrum by the model, the emission instead allows the model to reproduce the soft excess curvature between 0.6-1.5 keV from a wider range of column densities and ionization states, with lower velocity smearing.  We show this specifically by fitting the total smeared absorption/emission to the \\xmm~EPIC PN spectrum of PG1211+143, as part of a more complex model that accounts for complex narrow absorption systems and Compton reflection from an ionized accretion disk. This model results in a good fit to the data ($\\chi^{2}_{red}$=1.08) demonstrating that it is capable of reproducing even the strongest soft X-ray excesses observed. Furthermore, in comparison with the details discussed in GD04, the addition of the emission reduces the column density required to fit these data by a factor $\\sim$2 and reduces the smearing velocity from $\\sim$0.28c to $\\sim$0.2c. The accretion disk reflection component does contribute to the spectrum, but it does not produce a strong contribution to the soft excess. The reflection is smeared, but not dramatically so. This is in sharp contrast to the reflection dominated models for the soft excesses, which require extreme spin of the black hole.  However, there are significant problems with interpreting the absorbing material as an accretion disk wind. The interpretation as a UV line-driven disc wind, suggested in GD04, is especially problematic because such winds tend to be produced within $\\sim$25$^\\circ$ of the equatorial plane (PK04). Thus most type-1 AGN are expected to have lines of sight which would not intercept much of this material. This directly conflicts with the observations that most PG Quasars (type-1 AGN) show a soft excess. There is also a more generic problem in identifying this material as a wind. The required large velocity smearing implies a mass-loss rate of $\\dot M_{outflow}\\sim$190\\Msun~yr$^{-1}$, much larger than the accretion rate required to power the observed luminosity of $\\sim$1-2\\Msun~yr$^{-1}$ (see also Chevallier \\etal 2005).  We suggest an origin for this material as a `failed wind', {\\it i.e.}, that the central X-ray source is strong enough to over-ionize the wind, removing the acceleration through line absorption before the material reaches escape velocity. Thus the material is initially accelerated, but it cannot escape the system and eventually falls back to the disk. The unreasonably large mass-loss rate is then easily circumvented, since most of the material simply circulates, falling back and rejoining the disk at somewhat larger radii. Modelling such material is undoubtedly much more complex than described by our very simplistic assumptions on the density and velocity structure, but despite this we get a very good fit to the shape of the soft excess in PG1211+143 from a combination of absorption and emission in this material. Better models fit to a wide range of type-1 AGN will allow future studies to determine if this interpretation is physically plausible."
    },
    "0606/astro-ph0606668.txt": {
        "abstract": "It is well known that magnetic fields affect heat conduction in a different way in the direction parallel and perpendicular to the field. In this paper a formal derivation of this phenomenon and analytical expressions for the transport coefficients based in the Boltzmann equation are presented. Moreover, the Dufour effect or diffusion thermo effect is usually ignored in plasma transport theory. In this work, analytic expressions for the parallel and perpendicular thermal conductivities as well as the coefficients for both the thermal diffusion, or Soret, effect and the Dufour effect are formally derived. It is also shown how the heat conduction in the perpendicular direction decreases with increasing magnetic field and how in both directions the diffusion thermo effect is significant compared to thermal conduction. Other aspects of this work are also emphasized. ",
        "introduction": "The study of transport processes in plasmas is about to reach its seventieth birthday. Indeed, in 1936-1937 Landau published his pioneering paper by casting the Boltzmann collision integral in the form of a Fokker-Planck equation thus regarding that collisions in plasmas may be treated diffusively. Using very intuitive arguments he then arrived to what is today known as the collision integral in the logarithmic approximation for a gas with Coulomb interactions \\cite{key-1}. This method has become the basis of practically all approaches to the study of such processes in many applications. For instance, most of the work involving astrophysical problems use this method as further developed by Spitzer and collaborators \\cite{key-2,key-3} over fifty years ago, also later by other authors \\cite{key-4,key-5} and independently by Braginski \\cite{key-7,key-6}. The last effort along these lines is the apparently not fully recognized but magnificent work on transport theory in plasmas written by Balescu in 1988 \\cite{key-29}. In our opinion it is probably the best and most complete report written on the subject as of today. After Balescu's work, research on the kinetic theory of plasmas seems to have diverted in two directions. One is related to the kinetic problems which arise in fusion-related research. In tokamaks and other devices, the geometry and inhomogeneities present in the relatively large magnetic fields that prevail, affect the particle dynamics in a crucial way. This topic is referred to as neoclassical transport \\cite{key-30} and the literature here is abundant, a sample of which may be found in Refs. \\cite{key-9,key-10} and the works cited therein. The second application is characteristic in astrophysical problems where the Spitzer-Braginski approach seems to dominate. Typical situations are related to the problem of thermal conduction in tangled magnetic fields \\cite{key-39}, thermal conduction in clusters of galaxies \\cite{key-40} and similar problems \\cite{key-41}\\cite{key-42}. A few papers deal with traditional kinetic theory but they do not seem to proliferate. An example of a formal kinetic treatment of a two temperature gas can be found in Ref. \\cite{key-33} and of a multispecies plasma in Ref. \\cite{key-36}. However, a common denominator in most recent works in plasma kinetic theory is the use of the Landau-Fokker-Planck version of the kinetic equation and few have used the full Boltzmann equation to study this question. Indeed, in a recent authoritative monograph on the subject \\cite{key-43} it is clearly stated that Braginski's equations {}``are now being universally accepted and because of their completeness and correctness are considered standard''. We shall come back to these statements later on.  On the other hand, another approach to the study of transport phenomena in plasmas, using the full properties of the Boltzmann equation, was first attempted by Chapman and Cowling in 1939 in the first edition of their monumental work on kinetic theory \\cite{key-11}. It is very likely that since they had in mind mostly applications to the physics of metals, their method did not permeate to other branches of physics as smoothly as the Spitzer-Braginski ones did. In 1960 W. Marshall wrote a series of three reports entitled {}``The Kinetic Theory of an Ionized Gas'' in which he extended the preliminary results of Chapman and Cowling to the calculation of the transport coefficients of such a system when no magnetic field is present and in the presence of a magnetic field. Apparently these papers were never published in a regular journal and have thus been ignored in most work done in this subject \\cite{key-12}.  Another question that has been completely ignored in the several approaches to the transport processes, with the sole exception of Balescu's work \\cite{key-29}, is the derivation of the basic assumptions of Classical Irreversible Thermodynamics (CIT) from a kinetic model. This not just means discussions about the approach to equilibrium of a charged system, the H-theorem, and the explicit calculation of the entropy production exhibiting consistency with the second law of thermodynamics. It also includes the derivation of the complete equations relating the appropriate fluxes with their corresponding thermodynamic forces, the calculation of transport coefficients involved therein, specially those associated with the so-called cross-effects, and then showing explicitly that the matrix of such transport coefficients is a symmetric one, the well known Onsager reciprocity theorem. Curiously enough, even for an inert multicomponent mixture of dilute gases this latter property was obtained only recently \\cite{key-13}. This complete program has been accomplished for the so-called {}``vectorial'' transport processes in a dilute plasma, heat conduction, diffusion and electrical conduction. However the complete work, which includes viscomagnetism, is simply too bulky to attempt a publication in a single regular article \\cite{key-14}.  In this paper we would like to discuss with certain detail the problems of heat conduction and diffusion using the full Boltzmann equation as a formalism which not only involves the full two body dynamics of the interacting species but it also allows one to fulfill the program described in the previous paragraph. One reason for doing so is that these processes have become rather important in many astrophysical applications, specially in the enigmatic problem of the so-called cooling flows \\cite{key-15,key-17,key-40,key-41}. The main result which we want to emphasize here is that besides the heat transfer rates which arise from simple heat conduction due to thermal gradients, results which are essentially in agreement with those of Spitzer and Braginski, the influence of diffusion (mass transport) is rather surprising. With or without the presence of a magnetic field, the diffusion thermo effect or Dufour effect, as well as the pressure thermal effect, seem to play a rather important role in enhancing heat conduction. To our knowledge this particular result has never been published. It is as striking as the thermal analog of the well known Hall effect, the so-called Righi-Leduc effect, which we have reported in a separate paper \\cite{key-18}. Although similar results are also present in the case of electrical conduction we shall not deal with them here to keep the size of this paper within a reasonable length.  Yet, another reason for carrying out this work is to provide the kinetic foundations of magnetohydrodynamics within the framework of CIT which, as mentioned above, was partially done in Ref. \\cite{key-29} and later on extended to the relativistic regime \\cite{key-44}.  The structure of the paper is as follows. In Section II we briefly summarize the tenets of classical irreversible thermodynamics, Section III is devoted to the kinetic theory behind the calculation. In Section IV we compute and discuss the transport coefficients and Section V is left for some concluding remarks. ",
        "conclusions": "In this paper we have used the standard Boltzmann equation to study a fully ionized dilute plasma in a magnetic field. The main results obtained are explicit formulas for the transport coefficients associated with mass and heat flux as well as their cross effects. These latter ones have been completely ignored in the literature; here we stress their importance for different values of the magnetic field for a given particle density and temperature under the weak field approximation so that if $\\tau$ is the mean collision time, $\\omega_{e}\\tau$, where $\\omega_{e}$ is the Larmor frequency for electrons, is not too large. The full results are shown in the previous section. These results confirm Balescu's conjecture stating that the values of the transport coefficients are not very sensitive to the method used in obtaining them. Also, we obtain an enhanced heat conduction due to the Dufour cross-effect which leads to an effective thermal conductivity in all directions {[}Eq.\\,(\\ref{56})] even in the absence of a magnetic field.  We could have easily added an external field $\\vec{E}$ without any effort. This would lead to the study of electrical conductivity since in the expression for $\\vec{d}_{ij}$, Eq.\\,(\\ref{11}), the last term contains the full Lorentz force $\\vec{E}+\\vec{u}\\times\\vec{B}$. Since the electrical conduction flow is defined as\\begin{equation} \\vec{J}_{e}=\\sum_{i=a}^{b}\\frac{e_{i}}{m_{i}}\\vec{J}_{m}\\,,\\label{58}\\end{equation} and $\\vec{J}_{a}+\\vec{J}_{b}=0$,\\begin{equation} \\vec{J}_{e}=\\frac{m_{a}+m_{b}}{m_{a}m_{b}}e\\vec{J_{a}}\\,,\\label{59}\\end{equation} thus, using the results for the mass flow, $\\vec{J}_{e}$ can be readily computed from Eq.\\,(\\ref{59}). We avoided this discussion because of the intrinsic interest of electrical conduction which deserves a separate treatment. This has been discussed in another paper \\cite{key-45}.  The restriction here to a fully ionized hydrogen plasma is merely didactic. Taking $Z\\neq1$ and setting $n_{a}=\\mathcal{X}\\left(\\vec{r},t\\right)\\, n$ where $\\mathcal{X}\\left(\\vec{r},t\\right)<1$ is a trivial step, all the collision integrals in the Appendix B can be written in terms of $\\mathcal{X}\\left(\\vec{r},t\\right)$. In such a case $\\nabla\\frac{n_{a}}{n}\\neq0$ and besides the pressure diffusion coefficient there will be a true Fickian contribution. The reader may pursue this generalization with no problem whatsoever.  Finally, a word is indispensable concerning the results here obtained and those obtained by Balescu \\cite{key-29} and Braginski \\cite{key-7} using the Fokker-Planck (FP) equation. Balescu uses the Grad-like moment expansion method to compute the transport coefficients both for electrons and ions in a plasma for different values of $Z$. As we mentioned in the previous section the values for the thermal conductivities of the electrons is rather satisfactory. No comparison is possible with any cross coefficients since neither he nor Braginski calculate them. However, a word of caution must be here introduced. Moment methods, as has been clearly pointed out recently \\cite{key-34,key-35}, have to be dealt with carefully. On the one hand they may be subject to inconsistencies and secondly, when computing a transport coefficient, say with N (>5) moments, it is not clear which are the relevant contributions arising from different orders in the gradients. Although this problem may arise in the Chapman-Enskog expansion it is automatically taken care off by Knudsen's parameter. Therefore the different contributions in successive order in the gradients arise, Navier-Stokes, Burnett (second order), super-Burnett (third order) and so on. A fair comparison of the two methods would in principle require the introduction of such a parameter in the results derived from any moment expansion to classify the different contributions in terms of the gradient order. Despite its drawbacks, this method is still used in plasma kinetic theory calculations (see for example Ref. \\cite{key-38})  The results obtained by Braginski are carefully examined by Balescu so it is not necessary to repeat them here. Let us only add that in his result the relationship of the FP equation with the tenets of CIT is not pursued. The fact that the characteristic vector space to which the solution of the Boltzmann equation pertains is absent prevents the full discussion of the various cross effects when $\\vec{B}\\neq0$. The so-called frictional force introduced in the solution of the FP equation is artificial, the macroscopic (barycentric) velocity $\\vec{u}$ is not a thermodynamic force and according to Curie's principle \\cite{key-19} the only vectorial flows which may couple together are the heat, mass and charge fluxes. Thus we are convinced that our results are finer in their content than those arising from the FP equation. We will not pursue this discussion further since it would lead us away from the objective here acomplished.  \\appendix %dummy comment inserted by tex2lyx to ensure that this paragraph is not empty"
    },
    "0606/astro-ph0606127_arXiv.txt": {
        "abstract": "We present 105 stellar rotation periods in the young cluster IC\\,348, 75 of which are new detections, increasing the total number of known periods in this cluster to 143. The period distribution resembles that seen in the heart of the Orion Nebula Cluster by Herbst and colleagues.  Stars estimated to be less massive than 0.25\\,$M_{\\odot}$ show a unimodal distribution of fast rotators (P\\,$\\sim$\\,1--2~days) and a tail of slower rotators, while stars estimated to be more massive than 0.25\\,$M_{\\odot}$ show a bimodal distribution with peaks at $\\sim$\\,2 and $\\sim$\\,8 days.  We combine all published rotation periods in IC\\,348 with \\emph{Spitzer} mid-IR (3.6, 4.5, 5.8, and 8.0\\,\\micron) photometry, which provides an unprecedented efficient and reliable disk indicator in order to test the disk-braking paradigm. We find no evidence that the tail of slow rotators in low-mass stars or the long period peak in high-mass stars are preferentially populated by objects with disks as might be expected based on the current disk-braking model. Also, we find no significant correlation between period and the \\emph{magnitude} of the IR-excess, regardless of the mass range considered. Given the significant improvement of \\emph{Spitzer} observations over near-IR indicators of inner disks, our results \\emph{do not} support a strong correlation in this cluster between rotation period and the presence of a disk as predicted by disk-braking theory.  Rather, they are consistent with the suggestion that the correlation between period and the amplitude of the (I\\,--\\,K) excess reported in the past is a secondary manifestation of the correlation between the amplitude of near-IR excess and mass. By comparing our sample with a recent \\emph{Spitzer} census of IC\\,348, we find that the disk properties of our sample are indistinguishable from the overall disk properties of the cluster. We conclude that it is very unlikely that the lack of a correlation between rotation period and IR excess is due to a bias in the disk properties of our sample.  Finally, we find some indication that the disk fraction decreases significantly for stars with very short periods (P\\,$\\lesssim$\\,1.5 days). The fact that very fast rotators tend to have little or no excess has already been recognized by Rebull and colleagues for stars in the Orion Nebula Cluster.  This is the only feature of our sample that could \\emph{potentially} be interpreted as evidence for disk braking.  There is currently no alternative model to disk braking to explain the evolution of the angular momentum of pre-main-sequence stars.  It has been proposed that the observational signatures of disk braking might be significantly masked by the intrinsic breadth of the initial period distribution.  We argue that more rigorous modeling of angular momentum evolution and a quantitative analysis of the observational data are required before the disk-braking model can be regarded as inconsistent with observations. ",
        "introduction": "The evolution of the angular momentum of pre-main-sequence (PMS) stars is one of the longest-standing problems in star formation and is currently a controversial topic. As low-mass ($\\sim$0.1--1.2\\,$M_{\\odot}$) PMS stars evolve along their convective tracks, conservation of angular momentum dictates that they should spin up considerably.  Low-mass PMS stars are believed to contract by a factor of $\\sim$2--3 during their first 3\\,Myrs of evolution.  If these stars were to conserve their specific angular momentum, \\emph{j} (where \\emph{j}\\,$\\propto$\\,$R^{2}/P$), their rotation periods would be expected to decrease by a factor of $\\sim$\\,4--9. Using the models by \\citet[D98, hereafter]{dantona94, dantona98},\\footnote{Available at http://www.mporzio.astro.it/~dantona/prems.html} and assuming angular momentum conservation and homologous contraction, \\citet{herbst00} show that, given a starting period of 10 days,\\footnote{Adopting an initial period of 10 days is a conservative choice because recent observations \\citep[e.g.][]{covey05} show that deeply embedded PMS stars tend to rotate \\emph{faster} than optically revealed PMS stars.} all low-mass stars should rotate with periods shorter than $\\sim$2 days by an age of 2\\,Myrs.  However, observations show a large population of slow rotators (P\\,$\\sim$\\,8.0 days) in clusters that, according to the same evolutionary tracks, are $\\sim$2\\,Myrs old or older \\citep{luhman03}. Also, the \\emph{very} broad distribution of rotation periods in the zero-age main sequence (ZAMS) suggests a very broad range of angular momentum loss between the birth line and the ZAMS.  Proposed mechanisms to explain the presence of slow rotators come in the form of magnetic-field interaction between the PMS star and material in its inner disk, locking the star's rotation to the Keplerian velocity of the disk's inner edge \\citep{konigl91, shu94} and/or creating an accretion-driven wind \\citep{shu00}, thereby transferring angular momentum from the star to material from the disk, preventing it from spinning up, unchecked.  ``Disk locking'' and ``disk braking'' have become essential features of virtually every model of the angular momentum evolution of PMS stars.  The first order prediction of disk-locking and disk-braking models is that, as a group, stars with disks should rotate slower than stars that have already lost their disks and are free to spin up.  However, finding an observational correlation between the presence (or lack) of a circumstellar disk and the rotation period of PMS stars has proven a difficult task. It has been repeatedly argued that one of the main difficulties in finding this predicted correlation is the lack of an appropriate disk indicator \\citep{rebull01, herbst02}. \\emph{Spitzer}'s unprecedented mid-IR sensitivity allows, for the first time, an unambiguous identification of disks in a statistically significant sample of PMS stars with known rotation periods.  Here we report 75 previously unidentified rotation periods detected in the young cluster IC\\,348 which is situated in the East side of the Perseus Molecular Cloud. Our new periods, when combined with the periods in the literature, increase the total number of known periods in this cluster to 143. IC\\,348 is 315\\,pc away, young (2--3 Myrs), relatively compact ($\\sim$400 members, D\\,$\\sim$\\,20 arcminutes), and has low extinction (A$_{{\\rm V}}$\\,$<$\\,4) \\citep{luhman03}. We combine all published rotation periods in IC\\,348 with \\emph{Spitzer} photometry to investigate the disk-braking scenario. This paper is organized as follows. In Section~\\ref{previous}, we summarize previous observational attempts to identify a period--disk correlation.  In Section~\\ref{photometry}, we describe our time series photometry and the observed rotation period distribution.  In Section~\\ref{rotation}, we combine the rotation periods with the disk identification from \\emph{Spitzer} observations in order to search for the period--disk correlation predicted by the disk-braking paradigm.  In section 5, we discuss the implication of our results in the context of the disk-braking scenario. Finally, we summarize our results in Section 6. ",
        "conclusions": "We have obtained time series photometry of the young stellar cluster IC\\,348 and measured 75 new rotation periods. Our results increase the total number of known rotation periods in the cluster to 143.  We combined all published rotation periods in IC\\,348 with \\emph{Spitzer} photometry (3.6, 4.5, 5.8, and 8.0\\,\\micron) in order to test the disk braking paradigm, constructing a new, more reliable set of criteria for disk identification with IRAC data.  We find that the IC\\,348 rotation period distribution resembles that seen in the heart of the ONC: Stars estimated to be less massive than 0.25\\,$M_{\\odot}$ show a unimodal distribution dominated by fast rotators (P\\,$\\sim$\\,1--2 days) and a tail of slow rotators, while stars estimated to be more massive than 0.25\\,$M_{\\odot}$ show a bimodal distribution with peaks at $\\sim$2 and $\\sim$8 days. We find no evidence that the tail of slow rotators in low-mass stars or the long-period peak of high-mass stars are preferentially populated by stars with disks, a correlation which is predicted by the current disk-braking paradigm. Also, we find no significant correlation between period and the \\emph{magnitude} of the IR excess, regardless of the mass range considered.  Given the large improvement of IRAC observations over near-IR disk indicators, our results support the claim made by \\citet{littlefair05} that the correlation between period and (I--K) excess reported by several authors is a secondary manifestation of the correlation between near-IR excess and mass.  We find that the disk properties of our sample are indistinguishable from the disk properties of the cluster as a whole and conclude that it is very unlikely that the lack of a correlation between rotation period and IR excess is due to a bias in the disk properties of our sample.  Finally, we find some indication that the disk fraction might decrease significantly in stars with very short periods (P\\,$\\lesssim$\\,1.2 days). The fact that very fast rotators tend to have little or no excess has already been shown by \\citet{rebull05} for stars in the ONC and has been confirmed by a preliminary analysis of our ongoing work with objects in NGC\\,2264. The low disk fraction of these very fast rotators is the only feature of our sample that could \\emph{potentially} be interpreted as an evidence for disk braking.  However, the lack of evidence for disk braking (in the form of a correlation between PMS stellar rotation periods and IR excess) in all but the fastest rotators is not enough to rule out disk braking in PMS stars.  As shown by \\citet{rebull04}, current observational signatures of disk braking may be hidden by an initial large distribution of rotation periods in PMS stars.  Because there is currently no alternative mechanism for angular momentum loss in PMS stars and because there is evidence for star--disk interaction in very young stellar objects, a rigorous quantitative analysis of the effects of disk-braking parameters on current observational signatures is required to determine whether disk-braking may indeed play a significant role in the angular momentum evolution of these stars.  Simulations similar to those run by \\citet{rebull04}, with the improvements suggested in Section~\\ref{results} of this paper, can further constrain the importance of disk braking in the evolution of PMS stars.  We are in the process of testing such models and will present our results in a follow-up paper."
    },
    "0606/astro-ph0606311_arXiv.txt": {
        "abstract": "{} {Search for  Very High Energy $\\gamma$-ray  emission in the Kookaburra complex through  observations with the H.E.S.S. array.} {Stereoscopic imaging of Cherenkov light emission of the $\\gamma$-ray showers in the atmosphere is used for the reconstruction and selection of the events to search for  $\\gamma$-ray signals. Their spectrum is derived by a forward-folding maximum likelihood fit.} {Two extended $\\gamma$-ray sources with an angular (68\\%) radius of $3.3-3.4${\\arcmin}  are discovered at high ($>$13$\\sigma$) statistical significance: HESS~J1420-607 and HESS~J1418-609. They exhibit a flux above 1 TeV of ($2.97 \\pm 0.18_{\\rm stat}  \\pm 0.60_{\\rm sys}) \\times 10^{-12}$ and ($2.17 \\pm 0.17_{\\rm stat} \\pm 0.43_{\\rm sys}) \\times 10^{-12}$ cm$^{-2}$~s$^{-1}$, respectively, and similar hard photon indices $\\sim 2.2$. Multi-wavelength comparisons show spatial coincidence with the wings of the Kookaburra.  Two pulsar wind nebul{\\ae} candidates, K3/PSR J1420-6048 and the Rabbit, lie on the edge of the H.E.S.S. sources.} { The two new sources confirm  the non-thermal nature of at least parts of the two radio wings which overlap with the $\\gamma$-ray emission and establish their connection with the two X-ray pulsar wind nebul{\\ae} candidates. Given the large point spread function of EGRET, the unidentified source(s) 3EG~J1420$-$6038/GeV~J1417$-$6100 could possibly be related to either or both H.E.S.S. sources. The most likely explanation for the Very High Energy $\\gamma$-rays discovered by H.E.S.S. is inverse Compton emission of accelerated electrons on the Cosmic Microwave Background near the two candidate pulsar wind nebul{\\ae}, K3/PSR J1420-6048 and the Rabbit.  Two scenarios which could lead to the observed large ($\\sim$10 pc) offset-nebula type morphologies are briefly discussed. } ",
        "introduction": "\\begin{figure*}[t] \\centering \\includegraphics[width=0.5\\textwidth]{5511fig1.eps} \\caption{Smoothed excess map of the $1^{\\circ} \\times 1^{\\circ}$ field of view in the Kookaburra region. The unbinned excess map has been smoothed with a Gaussian of width 2{\\arcmin}. In the bottom left corner the point-spread function of this dataset smoothed in the same way is shown (along with the smoothing radius as a black dashed line). The white contours denote the 5$\\sigma$, 7.5$\\sigma$ and 10$\\sigma$ significance levels (with the outermost being the 5$\\sigma$ contour), derived with a point source integration radius of $0.1^{\\circ}$. The position of the pulsar PSR\\,J1420$-$6048 is marked with a star, the position of the rabbit (G313.3+0.1) is marked with a black and white triangle. The best fit positions of the two sources are marked with error crosses, the best-fit extensions are given as black circles.There is no significant evidence for a connecting bridge beyond what is expected from the Gaussian source shape convoluted with the point spread function of the instrument.} \\label{fig::skymap} \\end{figure*}   The complex of compact and extended radio/X-ray sources called Kookaburra, after the name of the Australian bird~\\citep{rrjg99}, spans over about one square degree along the Galactic plane around $l=313.4^{\\circ}$. It has been extensively studied in the search for counterparts to the unidentified EGRET source (or sources, see section 3.3) 3EG~J1420$-$6038/GeV~J1417$-$6100~\\citep{ht99,lm97}.  Radio images of this region have revealed a  large circular thermal shell with a broad wing to the North-East and a narrower wing to the South-West. Diffuse X-ray emission in the wings and point sources have been discovered through ASCA observations~\\citep{rrjg99,rrj01}, followed by higher resolution imaging with XMM-{\\it Newton} and {\\it Chandra}~\\citep{ng05}. A 3$^\\prime$  X-ray/radio nebula in the North-Eastern wing contains a young ($\\tau_c=1.3 \\times 10^4$ yr) and very energetic radio pulsar (${\\dot E} = 1.0\\times 10^{37}$erg/s), PSR J1420$-$6048~\\citep{dam01}, and has been proposed as a candidate pulsar wind nebula (PWN) counterpart to the EGRET source(s). A brighter nebula, G313.1+0.1, called the `Rabbit' \\citep{rrjg99}, lying in the South-Western wing and exhibiting extended hard X-ray emission has been proposed as a plausible PWN contributing also to the $\\gamma$-ray emission detected by EGRET.      In this paper, observations of the Kookaburra region with the H.E.S.S. (High Energy Stereoscopic System) telescopes and the discovery of two Very High Energy (VHE) $\\gamma$-ray sources coinciding with its wings are reported. H.E.S.S. is an array of four imaging atmospheric Cherenkov telescopes located in the Khomas Highland of Namibia~\\citep{HESS}. Each H.E.S.S. telescope has a mirror area of 107 m$^2$ ~\\citep{HESSOptics} and a total field of view of 5$^\\circ$\\citep{HESSCamera}. The system is run in a coincidence mode~\\citep{HESSTrigger} requiring at least two of the four telescopes to have triggered in each event. The H.E.S.S. instrument has an angular resolution of $\\sim 5^{\\prime}$ per event and a point-source sensitivity of $<2.0\\times\\,10^{-13}$ cm$^{-2}$s$^{-1}$ (1\\% of the flux from the Crab Nebula above 1 TeV) for a $5\\,\\sigma$ detection in a 25 hour observation.  Section 2 describes the  H.E.S.S. observations, data reduction and results. In section 3 multi-wavelength comparisons are made and in the following section  the interpretation of data together with possible associations for the new H.E.S.S. sources are discussed.   \\begin{figure*} \\centering \\includegraphics[width=0.75\\textwidth]{5511fig2.eps} \\caption{Reconstructed differential VHE $\\gamma$-ray spectrum of HESS\\,J1420$-$607 (full circles) and HESS\\,J1418$-$609 (open circles) along with power law fits to the two spectra. The flux points for HESS\\,J1418$-$609 have been scaled by a factor of 0.1 for ease of viewing. The fit to the HESS\\,J1420$-$607 data results in a photon index of $2.17\\pm0.06_{\\rm stat}\\pm 0.1_{\\rm sys} $ with a flux normalisation at 1 TeV of $(3.48 \\pm 0.20_{\\rm stat}\\pm 0.70_{\\rm sys}) \\times 10^{-12} \\mathrm{cm}^{-2} \\, \\mathrm{s}^{-1} \\, \\mathrm{TeV}^{-1}$. The power-law fit for HESS\\,J1418$-$609 yields a photon index of $2.22 \\pm 0.08_{\\rm stat}\\pm 0.1_{\\rm sys}$ and a flux normalisation of $(2.64 \\pm 0.20_{\\rm stat} \\pm 0.53_{\\rm sys})\\times 10^{-12} \\mathrm{cm}^{-2} \\, \\mathrm{s}^{-1} \\, \\mathrm{TeV}^{-1}$. The dashed line shows the Crab nebula spectrum as measured by H.E.S.S.~\\citep{HESSCrab}.} \\label{fig::spectrum} \\end{figure*} ",
        "conclusions": "Two extended VHE sources have been discovered in the wings of the Kookaburra complex, HESS~J1420$-$609 and HESS~J1418$-$607. They show similar $\\gamma$-ray angular extensions ($\\sim$ 3.3${\\arcmin}$-$3.4${\\arcmin}) and hard VHE $\\gamma$-ray spectra. This discovery confirms the non-thermal nature of at least parts of the wings of the Kookaburra which overlap with the VHE emission and establishes their connection with the two X-ray PWN candidates, K3  and the Rabbit. Within the limits of available multi-wavelength data, the SEDs of the two sources show also remarkable similarities and suggest analogous underlying objects and emission processes. The VHE $\\gamma$-ray emission could most plausibly be explained by IC emission from these PWNe in an offset-type configuration near the energetic PSR~J1420$-$6048, and the candidate pulsar in the Rabbit nebula. Given the poor spatial resolution of EGRET measurements, the unidentified source(s) 3EG~J1420$-$6038/GeV~J1417$-$6100 could possibly be related to either or both H.E.S.S. sources through a PWN-type emission above 100 MeV. The EGRET flux may also contain pulsed emission from the pulsar associated with K3 and the candidate pulsar in the Rabbit. Future GLAST observations of these objects should be able to establish the status of such pulsed emission, resulting in a more certain multiwavelength interpretation of the H.E.S.S. sources. The detection of the Rabbit pulsar, the morphology of the VHE and X-ray nebul{\\ae}, and the confirmation of a GeV-TeV connection are important considerations that need to be addressed through further investigations."
    },
    "0606/astro-ph0606582_arXiv.txt": {
        "abstract": "We present a photometric study of the star-forming region NGC 346 and its surrounding field in the Small Magellanic Cloud, using data taken with the Advanced Camera for Surveys (ACS) on board the Hubble Space Telescope (HST).  The data set contains both short and long exposures for increased dynamic range, and photometry was performed using the ACS module of the stellar photometry package DOLPHOT.  We detected almost 100,000 stars over a magnitude range of $V \\sim 11$ to $V \\sim 28$ mag, including all stellar types from the most massive young stars to faint lower main sequence and pre-main sequence stars. We find that this region, which is characterized by a plethora of stellar systems and interesting objects, is an outstanding example of mixed stellar populations. We take into account different features of the color-magnitude diagram of all the detected stars to distinguish the two dominant stellar systems: The stellar association NGC 346 and the old spherical star cluster BS 90. These observations provide a complete stellar sample of a field about 5\\arcmin\\ $\\times$ 5\\arcmin\\ around the most active star-forming region in this galaxy. Considering the importance of these data for various investigations in the area, we provide the full stellar catalog from our photometry. This paper is the first part of an ongoing study to investigate in detail the two dominant stellar systems in the area and their surrounding field. ",
        "introduction": "The investigation we present here deals with the stellar content of the star forming region in the vicinity of the stellar association NGC 346 in the Small Magellanic Cloud (SMC). NGC 346 is the largest stellar concentration in the SMC, related to the brightest {\\sc Hii} region of this galaxy, named LHA 115-N 66 or in short N 66 (Henize 1956), with an H{\\alp} luminosity almost 60 times higher than Orion's (Kennicutt 1984). Such active star forming regions are rare and can only be compared to star-bursts, like e.g., NGC 3603 in our Galaxy (Sher 1964).  As a consequence NGC 346 has been the subject of intense studies to establish the properties of the {\\sc Hii} region and of the bright early-type stellar content of the association in the low-metallicity environment of the SMC.  In a spectroscopic study of 42 bright main sequence stars Massey et al. (1989) confirmed 33 O-stars in the area of NGC 346. 22 of them are located in the central region of the association, and eleven are of type O6.5 or earlier. The stellar sample of Massey et al. is essentially complete down to $\\sim$ 10 M{\\solar} with six stars in the mass range 40 - 85 M{\\solar} and with a background field population of $\\sim$ 5 M{\\solar} stars.  An ultraviolet and optical spectral atlas of O stars in the SMC is compiled by Walborn et al. (2000) with data from the HST Space Telescope Imaging Spectrograph (STIS), AAT and ESO 3.6 m. Systematic phenomena in the sample include weak stellar-wind profiles on the O main sequence and in late-O giants, the absence of Si {\\sc iv} wind features throughout the O-giant sequence, low terminal velocities and enhanced He {\\sc ii} wind features in Of super-giants. The spectra of six O-type stars in NGC 346 from the Walborn et al. (2000) sample were modeled by Bouret et al. (2003) to determine their chemical abundances and wind parameters. The majority of the stars reveal CNO cycle-processed material at their surfaces during the main-sequence stage, thus indicating fast stellar rotation and/or very efficient mixing processes. The derived effective temperatures are lower than predicted from the widely used relation between spectral type and $T_{\\rm eff}$, resulting in lower stellar luminosities and lower ionizing fluxes.  However, Massey et al. (2005) argue that the data used by Bouret et al. (2003) were limited with respect to sky (nebular) subtraction. The effective temperatures of O3 - O7 dwarfs and giants in the SMC, as found by Massey et al. (2005) from their investigation of 33 O-type stars in the Magellanic Clouds, are about 4000 K hotter than for stars of the same spectral type in the Milky Way. Consequently, the effective temperature scale for O dwarfs in NGC 346 derived by these authors is significantly hotter than the Bouret et al. (2003) values. The differences decrease as one approaches stars of spectral type B0~V. In general, Massey et al. (2005) found that the winds momentum of O stars in the Magellanic Clouds scales with luminosity and metallicity as predicted by the theory for radiatively driven wind (Kudritzki 2002; Vink et al. 2001), supporting the use of photospheric analyses of hot luminous stars as a distance indicator for galaxies with resolved massive stellar populations.  A search for Be-stars in a field of 10\\arcmin\\ $\\times$ 10\\arcmin\\ centered on NGC 346 (Keller et al. 1999) found that only about 11\\% of the upper main sequence stars (with V $<$ 16 mag) located in the main body of the association show indeed strong H\\alp\\ emission. These authors adopt the most widely accepted explanation for the Be phenomenon, the rapid rotation of the B star (Bjorkman \\& Cassinelli 1993), due to which the stellar wind from a star is concentrated and confined to a disk around its equatorial plane. They relate, thus, the low fraction of Be stars in the cluster with the suggestion that it has been formed from low angular momentum gas.  The general area of NGC 346/N 66 hosts at least two supernova remnants. Not much is known about the first, named SNR 0056$-$7233, which has been identified by Ye et al. (1991) to the southwest at a projected distance of about 5.4 pc from the center of NGC 346. The second supernova remnant, named SNR 0057$-$7226, is located northeast close to the Wolf-Rayet luminous blue variable system HD 5980 in the vicinity of NGC 346. Koenigsberger et al. (2001) observed with HST/STIS interstellar and circumstellar absorption components along the line of sight toward HD 5980. They argued that SNR 0057$-$7226 had probably a progenitor, that was one of the brightest members of the association. An estimation of $T \\sim$ 40,000 K and a total mass between 400 and 1000 M{\\solar} for the supernova remnant shell was made by these authors. \\textit{Chandra} observations have shown that HD 5980 is surrounded by a region of diffuse X-ray emission from the supernova remnant, while NGC 346 itself shows only relatively faint X-ray emission, most of which seems correlated with the location of the brightest stars in the core of the association (Naz{\\'e} et al. 2003).  \\begin{figure*}[t!] \\centerline{\\hbox{ \\includegraphics[width=8.75truecm,angle=0]{f2a.ps} \\includegraphics[width=9.truecm,angle=0]{f2b.ps} }} \\caption{Typical uncertainties of photometry (left) and completeness factors (right) as derived by DOLPHOT from all datasets, for both bands. The bifurcation in the photometric errors is due to varying depth within the overall field.  The stars in the high-sigma branch are located in regions covered only by single exposures (dataset J92FA3).} \\label{fig_pherr-cmp} \\end{figure*}  Danforth et al. (2003) found with the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) strong {\\sc O vi} and {\\sc C iii} emission from a position at the edge of SNR 0057$-$7226. They determined the physical parameters of the interaction zone with N 66 and found that ionizing photons from massive stars of the association likely affect the ionization balance in the post-shock gas, hindering the production of lower ionization and neutral species. Danforth et al. (2003) proposed a physical relationship between SNR 0057$-$7226 and N 66, which suggests a slowly expanding bubble around NGC 346, powered by stellar winds and thermal pressure, while strong {\\sc O vi}, {\\sc C iii}, and X-ray emission arise from the shock interaction of the super-nova remnant with the denser, ionized gas on the rear side of N 66.  Rela{\\~n}o et al. (2002) computed photoionization models of the {\\sc H ii} region around NGC 346 with the CLOUDY 94 algorithm (Ferland et al. 1998) based on spectrophotometric data of Peimbert et al. (2000). These models showed that about 45\\% of the photons produced by the ionizing stars escape from NGC 346, implying that the system is a major source of ionizing flux for the surrounding diffuse interstellar medium. They also suggested that N 66 is a spherical density-bounded nebula.  The recent release of the {\\em Advanced Camera for Surveys} imaging toward the general area of NGC 346 from the {\\em HST Data Archive} offers a unique opportunity for the detailed photometric study not only of the bright massive stellar content, which is well known, but mostly of the fainter stellar members of the area down to the lower main sequence stars, which represent the majority of the stellar populations. In this paper we present our photometry of these observations, that cover a field about 5\\arcmin\\ $\\times$ 5\\arcmin centered on the association NGC 346. Our photometry detects magnitudes as faint as $V$ \\gsim\\ 26.5 mag, providing one of the richest stellar samples observed in the SMC. We also describe the stellar content of the area as it is seen in the derived color-magnitude diagram.  Specifically, in section 2 we present the observations and the photometry with software especially designed for ACS imaging. In the following section we present the stellar catalog and the color-magnitude diagram. We identify the different stellar populations coexisting in the area and we discuss their spatial distribution (\\S 3). A general description of the stellar content of the two main stellar concentrations in the area, the young stellar association NGC 346, and the old spherical cluster [BS95] 90 or in short BS 90 (Bica \\& Schmitt 1995) is given in section 4.  Final remarks and future prospects of the science that this catalog can provide are discussed in section 5.   \\begin{deluxetable}{ccccc} \\tablewidth{0pt} \\tablecaption{Log of the observations. Datasets refer to HST archive catalog. Exposure times (T$_{\\rm expo}$) per filter are given in seconds. \\label{tab1}} \\tablehead{\\colhead{Visit} & \\colhead{RA} & \\colhead{DEC} & \\multicolumn{2}{c}{T$_{\\rm expo}$}\\\\ \\colhead{Dataset} & \\multicolumn{2}{c}{(J2000)}& \\colhead{F555W} & \\colhead{F814W}} \\startdata J92F01&00:59:07.61&$-$72:09:39.8&1$\\times$456&1$\\times$484\\\\ &00:59:07.31&$-$72:09:40.0&1$\\times$456&1$\\times$484\\\\ &00:59:07.02&$-$72:09:40.1&1$\\times$456&1$\\times$484\\\\ &00:59:06.72&$-$72:09:40.3&1$\\times$456&1$\\times$484\\\\ &00:59:06.72&$-$72:09:40.3&2$\\times$3&2$\\times$2\\\\ \\tableline J92F02&00:59:05.58&$-$72:11:21.5&1$\\times$483&1$\\times$450\\\\ &00:59:05.28&$-$72:11:21.6&1$\\times$483&1$\\times$450\\\\ &00:59:04.99&$-$72:11:21.6&1$\\times$483&1$\\times$450\\\\ &00:59:04.68&$-$72:11:21.7&1$\\times$483&1$\\times$450\\\\ &00:59:04.68&$-$72:11:21.7&2$\\times$3&2$\\times$2\\\\ \\tableline J92FA3&00:59:06.18&$-$72:10:31.2&1$\\times$380&1$\\times$380\\\\ \\enddata \\end{deluxetable} ",
        "conclusions": ""
    },
    "0606/astro-ph0606330.txt": {
        "abstract": "We study the clustering properties of $K$-selected galaxies at $2<z<3.5$ using deep multiwavelength imaging in three fields from the MUSYC survey.  These are the first measurements to probe the spatial correlation function of $K$-selected galaxies in this redshift range on large scales, allowing for robust conclusions about the dark matter halos that host these galaxies.  $K$-selected galaxies with $K<21$ have a correlation length $r_0 \\sim 6 h^{-1} \\textrm{Mpc}$, larger than typical values found for optically-selected galaxies.  The correlation length does not depend on $K$-band magnitude in our sample, but it does increase strongly with color; the $J-K>2.3$ distant red galaxies (DRGs) have $r_0 \\sim 11 h^{-1} \\textrm{Mpc}$.  Furthermore, contrary to findings for optically-selected galaxies $K$-selected galaxies that are faint in the $R$-band cluster more strongly than brighter galaxies.  These results suggest that a color-density relation was in place at $z>2$; it will be interesting to see whether this relation is driven by galaxies with old stellar populations or by dusty star forming galaxies.  Irrespective of the cause, our results indicate that $K$-bright blue galaxies and $K$-bright red galaxies are fundamentally different, as they have different clustering properties.  Using a simple model of one galaxy per halo, we infer halo masses $\\sim 5 \\times 10^{12} M_\\odot$ for $K<21$ galaxies and $\\sim 2 \\times 10^{13} M_\\odot$ for DRGs.  A comparison of the observed space density of DRGs to the density of their host halos suggests large halo occupation numbers; however, this result is at odds with the lack of a strong small-scale excess in the angular correlation function.  Using the predicted evolution of halo mass to investigate relationships between galaxy populations at different redshifts, we find that the $z=0$ descendants of the galaxies considered here reside primarily in groups and clusters. ",
        "introduction": "\\label{sec-intro}  Optical surveys of the high-redshift universe have been very successful in finding relatively unobscured star-forming galaxies, primarily via the $U$-dropout technique.  These $z \\sim 3$ Lyman break galaxies (LBGs) typically have stellar masses $\\sim 10^{10} M_\\odot$, star formation rates of $10-100 M_\\odot \\textrm{yr}^{-1}$, and are thought to dominate the star formation density at that epoch \\citep{steidel03,shapley01,reddy05}.  However, it is becoming increasingly clear that substantial numbers of galaxies exist at these redshifts that have little rest-frame UV luminosity and are thus underrepresented in optical surveys.  Such galaxies may be detected in the near-infrared (NIR), which samples the rest-frame optical.  One criterion used to select galaxies in the NIR is $J-K>2.3$ \\citep{franx03,vandokkum03}.  These distant red galaxies (DRGs) typically have high star formation rates, $\\gtrsim 100 M_\\odot \\textrm{yr}^{-1}$, and dust obscuration $A_V > 1$, but must also have significant populations of evolved stars in order to explain their colors and spectra \\citep{forster04,papovich06,kriek06a}.  Some DRGs show little or no evidence of active star formation \\citep{labbe05, kriek06a,kriek06b,reddy06}.  DRGs must in general be very massive to account for their significant $K$-brightness; stellar population synthesis models imply masses $M_* \\sim ~10^{11} M_\\odot$.  Indeed, $~95\\%$ of galaxies with $M_* > 10^{11} M_\\odot$ at $2<z<3$ have $K<21.3$ \\citep{vandokkum06}.  Conversely, the median galaxy in this mass range has $R \\sim 25.9$, fainter than the limits typically reached by optical surveys.  The relationship between $K$-selected samples and optically-selected samples, not to mention present-day galaxies, remains unclear.  While typical optically-selected galaxies have different properties than $K$-selected galaxies, the $K$-bright subsample of optically-selected galaxies have stellar masses, star formation rates, and metallicities that are in approximate agreement with $K$-selected galaxies \\citep{shapley04,reddy05}.  Nevertheless, it is not clear whether the differences between these galaxies are transient (e.g. dust geometry, starbursts, mergers) or fundamental (e.g. age of the underlying old stellar populations, mass, environment).  An understanding of the nature of the differences between populations is essential to place them in evolutionary scenarios.  One way to investigate differences between galaxy populations is to measure their clustering properties.  As clustering measurements provide information that is independent of photometric properties, they can be used to distinguish between transient and fundamental differences between galaxy populations.  In the halo model of galaxy formation, the large-scale distribution of galaxies is determined by the distribution of dark matter halos.  The correlation function of galaxies can therefore be associated with the correlation function of the halos in which they reside.  Halo clustering, in turn, is a strong function of halo mass \\citep{mo96}, providing a means to study the relationship between galaxy properties and the mass of the dark matter halos.  Several studies have measured the dependence of clustering strength of high-redshift galaxies on color.  \\citet{daddi03} use the ultradeep imaging of the $4.5 \\textrm{arcmin}^2$ FIRES HDF-S field \\citep{labbe03} to study the clustering characteristics of $K$-selected galaxies at $2<z<4$.  Their most striking finding is that the correlation length increases strongly with $J-K$ color, with the reddest galaxies in their sample having correlation lengths $r_0 = 10-15h^{-1} \\textrm{Mpc}$, comparable to the most luminous red galaxies in the local universe.  The mass of dark matter halos with similar correlation lengths is $>10^{13} M_\\odot$, yet the galaxy number density is $\\sim 100$ times larger than that expected for such massive dark matter halos, implying that many galaxies must share the same halo.  More recently, \\citet{grazian06} measured a $r_0 = 13.4^{+3.0}_{-3.2} h^{-1} \\textrm{Mpc}$ for $z>2$ DRGs using the larger $135 \\textrm{arcmin}^2$ GOODS-CDFS field, also indicating that red galaxies are located in very massive halos.  The interpretation of these clustering measurements is complicated by the fact that, in order to derive information about the dark matter halos, the correlation function must be measured on large scales.  The correlation function has a contribution from galaxies that share halos (hereafter the ``1-halo'' term) and galaxies in separate halos (the ``2-halo'' term).  The shapes of these two contributions are shown with impressive detail in the correlation functions of large samples of $z \\sim 4$ LBGs presented by \\citet{lee06} and \\citet{ouchi05}.  In order to derive meaningful constraints on large-scale clustering properties, and thus on the host dark matter halos, it is important that both of these terms are taken into consideration.  If the correlation function is parameterized as a simple power-law then it should be measured on scales where the 2-halo term dominates the clustering signal.  In practice, a firm lower limit to the radial range in which the angular correlation function should be fitted is the halo virial radius.  At $z=3$, the virial radius $r_{200}$ of a $10^{13} M_\\odot$ halo corresponds to $22 \\arcsec$ \\citep[e.g.][]{mo02}.  The majority of the clustering signal from \\citet{daddi03} and \\citet{grazian06} is on scales $\\theta \\lesssim 30 \\arcsec$, which may lead to gross overestimates of the large-scale correlation length and the mass of the host dark matter halos.  In particular, \\citet{zheng04} shows that the measurements of \\citet{daddi03} are consistent with models in which the large-scale correlation length is as low as $r_0 \\sim 5 h^{-1}$ and the typical halo mass is $\\sim 10^{12} M_\\odot$.  This suggests that DRGs and LBGs may occupy similar halos, but that DRGs have higher occupation numbers.  The goal of this work is to study the clustering characteristics of a $K$-selected population of galaxies at $2<z<3.5$.  The increased field-of-view of our imaging allows for an improved determination of the clustering strength of $K$-selected galaxies at angular separations sufficiently large to investigate the large-scale distribution of galaxies, and thus to provide more meaningful estimates of the masses of the halos in which they reside. Secondarily, we wish to analyze the clustering results using models of halo clustering, and to use these models to shed light on evolutionary scenarios for $z>2$ galaxies.  These types of analyses have previously been performed for optically-selected samples \\citep[e.g.][]{moustakas02, ouchi04, adelberger05a}.  Throughout, we use the cosmological parameters $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70 h_{70} \\textrm{ km} \\textrm{ s}^{-1} \\textrm{ Mpc}^{-1}$, and $\\sigma_8=0.9$.  Results are given using $h_{70}=1$, except for correlation lengths which are scaled to units of $h=0.7$ in order to facilitate comparison to previous studies.  Optical magnitudes are given in AB and NIR magnitudes are given on the Vega system. ",
        "conclusions": "\\label{sec:discussion}  We have used $300 \\textrm{arcmin}^2$ of $UBVRIzJHK$ imaging from the MUSYC survey to study the angular and spatial correlation functions of $K$-selected galaxies with $2<z_{\\rm{phot}}<3.5$ and $K<21$.  The correlation length for this sample is $r_0 = 6.0^{+0.9}_{-1.1} h^{-1} \\textrm{Mpc}$, $\\sim 50\\%$ larger than optically-selected galaxies at similar redshifts \\citep{adelberger05a, lee06}.  The clustering of galaxies increases strongly with $J-K$ and $R-K$ color, and the $J-K>2.3$ population of distant red galaxies has $r_0 = 11.1^{+1.3}_{-1.4} h^{-1} \\textrm{Mpc}$.  Our results for DRGs are lower than previous results \\citep{daddi03,grazian06}.  This may be partially due to our smaller uncertainties.  Additionally, previous studies were only able to constrain the correlation function on small scales, where the signal may be strongly affected by galaxies that share the same halo \\citep{zheng04, ouchi05, lee06}.  In contrast, we perform fits to the correlation function on scales larger than the typical halo virial radius where this effect is reduced.  Nevertheless we confirm the basic trend indicated by previous studies, in which red galaxies are much more strongly clustered than optically-selected galaxies.  Moreover, this trend is not simply due to the fact that the red galaxies used in the previous studies were selected in the $K$-band: clustering increases strongly with color even within our $K$-selected sample, while there is no significant relationship between $K$-magnitude and color or between $K$-magnitude and correlation length.  These results suggest that a color-density relationship was in place at $z>2$, as is also indicated by the properties of optically-selected galaxies in different environments \\citep{steidel05}.  Whether this relationship is driven by galaxies that are red because of dust obscuration, or because of low specific star formation rates, remains to be seen.  \\citet{cucciati06} and \\citet{cooper06} investigate the fraction of red galaxies as a function of local galaxy density, finding that the color-density relation extends only to $z \\sim 1.3-1.5$.  This apparent contradiction with our results may be due to the difference in measurement techniques (fraction of red galaxies vs. correlation lengths).  These studies may also be dominated by galaxies that are less massive than the $z \\sim 1.5$ descendants of MUSYC galaxies, so there may not be a contradiction if the evolution of the color-density relationship is mass-dependent \\citep{cooper06}.  Finally, these studies are based on spectroscopy of galaxies that are selected in the optical, leaving the possibility that they are incomplete for the reddest galaxies at these redshifts.  A color-density relation at $z>2$ has implications for current galaxy selection techniques.  \\citet{adelberger05b} show that, among an optically-selected sample of $z \\sim 2.2$ BX objects, there is a strong correlation between clustering strength and $K$-band magnitude. Additionally, \\citet{shapley04} show that these galaxies are also massive, metal rich, and have high star formation rates.  Taken together, these results indicate that the $K$-bright galaxies found by optical surveys show similar properties to those uncovered by NIR surveys \\citep[e.g.][]{franx03, daddi03, daddi04, vandokkum04, vandokkum06, kriek06a}.  It has been argued that the primary difference between optically-bright massive galaxies and optically-faint massive galaxies may be that the former are observed during a chance period of unobscured star formation \\citep[e.g.][]{shapley05}.  However, the observed relationship between clustering and $J-K$ or $R-K$ color suggest that this is not the complete explanation.  For instance, among our $K$-selected sample, it is the \\emph{R-faint} galaxies \\citep[i.e.~those that could not make it into the sample of][]{adelberger05b} that cluster most strongly. Additionally, the very high correlation lengths measured by \\citet{adelberger05b}, $r_0 \\sim 10 h^{-1} \\textrm{Mpc}$ for galaxies with $K \\sim 20.5$, may not be representative because one of their four fields shows anomalously high clustering; the average of the other three fields is $r_0 = 5 \\pm 1 h^{-1} \\textrm{Mpc}$ (see their Fig.~2).  Applying the same $K < 20.5$, $R < 25.5$ selection criteria to the MUSYC sample, we find $r_0 = 4.5^{+2.7}_{-4.5} h^{-1}\\textrm{Mpc}$.  We conclude that the difference between massive optically-selected galaxies and massive red $K$-selected galaxies is probably not ``transient'', but is instead related to a fundamental difference in the host dark matter halos.  It is interesting to consider whether $r_0$ is primarily related to stellar mass, or some other property.  We use stellar population models to estimate the stellar masses, and do not find a significant relationship between mass and $r_0$.  This is interesting, as there is a strong relationship between color and mass; together, these results indicate that either low-mass red galaxies cluster very strongly or that massive blue galaxies cluster less strongly.  We find some evidence supporting both of these conclusions.  We discuss our findings in light of the recent results from optically-selected samples, finding that there is evidence that redder colors are associated with larger correlation lengths for both samples.  For optically-selected galaxies, the observed relationship between decreasing $R$-magnitude and increasing clustering strength for LBGs has been taken to suggest a positive relationship between halo mass and star formation rate at early times \\citep[e.g.][]{giavalisco01}. On the other hand, stellar population synthesis modelling has not uncovered a relationship between rest-frame UV luminosity and stellar mass for typical optically-selected galaxies \\citep[][but see \\citealt{papovich01} for a discussion of a fainter sample]{shapley01, shapley05}.  Similarly, \\citet{adelberger05b} show that the clustering of optically-selected BX galaxies is strongly related to $K$-brightness.  While there is a correlation between $K$ and stellar mass \\citep{shapley05}, there is also a relationship between $K$ and star formation rate \\citep{reddy05,erb06}.  So while firm conclusions cannot be drawn from this evidence, it appears that clustering may not be determined by stellar mass alone for optically-selected samples. The evidence presented here for $K$-selected galaxies also indicates that color, rather than stellar mass, may be the primary determinant of $r_0$.  We compared the observed clustering length of MUSYC galaxies to that of dark matter halos, assuming the simple case of one galaxy per halo above a halo mass threshold.  Galaxies with $2<z_{\\rm{phot}}<3.5$ and $K<21$ occupy halos with $M \\gtrsim 3 \\times 10^{12} M_\\odot$.  The number density of galaxies is similar to the number density of halos, indicating a mean halo occupation number of order $\\sim 1$.  Galaxies with redder colors reside in more massive halos, with DRGs residing in $M \\gtrsim 10^{13} M_\\odot$ halos.  However DRGs are found to be more numerous than these halos, suggesting a mean halo occupation number $N_{occup} \\approx 40^{+60}_{-30}$.  If such large numbers of galaxies occupy the same halos, then it is expected that the angular correlation function $w(\\theta)$ will show very large values on scales corresponding approximately to the halo virial radius.  We performed an independent estimate of the occupation numbers by comparing the small-scale values of $w(\\theta)$ to the values expected for dark matter halos; this yields occupation numbers $\\sim 1.5-2$ for all samples studied in this paper, regardless of color and stellar mass.  The cause of the discrepancy in the occupation number of red galaxies is unclear.  One possibility is that we have overestimated the correlation length of DRGs, although previous studies \\citep{daddi03,grazian06} have found correlation lengths even larger than the value presented here.  It is important to note that the high $r_0$ measured for DRGs is largely due to the integral constraint correction (\\S \\ref{sec:method}); neglecting this correction decreases the power law amplitude of $w(\\theta)$ by $\\sim 50\\%$, and decreases the correlation length from $r_0 = 11.1^{+1.3}_{-1.4} h^{-1} \\textrm{Mpc}$ to $r_0 = 7.7^{+1.6}_{-1.9} h^{-1} \\textrm{Mpc}$. Larger fields are necessary to reduce the effect of the integral constraint.  We also note that occupation numbers greater than unity indicate that more complex halo occupation distribution models--in which the number of galaxies per halo depends on the halo mass--are necessary in order to accurately quantify the relationship between galaxies and halos \\citep[e.g.][]{zheng04,lee06}.  Regardless of the cause of this discrepancy, the relationship between color and clustering in our sample has been established with $\\sim 96\\%$ significance.  Finally, we addressed the evolution of $z>2$ galaxies.  The descendants of our $K$-selected populations tend to occupy halos with masses $10^{13} - 10^{14} M_\\odot$, corresponding to the mass scales of groups.  The reddest samples, including the DRGs, may occupy cluster-scale halos, with masses $\\gtrsim 10^{14} M_\\odot$.  Even the descendants of the less clustered LBGs tend to reside in groups.  It appears that only a small subset of the $z>2$ galaxies that dominate current redshift surveys can be progenitors of typical $L_*$ field galaxies.  An important caveat is that each galaxy `population' will have rather heterogeneous properties, and it may be that discussing the evolutionary paths of population averages obscures important distinctions.  For instance, \\citet{adelberger05a} show that the correlation length of BX objects is $\\sim 4 h^{-1} \\textrm{Mpc}$, and argue from this that their descendants should be elliptical galaxies. \\citet{adelberger05b} show that the BX objects with $K \\lesssim 21$ contribute most strongly to the clustering measurement; $40\\%$ of the BX objects are at $K>21.5$, and have a correlation length $\\sim 2.5 h^{-1} \\textrm{Mpc}$.  Presumably these fainter galaxies will evolve into a much less clustered population by $z=0$.  A principle limitation of the preceding analysis is the heavy reliance on photometric redshifts.  Obtaining large numbers of spectroscopic redshifts for red $K$-selected galaxies has proven difficult on 6-10m telescopes.  While NIR spectroscopy yields a high success rate for determining redshift for bright galaxies \\citep{kriek06b}, the advent of multi-objects NIR spectroscopy will make the process more efficient.  Another limitation of our study is the small galaxy sample.  Recent clustering measurement of optically-selected galaxies are based on samples that are 1-2 orders of magnitude larger than the one presented here.  The next generation of NIR detectors will enable the imaging of significantly larger fields to comparable depth, allowing for more precise measurements of galaxy clustering."
    },
    "0606/astro-ph0606531_arXiv.txt": {
        "abstract": "We present predictions on the evolution of the Tully--Fisher (TF) relation with redshift, based on cosmological N-body/hydrodynamical simulations of disc galaxy formation and evolution. The simulations invoke star formation and stellar feedback, chemical evolution with non-instantaneous recycling, metallicity dependent radiative cooling and effects of a meta-galactic UV field, including simplified radiative transfer. At $z=0$, the simulated and empirical TF relations are offset by about 0.4~magnitudes (1~$\\sigma$) in the B and I bands. The origin of these offsets is somewhat unclear, but it may not necessarily be a problem of the simulations only.  As to evolution, we find a brightening of the TF relation between $z=0$ and $z=1$ of about 0.85~mag in rest--frame B band, with a non-evolving slope. The brightening we predict is intermediate between the (still quite discrepant) observational estimates.  This evolution is primarily a luminosity effect, while the stellar mass TF relation shows negligible evolution. The individual galaxies do gain stellar mass between $z=1$ and $z=0$, by a 50--100\\%; but they also correspondingly increase their characteristic circular speed. As a consequence, individually they mainly evolve {\\it along} the stellar mass TF relation, while the relation as such does not show any significant evolution. ",
        "introduction": "\\label{sect:introduction} The Tully--Fisher (TF) relation is a most important distance indicator, but also a key to understand the structure and evolution of disc galaxies. Its slope is a standard test--bench for cosmologically motivated galaxy formation models (Dalcanton, Spergel \\& Summers 1997; Mo, Mao \\& White 1998; Sommer--Larsen, G\\\"otz \\& Portinari 2003, henceforth SGP03) and its absolute luminosity level is also a crucial constraint, suggesting low stellar mass--to--light ratios and  ``bottom--light'' Initial Mass Functions in disc galaxies (Portinari, Sommer--Larsen \\& Tantalo 2004; henceforth PST04).  \\begin{table*} \\label{tab:literature} \\caption{Overview of literature results on the evolution of the TF relation; we indicate its brightening as $\\Delta M_B = M_B (V_c,0)-M_B (V_c,z)$. Then we report the redshift range and median redshift of the sample, the number of galaxies and the reference paper.} \\begin{tabular}{c c c l p{4truecm}} \\hline $\\Delta M_B$ & $z$ range & $<z>$ & $N_{gal}$ & reference \\\\ \\hline \\smallskip 1.5        & 0.15--0.35 & 0.25 & 24 & Rix \\etal (1997) \\\\ \\smallskip 1.5--2     & 0.25--0.45 & 0.35 & 12 & Simard \\& Pritchet (1998) \\\\ \\smallskip 0.4        &  0.1--1    & 0.5  & 16 & Vogt \\etal (1996, 1997) \\\\  0.2        &  0.2--1    & 0.5  & $\\sim$100 & Vogt (1999) \\\\ \\smallskip 1.1        &  0.5--1.5  & 0.9  & 22 & Barden \\etal (2003) \\\\ \\smallskip 1.6 $z$    & 0.15--0.9  & 0.4  & 19 (field) & Milvang--Jensen \\etal (2003) \\\\ $\\sim 0.8^a$  & 0.05--1 & 0.45 & $\\sim120$ & B\\\"ohm \\etal (2004), \\\\ \\smallskip (1.22 $z$)$^a$ &         &      &           & B\\\"ohm \\& Ziegler (2006)\\\\ \\smallskip 0.3 & \\multicolumn{3}{c}{revised analysis of B\\\"ohm \\etal's data} & Kannappan \\& Barton (2004) \\\\ \\smallskip 1.0 $z$    &  0.1--1    & 0.33 & 89 & Bamford \\etal (2006) \\\\ \\smallskip 1.3 $z$    & 0.19--0.74 & 0.39 & 14 & Nakamura \\etal (2006) \\\\ \\smallskip $\\gsim 0$   &  0.4--0.75 & 0.53 & 11 (non-disturbed) & Flores \\etal (2006) \\\\ \\hline \\end{tabular}  {\\footnotesize $^a$ For the sake homogeneous comparison to other data, we report in this table the B\\\"ohm \\etal results for constant TF slope, \\hfill  although the authors favour a slope evolution scenario to interpret their data (see text). \\hfill} \\end{table*}  In the last decade, the increasing amount of observational data at intermediate and high redshifts have allowed to study also the evolution of the TF relation, as well as of other galaxy scaling relations. However, there is as of yet no convergence on the results (see the summary in Table~\\ref{tab:literature}). Rix \\etal (1997) and Simard \\& Pritchet (1998) argued for strong luminosity evolution of disc galaxies (1.5--2~mag in rest--frame B band) already at $z=0.3$. Conversely, Vogt \\etal (1996, 1997) found at most moderate evolution ($\\Delta M_B \\lsim$0.4~mag) for a sample of about 16 spiral galaxies at a median redshift of $z \\sim 0.5$. Later on, basing on a much larger sample of 100 galaxies between $0.2<z<1$, Vogt (1999) further reduced the estimated evolution to  $\\Delta M_B \\lsim$0.2~mag, almost entirely driven by strong brightening of a few dwarf galaxies at $z>0.5$. Most recently, from the analysis of the DEEP1 Groth Strip survey, Vogt (2006) reports no evolution at all (within $\\pm$0.3~mag) in the TF relation out to $z \\sim$1, in the sense that all of the apparent evolution can be entirely accounted for by observational and selection effects. Other groups instead find, out to $z \\sim$1, a brightening up to 1~mag or more (Rigopoulou \\etal 2002; Barden \\etal 2003; Milvang--Jensen \\etal 2003; B\\\"ohm \\etal 2004; Bamford, Arag\\'on--Salamanca \\& Milvang--Jensen 2006; Nakamura \\etal 2006). Significant brightening is also indicated by the line width -- luminosity relation at even higher redshifts: van Starkenburg \\etal (2006) report a 2~mag offset for 3 galaxies at $z \\sim$1.5. However, at so early epochs, one major concern is how to interpret the kinematics of galaxies, which may not have settled onto a well developed rotation curve comparable to those of local spirals.  Limiting our discussion to galaxies up to $z \\sim$1, the discrepancy between different results shows that these are plagued by observational difficulties and selection effects. Smith \\etal (2004) observed a disc galaxy at $z \\sim$0.8 in integral field spectroscopy, deducing for this one object a far more negligible brightening with respect to the local TF relation, than found by Barden \\etal with long--slit spectroscopy; this highlights the difficulties of measuring the kinematics of high redshift galaxies. Kannappan \\& Barton (2004) further suggest that a significant contribution to the apparent evolution may be due to disturbed kinematics biasing the measurements of rotation velocities in high redshift samples, where kinematic anomalies are both more frequent (in the hierarchical clustering scenario) and more difficult to assess; kinematic disturbances might also contribute to the discrepancies between results obtained from selecting large undisturbed discs (e.g.\\ the Vogt paper series) and from other less selective samples. Indeed recently, Flores \\etal (2006) have used 3D integral field spectroscopy to recover the velocity field of disc galaxies at $z=0.4-0.75$ and showed that the intermediate redshift TF relation of undisturbed rotating discs shows much reduced scatter and evolution than a generic sample including perturbed or complex kinematics.  It seems therefore that the results on the TF evolution are still dominated by observational difficulties, selection effects and systematics. A recalibration of the zero--point of the local comparison TF relation by about 0.5~mag (Tully \\& Pierce 2000 vs.\\ Pierce \\& Tully 1992; see Fig.~\\ref{fig:TFz0}a) adds further confusion to the issue.  As to the slope of the TF relation, within the large scatter of the high redshift samples the slope is not very well constrained directly so the ``minimal assumption'' is usually adopted of an evolution at constant slope, i.e.\\ of a uniform luminosity offset. However, there are some indications that evolution might be mass dependent, and that sub--$L_*$ objects drive the brightening (Simard \\& Prichet 1998; Vogt 1999); differential mass evolution might partly explain the discrepancies between the results from different samples. In their recent VLT sample B\\\"ohm \\etal (2004) and B\\\"ohm \\& Ziegler (2006) find a brightening of about 0.8~mag if a constant TF slope is imposed, but their data are better fit by a changing slope implying much stronger brightening for dwarf galaxies. Kannappan \\& Barton (2004) suggest though that this differential behaviour might be driven by kinematic anomalies, namely that the objects with stronger evolution are likely kinematically disturbed rather than genuinely brighter; comparing the B\\\"ohm \\etal data to a low redshift sample with similar selection as the high redshift one, they revise the estimated TF evolution down to a modest $\\Delta M_B$=0.3~mag, with no change in slope. Bamford \\etal (2006) further point out that a spurious change in slope may be due to intrinsic coupling between scatter in rotation speed, TF residuals and limited magnitude range of the sample; also Weiner \\etal (2006) find no slope evolution in a revised analysis of the B\\\"ohm \\etal (2004) data. Finally, very recently Weiner \\etal (2006), from a large sample of TKRS/GOODS galaxies with $z$=0.4--1.2, argue for an evolution in TF slope {\\it in the opposite sense}, with brighter, massive galaxies fading more than low mass galaxies. The issue of slope evolution remains quite controversial.  However quantitatively significant the luminosity evolution of the TF relation, there is some evidence that it must be largely a luminosity dimming effect of the aging stellar populations, while in terms of the stellar mass there is hardly any offset, at fixed $V_c$, between the TF relation at $z=0$ and at $z=1$ (Conselice \\etal 2005; Flores \\etal 2006).  In this paper we present our predictions on the evolution of the TF relation with redshift, based on cosmological+hydrodynamical simulations of disc galaxy formation. In Section~\\ref{sect:simulations} we describe our simulations and their analysis. In Section~\\ref{sect:TFz0} we discuss the theoretical TF relation at $z=0$ and the effects of numerical resolution and softening length on the results. In Section~\\ref{sect:TFevolution} we illustrate the evolution of the TF relation as predicted by our simulations. A special experiment with no infall of hot halo gas after $z=1$ is discussed in Section~\\ref{sect:noinfall}. Finally, Section~\\ref{sect:conclusions} outlines our conclusions. ",
        "conclusions": "\\label{sect:conclusions} In this paper we discuss the TF relation and its evolution as predicted by cosmological simulations of disc galaxy formation including hydrodynamics (Tree-SPH) and all the relevant baryonic physics, such as star formation, chemical and photometric evolution, metal--dependent cooling, energy feedback.  The resulting disc galaxies at $z=0$ have scalelengths compatible with the observed ones in the low--mass range ($V_c \\sim$100--120~km/sec), and somewhat shorter, yet within a factor of 2, for Milky Way--type or more massive objects. Though still somewhat plagued by the angular momentum problem, our physical recipe and implementation of strong early feedback goes a long way in creating realistic galaxies (SGP03).  At $z=0$, the TF relation defined by our simulated galaxies is offset by about 0.4~mag in luminosity, and by about 40\\% in stellar (and baryonic) mass, with respect to observational TF relations. This is a well known problem of galaxy formation simulations (Navarro \\& Steinmetz 2000) and the origin of the offset remains somewhat unclear. We demonstrate that it is likely not due to the dark matter cusp in the central regions, and that it cannot be cured by simply avoiding adiabatic contraction. A different choice of cosmological parameters (such as $\\sigma_8$) and/or of stellar IMF in disc galaxies may concur to solve the problem; some mechanism for adiabatic expansion of the halo from the baryons has also been recently invoked in this respect (Dutton \\etal 2006).  However, we remark that the location of the Milky Way with respect to the TF relation is also offset (Flynn \\etal 2006), by as much as our simulations. The fact that our Galaxy lies on the same TF locus as our simulations adopting a Solar Neighbourhood--like IMF (Fig.~\\ref{fig:TFz0}), suggests the puzzling possibility that the problem may not lie with the simulations or the cosmology, but rather with the luminosity zero--point of the observed TF relation of external galaxies, or with the IMF of the Milky Way which may not be representative for all disc galaxies.  As to the luminosity evolution of the TF relation with redshift, there is presently no consensus observationally, with estimates ranging from negligible evolution to $\\Delta M_B >$1~mag by $z \\sim$1. Also, there are claims that the evolution may be mostly in the slope (with a stronger brightening for lower mass objects) rather than a systematic offset, but this also is still debated. There is however some preliminary consensus that the stellar mass TF relation hardly evolves out to $z \\sim$1.  Our simulations predict a significant B--band luminosity evolution ($\\Delta M_B$=0.7~mag by $z$=0.7 and $\\Delta M_B$=0.85~mag by $z$=1), while the TF slope remains nearly constant (systematic zero--point evolution). The redshift evolution of the TF relation is a result of the correlated increase in stellar mass and $V_c$ in the individual galaxies, combined with luminosity fading and reddening caused by the aging of the host stellar populations. At fixed $V_c$, the net effect is a systematic decrease in B band luminosity, while in terms of stellar mass the TF relation hardly evolves at all, in agreement with the present available observational evidence.  This does not mean that the stellar mass content of the individual galaxies does not evolve: each galaxy grows typically by 50--100\\% in mass between $z=1$ and $z=0$, but as its baryonic mass grows, its circular velocity also increases so that the evolution of the individual object occurs largely {\\it along} the TF relation. The relation as such does not show any significant evolution.  We also consider as a special experiment a simulation where late gas infall on the disc is precluded, by artificially stripping the galaxy of its hot, dilute halo gas at $z=1$. This is useful to assess alternative scenarios of disc galaxy formation, where late disc growth is driven by other processes (like satellite or cold gas accretion), as well as a possible scenario of spiral evolution in clusters, where ram pressure stripping from interaction with the ICM is a possible driver of morphological transformations. We find that, from the point of view of the TF properties, the resulting object is indistinguishable from early type spirals with little recent star formation and red colours --- henceforth the TF relation for cluster galaxies at $z=0$ should not be markedly different from the field one, once colour--corrections are applied. Its TF evolution is however more significant than that of the ``normal'' simulations, suggesting that, if halo gas infall had no major role in disc galaxy evolution at late times, the TF relation would show a much stronger luminosity evolution. Unfortunately the present observational uncertainty does not allow to discuss if either scenario is really favoured. This experiment, artificially creating an early--type galaxy, also underlines that the amount of TF evolution likely depends on Hubble spiral type, being more significant for early types; this effect might contribute to the discrepancy between different observational results, and should be taken into account when selecting the local reference TF sample."
    },
    "0606/astro-ph0606707_arXiv.txt": {
        "abstract": "We calculate the broadband radio--X-ray spectra predicted by microblazar and microquasar models for Ultra-Luminous X-ray sources (ULXs), exploring the possibility that their dominant power-law component is produced by a relativistic jet, even at near-Eddington mass accretion rates. We do this by first constructing a generalized disk--jet theoretical framework in which some fraction of the total accretion power $\\Pa$ is efficiently removed from the accretion disk by a magnetic torque responsible for jet formation. Thus, for different black hole masses, mass accretion rates and magnetic coupling strength, we self-consistently calculate the relative importance of the modified disk spectrum, as well as the overall jet emission due to synchrotron and Compton processes. In general, transferring accretion power to a jet makes the disk fainter and cooler than a standard disk at the same mass accretion rate; this may explain why the soft spectral component appears less prominent than the dominant power-law component in  most bright ULXs. We show that the apparent X-ray luminosity and spectrum predicted by the microquasar model are consistent with the observed properties of most ULXs. We predict that the radio synchrotron jet emission is too faint to be detected at the typical threshold of radio surveys to date. This is consistent with the high rate of non-detections over detections in radio counterpart searches. Conversely, we conclude that the observed radio emission found associated with a few ULXs cannot be due to beamed synchrotron emission from a relativistic jet. ",
        "introduction": "Ultra-Luminous X-ray sources (ULXs) are among the most intriguing sources to have been discovered by X-ray satellites. They are defined as point-like X-ray sources which, if located within the galaxy they are associated with on the sky, are off-nuclear and have X-ray luminosities $\\Lx \\gtapprox \\, 2 \\times 10^{39}\\, {\\rm erg \\, s}^{-1}$ in the $0.5 - 8.0\\,$keV bandpass \\citep*{Irwin04}. This exceeds the theoretical Eddington limit for a neutron star or stellar-mass black hole, $L_{\\rm Edd} = 4\\pi G\\Mbh \\mu m c/\\sigmaT$, where $\\Mbh$ is the black hole mass and $\\mu$ is the mean molecular weight. It also exceeds the peak in observed X-ray luminosities of known Galactic X-ray binaries (XRBs). Spectroscopic studies have revealed that some ULXs can be identified with supernova remnants or background AGN \\citep[see e.g.][and references therein]{GutLopez05}. Of the remaining unidentified sources, however, rapid X-ray variability \\citep*[e.g.][]{StroMush03,SoriaMotch04,Krauss05} and spectral state transitions \\citep*{Kubota01,Kubota02,LaParola01,MizKubMak01,Strickland01} strongly suggest that we are dealing with a population of close interacting binary systems involving accretion onto a black hole.   Perhaps the most contentious issue surrounding ULXs is the mass of the black hole, $\\Mbh$, which determines $L_{\\rm Edd}$. There are at least three possible explanations for why many ULXs are apparently $\\sim 20-30$ times  more luminous than typical Galactic black hole XRBs: 1. The emission is Doppler-boosted by a relativistic jet pointing towards us (microblazar scenario: e.g. \\citealt*{FabMesch01,Kording02}). In this case, $\\Mbh$ need not be more than $\\sim \\, \\mbox{a few} \\, M_\\odot$, similar to typical stellar-mass black holes in the Galaxy. 2. The brightness enhancement is due to a combination of various factors: mild beaming, $\\Mbh$ more massive than typical Galactic sources, and super-Eddington accretion. A factor of $\\sim 3$ for each of these terms is physically plausible and would suffice for most ULXs \\citep[e.g.][]{KingDehnen05}. In particular, mild geometric beaming could result from a sub-relativistic, radiatively-driven disk outflow (\\citealt{King01,King04}). Alternatively, it could be the result of a relativistic jet pointing slightly away from us (microquasar scenario). 3. The X-ray emission is isotropic and Eddington-limited; the high luminosity arises from a more massive black hole (an intermediate-mass black hole, IMBH), with $\\Mbh \\sim 10^{2-3} M_\\odot$  \\citep*{ColMush99,Makishima00,Miller03}. In this case the X-ray source is truly ultra-luminous, and hence it requires an accretion rate higher than that for a beamed scenario. If ULXs were proven to be powered by an accreting IMBH, the implications would be far-reaching, impacting on many other fields, from star-formation to cosmology (see \\citealt{Miller04} for a review).  In principle, phase-resolved optical spectroscopic and photometric studies are the most direct way to determine the binary system parameters (viewing angle, orbital period, radial velocity curve and hence, mass function -- see e.g. \\citealt{Charles98}) and constrain the nature of ULXs. This is how Galactic black holes were first identified. However, optical counterpart searches and phase-resolved studies have proven exceedingly difficult for ULXs because even the nearest ones are located at distances of a few Mpc. Moreover, the brightest ULXs tend to be preferentially located in crowded star forming regions  \\citep{Irwin04}, making unambiguous optical identifications often impossible, even within the \\textit{Chandra} error circle. When optical counterparts are found, they are generally consistent with O or B0 stars \\citep{Liu02,Liu04,Kaaret04,Zampieri04}. However, this could be a selection effect, since smaller donors (perhaps lower-mass stars evolving through their giant phase) would be too faint to be detected. Hence, it is important to identify the discriminating signatures of ULXs in other wavebands.   Radio observations may provide a more effective means of differentiating the isotropic IMBH and beaming models. By analogy with AGN, we can expect a ULX microblazar to produce strong synchrotron radio emission from a compact, unresolved core, while more extended, and possibly elongated and resolved radio structure may arise from a ULX microquasar, seen at larger inclination. To date, radio counterparts have been detected for only a few ULXs: two in M82 \\citep*{KronSram85,KordColFalck05}; one in NGC\\,5408 \\citep*{Kaaret03,Soria06a}; one in Holmberg\\,II \\citep*{TonWest95,MillMushNeff05}; one in NGC\\,7424 \\citep*{Soria06b}; and one in NGC\\,6946 (van Dyk et al. 1994; Blair, Fesen \\& Schlegel, 2001; Swartz et al. 2006, in prep.). It is still unclear whether the emission in those radio sources is due to an underlying supernova remnant or is directly produced by ULX jet activity, either from the core or the lobes. Hence, it is important to determine the radio signatures of different ULX models.  In the X-ray band, ULX spectra are generally dominated by a power-law, particularly above 1~keV. A cool thermal component is present in some sources, but contributes at most $10 - 20\\,$\\% of the flux \\citep*[e.g.][]{Stobbart06}. A power-law component is also dominant in the X-ray spectra of Galactic black holes in two of the canonical spectral states \\citep{McClinRem06}: the low/hard state, when the photon index is $\\Gamma \\sim 1.5 - 2$ and the X-ray luminosity is $\\Lx \\ltapprox 10^{-2} L_{\\rm Edd}$; and the very high (or steep power-law) state, when $\\Gamma \\gtapprox 2.5$ and $\\Lx \\sim L_{\\rm Edd}$. In the low/hard state, the accretion disk is thought to be truncated far from the innermost stable circular orbit (ISCO), and replaced by a radiatively inefficient flow with a steady jet \\citep[see e.g.][]{Esin97}. The very high state, on the other hand, is characterized by a high radiative efficiency, quasi-periodic X-ray variability, and absence of a steady jet, although flares or sporadic radio ejections are associated with this state. It is not clear whether ULXs can be classified into either one of these power-law dominated XRB states; their photon index is somewhat  intermediate, with $\\Gamma \\approx 1.5 - 2.5$ (e.g. \\citealt{Swartz04,Stobbart06,Winter06}). If they are in a classical low/hard state, their apparent X-ray luminosities suggest masses $\\gtapprox \\mbox{a few} \\times 10^3 M_\\odot$. What seems clear is that ULXs are not in a high/soft state, when the X-ray spectrum is dominated by soft disk blackbody emission and the power-law is weaker.  Regardless of which spectral state ULXs belong to, the power-law component could be produced either by thermal Comptonization in a hot, diffuse corona, or by nonthermal synchrotron and Compton processes in a relativsitic jet, or possibly a combination of both \\citep*[see][for a discussion]{Markoff03}. Although it is widely accepted that relativistic jets are responsible for extragalactic radio sources, analogous jet models for Galactic XRBs have only been considered relatively recently \\citep*[see][for a review]{Fender05}. Existing disk-jet models for binaries, however, either neglect the radio synchrotron properties altogether \\citep*[e.g.][]{GeorgAharKirk02} or are only appropriate for the low/hard state \\citep*[e.g.][]{MarkFalckFend01,Markoff03,FendBellGall04}. It is important to explore the possibility of a jet coexisting with (and coupled to) a bright disk extending all the way down to the ISCO\\footnote{Recent X-ray/radio studies of FRII radio galaxies and quasars (e.g. \\citealt{PunslyTingay05,BallFab05}) provide examples of powerful nuclear jets co-existing with a high luminosity spectral state ($\\LX \\sim L_{\\rm Edd}$).}.   In this paper, we investigate both the radio and X-ray properties of relativistic beaming models for ULXs using an accreting black hole framework generalized to include coupled disk-jet spectral components. Specifically, the radio and hard X-ray emission are produced in a relativistic jet that is magnetically coupled to an accretion disk (which can contribute to the soft X-ray emission). The jet drains a substantial fraction of the total accretion power from the disk and thus modifies the multi-colour disk spectrum. In order to calculate the spectral energy distribution, we explicitly model the mechanism which magnetically couples the disk and jet. Thus, our generalized model calculates both the jet and disk spectra \\textit{self-consistently} for different values of the black hole mass $\\Mbh$, mass accretion rate $\\Mdota$, and coupling strength. This leads to two unique features of the model that have important implications not just for ULXs, but for all accreting systems: 1. The nett accretion power can exceed the Eddington limit because the jet efficiently channels energy from the disk, which remains sub-Eddington; and 2. The presence of a jet can be inferred indirectly from its effects on the disk, which consequently emits a spectrum that can be considerably different from that predicted by the standard Shakura-Sunyaev \\citep{SS73} disk in the EUV/soft-X-ray bandpass for XRBs. Our coupled disk-jet theory is outlined in Sec.~\\ref{s:theory}; results for beamed ULX models are presented in Sec.~\\ref{s:results}; and a discussion and conclusions are given in Secs.~\\ref{s:discussion} and \\ref{s:conclusions}, respectively. ",
        "conclusions": "\\label{s:conclusions}  Current observational evidence suggests that most of the X-ray emission from ULXs does not come directly from the accretion disk, which is either not directly visible, or colder than in bright stellar-mass XRBs. We have used a coupled disk+jet theoretical framework to explain why most of the energy is not directly radiated from the disk, and to test whether a relativistic beaming scenario (microblazar or microquasar) is consistent with the observed X-ray and radio spectra. Our main conclusions can be summarized as follows: \\begin{list}{}{\\itemsep=0pt} \\item[1.] The accretion disk can be substantially modified by the presence of a magnetized jet; this gives rise not only to a modified disk spectrum, but also to a decrease in the peak colour temperature ($T_{\\rm in}$) of the inner disk for a fixed mass accretion rate. The modified disk+jet model thus predicts an anticorrelation between the relative importance of hard (jet) and soft (disk) X-ray spectral components. \\item[2.] We confirm that both the stellar-mass microblazar and microquasar scenarios can produce apparent X-ray luminosities in the ULX regime ($\\Lx \\gtapprox 2 \\times 10^{39} \\, {\\rm erg \\, s^{-1}}$). A general feature of our disk+jet model is that it can produce hard, nonthermal X-ray spectra that are analogous to the classic low/hard spectral state in XRBs, but applicable to high mass accretion rates ($\\Mdota \\sim \\dot M_{\\rm Edd}$) and, therefore, high luminosity sources. In our model, a magnetic torque provides the mechanism for transferring a substantial fraction of the total accretion power $\\Pa \\sim L_{\\rm Edd}$ from the disk to the jet (i.e. $\\Pj \\sim L_{\\rm Edd}$). Our spectral model for the microblazar predicts a substantially hardened X-ray spectrum due to strong beaming and Comptonization effects. The microquasar model, on the other hand, produces a featureless power-law spectrum. \\item[3.] Both the microblazar and microquasar models predict that, despite the beaming effects, ULXs are intrinsically weak radio emitters. This result is compatible with the overwhelming excess of non-detections over detections in ULX radio observations to date. Neither of the beaming models are consistent with the few ULX radio counterpart detections reported to date. We therefore conclude that the observed radio emission cannot be attributed to beamed synchrotron emission in a relativistic jet. We suggest it is likely due to radio lobes or a radio SNR. We predict that deeper observations around 20\\,GHz should in principle detect a change in radio spectral index, thus resolving the direct jet contribution from the lobe emission. \\end{list}  In this paper, we have restricted ourselves to modelling beamed emission from accreting black holes with masses in the known stellar-mass range ($5-20 \\, M_\\odot$). If IMBHs exist, it is legitimate to assume (based on the stellar-mass and supermassive black hole analogies) that  some IMBHs can also possess relativistic jets. If some IMBHs have beamed emission, then we would expect to find at least a few milliblazars and milliquasars with apparent X-ray luminosities in excess of $10^{41} \\, {\\rm erg \\, s^{-1}}$. Such sources have not yet been seen.   The disk+jet model we have proposed here requires that more gravitational energy is channelled into a jet than is radiated by a disk, at least during some phases of accretion. At low or moderate accretion rates, new evidence of significant feedback onto the ISM surrounding accreting black holes indicates that jets can indeed carry away much more energy than is radiated by the disk \\citep[e.g.][]{Owen00,Churazov02}. At accretion rates approaching the Eddington rate, however, observations suggest that jets are quenched in Galactic XRBs \\citep[e.g.][]{FendBellGall04} and in low-mass AGN \\citep[e.g.][]{GreeHoUlv06}, but can persist in AGN with high black hole masses \\citep[e.g.][]{Laor00,McLureJarvis04}. Our model predicts that the fractional accretion power chanelled into jets scales with black hole mass and is independent of the mass accretion rate. It therefore predicts that jets can be the dominant energy carrier, even at high ($\\sim$ Eddington) accretion rates. If this is correct, it should be possible to find evidence of relativistic jet emission also in some other accreting black hole sources in which the hard X-ray component has perhaps been misinterpreted as, for example, thermal Comptonization in a corona."
    },
    "0606/astro-ph0606477_arXiv.txt": {
        "abstract": "We present a theoretical study of void spins and their correlation properties. The concept of the spin angular momentum for an unbound void is introduced to quantify the effect of the tidal field on the distribution of matter that make up the void. Both the analytical and numerical approaches are used for our study. Analytically, we adopt the linear tidal torque model to evaluate the void spin-spin and spin-density correlations, assuming that a void forms in the initial region where the inertia momentum and the tidal shear tensors are maximally uncorrelated with each other. Numerically, we use the Millennium run galaxy catalog to find voids and calculate their spin statistics. The numerical results turn out to be in excellent agreement with the analytic predictions, both of which consistently show that there are strong spatial alignments between the spin axes of neighbor voids and strong anti-alignments between the void spin axes and the directions to the nearest voids. We expect that our work will provide a deeper insight into the origin and properties of voids and the large scale structure. ",
        "introduction": "Recent galaxy redshift surveys \\citep[e.g., 2dFGRs,][]{col-etal01} have allowed us to map the universe on the largest scale ever and study systematically the detailed properties of the cosmic large-scale structures. It is confirmed by these surveys that the galaxies in the universe are not distributed evenly but develop pattern of network connected by filaments and sheets which in turn are separated by immense voids.  Voids are defined observationally as large regions with very low galaxy number density and shown to occupy approximately $40\\%$ of the cosmic volume \\citep{hoy-vog04}. The striking vastness of the observed void volumes inspired a flurry of research on the origin and properties of cosmic voids \\citep{ela-pir97,sch-etal01,pee01,mat-whi03,ben-etal03,she-van04,hoy-vog04, roj-etal04,col-etal05,hoy-etal05,got-etal05,sha-etal06}. It is now generally believed that voids are originated from the local minima of the primordial density fluctuations and expand faster than the rest of the universe.  The very fact that voids are extremely underdense and expand faster has led many authors to assume somewhat naively that the shapes of voids should be spherical. In fact previous analytical and numerical works on voids are based on this assumption \\citep[e.g.,][]{dub-etal93,van-van93,fri-pir01,she-van04}.  Very recently, however, \\citet{sha-etal06} investigated systematically the shapes of voids by applying the excursion set approach to the density field in N-body simulations and demonstrated that the intrinsic shapes of voids are far from spherical symmetry. They claimed that voids are most vulnerable to external shears due to their low-density, so that their shapes are disturbed to be nonspherical in spite of their faster expansion.  Although \\citet{sha-etal06} focused on the shapes and morphology of voids, we note here that if the tidal shear plays a significant role in the case of voids, then it should cause not only the shapes of voids to be nonspherical but also voids to acquire the spin angular momentum. Our goal here is to construct an analytic model for the spin angular momentum of voids originated from the tidal shear effect and test the analytic predictions against the data from the recent Millennium run simulation.  The organization of this paper is as follows. In $\\S 2$, the analytic model for the void spin angular momentum and the void spin correlation properties are presented. In $\\S 3$, the Millennium run galaxy catalog is analyzed to determine void spin statistics, and the numerical results are compared with the analytic predictions. In $\\S 4$, the achievements and caveats of our works are discussed, and a final conclusion is drawn. ",
        "conclusions": "In the light of the work of \\citet{sha-etal06} who claimed that the tidal shear must play a strong role in case of voids, we have proposed a scenario that voids are originated from the regions in the initial density field whose inertia tensors are completely uncorrelated with the local tidal tensors. Since the alignment between the principal axes of the inertia and tidal tensors plays a role of reducing the tidal shear effect according to the linear tidal torque theory, the hypothesis implies that voids should have very high spin angular momentum. We define the concept of the spin angular momentum for an unbound void to quantify the effect of the tidal field on the distribution of matter that make up voids.  Our analytic model based on the linear tidal torque theory have predicted that there is a strong spatial alignment between the spin axes of neighbor voids and a strong anti-alignment between the spin axes of voids and the directions to the nearest voids.  Our predictions have been tested against the Millennium run galaxy catalog where $13507$ large voids with more than $30$ galaxies are identified. We have measured numerically the distribution and correlation functions of void spins, and found that the predictions of our analytic model are in excellent agreement with the numerical results.  The success of our work is, however, subject to a couple of caveats. The first caveat comes from our implicit assumption that the matter voids are same as the galaxy voids. In our analytic model, we consider the matter voids. While in the numerical analysis we deal with the galaxy voids. Although the two types of voids are not necessarily same, they should be at least closely correlated with each other.  The second caveat lies in the fact that there is no standard way to identify voids. Unlike the case of halo-finding, the use of different algorithms could produce different results for some properties of voids. Especially, the shapes and volumes of voids can differ significantly by the underlying criterion of the void-finding algorithm. If one use an algorithm which regards a void as a spherical region of low density, then all voids found by that algorithm would be spherical. In spite of this ambiguity in finding voids,  we expect that our result on the correlation properties of the void spin angular momentum would not change significantly by the different choice of the void finding algorithm, since the correlation properties are basically determined not by the shapes of voids but by the galaxies in voids.  Another issue we would like to discuss here is a possibility of testing numerically our key assumption that a void forms in the initial underdense region where the misalignment between the principal axes of the inertia and the tidal tensors becomes maximum. Although \\citet{lee06} has already shown analytically that our assumption is true,  it would be highly desirable to test it against numerical data from the Millennium run simulation. For this, it would be necessary to find the Lagrangian positions of all the particles that make up each void and to determine the inertia momentum tensor of each Lagrangian proto-void region in the principal axis frame of the initial tidal field. As the initial density field of the Millennium run simulation is not available to the public yet, the numerical test of our assumption is postponed. We hope to report the test result elsewhere in the future.  Finally, we conclude that our work on void spins will provide a deeper insight into the origin and properties of the large scale structure in the universe."
    },
    "0606/astro-ph0606194_arXiv.txt": {
        "abstract": "We report results from a systematic study of the spectral energy distribution (SED) and spectral evolution of XTE~J1550--564 and H~1743--322 in outburst. The jets of both sources have been directly imaged at both radio and X-ray frequencies, which makes it possible to constrain the spectrum of the radiating electrons in the jets. We modelled the observed SEDs of the jet `blobs' with synchrotron emission alone and with synchrotron emission plus inverse Compton scattering. The results favor a pure synchrotron origin of the observed jet emission. Moreover, we found evidence that the shape of the electron spectral distribution is similar for all jet `blobs' seen. Assuming that this is the case for the jet as a whole, we then applied the synchrotron model to the radio spectrum of the total emission and extrapolated the results to higher frequencies. In spite of significant degeneracy in the fits, it seems clear that, while the synchrotron radiation from the jets can account for nearly 100\\% of the measured radio fluxes, it contributes little to the observed X-ray emission, when the source is relatively bright. In this case, the X-ray emission is most likely dominated by emission from the accretion flows. When the source becomes fainter, however, the jet emission becomes more important, even dominant, at X-ray energies. We also examined the spectral properties of the sources during outbursts and the correlation between the observed radio and X-ray variabilities. The implication of the results is discussed. ",
        "introduction": "There is growing evidence that the central engine of microquasars is qualitatively similar to that of active galactic nuclei (AGN) or of gamma-ray bursts (GRBs). Collectively, these systems provide an excellent laboratory for studying particle acceleration in the jets of black holes over a vast range of physical scales (Mirabel 2004; Cui 2006). As has been well demonstrated in the study of AGN and GRBs, modelling the broadband spectral energy distribution (SED) of microquasars can be a very effective way to cast light on emission mechanisms and thus on the nature of the central engine. Unlike the cases of AGN and GRBs, however, the availability of such data is highly limited for microquasars, due to the combination of the transient nature of most such sources and the rarity of their outbursting activities (during which they can usually be seen). In spite of the difficulty, much has been learned from the studies that have been carried out (e.g., Markoff, Falcke \\& Fender 2001; Hynes et al. 2002; Ueda et al. 2002; Chaty et al. 2003; Fuchs et al. 2003; Markoff \\& Nowak 2004; Markoff, Nowak \\& Wilms 2005; Yuan, Cui \\& Narayan 2005).  The broadband SED of microquasars is usually composed of several distinct components. Very roughly, at radio frequencies, the spectrum can be described by a power law, which may have either positive or negative frequency dependence. The emission is generally thought to be synchrotron radiation from relativistic electrons in the jets, which may be optically thin or thick to synchrotron self-absorption (Hjellming \\& Han 1995; Fender 2006). Such a synchrotron spectrum could extend up to infrared/optical frequencies in some cases (e.g., Chaty et al. 2003; Migliari et al. 2007). From optical to soft X-ray frequencies, the spectrum takes on a blackbody-like shape, which is usually viewed as the signature of an optically thick, geometrically thin accretion disc (review by Liang 1998). It is seen to peak at different frequencies for different sources or even for a given source at different fluxes; the latter has been interpreted as being related to the movement of the inner edge of the disc (e.g., Esin, McClintock \\& Narayan 1997). From hard X-ray to soft gamma-ray frequencies, the spectrum can, once again, be roughly described by a power law, which rolls over at some characteristic frequency under certain circumstances. It is generally modelled as Compton upscattering of soft photons by energetic electrons in an optically-thin configuration (Liang 1998).  The physical origin of Comptonizing electrons is still not entirely clear. They may be thermal electrons associated with advection dominated accretion flows (ADAF; see, e.g., Narayan, Mahadevan \\& Quataert 1998) or magnetic flares (e.g., Poutanen \\& Fabian 1999) above the thin disc. Over the years, intense efforts have been made to fit the X-ray spectrum of microquasars in the so-called low-hard state with models that include optically-thick emission from the thin disc and a Comptonized emission from some `corona'. The most successful and physically self-consistent models all seem to prefer a geometry that consists of an inner (roughly) spherical corona plus an outer thin disc (e.g., Dove, Wilms \\& Begelman 1997; Esin et al. 1997). Alternatively, the Comptonizing electrons may be non-thermal in nature (e.g., Zdziarski et al. 2001; Titarchuk \\& Shrader 2002). This is thought to be particularly relevant to the so-called high-soft state of microquasars (Grove et al. 1998). A promising place to accelerate electrons to the required energies is the jets, as in the case of AGN and GRBs, although the electrons could also be associated with the accretion flows. The broadband SED has proven to be a valuable tool for assessing the roles of jets and accretion flows in microquasars, as well as the coupling between the two (Yuan et al. 2005).  In this work, we examined two microquasars, XTE J1550--564 and H 1743--322, whose jets have been directly imaged at both radio and X-ray frequencies (Corbel et al. 2002, 2005), mainly to gain insights into the roles of the jets and accretion flows in microquasars. The broadband SED of the jets allowed us to assess the contribution of jets to the overall SED in a relatively model-independent manner. Moreover, modelling the jet SED helped constrain the properties of radiating electrons in the jets and thus reduced degeneracy in the modelling of the overall SED of the source. We also studied the spectral evolution of the sources during outburst and the correlation between the observed radio and X-ray variabilities. ",
        "conclusions": "We have carried out a systematic study of the SED of XTE~J1550--564 and H 1743--322, and their spectral evolution during a number of major outbursts. The main conclusions of the work are summarized as follows:  \\begin{enumerate} \\item The results from physical modelling of the broadband SED of the resolved components of the jets support a pure synchrotron origin of the observed emission, from radio to X-ray frequencies, in XTE J1550--564 and H 1743--322. The effects of inverse Compton scattering in the jets were examined and found to be negligible. This is at variance with some of the published works (e.g., Markoff \\& Nowak 2004; Giannios 2005; Markoff, Nowak \\& Wilms 2005), which argue for a dominant role of inverse Compton scattering in the production of X-rays from microquasars based on modeling the unresolved (jets plus core) emission of the sources.  \\item We found that the synchrotron radiation from the jets can account for 100\\% of the observed radio emission but seems to contribute little to the observed X-ray emission from XTE J1550--564 and H 1743--322, when the sources are relatively bright. We think that the bulk of the X-ray emission comes from the accretion flows, which is always thought to be the case. We have found observational evidence to show that the jet contribution at X-ray energies increases as the sources become fainter, and might eventually dominate in or close to the quiescent state.  \\item We found it straightforward to define the spectral states based on the shape of SEDs. The presence of a dominant soft (disc) component distinguishes the HSS from transitional states. In the context of the adopted definitions of the spectral states, we presented clear evidence for spectral hysteresis associated with the LHS-to-HSS and HSS-to-LHS transitions associated with the rising and decaying phases of an outburst. We also showed additional evidence to support a previous suggestion that the mass accretion rate alone cannot uniquely determine spectral states.  \\item There is a general positive correlation between the X-ray and radio fluxes in XTE J1550--564 even during the transitional states. We found evidence for a possible frequency dependence of the correlation.  \\end{enumerate}"
    },
    "0606/astro-ph0606641_arXiv.txt": {
        "abstract": "The {\\it Gamma-Ray Large Area Space Telescope (GLAST)}, to be launched in the fall of 2007, will measure the spectra of distant extragalactic sources of high energy \\grays, particularly active galactic nuclei and \\gray\\  bursts.   {\\it GLAST}  can  look  for energy-dependent  \\gray\\ propagation  effects  from  such   sources  as  a  signal  of  Lorentz invariance  violation.  These sources  should  also  exhibit the  high energy   cutoffs  predicted   to  be   the  result   of  intergalactic annihilation interactions with low  energy photons having a flux level as   determined   by    various   astronomical   observations.    Such annihilations  result  in electron-positron  pair  production above  a threshold energy given by $2m_{e}$  in the center of momentum frame of the  system, assuming  Lorentz  invariance. If  Lorentz invariance  is violated (LIV),  this threshold can be  significantly raised, changing the  predicted absorption  turnover in  the observed  spectrum  of the sources.   Stecker and  Glashow  have shown  that  the existence  such absorption  features  in the  spectra  of  extragalactic sources  puts constraints on LIV.  Such  constraints have important implications for some  quantum  gravity  and   large  extra  dimension  models.  Future spaceborne  detectors dedicated to  measuring \\gray\\  polarization can look for  birefringence effects as  a possible signal of  loop quantum gravity.  As  shown by Coleman  and Glashow,  a much  smaller amount  of Lorentz invariance   violation  has   potential   implications  for   possibly supressing  the   ``GZK  cutoff''  predicted  to  be   caused  by  the interactions of  cosmic rays having multi-Joule  energies with photons of  the 2.7  K  cosmic background  radiation  in intergalactic  space. Owing  to the  rarity  of  such ultrahigh  energy  cosmic rays,  their spectra are  best studied by  a UV-sensitive satellite  detector which looks down  on a large volume  of the Earth's atmosphere  to study the nitrogen fluorescence  tracks of  giant air-showers produced  by these ultrahigh  energy  cosmic rays.   We  discuss  here  in particular,  a two-satellite mission called {\\it OWL},  which would be suited to make such studies. ",
        "introduction": "The theory of  relativity is one of the  fundamental pillars of modern physics.  However, because  of  the problems  associated with  merging relativity with quantum theory, it  has long been felt that relativity may  have to  be modified  in some  way.  It  has been  suggested that relativity,  {\\it  i.e.}  Lorentz  invariance  (LI), may  be  only  an approximate   symmetry  of  nature\\cite{sa72}.    There  has   been  a particular interest in the  possibility that a breakdown of relativity may be associated  with the Planck scale of  $M_{QG} \\sim 10^{19}$ GeV where  quantum   effects  are   expected  to  become   significant  in gravitational  theory.  Although no  true  quantum  theory of  gravity exists, it  was independently  proposed that LI  might be  violated in such  a   theory  with  astrophysical   consequences\\cite{ac98}  being manifested at an  energy scale $<< M_{QG}$. The  subject of this paper is the  potential use  of observations of  high energy  phenomena from satellite detectors  to search for  the possible breakdown  of Lorentz invariance. ",
        "conclusions": ""
    },
    "0606/astro-ph0606084_arXiv.txt": {
        "abstract": "With close pairs of quasars at different redshifts, a background quasar sightline can be used to study a foreground quasar's environment in \\emph{absorption}. We used a sample of 17 Lyman limit systems with column density $\\mnhi > 10^{19}~{\\rm cm^{-2}}$ selected from 149 projected quasar pair sightlines, to investigate the clustering pattern of optically thick absorbers around luminous quasars at $z\\sim 2.5$. Specifically, we measured the quasar-absorber correlation function in the transverse direction, and found a comoving correlation length of $r_0=9.2^{+1.5}_{-1.7}~\\hMpc$ (comoving) assuming a power law correlation function, $\\xi\\propto r^{-\\gamma}$, with $\\gamma=1.6$. Applying this transverse clustering strength to the line-of-sight, would predict that $\\sim 15-50\\%$ of \\emph{all} quasars should show a $\\mnhi > 10^{19}~{\\rm cm^{-2}}$ absorber within a velocity window of $\\Delta v < 3000~\\kms$. This overpredicts the number of absorbers along the line-of-sight by a large factor, providing compelling evidence that the clustering pattern of optically thick absorbers around quasars is highly anisotropic.  The most plausible explanation for the anisotropy is that the transverse direction is less likely to be illuminated by ionizing photons than the line-of-sight, and that absorbers along the line-of-sight are being photoevaporated.  A simple model for the photoevaporation of absorbers subject to the ionizing flux of a quasar is presented, and it is shown that absorbers with volume densities $n_{\\rm H} \\lesssim 0.1~{\\rm cm}^{-3}$ will be photoevaporated if they lie within $\\sim 1~{\\rm Mpc}$ (proper) of a luminous quasar.  Using this simple model, we illustrate how comparisons of the transverse and line-of-sight clustering around quasars can ultimately be used to constrain the distribution of gas in optically thick absorption line systems. ",
        "introduction": "\\label{sec:intro}  \\begin{figure} \\centerline{\\epsfig{file=f1.eps,bb=100 80 600 680, width=0.50\\textwidth}} \\caption{Distribution of foreground quasar redshifts, transverse separations, and ionizing flux ratios probed by the background quasar sightlines.  The upper plot shows the ratio of the ionizing flux to the UV background $g_{\\rm UV}$, note the general $R^{-2}$ trend.  The lower plot shows foreground quasar redshift versus transverse comoving separations.  The (blue) squares have a Keck spectrum of the background quasar, (red) triangles have Gemini background spectra, (magenta) upside down triangles have MMT background spectra, and (green) circles have SDSS background spectra. Filled symbols outlined in black have a super-LLS within $|\\Delta v| < 1500\\kms$ at the foreground quasar redshift (see Table~\\ref{table:slls}) and open symbols have no absorber.  The region to the left of the dotted line is excluded by the SDSS fiber collision limit of $\\theta=55\\arcsec$, which explains the paucity of SDSS background spectra there.  The Keck/Gemini/MMT spectra probe angular separations an order of magnitude smaller than the fiber collision limit, allowing us to study the foreground quasar environment down to down to $70~\\hkpc$ where the ionizing flux is $\\sim 10,000$ times the UV background.\\label{fig:scatter}} \\end{figure}  Although optically thick absorption line systems, that is Lyman Limit Systems (LLSs) and damped Lyman-$\\alpha$ systems (DLAs), are detected as the strongest absorption lines in quasar spectra, the two types of objects, quasars and absorbers, play rather different roles in the evolution of structure in the Universe.  The hard ultraviolet radiation emitted by luminous quasars gives rise to the ambient extragalactic ultraviolet (UV) background \\citep[see e.g.][]{HM96,Miralda03,Meiksin05} responsible for maintaining the low neutral fraction of hydrogen ($\\sim 10^{-6}$) in the intergalactic medium (IGM), established during reionization. However, high column density absorbers represent the rare locations where the neutral fractions are much larger.  Gas clouds with column densities $\\log N_{\\rm HI}>17.2$ are optically thick to Lyman continuum ($\\tau_{\\rm LL} \\gtrsim 1$) photons, giving rise to a neutral interior self-shielded from the extragalactic ionizing background.  In particular, the damped Ly$\\alpha$ systems dominate the neutral gas content of the Universe \\citep{phw05}, which is the primary reservoir for the formation of stars observed in local galaxies.  One might expect optically thick absorbers to be absent at small separations from luminous quasars.  For example, a quasar at $z=2.5$ with magnitude $r=19$, emits a flux of ionizing photons that is 130 times higher than that of the extragalactic UV background at an angular separation of $60\\arcsec$, corresponding to a proper distance of 340~$\\hkpc$.   Indeed, the decrease in the number of \\emph{optically thin} absorption lines ($\\log N_{\\rm HI}<17.2$ hence $\\tau_{\\rm LL}\\lesssim 1$), in the vicinity of quasars, known as the \\emph{proximity effect} \\citep{BDO88}, provides a measurement of the UV background \\citep{Scott00}. If Nature provides a nearby background quasar sightline, one can also study the \\emph{transverse proximity effect}, which is the expected decrease in absorption in a \\emph{background} quasar's Ly$\\alpha$ forest, caused by the transverse ionizing flux of a \\emph{foreground} quasar. The transverse effect has yet to be detected, in spite of many attempts \\citep[][but see Jakobsen \\etal 2003]{Crotts89,DB91,FS95,LS01,Schirber04,Croft04}.  On the other hand, it has long been known that quasars are associated with enhancements in the distribution of galaxies \\citep{BSG69,YG84,YG87,BC91,SBM00,BBW01,Serber06,Coil06}, although these measurements of quasar galaxy clustering are mostly limited to low redshifts $\\lesssim 1.0$. Recently, \\citet{AS05}, measured the clustering of Lyman Break Galaxies (LBGs) around luminous quasars in the redshift range ($2 \\lesssim z \\lesssim 3.5$), and found a best fit correlation length of $r_0=4.7~\\hMpc$ ($\\gamma=1.6$), very similar to the auto-correlation length of $z\\sim 2-3$ LBGs \\citep{Adel03}. \\citet{Cooke06} recently measured the clustering of LBGs around DLAs and measured a best fit $r_0=2.9~\\hMpc$ with $\\gamma=1.6$, but with large uncertainties \\citep[see also][]{Gawiser01,Bouchet04}.  If LBGs are clustered around quasars, and LBGs are clustered around DLAs, might we expect optically thick absorbers to be clustered around quasars?  This is especially plausible in light of recent evidence that DLAs arise from a high redshift galaxy population which are not unlike LBGs \\citep{Schaye01,Moller02}.  Clues to the clustering of optically thick absorbers around quasars come from \\emph{proximate DLAs}, which have absorber redshifts within $3000\\kms$ of the emission redshift of the quasars \\citep[see e.g.][]{Moller98}. Recently, \\citet{REB06} \\citep[see also][]{Ellison02}, compared the number density of proximate DLAs per unit redshift to the average number density of DLAs in the the Universe \\citep{phw05}. They found that the abundance of DLAs is enhanced by a factor of $\\sim 2$ near quasars, which they attributed to the clustering of DLA-galaxies around quasars.  Projected pairs of quasars with small angular separations ($\\theta \\lesssim 5\\arcmin$) but different redshifts, can also be used to study the clustering of absorbers around quasars.  In this context, clustering is manifest as an excess probability, above the cosmic average, of detecting an absorber in the neighboring background quasar spectrum, near the redshift of the foreground quasar.  In the first of this series of four papers on optically thick absorbers near quasars \\citep[][henceforth QPQ1]{LLS1}, we used background quasar sightlines to search for optically thick absorption in the vicinity of foreground quasars: 149 projected quasar pairs were systematically surveyed for Lyman limit Systems (LLSs) and Damped Ly$\\alpha$ systems (DLAs) in the vicinity of $1.8 < z < 4.0$ luminous foreground quasars. A sample of 27 new absorbers were uncovered with transverse separations $R < 5~\\hMpc$ (comoving) from the foreground quasars, of which 17 were super-LLSs with $\\mnhi > 10^{19}~{\\rm cm^{-2}}$.  The distribution of foreground quasar redshifts, transverse separations, and ionizing flux ratios probed by the projected pair sightlines studied in QPQ1 are illustrated in Figure~\\ref{fig:scatter}.  The filled symbols outlined in black indicate the sightlines which have an absorption line system with $\\mnhi > 10^{19}~{\\rm cm}^{-2}$ within a velocity interval of $\\left|\\Delta v\\right| = 1500~\\kms$ of the foreground quasar redshift, and the open symbols represent sightlines with no such absorber. The line density of absorbers per unit redshift at $z\\sim 2.5$ with column densities $\\mnhi > 10^{19}~{\\rm cm}^{-2}$ is $dN\\slash dz \\simeq 0.8$\\citep{Omeara06}, which implies that the expected number of \\emph{random} quasar-absorber coincidences within $|\\Delta v|=1500\\kms$ is $\\sim 3\\%$.  The expected number of random quasar-absorber associations in the sample of 149 sightlines is $\\langle N \\rangle = 4.4$; whereas, 17 systems were discovered in QPQ1 -- compelling evidence for clustering.  This clustering is particularly conspicuous on small scales: out of 8 sightlines with $R < 500~\\hkpc$, 4 quasar-absorber pairs were discovered, which implies a small scale covering factor of $\\sim 50\\%$, whereas $\\langle N\\rangle = 0.18$ would have been expected at random. This excess of $\\sim 25$ over random may not come as a surprise when one considers that galaxies are strongly clustered around quasars on small scales. However, the light from \\emph{every} isolated quasar in the Universe traverses these same small scales on the way to Earth.  Naively, we would then expect a comparable covering factor of $\\mnhi > 10^{19}~{\\rm cm}^{-2}$ super-LLSs along the line-of-sight (LOS) -- which is definitely not observed to be the case!  How can we quantify the clustering of absorbers around quasars? How do we compare the transverse clustering pattern to the number of proximate absorbers observed along the LOS?  What can the quasar-absorber clustering pattern teach us about quasars and absorbers?  These questions are addressed in this second paper of the series. We briefly review the QPQ1 survey of projected quasar pairs and discuss the quasar-absorber sample in \\S~\\ref{sec:sample}.  A formalism for quantifying the line-of-sight (LOS) and transverse clustering of absorbers around quasars is presented in \\S~\\ref{sec:formalism}.  In \\S~\\ref{sec:maxL} we introduce a maximum likelihood technique for estimating the quasar-absorber correlation function, which we use to measure the clustering of our quasar-absorber sample in \\S~\\ref{sec:cluster}. Our measurement of the transverse clustering is used to predict the expected number of proximate absorption line systems which should be observed along the line-of-sight in \\S~\\ref{sec:proximity}. In \\S~\\ref{sec:photo}, we introduce a simple analytical model for the photoevaporation of optically thick absorbers by quasars, which is used to illustrate how comparisons of the line-of-sight and transverse clustering can constrain the distribution of gas in optically thick absorbers. We summarize and discuss our results in \\S~\\ref{sec:conc}  Paper III of this series \\citep{LLS3} investigates fluorescent Ly$\\alpha$ emission from our quasar-absorber pairs, and echelle spectra of several of the quasar-LLS systems published here are analyzed in Paper IV \\citep{LLS4}.  Throughout this paper we use the best fit WMAP (only) cosmological model of \\citet{Spergel03}, with $\\Omega_m = 0.270$, $\\Omega_\\Lambda =0.73$, $h=0.72$.  Unless otherwise specified, all distances are comoving\\footnote{Note that in QPQ1 proper distances were quoted.}.  It is helpful to remember that in the chosen cosmology, at a redshift of $z=2.5$, an angular separation of $\\Delta\\theta=1\\arcsec$ corresponds to a comoving transverse separation of $R=20~\\hkpc$, and a velocity difference of $1500\\kms$ corresponds to a radial redshift space distance of $s=15~\\hMpc$.  For a quasar at $z=2.5$, with an SDSS magnitude of $r=19$, the flux of ionizing photons is 130 times higher than the ambient extragalactic UV background at an angular separation of $60\\arcsec$ (comoving $R=1.2~\\hMpc$).  Finally, we use the term optically thick absorbers and LLSs interchangeably, both referring to quasar absorption line systems with $\\log N_{\\rm HI}>17.2$, making them optically thick at the Lyman limit ($\\tau_{\\rm LL}\\gtrsim 1$). ",
        "conclusions": ""
    },
    "0606/astro-ph0606567_arXiv.txt": {
        "abstract": "With a redshift of $z=6.295$, GRB 050904 is the most distant gamma-ray burst ever discovered. It was an energetic event at all wavelengths and the afterglow was observed in detail in the near-infrared bands. We gathered all available optical and NIR afterglow photometry of this GRB to construct a composite NIR light curve spanning several decades in time and flux density. Transforming the NIR light curve into the optical, we find that the afterglow of GRB 050904 was more luminous at early times than any other GRB afterglow in the pre-\\emph{Swift} era, making it at these wavelengths the most luminous transient ever detected. Given the intrinsic properties of GRB 050904 and its afterglow, we discuss if this burst is markedly different from other GRBs at lower redshifts. ",
        "introduction": "The detection of an extremely bright prompt optical flash accompanying GRB 990123 \\citep{Akerlof1999} highlighted the possibility of using GRB afterglows as backlighting sources to probe the high-redshift universe and the reinonization era \\citep{Lamb2000, Inoue2004, Totani2006}. This promise was finally fulfilled with the discovery of GRB 050904 by the \\emph{Swift} satellite \\citep{Cusumano2006a, Cusumano2006b}, lying at a redshift of $z=6.295$ \\citep{Kawai2006}. Not only was a bright NIR afterglow discovered for this GRB \\citep{Haislip2006, Tagliaferri2005}, a prompt flash contemporaneous to the $\\gamma$-ray emission was seen by the robotic 25cm TAROT telescope \\citep{Boer2006}.  Had the prompt optical flash of GRB 990123 been at a redshift of $z=1$ and had it been unextinguished by any dust, it would have peaked at an apparent magnitude\\footnote{In this paper, we use WMAP concordant cosmology with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$, $\\Omega_\\Lambda=0.73$ \\citep{Spergel2003}. For $z=6.295$, this leads to a distance modulus of $44.11$ mag.} $R=7.6^{+0.04}_{-0.05}$ \\citep[][henceforth K06] {PaperIII}. This corresponds to an absolute $R$-band magnitude $M_R=-36.5$ \\citep[$2.3\\cdot10^{16}L_{\\odot R}$, assuming the $V-R$ color for a G2V star from][]{Ducati2001}, making this the most luminous optical source ever detected at that time. In this paper, we compile all available data on the afterglow of GRB 050904 and construct a synthetic light curve for a hypothetical perfectly ionized universe showing no Lyman dropout up to the redshift of the GRB. We then derive light curve parameters and use the spectral energy distribution (SED) of the afterglow to extrapolate the afterglow light curve into the $R$ band assuming $z=1$. We find that the prompt flash of GRB 050904 was the most luminous optical/NIR transient ever detected, even exceeding the $R$-band luminosity of GRB 990123 at peak by more than one magnitude.   \\begin{deluxetable*}{lccccccccccc} \\tablecolumns{11} \\tabletypesize{\\scriptsize} \\tablecaption{Results of the composite $YJHK$ light curve fitting} \\tablehead{ \\colhead{fit\\tablenotemark{a}}   & \\colhead{$\\chi^2$} & \\colhead{d.o.f.}  & \\colhead{$m_c$\\tablenotemark{b}}  & \\colhead{$\\alpha_0$}  & \\colhead{$\\alpha_1$}  & \\colhead{$\\alpha_2$}  & \\colhead{$t_{b1}$\\tablenotemark{c}}  & \\colhead{$t_{b2}$}  & \\colhead{$n_1$\\tablenotemark{d}}  & \\colhead{$n_2$}} \\startdata 1 & 59.2 & 17 & $18.80\\pm0.47$ & $1.39\\pm0.13$ & $0.92\\pm0.05$ & \\nodata      & $0.35\\pm0.14$ & \\nodata      & $-10$    & \\nodata         \\\\ 2 & 59.7 & 17 & $20.65\\pm0.17$ & \\nodata      & $0.85\\pm0.08$ & $2.45\\pm0.18$ & \\nodata      & $2.63\\pm0.37$ & \\nodata & $1.82\\pm1.02$   \\\\ H06 & \\nodata & \\nodata & \\nodata & $1.36^{+0.06}_{-0.07}$  & $0.82^{+0.08}_{-0.21}$ & \\nodata & $\\approx0.5$      & \\nodata & \\nodata & \\nodata \\\\ T05 & \\nodata & \\nodata & \\nodata & \\nodata  & $0.72^{+0.15}_{-0.20}$ & $2.4\\pm0.4$ & \\nodata      & $2.6\\pm1.0$ & \\nodata & \\nodata \\\\ \\enddata \\tablenotetext{a}{The first fit uses data from 0.1 to 3 days after the burst, while the second fit uses only data after 0.3 days. H06 denotes the results from \\cite{Haislip2006}, and T05 those of \\cite{Tagliaferri2005}.} \\tablenotetext{b}{The apparent $J$-band magnitude at the respective break time is denoted as $m_c$ (cf. Z06, Eq. 1).} \\tablenotetext{c}{The break at which the forward shock afterglow begins to dominate over the prompt flash emission is $t_{b1}$, while $t_{b2}$ denotes the jet break time. Break times are given in units of days after the burst trigger.} \\tablenotetext{d}{$n_1$ and $n_2$ are the break smoothness parameters of the first and second breaks, respectively.} \\label{LC} \\end{deluxetable*} ",
        "conclusions": "\\label{Sec5}  We have shown that the early flare of GRB 050904 detected by the TAROT telescope is the most luminous optical/NIR transient ever detected, exceeding even the peak magnitude of the prompt optical flash of GRB 990123. While GRB 050904 is not markedly different from other GRBs in any single aspect, it is among the most extreme GRBs ever detected.  So far, only a single GRB at the end of the reionization era has been detected. It is thus not possible to decide whether GRB 050904 is typical for GRBs at a very high redshift, or if it is simple observational bias that the first GRB to be found at $z>6$ is one of the most extreme events ever detected. More typical events would not have sufficient follow-up to determine the necessary parameters, especially the redshift. The \\emph{Swift} satellite, with its high sensitivity in an energy band that is typical for highly redshifted hard bursts and its possibility of long-duration image triggers, is a very powerful instrument for detecting these very distant events. Probably, more very high redshift GRBs will be detected during the lifetime of \\emph{Swift}, possibly even breaking the redshift records for quasars \\citep{Fan2003} and galaxies \\citep{Iye2006}. Indeed, a GRB at a higher redshift than GRB 050904 may have already been found, GRB 060116 \\citep{Grazian2006}\\footnote{However, see also \\cite{Piranomonte2006} and \\cite{Tanvir2006}.}.  While the detection of a GRB by \\emph{Swift} is the first and most important step, most parameters characterizing the event can only be obtained by ground-based telescopes. The determination of the true burst energetics depends on the measurement of the redshift and the jet break time, and GRBs at $z>6$ are bright only in the NIR. While optical spectroscopy is still feasible to $z\\approx7$, any higher redshifts will move the metal lines that are used for an exact determination of $z$ beyond the typical limit of optical spectroscopy (and the region beyond $0.7\\;\\mu$m is already strongly affected by sky lines). NIR spectroscopy even at large facilities needs bright sources, and thus far, NIR spectroscopy of a GRB afterglow has never been successful in detecting any lines. The dilemma is that observing time at large facilities is precious and observations would only be triggered in case there is already a strong suspicion that the event lies at a very high redshift. This implies the need for rapid photometric redshift determinations that use the Lyman $\\alpha$ trough, as it has been successfully applied to GRB 050904 \\citep{Haislip2006, Tagliaferri2005, Price2006}. One telescope that is already in use and has the capability for rapid multi-color photometry is the MAGNUM telescope, which also observed GRB 050904 \\citep{Price2006}. A key element would be the deployment of more moderately large robotic NIR telescopes such as PAIRITEL \\citep{Bloom2005} and REM \\citep{Zerbi2001}, which also has an optical camera.  More rapid follow-up of high redshift bursts will finally allow GRBs to fulfill their promise as the ultimate probes of the very early universe."
    },
    "0606/physics0606136_arXiv.txt": {
        "abstract": "An increasing number of numerical simulations and experiments describing the turbulent spectrum of Rayleigh-Taylor (RT) mixing layers came to light over the past few years. Results reported in recent studies allow to rule out a turbulence {\\it \\`a la Kolmogorov} as a mechanism acting on a self similar RT turbulent mixing layer. A different mechanism is presented, which complies with both numerical and experimental results and relates RT flow to other buoyant flows. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606421_arXiv.txt": {
        "abstract": "We apply an inverse Compton $e^\\pm$ pair cascade model for $\\gamma$-ray production in the massive binary system LSI +61$^{\\rm o}$ 303 assuming that electrons are accelerated already inside the inner part of the jet launched by the compact object. $\\gamma$-ray spectra, affected by the cascade process, and lower energy spectra, from the synchrotron cooling of the highest energy electrons in the jet, are calculated as a function of the phase of this binary system. $\\gamma$-ray spectra expected in such model have different shape than those ones produced by electrons in the jet directly to observer. Moreover, the model predicts clear anti-correlation between $\\gamma$-ray fluxes in the GeV (1-10 GeV) and TeV ($>200$ GeV) energy ranges with the peak of the TeV emission at the phase $\\sim$0.5 (the peak half width ranges between the phases $\\sim$0.4-0.9 for the inclination of the binary system equal to $60^{\\rm o}$, and $\\sim$0.4-0.1 for $30^{\\rm o}$). The fine features of TeV $\\gamma$-ray emission (fluxes and spectral shapes) as a function of the phase of the binary system are consistent with recent observations reported by the MAGIC collaboration. Future simultaneous observations in the GeV energies (by the GLAST and AGILE telescopes) and in the TeV energies (by the MAGIC and VERITAS telescopes) should test other predictions of the considered model supporting or disproving the hypothesis of acceleration of electrons already in the inner part of the microquasar jets. ",
        "introduction": "The binary system LSI +61$^{\\rm o}$ 303 contains a massive Be type star, with a radius  of $r_\\star$=13.4 R$_\\odot$ and a surface temperature $T_{\\rm s} = 2.8\\times 10^4$ K, and a compact object on a close  orbit characterised by a semi-major axis $a = 5.3 r_\\star$, an ellipticity $e = 0.72$, and an azimuthal angle of the observer in respect to the periastron passage $\\omega = 70^{\\rm o}$ (Casares et al.~2005). The inclination angle of the binary system is not well constrained by the observations and ranges from  $25^{\\rm o} < i < 60^{\\rm o}$ for the case of a neutron star and $\\theta< 25^{\\rm o}$ for a black hole. In the radio wavelengths, the observed asymmetric extended structure is usually interpreted as due to relativistic radio jets, with an estimated  speed $\\sim$0.6c (Massi et al. 2004). This radio source has been related by Gregory \\& Taylor~(1978) to the  COS B $\\gamma$-ray source CG135+01 (Hermsen et al. 1977). CG135+01 has been also detected by the EGRET instrument above 100 MeV (3EG J0241+6103, Hartman et al. 1999) with a hard spectrum (with a photon spectral index $2.05\\pm 0.06$, Kniffen et al. 1997), and by the COMPTEL detector in the energy range $\\sim$0.75 to 30 MeV (with a photon spectral index $1.95^{+0.2}_{-0.3}$, van Dijk et al. 1996). The analysis of different EGRET observations shows evidences of variability (Tavani et al. 1998, confirmed by Wallace et al. 2000) presenting a modulation with timescales similar to the orbital period of LS I +61$^{\\rm o}$ 303 and the emission maxima likely near the periastron passage and phase 0.5 (Massi 2004). The Whipple group (Hall et al.~2003, Fegan et al.~2005) set an upper limit on the TeV flux from this source, $0.88\\times 10^{-11}$ cm$^{-2}$ s$^{-1}$ above 500 GeV and $1.7\\times 10^{-11}$ cm$^{-2}$ s$^{-1}$ above 350 GeV, which is clearly below an extrapolation of the EGRET spectrum. Very recently, the MAGIC telescope observed TeV $\\gamma$-ray emission at energies above $\\sim$200 GeV with the power law spectrum (spectral index -2.6) and extending up to a few TeV (Albert et al.~2006). This emission is observed after the periastron passage of the compact object. $\\gamma$-ray emission from LSI +61$^{\\rm o}$ 303 has been interpreted in terms of the inverse Compton (IC) scattering model in which stellar radiation is scattered by electrons accelerated on the shock wave either created in collisions of the stellar and pulsar winds (e.g. Maraschi \\& Treves 1981), or by electrons accelerated in the jet and scattering of synchrotron and stellar radiation (e.g. Bosch-Ramon \\& Paredes~2004). The alternative hadronic model for this source has been also discussed by Romero, Christiansen \\& Orellana~(2005). The variable GeV gamma-ray emission reported from LSI +61$^{\\rm o}$ 303 by EGRET can be naturally explained provided that the production region of $\\gamma$-rays lays inside the binary system, i.e. not far away from the massive star. Then, the optical depths for accelerated electrons (and produced $\\gamma$-rays) are large enough for developing the IC $e^\\pm$ pair cascade in the radiation field of the massive star (Dubus 2006, Bednarek~2006 (B06)). For the geometry of the considered binary system, the radiation field is strongly anisotropic if considered with respect to the injection place of the electrons. Such general type of anisotropic IC $e^\\pm$ pair cascades in the radiation of the massive stars have been considered in the past by Bednarek~(1997, 2000) and Sierpowska \\& Bednarek~(2005), and recently applied to the binary systems LS 5039 and LSI  +61$^{\\rm o}$ 303 by Bednarek~(B06). Aharonian et al.~(2006) also consider anisotropic cascades with the application to LS 5039. The basic features of such cascades differ significantly from the well-known IC $e^\\pm$ pair cascades  in the isotropic radiation fields studied in detail in the past (e.g. Protheroe~1986, Honda~1989, Protheroe \\& Stanev~1993, Biller~1995).  In our previous paper (B06), we considered $\\gamma$-ray production in the cascade model assuming that electrons are injected isotropically at some distance from the massive star. Based on those calculations, we were able to predict at which orbital phases the largest fluxes of TeV $\\gamma$-rays should be expected. In fact, these predictions have been successfully tested by the MAGIC observations (Albert et al. 2006). In this paper, based on the anisotropic cascade calculations (B06), we discuss a more specific model for the $\\gamma$-ray production  in LSI +61$^{\\rm o}$ 303. Here we assume that electrons are accelerated on shocks waves propagating along the jet launched from the vicinity of the compact object. They are injected along the jet with specific injection efficiency which does not depend on the phase of the binary system. We do not consider the effects of possible changes of the acceleration efficiency related to, e.g. variations in the accretion rate onto the compact object and neglect the possible radiation from the equatorial disk around the massive Be star in respect to its emission from the surface. In the present code we include the synchrotron losses of electrons in the jet which were not considered in previous work (B06). ",
        "conclusions": "The massive binary systems provide unique conditions for investigating of the basic problem of astrophysics, i.e. the acceleration process of particles, due to the presence of a well-defined target (the radiation field of the massive star). However in order to derive useful constraints the details of the sometimes complicated radiation processes by accelerated particles (e.g. cascades in the anisotropic radiation of the massive star) have to be considered. Calculations of the type described in this paper predict observational features of the $\\gamma$-ray emission (in the GeV to TeV energy range) which, if observed, strongly indicate on a very efficient acceleration of electrons already close to the base of the microquasar jets and allows to constrain the efficiency of this acceleration process. In fact, the $\\gamma$-ray emission from the microquasar type massive binary LS 5039 has been already observed in the TeV energies (Aharonian et al.~2005) and in the GeV energies (Paredes et al.~2000), and from LSI +61$^{\\rm o}$ 303 is observed in the TeV energies (Albert et al.~2006) and  in the GeV energies (Kniffen et al.~1997). Our cascade IC $e^\\pm$ pair model predicts the $\\gamma$-ray light curve in the GeV energy range from LSI +61$^{\\rm o}$ 303 (Fig.~3) which shows maximum close to the periastron as suggested by Massi~(2004). In contrast, the TeV $\\gamma$-ray light curve peaks at the phase $\\sim 0.5$ showing clear anti-correlation with the GeV emission (see also earlier calculations of the propagation of $\\gamma$-rays inside LSI +61$^{\\rm o}$ 303 by Bednarek~(B06)). This cascade model also predicts the flattening of the $\\gamma$-ray emission above a few hundred GeV and lower $\\gamma$-ray flux at phases close to the periastron.  \\begin{figure} \\vskip 4.5truecm \\special{psfile=lsi-fig4.eps voffset=-30 hoffset=0 hscale=55 vscale=55} \\caption{The comparison of the $\\gamma$-ray spectrum produced by electrons in the jet directly to the observer (dotted curve) with the $\\gamma$-ray spectrum escaping to the observer in the cascade process considered in this paper (full curve). The parameters of the example model are the following: $\\xi = 0.1$ and $\\eta = 0.1$ and the electron spectrum $\\propto E^{-2}$, inclination of the binary system $i = 60^{\\rm o}$ and the compact object is at the apastron passage.} \\label{fig4} \\end{figure}  The $\\gamma$-ray spectra shown in Figs.~2 and 3 are calculated only from the part of the jet extending below $10r_\\star$ (in order to fulfil the condition of local, efficient cooling of electrons) and above $0.1r_\\star$. Note, however, that the considered region of the jet is broad enough as to provide the dominant contribution to the escaping $\\gamma$-rays (see also Bosch-Ramon et al.~2006). The region below $0.1r_\\star$ might contribute only to the GeV part of the spectrum. Moreover, the IC model in which electrons in the inner part of the jet scatter soft radiation coming from the accretion disk should predict different $\\gamma$-ray light curves. For example, in our model the $\\gamma$-ray flux depends on the relative orientation of the jet, massive star and the observer predicting the peak of TeV emission at phase $\\sim$0.5. These features can be tested by the future GeV and TeV $\\gamma$-ray observations. As we mentioned in Sect.~1, the IC  model in which collimated electrons scatter stellar radiation has been  already considered in a few papers, e.g. recently by Bosch-Ramon et al.~(2006) , Dermer \\& B\\\"ottcher~(2006). These previous works do not take into account the  IC $e^\\pm$ pair cascade effects which, as we have show above, can play an important role. In order to quantify this effect we compare in Fig.~4 the $\\gamma$-ray spectra produced by electrons in the jet directly toward the observer and the $\\gamma$-ray spectra which are formed in the cascade process for the apastron passage where differences between these two spectra are expected to be the smallest. Although the $\\gamma$-ray spectra above a few hundred GeV are almost this same, they differ significantly at lower energies. Therefore, the broad band shape of the $\\gamma$-ray spectra are clearly different. Note that similar effects are expected for the inclination angle $i=30^{\\rm o}$ since differences between the spectra at the apastron passage for these two inclination angles are small (see the light curves for these two inclination angles in Fig.~3 for the phase of $0.73$). The cascade process considered in this paper significantly re-distribute the energy in the $\\gamma$-ray spectrum toward the observer, influencing e.g.  the normalization of the spectrum  at lower energies. In the works mentioned above the contribution of $\\gamma$-ray emission from electrons  scattering of the disk radiation in the jets of microquasars have also been included. It was shown that this contribution can not dominate over the $\\gamma$-rays produced  by scattering the stellar radiation.  On the other hand, electrons accelerated in the jet above $10r_\\star$ might contribute to the highest energy part of the $\\gamma$-ray spectrum and to the low energies (X-rays) as a result of synchrotron process and synchrotron self-Compton mechanism since synchrotron energy losses of electrons might dominate far away from the massive star (e.g. Atoyan \\& Aharonian~1999). However, significant variability of the TeV flux with timescales similar to the period of the binary system has been recently detected by the MAGIC telescope (Albert et al.~2006). So then, the contribution to the observed $\\gamma$-ray spectrum from the outer parts of the jet should be rather minor in case particle injection were constant along the orbit. The problem of the phase dependent accretion onto the compact object has not been considered in the present work (for possible indications how it can occur see e.g. Bosch-Ramon et al.~2006, Christiansen et al.~2006). We concentrated only on the radiation processes occurring during uniform injection of energy along the jet (independent on the phase).  The electrons cooled up to a few GeV in the inner jet are advected to the outer parts of the jet. They can be responsible for the lower energy emission observed from some microquasars. Based on the considered here cascade model, we estimate that a significant fraction of the particle energy ($\\sim 10^{-3}$) survives the $0.1-10r_\\star$ region to allow the formation of radio jets, but it's treatment is out of the scope of the paper."
    },
    "0606/astro-ph0606492_arXiv.txt": {
        "abstract": "We present the diffuse X-ray emission identified in \\chandra\\ observations of the young, massive Galactic star cluster Westerlund 1. After removing point-like X-ray sources down to a completeness limit of $\\approx$$2\\times10^{31}$ \\ergs, we identify $(3\\pm1)\\times10^{34}$ \\ergs\\ (2--8 keV) of diffuse emission. The spatial distribution of the emission can be described as a slightly-elliptical Lorentzian core with a half-width half-maximum along the major axis of 25\\arcsec$\\pm$1\\arcsec, similar to the distribution of point sources in the cluster, plus a 5\\arcmin\\ halo of extended emission. The spectrum of the diffuse emission is dominated by a hard continuum component that can be described as a $kT$$\\ga$3 keV thermal plasma that has a low iron abundance ($\\la$0.3 solar), or as non-thermal emission that could be stellar light that is inverse-Compton scattered by MeV electrons. Only 5\\% of the flux is produced by a $kT$$\\approx$0.7 keV plasma. The low luminosity of the thermal emission and the lack of a 6.7 keV iron line suggests that $\\la$40,000 unresolved stars with masses between 0.3 and 2\\msun\\ are present in the cluster. Moreover, the flux in the diffuse emission is a factor of two lower than would be expected from a supersonically-expanding cluster wind, and there is no evidence for thermal remnants produced by supernovae. Less than $10^{-5}$ of the mechanical luminosity of the cluster is dissipated as 2--8 keV X-rays, leaving a large amount of energy that either is radiated at other wavelengths, is dissipated beyond the bounds of our image, or escapes into the intergalactic medium. ",
        "introduction": "Sensitive X-ray observations are an increasingly-important tool for studying young star clusters, particularly now that the {\\it Chandra X-ray Observatory} and the {\\it Newton} X-ray Multi-Mirror Mission have made harder X-rays (2--10 keV) available for study. Young stars of all types are strong X-ray sources, with low-mass ($M$$<$3\\msun) pre-main-sequence stars producing X-rays in their active magnetic coronae \\citep{pf05,fei05}, and massive OB stars ($M$$\\ga$8\\msun) producing X-rays through shocks in their stellar winds \\citep{cg91,ber97,ski02a}. Therefore, using observations of local star forming regions (e.g., Orion) as templates, measurements of the integrated X-ray luminosities of more distant clusters can be used to constrain their total stellar population, including the numbers of young stars that may be unobservable in the optical and infrared because of extinction or source confusion \\citep{fei05,ns05}.  X-ray observations of clusters of massive stars also reveal diffuse X-ray emission that is produced as stellar winds encounter each other and the surrounding interstellar medium \\citep[ISM; e.g.,][]{sh03,tow03,lyz04,tow05,tow06}. % Learning the fate of the energy carried by these winds, and eventually by supernovae, would provide insight into how galaxies evolve. If the energy is transferred to the ISM, it might at first trigger future generations of star formation, but a sufficiently large input of energy could clear away the ISM and halt star formation. Alternatively, if the energy escapes a galaxy, stellar winds and supernovae would enrich the intergalactic medium with metals. To determine the fate of that energy, it is necessary to obtain X-ray observations of clusters that have a range of ages and populations, and that are surrounded by ISM with a variety of densities \\citep[e.g.,][]{tow03,tow05}.  \\begin{figure*}[htb] \\centerline{\\epsfig{file=diffuse_img_wide.ps,width=0.5\\linewidth,angle=90}} \\caption{ Images of the 10\\arcmin\\ by 10\\arcmin\\ field around the center of Westerlund 1. The {\\it left panel} displays the unbinned image, with the point-like X-ray sources described in Clark \\etal\\ (in prep.) marked with circles. The {\\it right panel} displays an image of the diffuse X-ray emission in which the point sources have been excised and image has been adaptively binned to a signal-to-noise of 10 using the Weighted Voronoi Tessellation algorithm \\citep{ds06}. The location of the 10.6 s X-ray pulsar \\citep{mun06} is also indicated. } \\label{fig:diffimg} \\end{figure*}  In this paper, we report on \\chandra\\ observations of the diffuse X-ray emission from the young Galactic star cluster Westerlund 1. The cluster contains 24 Wolf-Rayet (WR) stars, more than 80 blue super-giants, at least 3 red super-giants, a luminous blue variable, and an amazing 6 yellow hyper-giants, only 6 of which are known in the entire rest of the Galaxy \\citep{wes87, cn02, cn04, nc05, cla05}. Assuming a standard initial mass function \\citep{krou02}, Westerlund 1 could be as massive as $10^5$ \\msun, making it several times larger than the well-known, young Galactic clusters the Arches, Quintuplet, and NGC 3603. Westerlund 1 is also located only $\\approx$5 kpc away \\citep{cn02,cla05}, so it is one of the closest young, dense star clusters. Therefore, Westerlund 1 is a crucial object for understanding the evolution of star clusters, and their impact on the ISM of their host galaxies.  This is one of several papers describing the \\chandra\\ observations. In \\citet{mun06}, we reported the detection of a slow X-ray pulsar in Westerlund 1. \\citet{ski06} examined the X-ray emission from the Wolf-Rayet (WR) stars in the cluster, as well as a subset of the OB supergiants that are brightest in X-rays. In Clark \\etal\\ (in prep.), we will report the spectroscopic identification of further optical counterparts to the X-ray sources, and discuss the origin of the X-ray emission from these stars. In this paper, we describe the spatial distribution (\\S\\ref{sec:dist}) and spectrum (\\S\\ref{sec:spec}) of the diffuse X-rays. We compare the emission seen from Westerlund 1 to that of other massive star clusters in the Local Group (\\S3). We suggest that Westerlund 1 is one of only a few star clusters to produce mostly non-thermal X-rays. We discuss the contraints that this places on the contributions of pre-main-sequence stars (\\S3.1), stellar winds (\\S3.2), and supernovae (\\S3.3) to the diffuse emission, and examine what could be causing non-thermal emission (\\S3.4). ",
        "conclusions": "The origin of the diffuse X-ray emission from clusters of massive young stars is currently under debate. Several authors \\citep[e.g., Cant\\'{o}, Raga, \\& Rodr\\'{i}guez 2000;][]{sh03} have modeled the diffuse X-rays from the most massive clusters in the Local Group as a cluster wind. Under this model, the winds of individual stars collide, thermalize, and form a pressure-driven bulk flow that expands supersonically into the interstellar medium \\citep{cc85}. \\citet{sh03} tabulated results from studies of R136, NGC 3603 \\citep{mof02}, NGC 346 \\citep{naz02}, the Rosette \\citep{tow03}, and the Arches \\citep{yz02}, and showed that the luminosities of their diffuse X-ray emission ($[1-6]\\times10^{34}$ \\ergs) were considerably larger than would be expected from the standard cluster wind model. The large X-ray luminosities can be explained several ways: the densities of the cluster winds could be higher than expected because the stellar winds entrained cooler material or because radiative losses decreased the temperature of the shocked plasma \\citep{sh03}; the wind energy could be dissipated through heat conduction where it encounters nearby molecular clouds \\citep{dm87}; or the wind could be confined by the surrounding ISM \\citep{chu95}. Alternatively, the large X-ray luminosities might partly result from the fact that unresolved pre-main-sequence stars should contribute significantly to the luminosity of the (apparently) diffuse emission, especially for more distant clusters \\citep[e.g.,][]{tow06}.  Westerlund 1 is at least as massive as NGC 3603, R136, and the Arches \\citep{cla05}, so from an observational standpoint the luminosity of its diffuse X-rays ($[3\\pm1]\\times10^{34}$ \\ergs; 2--8 keV) is understandable. However, the spectrum and spatial distribution of the diffuse X-ray emission from Westerlund 1 presents more of a puzzle. First, the spectrum lacks the line emission from He-like Fe that would be expected from a thermal plasma given the hard continuum flux. % In contrast, the spectra of pre-main sequence stars exhibit prominent lines from He-like and H-like Si, S, Ar, and Fe that imply metal abundances up to ten times the solar value \\citep[e.g.,][]{fei05}. However, X-ray spectra of O and WR stars often exhibit weak Fe emission that imply abundances $\\la$0.3 solar \\citep[e.g.,][]{ski01,ski02a,sch03,ski05}, so it is possible that the diffuse flux is dominated by plasma from the O and WR star winds, with little contribution from pre-main-sequence stars. It is also possible that the diffuse emission is non-thermal, by analogy with similar interpretations for the hard flux from a handful of young stellar associations, including RCW 38 \\citep{wol02}, DEM L192 \\citep[N51D;][]{coo04}, 30 Dor C \\citep{bam04}, and possibly the Arches cluster \\citep{lyz04}. In most of the above cases, the non-thermal emission has been interpreted as synchrotron emission from supernova remnants. We will examine this hypothesis for Westerlund 1 in \\S\\ref{sec:nonthermal}.  The second surprise is that the the diffuse X-ray emission from Westerlund 1 seems to extend far beyond the core of the cluster. Within the inner 2\\arcmin, the surface brightness of the diffuse emission falls off with a half-width half-maximum of 0.5 pc (Fig.~\\ref{fig:diffimg} and \\ref{fig:diffdist}), which is identical to the distribution of point-like X-ray sources (Clark \\etal, in prep). This core of diffuse emission could be produced either from the cluster wind, which radiates X-rays mostly in the region where the colliding winds are thermalized \\citep{sh03}, or from an unresolved population of pre-main sequence stars. However, between 2\\arcmin\\ and 5\\arcmin\\ (3--7 pc) from the cluster core the diffuse X-ray flux attains a constant level $\\approx$$7\\times10^{-14}$ \\ergcms\\ arcmin$^{-2}$ (Tab. \\ref{tab:spec}; Fig. \\ref{fig:diffimg}), which is 2--3 times larger than is expected from the Galactic plane \\citep[e.g.,][]{han04,ebi05}. An expanding thermal plasma would exhibit a rapidly-declining temperature profile, yet the spectrum of this halo of diffuse emission is quite hard and lacks the line emission expected from a cooling plasma. This makes it tempting to interpret the diffuse halo as non-thermal particles that are accelerated in a large-scale outflow.  Therefore, although the luminosity of the diffuse X-ray emission from Westerlund 1 is not surprising, the lack of line emission in the spectrum and broad spatial distribution of the diffuse X-rays is. To address this, in the following sections we quantify the probable contributions of pre-main-sequence stars, stellar winds, and supernovae to the X-ray emission from Westerlund 1.  \\subsection{Unresolved Low-Mass Stars}  The non-thermal spectrum of the diffuse X-ray emission from Westerlund 1 puts interesting constraints on the population of low-mass stars in the cluster. The average spectrum of the lightly-absorbed pre-main-sequence stars in Orion can be qualitatively described as a two-temperature plasma with $kT_1$=0.5 keV and $kT_2$=3.3 keV, and with metal abundances of up to 10 times solar for S, Ar, Ca, and Fe \\citep{fei05}. In contrast, in our models for the diffuse emission from Westerlund 1, a $kT$$\\la$1 keV plasma contributes only 5\\% of the 2--8 keV diffuse X-ray flux, and the remaining hard flux does not exhibit the expected He-like Fe line at 6.7 keV, placing an upper limit on the Fe abundance of $\\la$0.3 solar (Tab.~\\ref{tab:spec}). To obtain a conservative estimate of the number of low mass stars in the cluster, we assume that they have solar Fe abundances (i.e., much lower abundances than the stars in Orion). We find that implies that they they produce $\\la$30\\% of the diffuse flux from Westerlund 1, or $\\approx$$9\\times10^{33}$ \\ergs (2--8 keV).  We use the results of the \\chandra\\ Orion Ultradeep Project (COUP) to convert this luminosity into a number of low-mass stars, taking into account the difference in ages between the two clusters. For Orion, the 1398 stars later than B4 in the COUP observations have an integrated, de-absorbed 2--8 keV luminosity of $1.2\\times10^{33}$ \\ergs\\ \\citep{fei05}. Most of this emission is produced by stars with 0.3$<$$M$$<$3 \\msun.\\footnote{The COUP sample of X-ray sources is complete to $\\approx$0.2\\msun, and excluding the singularly bright O6 star $\\theta^1$ Ori C (with a luminosity of $3\\times10^{32}$ \\ergs\\ [2--8 keV], it could be detected as a point source in our observations of Westerlund 1) the $>$3\\msun\\ stars in Orion produce only $\\approx$6\\% of the integrated X-ray luminosity.} However, the stars in Westerlund 1, with ages of $\\approx$4 Myr \\citep{cla05}, are significantly older than the 1-Myr-old population in Orion. To take this into account, we first note that when 2--3\\msun\\ stars reach an age of $\\approx$4 Myr, they become fully radiative and their X-ray luminosities drop by an order of magnitude \\citep{fla03}. Even though they are only 5\\% of low-mass stars by number, in Orion these 2--3\\msun\\ stars produce $\\approx$30\\% of the flux from 0.3$<$$M$$<$3 \\msun\\ stars \\citep[see Fig 4.\\ in][]{fei05}, or $\\approx$$8\\times10^{32}$ \\ergs. Second, \\citet{pf05} find that the X-ray luminosities of young stars with 0.5$<$$M$$<$2\\msun\\ fall off with time $\\tau$ as $L_{\\rm X} \\propto \\tau^{-0.75}$, so at 4~Myr the stars in Orion should be $\\approx$3 times fainter. Therefore, if Orion were 4 Myr old, we would expect its $\\approx$1400 stars with 0.3$<$$M$$<$2 \\msun\\ to have a luminosity of $3\\times10^{32}$ \\ergs\\ (2--8 keV). Our upper limit to the integrated X-ray luminosity of low-mass stars in Westerlund 1 is $\\la$$9\\times10^{33}$ \\ergs, so we infer that Westerlund 1 contains $\\la$40,000 stars with masses between 0.3$<$$M$$<$2\\msun.  This number of low-mass stars is smaller than one would expect if one were to extrapolate from the number of massive, post-main-sequence stars in the cluster using a standard initial mass function \\citep{krou02}. There are $\\approx$150 stars brighter than $V$=21 within 5\\arcmin\\ of the center of Westerlund 1, the faintest of which have recently been identified as O7 main sequence stars that would have initial masses $\\ga$30 \\msun\\ (J. S. Clark \\etal\\ in prep). The maximum initial mass of the stars remaining in the cluster is uncertain because there is no precise means to determine the initial masses of the supergiants and WR stars, but \\citet{cla05} argue that it is probably in the range of 40--50 \\msun. If we assume that the initial mass function can be described as a broken power law of the form $dN \\propto M_i^{-\\alpha} dM_i$, where $\\alpha$=2.3 for $M_i$$>$0.5\\msun\\ and $\\alpha$=1.4 for 0.3$<$$M_i$$<$0.5 \\msun\\ \\citep{krou02}, then if there are 150 stars with 30$<$$M_i$$<$50 \\msun, we would expect 100,000 stars with 0.3$<$$M$$<$2 \\msun. This inferred lack of low-mass stars can be explained several ways. The slope of the initial mass function could be flat ($\\alpha$$\\la$2.1), as has been inferred for NGC 3603 and the Arches cluster based on infrared star counts \\citep[e.g.,][]{eis98,sto05}. The mass function could be truncated at low masses ($M$$<$0.6\\msun), by analogy with the fact that $M$$<$7\\msun\\ stars appear to be depleted in the Arches cluster \\citep{sto05}. Finally, the initial masses of the post-main sequence stars in Westerlund 1 could span a much wider range of masses (20--60\\msun) than assumed in \\citet{cla05}.  If we assume that the mass function is truncated at low masses, the total mass of the cluster would not differ significantly from the estimate of $10^{5}$ \\msun\\ in \\citet{cla05} based on the un-modified Kroupa form. However, if the mass function is flat, or the optically-detected stars had a wider range of initial masses, the total mass of the cluster would be only 40,000-70,000 \\msun. Obviously, an accurate measurement of the mass function, and consequently the total mass, requires direct infrared observations of the low-mass stars in Westerlund 1. However, these X-ray observations provide a useful starting point.  \\subsection{Stellar Winds}  It is not clear whether stellar winds or supernovae are the dominant source of the diffuse X-ray emission from Westerlund 1, because it is at an age when both should contribute equally to its mechanical output (e.g., Leitherer, Robert, \\& Drissen 1992)\\nocite{lrd92}. Individual stellar winds carry a power of \\begin{equation} L_{\\rm w} = 3\\times10^{35} \\left( \\frac{\\dot{M}}{10^{-6}M_\\odot/{\\rm yr}} \\right) \\left( \\frac{v_{\\rm w}}{10^3~{\\rm km/s}} \\right)^2~{\\rm erg~s}^{-1}, \\end{equation} where $\\dot{M}$ is the mass loss rate, and $v_{\\rm w}$ is the wind velocity. The WR stars dominate the mechanical output from stellar winds, with typical $\\dot{M} \\approx 6\\times10^{-5}$ \\msun\\ yr$^{-1}$ and $v_{\\rm w} \\approx 1700$ km s$^{-1}$, so that $L_{\\rm w} \\approx 5\\times10^{37}$ \\ergs\\ \\citep{lrd92}. With at 24 WR stars in the cluster, the total mechanical energy output from winds is $>1\\times10^{39}$ \\ergs.  We can estimate the X-ray luminosity of the resulting cluster wind using the analytic solutions to the density and temperature that \\citet[][see also Stevens \\& Hartwell 2003]{crr00} derived for a wind expanding supersonically into the interstellar medium \\citep{cc85}. Within the radius of the cluster where the stars input their energy, the density is given by \\begin{equation} n_0 = 0.1 N \\left( \\frac{\\dot{M}_{\\rm w}}{10^{-5} M_\\odot/{\\rm yr}} \\right) \\left(\\frac{v_{\\rm w}}{10^3~{\\rm km/s}} \\right)^{-1} \\left( \\frac{R_c}{{\\rm pc}} \\right)^{-2}{\\rm cm^{-3}} \\end{equation} (note that we used the supersonic solution with an adiabatic index $\\gamma = 5/3$ for the original equation) and the temperature by \\begin{equation} kT_0 = 1.3\\left( {{v_{\\rm w}}\\over{1000~{\\rm km/s}}} \\right)^2~{\\rm keV}, \\end{equation} where $N$ is the number of stars, $\\dot{M}_{\\rm w}$ is the average mass loss rate of the stars, and $R_c$ is the radius within which the stars are contained and the winds are thermalized. If we use the values above for WR stars, assume $N$=24, and take the radius of the cluster to be $R_c$=4 pc ($\\approx$3\\arcmin at 5 kpc), then we find $n_0$=0.6 cm$^{-3}$ and $kT_0$=4 keV. Using any of the standard plasma models in XSPEC \\citep[e.g.,][]{mlo86,log95}, and converting $n_0$ and $R_c$ to an emission measure (i.e., $K_{\\rm EM} = \\frac{4}{3} \\pi R_c^3 n_0^2$) we find a predicted $L_{\\rm X} = 3\\times10^{34}$ \\ergs\\ (2--8 keV). Therefore, a cluster wind could in principle account for all of the diffuse X-rays from Westerlund 1.  However, the cluster wind model is not able to account for the spatial distribution of the diffuse X-rays from Westerlund 1. A cluster wind would produce almost all of the diffuse X-ray emission within the core radius $R_c$ \\citep{sh03}, whereas at 70\\% of the diffuse emission from Westerlund 1 is part of a halo that extends out to at least 5\\arcmin\\ (Fig. \\ref{fig:diffdist}). If consider only the core of the diffuse emission as originating from stellar winds, then the cluster is underluminous by a factor of two.  This result is particularly surprising given that the standard cluster wind model applied by \\citet{sh03} to NGC 3603, R136, and NGC 346 predicted significantly {\\it less} flux than is observed. In those cases, \\citet{sh03} favored the hypothesis that cold material was being entrained in the wind. To reconcile our results for Westerlund 1 to the cluster wind model, we would have to assume that enough cold material is entrained that the cluster wind no longer emits in the 2--8 keV bandpass. This requires that the plasma be cooled by a factor of $\\approx$10, to $\\la$0.4 keV. Based on Figure~1 in \\citet{sh03}, we estimate that this would require that the mass of cold material input into the wind is twice that of the hot material, or roughly $3\\times10^{-3}$ \\msun\\ yr$^{-1}$. The mass loss rates from the $\\approx$10 luminous blue variables, red supergiants, and yellow hypergiants, which should be $<$$10^{-4}$ \\msun\\ yr$^{-1}$ each \\citep[e.g.,][]{jk90,lrd92,svk04}, probably could not account for this large amount of cool mass. Therefore, either there is a currently-unseen source of mass in Westerlund 1, or the stellar winds are not thermalized within the cluster and escape without radiating much.   \\subsection{Supernovae}  The presence of an isolated X-ray pulsar in Westerlund 1 \\citep{mun06} confirms that supernovae have occurred there. If we extrapolate an initial mass function with slope $\\alpha$=1.8--2.7 for $M$$>$30\\msun\\ to higher masses, we expect that the cluster originally contained $\\approx$80--150 stars with initial masses $>$50\\msun\\ that have already undergone supernova. For the most massive stars, this would have started when the cluster was about 3~Myr old, so the average supernova rate over the last 1 Myr should be on order one every 7,000--13,000 yr. If each supernova had a kinetic energy of $10^{51}$ erg, then the average power released by supernovae is $\\sim (2-5)\\times10^{39}$ \\ergs.  No obvious supernova remnant is present in our \\chandra\\ image of Westerlund 1, but this is not surprising. We have examined images from the {\\it Spitzer} GLIMPSE program (R. Indebetouw, private communication), and there is no evidence that dense gas or dust still surrounds Westerlund 1. Therefore, Westerlund 1 appears to have cleared away the ISM for parsecs around. When a supernova occurs in such an evacuated cavity, a typical radio and X-ray remnant is not expected until the remnant encounters the boundaries of the bubble blown by the cluster \\citep[e.g.,][]{cde89}.  Whether the hard, possibly non-thermal emission from Westerlund 1 is produced by supernovae is unclear. Most supernova remnants produce thermal X-ray emission with strong lines, but a few are also non-thermal X-ray sources. For example, RCW 86 and SN 1006 exhibit non-thermal filaments near the outer boundary of the shock, and thermal emission in the interior \\citep[e.g.,][]{drb04,rho02}. AX J1843.8$-$0352 and G346.3-0.5 exhibit non-thermal emission almost exclusively throughout the remnant \\citep[e.g.,][]{uen03,laz05,hir05}. Unfortunately, there is not a satisfactory explanation as to why a small fraction of supernova remnants produce non-thermal emission, so the issue remains unresolved for Westerlund 1.  \\subsection{Non-Thermal Particles\\label{sec:nonthermal}}  In principle, non-thermal particles can be produced either in supernova remnants \\citep[e.g.,][]{lp04} or in colliding stellar winds \\citep[e.g.,][]{eu93}. Once they are produced, inverse-Compton scattering should dominate synchrotron losses by a large factor in Westerlund 1 \\citep[see, e.g.,][]{rl79}. The ratio of the energy-loss rates is given by the ratio of the background radiation to the magnetic energy density. The energy density of the stellar light from the OB and WR stars in Westerlund 1 is approximately \\begin{eqnarray} \\nonumber U_{\\rm phot} & = & {{L_{\\rm stars}}\\over{4\\pi c D^2}} \\\\ & = & 5.5\\times10^{-9} \\left({{L_{\\rm stars}} \\over {10^{7} L_\\odot}}\\right) \\left(\\frac{d}{1~{\\rm pc}}\\right)^{-2}~{\\rm erg~cm}^{-3} \\end{eqnarray} where $L_{\\rm stars} \\sim 10^7 L_\\odot$ the luminosity of the cluster. For synchrotron losses to be important, magnetic fields would have to have an energy of $B^2/8\\pi \\ga U_{\\rm phot}$, which corresponds to $B\\ga$0.4 mGauss. This is much stronger than the microGauss fields generally assumed for the interstellar medium \\citep[e.g.,][]{beck01}, so inverse-Compton scattering is probably the dominant loss mechanism for non-thermal particles.  If the non-thermal X-ray emission is produced by inverse-Compton scattering, the energy requirements are modest. Non-thermal particles would only need to be replenished at a rate sufficient to balance the X-ray luminosity, $3\\times10^{34}$ \\ergs. Furthermore, inverse-Compton scattering photons from optical and UV energies ($E_{\\rm in} = 2-20$ eV) into the X-ray band ($E_{\\rm out} \\approx 3$ keV) only requires electrons with $\\gamma^2 \\sim E_{\\rm out}/E_{\\rm in}$, or energies of 6--20 MeV. These particle energies are rather small. For comparison, if the magnetic field in Westerlund 1 has a strength of only 10 $\\mu$Gauss, producing non-thermal synchrotron radio emission like that seen from 30 Dor C \\citep{bam04} requires electrons with energies of a few GeV. Therefore, detecting diffuse, non-thermal radio emission from a star cluster like Westerlund 1 \\citep[e.g.,][]{yz03} would provide a much more interesting constraint on the maximum energies of particles than detecting non-thermal X-rays. The interferometric radio observations in the literature \\citep{cla98} would have resolved out arcminute-scale diffuse radio emission, so single-dish observations are necessary to determine whether higher-energy particles are also produced by the cluster."
    },
    "0606/astro-ph0606171_arXiv.txt": {
        "abstract": "A two-component fluid representing dark energy is studied. One of the components has a polytropic form, while the other has a barotropic form. Exact solutions are obtained and the cosmological parameters are constrained using supernova type Ia data. In general, an open universe is predicted. A big rip scenario is largely preferred, but the dispersion in the parameter space is very high. Hence, even if scenarios without future singularities can not be excluded with the allowed range of parameters, a phantom cosmology, with an open spatial section, is a general prediction of the model. For a wide range of the equation of state parameters there is an asymptotic de Sitter phase. ",
        "introduction": "Several cosmological observables indicate that the present Universe is in a state of accelerated expansion. The first evidence in this sense came in the end of the last decade, when two independent observational projects \\cite{SN}, using type Ia supernovae luminosity distance-redshift relation, provided an estimation of the deceleration parameter $q = - a\\ddot a/\\dot a^2$. With a catalogue of about $50$ type Ia supernova with low and high redshifts, the analysis revealed a negative $q$, indicating an accelerated Universe. Today, more than $300$ type Ia supernovae have been identified, with high redshift, and the conclusion that $q$ is negative remains \\cite{tonry}. Since then, there has been an extensive discussion on the quality of the data. This led to a restricted sample of $157$ supernova, called the \"gold sample\" \\cite{riess}. A more recent survey led to the so-called \"legacy sample\", of about $100$ supernova, with high quality data \\cite{astier}. Even if the precise estimation of the cosmological parameters, like the matter density, Hubble parameter, etc, depends quite strongly on the choice of the sample, the conclusion that the Universe is accelerating has remained. Hence, a large part of the community of cosmologist accepts the present acceleration of the Universe as a fact. \\par A combination data from CMB anisotropies of the cosmic microwave background radiation \\cite{verde}, large scale structure \\cite{tegmark1} and type Ia supernovae data \\cite{colistete}, indicates an almost flat Universe, $\\Omega_T \\sim 1$ and a matter (zero effective pressure) density parameter of order $\\Omega_m \\sim 0.3$ \\cite{tegmark2}. Since an accelerated expansion can be driven by a repulsive effect, which can be provided by an exotic fluid with negative pressure, it has been concluded that the Universe is also filled by an exotic component, called dark energy, with density parameter $\\Omega_c \\sim 0.7$. This exotic fluid leads to an accelerated expansion, remaining at same time smoothly distributed, not appearing in the local matter clustering. \\par The first natural candidate to represent dark energy is a cosmological constant, which faces, however, many well-known problems. More recently, other candidates have been studied in the literature: quintessence, k-essence, Chaplygin gas, among many others. For a review of these proposals, see reference \\cite{sahni}. There are also claims that a phantom field (fields with a large negative pressure such that all energy conditions are violated) leads to the best fit of the observational data \\cite{fantasma}. A phantom field implies a singularity in a finite future proper time, which has been named big rip, where density and curvature diverge. This is of course an undesirable feature, but more detailed theoretical and observational analyses must be made in order to verify this scenario. Some authors state, based on considerations about the evaluation of the cosmological parameters, that there is no such {\\it phantom menace} \\cite{pad}. This is still an object of debate. \\par Most of the studies made until now lay on the assumption of a simple relation between pressure and density expressed generically, in a hydrodynamical representation, by $p = w\\rho^\\alpha$. Quintessence, like others dark energy candidates, imply that $w$ varies with the redshift, not being a constant. Chaplygin gas models (generalized or not) \\cite{chaplygin} imply a general value for $\\alpha$, but typically negative, and $w < 0$. Phantom fields could be represented by $\\alpha = 1$, $w < - 1$. Due to the high speculative nature of the dark energy component, many possibilities have been considered in the literature, both from fundamental or phenomenological point of views. \\par In the present work we intend to exploit a more generic relation between pressure and density with respect to those cases normally considered. The main idea is to use a double component equation of state. The relation between pressure and density may be written as \\begin{equation} \\label{eos} p_e = - k_1\\rho_c^\\alpha - k_2\\rho_c \\quad . \\end{equation} where $\\alpha$, $k_1$ and $k_2$ are constants, and the subscript $c$ indicates that such relation concerns the dark energy component of the matter content of the Universe. Let us call the component labelled by $k_1$ as the {\\it polytropic component}, and that one labelled by $k_2$ as the {\\it barotropic component}. This kind of equation of state has been, for example, studied in a theoretical sense in reference \\cite{indianos}. \\par The equation of state (\\ref{eos}) may be also seen as a realisation of the so-called {\\it modified Chaplygin gas} \\cite{benaoum,chakra1,chakra2}. The usual Chaplygin gas model, generalised or not, has been introduced in order to obtain an interpolation between a matter dominated era and a de Sitter phase \\cite{chaplygin}. The modified Chaplygin gas model allows to obtain an interpolation between, for example, a radiative era and a $\\Lambda CDM$ era. In general, in this case, a negative value for $\\alpha$ is considered. But, as it will be seen later, even for a positive value of $\\alpha$, such an interpolation is possible. From the fundamental point of view, the equation of state (\\ref{eos}) can be obtained in terms of self-interacting scalar field. In references \\cite{benaoum,chakra1,chakra2} a connection with the rolling tachyon model has been established. This allows to consider the equation of state (\\ref{eos}) as a phenomenological realisation of a string specific configuration. \\par In reference \\cite{alcaniz1}, a structure similar to ({\\ref{eos}) has been analysed, using observational data, but fixing $k_2 = 1$, with the conclusion that the fitting of the supernova data are quite insensitive to the parameter $\\alpha$. Here, we follow another approach: we will fix $\\alpha = 1/2$, leaving $k_2$ free. This has the advantage of leading to explicit analytical expressions for the evolution of the Universe. Moreover, and perhaps more important, this may lead to interesting scenarios where, for example, the Universe evolves asymptotically as in a de Sitter phase, even if the equation of state is not characteristic of the vacuum state, $p = - \\rho$. \\par We will test the equation of state (\\ref{eos}) against type Ia supernovae data. We will span a four dimensional phase space, using as free parameters the dark matter density parameter $\\Omega_{m0}$, the exotic fluid density parameter $\\Omega_{c0}$ (or alternatively, the curvature parameter $\\Omega_{k0}$), the Hubble parameter $H_0$, and the equation of state parameter $k_1$ or $k_2$. The subscript $0$ indicates that all these quantities are evaluated today. Using the gold sample, we will show that the preferred values indicate $k_2 \\sim 6$, $\\Omega_{m0}$ and $\\Omega_{c0}$ $\\sim 0.3$, and $H_0 \\sim 67$. An open universe is a general prediction for this model. A phantom behaviour is largelly favoured. However, the dispersion is very high, and an asymptotic cosmological constant phase can not be discarded. \\par The use of other observables, like the spectrum of the anisotropy of the cosmic microwave background radiation (CMB) and the matter power spectrum, can in principle restrict more severely the parameter space. However, we postpone this evaluation to a future study because, in both cases, a perturbative analysis of the model is necessary. In this case we must replace the hydrodynamical representation presented above by a fundamental description of the fluid, for example, in terms of self-interacting scalar fields. We note {\\it en passant} that the hydrodynamical representation employed here may lead, at perturbative level, to instabilities at small scales due to a imaginary effective sound velocity, instabilities that can be avoided with a fundamental representation \\cite{jerome}. There are many different ways to implement this more fundamental description, which can lead to different results. The supernova data, on the other hand, test essentially the background, which is somehow independent of the description of the fluid. \\par The paper is organised as follows. In the next section, we obtain some analytical expressions for the evolution of the Universe, and derive the luminosity distance relation for the model. In section $3$, we make the comparison between the theoretical model and the observational data. In section $4$, we present our conclusions. ",
        "conclusions": "In this work, we have explored the possibility that the dark energy has an equation of state given by (\\ref{eos}). This proposal has already been explored in reference \\cite{alcaniz1}, but restricting one of the components to behave like a cosmological constant: a two-component fluid, which is a \"variation\" around a cosmological constant, has been analysed. Here, we alleviate this restriction, but introducing another one: the linear component can have any barotropic index, but the second component must vary as the square root of the density. This allows us to obtain an analytical expression for the evolution of the Universe, at least for a flat spatial section. This analytic expression reveals that it is possible to have an asymptotically de Sitter phase, for $p < 0$ and $k_2 < 1$, even if only $k_2 = 1$ represents the cosmological constant. This is an intriguing aspect of the model. For $k_2 > 1$ there is always a big rip. \\par The restriction in the polytropic factor $\\alpha$ seems not to be so relevant in view of the results of reference \\cite{alcaniz1}: if the second component obeys a polytropic power law, there is a strong degeneracy on the polytropic factor, and almost any value of the power is allowed. \\par Our results indicate an open universe with a matter density parameter around $\\Omega_{m0} \\sim 0.3$, with a similar estimation for the dark energy component. The $\\Lambda CDM$ model favours a closed universe. On the other hand, the predicted value for the Hubble parameter is more consistent with other observational tests, like CMB, that is, $H_0 \\sim 67\\,km/s\\cdot Mpc$ \\cite{spergel}. For the barotropic index in the two-component fluid $k_2$, the results indicate that $k_2 > 1$ is highly favoured. The dispersion is very high, but tends still to favour a phantom scenario. As the component index $k_1$ increases, the preferred value of $k_2$ decreases slightly. \\par It must be remarked that the model described here exhibits a $\\chi^2$ slightly smaller than for the $\\Lambda CDM$ model. This shows that the model is quite competitive. In our opinion, even if a phantom scenario is clearly preferred, the fact of predicting a low density universe is interesting in its own, mainly when compared with the analysis of clustering of matter in the universe.   {\\bf Acknowledgments:} We thank the anonymous referee who has suggested us, besides other remarks, to consider the non-flat case, what has led to new interesting results with respect to the previous version of the paper. We thank also CNPq (Brazil) and CAPES (Brazil) for partial financial support. F.C. and J.F.V.R. thank also FAPERJ (Brazil) for partial financial support. We thank Roberto Colistete Jr. and Martin Makler for their critical remarks."
    },
    "0606/astro-ph0606347_arXiv.txt": {
        "abstract": "We have used the SHARC II camera at Caltech Submillimeter Observatory to make 350 $\\mu$m and 450 $\\mu$m images of the Vega dust disk at spatial resolutions (FWHM) of 9\\farcs7 and 11\\farcs1, respectively.  The images show a ring-like morphology (radius $\\sim100$ AU) with inhomogeneous structure that is qualitatively different from that previously reported at 850 $\\mu$m and longer wavelengths.  We attribute the 350/450 $\\mu$m emission to a grain population whose characteristic size ($\\sim1$ mm) is intermediate between that of the cm-sized grains responsible for emission longward of 850 $\\mu$m and the much smaller grains ($\\stackrel{<}{_\\sim}18$ $\\mu$m) in the extensive halo, visible at 70 $\\mu$m, discussed by \\citet{su05}. We have combined our submillimeter images with {\\em Spitzer\\/} data at 70 $\\mu$m to produce 2-d maps of line-of-sight optical depth (relative column density). These ``tau maps\" suggest that the mm-sized grains are located preferentially in three symmetrically-located concentrations.  If so, then this structure could be understood in terms of the \\cite{wya03} model in which planetesimals are trapped in the mean motion resonances of a Neptune-mass planet at 65 AU, provided allowance is made for the spatial distribution of dust grains to differ from that of the parent planetesimals. The peaks of the tau maps are, in fact, located near the expected positions corresponding to the 4:3 resonance.  If this identification is confirmed by future observations, it would resolve an ambiguity with regard to the location of the planet. ",
        "introduction": "The observable behavior of circumstellar dust particles under the influence of gravity and radiation pressure can provide information on the locations and masses of unseen planets (see, for example, \\citet{wya03,wya06,del05}). A particularly suitable object for study is the A0 star Vega at 7.76 pc whose disk is seen nearly face-on, as suggested by the low inclination ($5^\\circ$) of the stellar rotation axis \\citep{gul94}. Images at 850 $\\mu$m and 1.3 mm \\citep{hol98,wil02,koe01} indicate a partial dust ring ($r\\sim100$ AU) dominated by two unequal clumps interpreted as collisional debris from resonantly-trapped planetesimals. More recent observations with the Multiband Imaging Photometer for {\\em Spitzer\\/} (MIPS) \\citep{su05} have revealed an extensive halo representing small ($\\sim2$ $\\mu$m) and medium-sized ($\\sim18$ $\\mu$m) grains blown out from the ring by radiation pressure, the ring itself being dominated by large grains of radii $a>180$ $\\mu$m.  Dynamical modeling shows the spatial distribution of the ring to be consistent with the 2:1(u) resonance of a Neptune-mass planet which has undergone migration \\citep{wya03}. Such a model predicts the existence of other populated resonances (for example 3:2 and 4:3), the detection of which would help confirm the model and provide new constraints.  We present new observations at 350 $\\mu$m and 450 $\\mu$m which bear on this issue. ",
        "conclusions": ""
    },
    "0606/astro-ph0606037_arXiv.txt": {
        "abstract": "A double-layer silicon detector consisting of two 500$\\mu$m-thick silicon PIN photodiodes with independent readouts was mounted in a vacuum chamber and tested with X-ray sources.  The detector is sensitive from 1--30 keV with an effective area of 6 mm$^2$.  The detector performs best at $-$35 $^\\circ$C with an energy resolution of 220~eV (FWHM, full width at half maximum) at 5.9~keV, and is able to operate at room temperature, $+$25 $^\\circ$C, with moderate resolution around 760~eV (FWHM).  The response of the top layer sensor is highly uniform across the sensitive area.  This large-format silicon detector is appropriate for future X-ray timing missions. ",
        "introduction": "New detectors with large effective areas, rapid timing response, and good sensitivity up to $\\sim$30~keV  are required to fulfill the scientific goals for a future X-ray timing observatory \\cite{kaa04} beyond the Rossi X-ray Timing Explorer (RXTE) \\cite{bra93}.  RXTE has been operating for ten years with fruitful scientific results obtained from its Proportional Counter Array (PCA), which is close to the end of its lifetime.  Thick silicon detectors are, perhaps, the best candidate detectors for future timing missions.  Silicon PIN photodiodes are simple to fabricate, with mature applications in industry and science \\cite{mit04}, and could be manufactured in large volumes at moderate cost with existing technology.  Large-format thick silicon PIN diode arrays sensitive in 1-30 keV would give spectral resolutions of a few hundred eV (FWHM) at 6~keV, which is better than that of PCA.   Also, solid state detectors should have good long-term stability.  This paper describes the structure and test of a double-layer silicon PIN detector.  We present measurement of the spectral resolution and its variation with temperature, and the uniformity of the detector efficiency, charge collection, and resolution.  We conclude with a summary of the performance of the detector and the important characteristics for its application in X-ray astronomy. ",
        "conclusions": "We have demonstrated a large-format silicon PIN photodiode which presents good performance in sensitivity, spectroscopy and timing.  The detector has high quantum efficiency from 1-30 keV, fine spectral resolution of 220~eV at $-$35 $^\\circ$C and 760~eV at 25 $^\\circ$C, and good time resolution of several tens of microseconds or better.  The intended purpose of the double layer structure was to test background rejection by coincidence in both layers.  Spectra taken using a $^{60}$Co source to simulate the radiation environment on-orbit show that only a small fraction of the background events deposit energy on both layers simultaneously.  However, the Au foil incorrectly inserted between the two detector layers, which gave bad uniformity for the bottom layer, decreased the efficiency for capturing coincident events.  The on-orbit background counting rate for Si detectors and means for its reduction are an important topic for future study.  The top layer of the detector shows a excellent uniformity of efficiency, charge collection and spectral resolution.  This suggests that we would obtain similar excellent uniformity for Si PIN photodiode detectors consisting of arrays of small pixels and operated as a position sensitive detector."
    },
    "0606/astro-ph0606201_arXiv.txt": {
        "abstract": "Recent GLIMPSE data have further confirmed the hypothesis of the existence of an in-plane long bar different from the bulge of the Milky Way with the same characteristics as emphasized some years ago by our team. In this paper, we present two new analyses that corroborate recent and earlier claims concerning the existence in our Galaxy of a long flat bar with approximate dimensions 7.8 kpc $\\times $1.2 kpc $\\times $0.2 kpc and a position angle of approximately 43$^\\circ $: 1) star counts with {\\itshape 2MASS} All-Sky Release and {\\itshape MSX} data, which give an excess in the plane region along $0<l<30^\\circ $ compared with $-30^\\circ<l<0$, and which cannot be due to the bulge, spiral arms, a ring, or extinction; 2) new data on the distance of the long bar using the red clump method, together with recent observations of our own that are compared with our model, and that are in agreement with the long-bar scenario. ",
        "introduction": "Among the different stellar components of our Galaxy are a thin and thick disc, a halo, spiral arms, possibly a stellar ring and the bulge/bar. The bar/bulge components (bulge, or bar, or bulge + bar) are perhaps the most difficult to disentangle from other components because of extinction and our perspective of them due to our position in the Galactic plane  (Zhao 2000); they have been the topic of recent controversy concerning their morphology (e.g.\\ Sevenster et al.\\ 1999; Garz\\'on 1999; Merrifield 2004, and references therein). Our hypothesis in this debate has been that there are indeed two components (Hammersley et al.\\ 2001): a bar which is long and in the plane ($-14^\\circ <l<+30^\\circ $, $|b|<\\approx 1.5^\\circ $), and a triaxial bulge, which is short and much wider in latitude (observable at $|l|<\\approx 15^\\circ $, $|b|<\\approx 10^\\circ $). Their position angles (angle between the major axis and the solar vector radius) might not be coincident, although the latest measurements give values that do not differ greatly: 43$^\\circ \\pm 7^\\circ $ for the long bar (Hammersley et al.\\ 2000; hereafter H00); and 28$^\\circ \\pm \\sim 8^\\circ $(syst.) for the triaxial bulge (L\\'opez-Corredoira et al.\\ 2005).  The story of the Galactic bar really begins with the discovery of large non-circular velocities  in the 3 kpc and  135 km/s spiral arms (Rougoor \\& Oort 1960; Rougoor 1964). De Vaucouleurs (1964) was the first to attempt to explain the motions of these arms in terms of a non-axisymmetric potential (a bar) at the centre of the Galaxy. He later predicted that the Milky Way was of type SAB(rs)bc on the grounds of the Galaxy's high spiral arm multiplicity, broken ring structure and non-circular H\\ {\\sc i} motions (de Vaucouleurs 1970). The non-axisymmetry hypothesis, or ``field model'', was not well received at the time by the astronomical community, which favoured Oort's (1977)  ``ejection model'', advocating the expulsion of gas from the centre of the Galaxy, even though it would require energies of the order 10$^{55}$ erg to power such an expansion of the spiral arms (van der Kruit 1971). Kerr (1967), however, explained the 3 kpc and 135 km/s arms in terms of a bar. Peters (1975) modelled the H~{\\sc i} distribution derived from $l$--$v$ diagrams in the inner 4 kpc of the Galaxy in terms of concentric elliptical orbits. He found that an orientation of 45 deg w.r.t.\\ to the solar radius vector fitted the maximum number of radio features in the $l$--$v$ diagrams, including the 3 kpc and 135 km/s arms. With regard to molecular gas, Nakai (1992) compared CO ($J = 1\\rightarrow 0$) emission maps of the barred spiral galaxies Maffei 2, NGC 2903 and NGC 253 with that of the Milky Way Galaxy (derived from the CO radial velocity at a sample of galactic longitudes). All four galaxies show a similar radial CO distribution with a sharp central peak and a secondary hump (corresponding to the bar--spiral arm transition). The radio evidence, then, gave strong indications of: a) a non-axisymmetric distribution of gas orbits and b) a CO intensity distribution that corresponded with other barred galaxies.  Our hypothesis concerning the existence of a long bar originate from near-infrared Two Micron Galactic Survey data (TMGS, Garz\\'on et al.\\ 1993). It began with an analysis of $K$-band TMGS star counts that led Hammersley et al.\\ (1994) to posit the existence of a bar of radius 4 kpc, with one of the tips at $l=27^\\circ $, and the other one at $l=-22^\\circ $, which would give an angle of 75 degrees \\footnote{This was later retracted by H00, who realized that the negative tip is in fact at $l\\approx -12^\\circ $, so that the angle of the bar is instead $\\approx 43^\\circ $.}. Calbet et al.\\ (1995) presented ``tomographs'' of the Galactic plane (derived from TMGS star counts and using the rather inaccurate assumption of a Dirac delta luminosity function, which could serve as a zeroth order approximation of the structure of the Galaxy), and the long bar was evident (figure 14 of Calbet et al.\\ 1995), extending over a radius larger than 3--3.5 kpc, although the method was not accurate enough to provide information about its morphology. Calbet et al.\\ (1996) also proposed the existence of a dust lane preceding the bar at negative longitudes, which would explain the higher extinction observed in this region. Considerable progress was made by H00, who measured the distance of several points along the bar, including $l=27^\\circ $, and, using a technique of analysis of red-clump stars in near-infrared colour--magnitude diagram, concluded that the angle was $43^\\circ \\pm 7^\\circ $. The same analysis with the same data, only using near-infrared filters, was presented by Picaud et al.\\ (2003), who also compared the data with the Besan\\c con model (Robin et al.\\ 2003), and they obtained a similar result ($45^\\circ \\pm 9^\\circ $). Other papers of our team (Garz\\'on et al.\\ 1997; L\\'opez-Corredoira et al.\\ 1999, 2001b) have produced further evidence, as will be described in Section \\ref{.counts}.  Concerning the stellar populations, Ng (1998) finds two populations towards the Galactic Centre: an old metal rich population ($Z=0.005$--0.08, $t$ = 13--15 Gyr), presumably the old Galactic bulge, and a younger less metal-rich population ($Z=0.003$--0.03, $t=8$--9 Gyr), perhaps the long bar. Cole \\& Weinberg (2002) also showed a non-axisymmetric structure that is likely to have been formed more recently than 3 Gyr ago and must be younger than 6 Gyr. However, they are at $|b|>2^\\circ $ so they were possibly  measuring the suburbs of the long bar, some anomalously young population in the bulge, or some contamination of the disc.  It is supposed that bars form as the result of an instability in differentially rotating discs (Sellwood 1981) whereas bulges are a primordial galactic component, so this distinction makes sense and it is not merely an artefact invented to fit certain morphological features. Athanassoula \\& Misiriotis (2002), Athanassoula (2005) and Athanassoula \\& Beaton (2006) use $N$-body simulations to produce long bars, which are broadened in at their centres into boxy bulges (the ratio of sizes between the semi-axis of the long bar and the major axis of the boxy bulge is approximately 1.5). Boxy bulges are in fact just a part of the long bar (Athanassoula 2005). From a comparison of their $N$-body simulations and near-infrared photometry Athanassoula \\& Beaton (2006) argue that the boxy bulge of M31 extends into a long bar. There is a striking parallel with our own Galaxy, which also has a boxy bulge (Picaud \\& Robin 2004; L\\'opez-Corredoira et al.\\ 2005) that extends into the long bar described in this and previous papers. Bars might be vertically extended (Kuijken 1996) but they might be constrained to the plane too. Indeed, many other galaxies exhibit both a bar and a bulge; e.g.\\ NGC 3351, NGC 1433 and NGC 3992, whose images in the Palomar plates, for instance, show the existence of both structures.  Nevertheless, many authors (reviewed by Merrifield 2004) have for many years ignored the bar + bulge hypothesis and preferred to speak about a unique non-axisymmetric structure, namely a bar (i.e.\\ bulge) that is triaxial, short and fatter (what we call this component the bulge). These authors did not consider the possibility of a long bar constrained in the Galactic plane. Although they were more or less successful in fitting the stars in off-plane regions, the abundance of new data now indicates that some new component in the Galactic plane at longitudes $+15^\\circ <l<+30^\\circ $ is being observed that is unexplainable in terms of a disc + bulge + extinction model. This fact will be corroborated in \\S \\ref{.counts} with All-Sky {\\itshape 2MASS} and {\\itshape MSX} data. In addition, in \\S \\ref{.redclump} we use red clump star data from the recent literature (including new observations of our own: Cabrera-Lavers et al.\\ 2006b) to reinforce the old picture of our hypothesis that there is a long bar apart from the bulge. Furthermore, the position angle obtained in the past is confirmed by the latest data (Benjamin et al. 2005). ",
        "conclusions": "Our conclusion is therefore that recent available data on the long bar corroborate H00's hypothesis. The Milky Way has a long thin bar with approximate dimensions 7.8 kpc $\\times $1.2 kpc $\\times $0.2 kpc; a position angle of the major axis with respect to the line Galactic center-Sun of approximately 43$^\\circ $; density of stars (brighter in K-band than $M_K=-6.1$) around $7\\times 10^{-5}$ pc$^{-3}$, and approximately the third part of stars of the thick bulge if we assume similar stellar populations for both components. GLIMPSE data (Benjamin et al.\\ 2005) explicitly supported this model, and other data on the distance of the long bar derived with the red clump method agree in general with this scenario."
    },
    "0606/astro-ph0606215_arXiv.txt": {
        "abstract": "We present an X-ray analysis of 54 normal elliptical galaxies in the {\\it Chandra} archive and isolate their hot gas component from the contaminating point source emission, allowing us to conduct, for the first time, a morphological analysis on the gas alone. A comparison with optical images and photometry shows that the hot gas morphology has surprisingly little in common with the shape of the stellar distribution. We observe no correlation between optical and X-ray ellipticities in the inner regions where stellar mass dominates over dark matter. A shallow correlation would be expected if the gas had settled into hydrostatic equilibrium with the gravitational potential. Instead, observed X-ray ellipticities exceed optical ellipticities in many cases. We exclude rotation as the dominant factor to produce the gas ellipticities. The gas appears disturbed, and hydrostatic equilibrium is the exception rather than the rule. Nearly all hydrostatic models can be ruled out at 99\\% confidence, based of their inability to reproduce the optical-X-ray correlation and large X-ray ellipticities. Hydrostatic models not excluded are those in which dark matter either dominates over stellar mass inside the inner half-light radius or has a prominently cigar-shaped distribution, both of which can be ruled out on other grounds. We conclude that, even for rather X-ray faint elliptical galaxies, the gas is at least so far out of equilibrium that it does not retain any information about the shape of the potential, and that X-ray derived radial mass profiles may be in error by factors of order unity. ",
        "introduction": "The launch of {\\it Chandra} and {\\it XMM-Newton} opened a new era in our understanding of the hydrodynamic histories of galaxies and clusters. Early spectral evidence from {\\it ROSAT} and {\\it ASCA} \\citep[e.g.][]{FabbianoROSATspectra, BuoteASCA} already suggested soft diffuse gas and harder stellar point sources as the two dominant components of the X-ray emission of normal elliptical galaxies. But {\\it Chandra} made it possible for the first time to spatially resolve a significant fraction of the point source component into individual sources, which are now believed to consist mainly of low-mass X-ray binaries \\citep[LMXBs, e.g.][and references therein]{SarazinLMXBGC}. This has significantly contributed to the understanding of the correlation between X-ray and blue luminosity. The $L_{\\rm X}$--$L_{\\rm B}$ diagram shows a steep $L_{\\rm X}\\propto L_{\\rm B}^2$ relation at the gas-dominated group and cluster scale, and gets shallower toward low X-ray luminosities for normal elliptical galaxies due to the increasing importance of the point source component \\citep{OSullivanLxLb}. For these galaxies, the LMXB emission severely contaminates the diffuse hot gas emission and complicates efforts to reveal its spatial structure.  It has long been assumed that the hot interstellar medium (ISM) in elliptical galaxies is in hydrostatic equilibrium with the underlying gravitational potential \\citep[e.g.][]{Forman1985}. The desire to make this assumption is natural since it then gives us a powerful tool to probe the host galaxy's mass distribution. As such, radial mass profiles derived from observed X-ray pressure profiles are among the strongest providers of evidence for the existence of massive dark matter halos surrounding normal elliptical galaxies \\citep[e.g.][]{Forman1985,KilleenEinsteindarkmatter,FabbianoNGC507, HumphreyDarkmatter,FukazawaMassprofiles} and have been found to be consistent with those derived from gas kinematics in spiral galaxies \\citep[][and references therein]{SofueSpiralDarkmatter} and stellar kinematics in ellipticals \\citep[e.g.][]{StatlerDarkmatter}, implying dark halos with mass profiles similar to those of isothermal spheres. Unfortunately, stellar kinematics can only probe the dark matter content out to a few optical radii. Farther out, one has to rely on X-ray mass profiles, gravitational lensing \\citep{KeetonDarkmatter} or kinematical planetary nebula (PN) data \\citep{NapolitanoPN}. A recent study by \\citet {RomanowskyPNDarkmatter} using PN data advocates a rather low dark matter contribution around some elliptical galaxies, at odds with standard $\\Lambda$-CDM simulations \\citep[e.g.][]{SpringelLargescalestruc}. This discrepancy may be reconciled by appealing to radial orbits of the halo stars \\citep{DekelLostFound}, which underlines the difficulty of interpreting PN kinematics.  The assumption of hydrostatic equilibrium in elliptical galaxies has always been considered well founded and understood. That radial mass profiles consistent with other methods can be derived from it indicates that it is not seriously wrong. If hydrostatic equilibrium holds precisely, the isophotes of the hot gas emission should very nearly trace the projected potential isophotes, and may be used to trace the total mass distribution. Buote and collaborators use this approach to argue for a highly flattened triaxial dark matter halo in the elliptical galaxy NGC~720 \\citep{BuoteGeomtest,BuoteNGC720}. However, no test has been devised to verify that hydrostatic equilibrium holds precisely enough all the way through a galaxy to make this kind of inference valid.  Indeed, substantial evidence is accumulating to the contrary. Combining stellar kinematic information with X-ray observations of NGC~4472, both \\citet{CiottiXraymassproblems} and \\citet{MathewsNGC4472orbits} independently argue for a lack of hydrostatic equilibrium in this particular galaxy. \\citet{NGC1700} argue that the extreme X-ray flattening of NGC~1700 provides a counterexample of an object that cannot be in hydrostatic equilibrium, and suggest that it is rotationally flattened. Rotational flattening might be expected for a variety of reasons \\citep{MathewsReview}. A large fraction of the hot gas in elliptical galaxies is thought to come from stellar mass loss, and should carry the stellar angular momentum. A fraction of the hot gas may also be acquired externally during mergers, stripped off during close encounters, or fall in from a circumgalactic gas reservoir \\citep{BrighentiCircumgas}. In each case, the gas should contain significant angular momentum. In the standard cooling flow scenario the gas should slowly flow inward, conserving angular momentum and settling into a rotationally supported cooling disk \\citep{BrighentiDisks,BrighentiDisks2,KleyXraydisks}. However, \\citet{BregmanRotating} demonstrate, with {\\it ROSAT} and {\\it Einstein} data for 6 elliptical galaxies, a lack of gas disk signatures. They find ellipticities that generally do not exceed values of $\\sim 0.2$, whereas disk models predict ellipticities larger than $\\sim 0.5$.  Detailed {\\it Chandra} and {\\it XMM} studies of over two dozen individual early-type galaxies are now published in the literature. Many of these observations reveal systems that are morphologically disturbed, with a large variety of suggested causes. For example, \\citet{Jon01} argue for the presence of shocks in NGC~4636; \\citet{Fin01} find evidence for interactions with the central radio source in NGC~4374; and \\citet{Machacek} suggest that NGC~1404 is moving relative to the Fornax intracluster gas. Only a few objects appear round and quiescent, such as NGC~4555 \\citep{OSullivanNGC4555} or NGC~6482 \\citep{PonmanNGC6482}. Nonetheless, the use of hydrostatic equilibrium as a fundamental assumption about the majority of early-type galaxies persists \\citep[e.g.][]{FukazawaMassprofiles}.  The main obstacle to understanding the true physical state of the hot ISM in elliptical galaxies is the separation of the gas and unresolved LMXB contributions to the diffuse emission, which is particularly important for X-ray faint galaxies. So far, two main approaches have been employed. (1) Resolved point sources are removed, and the residual diffuse image is adaptively smoothed. It is then argued or hoped that the contribution of unresolved point sources can be neglected while interpreting the diffuse emission and that what is shown is close to the gas morphology \\citep[e.g.][]{BuoteNGC720}. (2) The observation is separated into broad radial bins with sufficient signal to extract and fit a spectrum with a two component model, thus spectrally separating gas and LMXBs in the radial profile \\citep[e.g.][]{HumphreyDarkmatter}. Method (1) has the obvious disadvantage that one does not know what one is looking at. One gains no information on gas morphology in X-ray faint systems which are known to be LMXB dominated. Method (2) allows some insight into the radial distribution and extent of the gas but reveals nothing about the true elliptical shape or asymmetric features.  This is the first paper in a series that analyzes data on 54 elliptical galaxies in the {\\it Chandra} public archive. We homogeneously reanalyze the observations and introduce a new technique to isolate the gas from LMXBs. An adaptive binning technique \\citep{DiehlWVT,Cappellari} is then used to reveal the morphology of the gas {\\it alone}. We present a gas gallery and a quantitative morphological analysis. For the first time, we are able to test whether hydrostatic equilibrium generally holds in normal ellipticals, by comparing gas and stellar isophotal ellipticities in the region where the stellar mass is expected to dominate the potential. We show that the archive sample fails the test. We also examine the evidence for rotational support of the gas, and find it to be extremely weak. The subsequent papers in this series will take up the origin of the observed gas morphology \\citep[][hereafter Paper II]{DiehlAGN} and address the importance of central active galactic nuclei (AGN) in reheating the gas \\citep[][hereafter Paper III]{DiehlEntropy}.  The remainder of this paper is organized as follows. In \\S\\ref{s3.data} we describe the details of the data reduction pipeline, the new isolation technique and the method used to derive ellipticity profiles. In \\S\\ref{s3.results} we present our comparison between optical and X-ray properties and address the evidence for hydrostatic equilibrium and rotational support. We discuss the implications of our findings in \\S\\ref{s3.discussion}, before we summarize in \\S\\ref{s3.conclusions}. ",
        "conclusions": "We have analyzed the X-ray emission of 54 normal elliptical galaxies in the {\\it Chandra} archive and isolated their diffuse hot interstellar gas emission from the emission of discrete stellar point sources for the first time. We qualitatively and quantitatively compare the morphology of the hot gas to the shape of the stellar distribution and find that they have very little in common, despite the known $L_{\\rm X}$--$L_{\\rm B}$ relation. We compute ellipticity and position angle profiles for the X-ray gas and compare them to published optical profiles. In particular, we do not find a correlation between optical and X-ray ellipticities measured in the regions where stellar mass dominates the potential, suggesting that the gas is at least far enough out of hydrostatic equilibrium that the information about the shape of the underlying potential is lost. We also argue that X-ray derived radial mass profiles may be in error by factors of as much as a few, without the necessity for the galaxy to drive a global wind.  Although we find large X-ray gas ellipticities to be common, the gas morphology is generally inconsistent with rotationally flattened disks, and a comparison with stellar rotational velocities yields no evidence for significant rotational support.  The fact that neither the shape of the underlying potential nor rotational support determine the overall distribution of the X-ray emitting gas, combined with its general disturbed appearance, suggests the involvement of another major component: the central AGN. We assess the importance of the AGN in Paper II, where we draw a connection between gas morphology and AGN luminosity. These new findings are consistent with the AGN constantly stirring up the interstellar medium by inflating buoyant bubbles, which may also play a role in redistributing the angular momentum of the hot gas through entrainment. These intermittent AGN outbursts could also be responsible for disrupting or masking the signatures of cooling disks in the central regions of elliptical galaxies."
    },
    "0606/astro-ph0606023_arXiv.txt": {
        "abstract": "While conducting a near-infrared (NIR) survey of ``Digel Clouds'', which are thought to be located in the extreme outer Galaxy (EOG), Kobayashi \\& Tokunaga found star formation activity in ``Cloud 2'', a giant molecular cloud at the Galactic radius of $\\sim$ 20 kpc.  Additional infrared imaging showed two embedded young clusters at the densest regions of the molecular cloud. Because the molecular cloud is located in the vicinity of a supernova remnant (SNR) HI shell, GSH 138-01-94, it was suggested that the star formation activity in Cloud 2 was triggered by this expanding HI shell.  We obtained deep J (1.25$\\,\\mu$m), H (1.65$\\,\\mu$m) and K (2.2$\\,\\mu$m) images of one of the embedded clusters in Cloud 2 with high spatial resolution (FWHM $\\sim 0\\farcs 3$) and high sensitivity (K $\\sim$ 20 mag, 10$\\sigma$). We identified 52 cluster members. The estimated stellar density ($\\sim$ 10 pc$^{-2}$) suggests that the cluster is a T-association. This is the deepest NIR imaging of an embedded cluster in the EOG.  The observed K-band luminosity function (KLF) suggests that the underlying initial mass function (IMF) of the cluster down to the detection limit of $\\sim 0.1 M_\\odot$ is not significantly different from the typical IMFs in the field and in the near-by star clusters.  The overall characteristics of this cluster appears to be similar to those of other embedded clusters in the far outer Galaxy. The estimated age of the cluster from the KLF , which is less than 1 Myr, is consistent with the view that the star formation was triggered by the HI shell whose age was estimated at 4.3 Myr (Stil \\& Irwin).  The 3-dimensional geometry of SNR shell, molecular cloud and the embedded cluster, which is inferred from our data, as well as the cluster age strongly suggest that the star formation in Cloud 2 was triggered by the SNR shell. ",
        "introduction": "\\label{sec:INTRO}  The extreme outer Galaxy (EOG), whose Galactic radius ($R_g$) is more than 18 kpc, is known to be a very low-density and low-metallicity region.  Because the distribution limit of population I and II stars are 18-19 kpc and 14 kpc, respectively, the star forming process in such environment has not been studied as well as near-by star forming regions or star forming regions in the inner Galaxy due to the faintness of the sources. The EOG also serves as an excellent laboratory for the study of star forming process because there is no complex star formation history as in the inner Galaxy.  ``Cloud 2'' is one of the molecular clouds in the outermost Galaxy region discovered by \\cite{Digel1994} based on CO observations of distant HI peaks in the Maryland-Green Bank survey \\citep{WW1982}. % It is located at a very large Galactic radius ($R_g$ = 18$-$28 kpc) in the second quadrant of the Galaxy ($l \\sim 130^\\circ$). It is the largest molecular cloud (M $\\sim$ 1 $\\times$ 10$^4$ M$_\\odot$; \\citet{Digel1994}) among those in the outermost region. The Galactic radius of Cloud 2 has been estimated by various methods: $R_g = 28$ kpc (heliocentric distance: $D=$ 21 kpc) from the kinematic distance of CO \\citep{Digel1994}, $R_g =$ 23.6 kpc ($D = 16.6$ kpc) from the latest HI observation \\citep{Stil2001}, and $R_g \\sim$ 15$-$19 kpc ($D=$ 8.2$-$12 kpc) from the optical spectroscopy of a B-type star, MR-1 \\citep{Muzzio1974, Smartt1996}, which is thought to be associated with Cloud 2 \\citep{de Geus}. Throughout this paper we adopt $R_g$=19 kpc ($D=$12 kpc) because it is about the mean value of all the estimated distances and because spectroscopic distance of stars should be more accurate than kinematic distance. Metallicity at the Galactic radius of Cloud 2 is estimated at $\\sim$ -1.0 dex from the standard metallicity curve by \\citet{Smartt1997}. In fact, the metallicity of the B-type star, MR-1, is measured at $-$0.5 to $-$0.8 dex % \\citep{Smartt1996, Rolleston2000} and that of the Cloud 2 itself, is measured at $\\sim$ -0.7 dex from the radio molecular emission lines \\citep{Lubowich}. These metallicity values are comparable to that of LMC ($\\sim -0.5$ dex) and SMC ($\\sim -0.9$ dex) \\citep{Arnault}. Also, Cloud 2 was found to be located in the vicinity of a large SNR HI shell (GSH 138-01-94) which has an almost complete spherical shape with a radius of 180 pc \\citep{Stil2001}.   \\citet{KT2000} found young stellar objects in Cloud 2 with a wide field near-infrared survey. \\citet{NK2006} also found two young stellar clusters at the northern and southern  CO peaks of Cloud 2 with deeper near-infrared imaging with UH 2.2m telescope. They suggested from the geometry of the clusters and SNR HI shell that the star formation activity in Cloud 2 was triggered by the SNR HI shell. Cloud 2 may be the clearest example of SN triggered star formation, which makes use of the advantage of the ``star formation laboratory'' in  the EOG. As a next step, we obtained deep high-resolution images of the stellar clusters in Cloud 2 with the Subaru 8.2 m telescope for studying the details of the SN triggered star formation. Here we present the results for the stellar cluster in the northern CO peak of Cloud 2 (hereafter ``Cloud2-N cluster''). ",
        "conclusions": "\\label{sec:DIS}  \\subsection{Age of the Cloud2-N cluster} \\label{sec:Age} Because KLFs of different ages are known to have different peak magnitudes, the age of the young clusters can be estimated. The comparison of the observed KLF with the model KLFs in Figure~\\ref{fig:KLF} suggests that the age of the Cloud2-N cluster is less than 1 Myr and most likely no more than 0.5 Myr. It is difficult to estimate the age of the cluster with an accuracy of 0.1 Myr because isochrone models for ages of less than 1 Myr is thought to be uncertain \\citep{Baraffe}. However, we can at least conclude that the age of the Cloud2-N cluster is no more than 1 Myr from the above comparison (Fig.~\\ref{fig:KLF}).   There are still two uncertainties which have not been considered in the model KLFs. First, the distance to Cloud 2 has a moderate uncertainty around the most likely value $D = 12$ kpc  (distance modulus (DM) = 15.4 mag) (see section~\\ref{sec:INTRO}). However, we found that the estimated age from model KLFs is still $0.5-1.0$ Myr for the nearest possible distance ($D=8.5$ kpc, DM = 21.1 mag) and the age is even younger than 0.1 Myr for the most distant case ($D=16.6$ kpc, DM = 14.6 mag). The second uncertainty comes from the possible large difference of IMF from the typical IMFs because of the special environment of the EOG. In this case, the age of the cluster could be much older than 1 Myr. The comparison of the observed KLF and the model KLFs for ages more than 2 Myr suggests that the IMF could have been significantly weighted to higher mass compared to the typical IMFs if the age of the cluster were more than 2 Myr. To check this we tried to fit the slope of the observed KLF in $m_K$ = 16 $-$ 19 range assuming the age of 2 Myr (see the dot-dashed line in Figure \\ref{fig:KLF}d). The necessary IMF slope for this fitting required a very unrealistic slope of the IMF at the stellar mass $>$ 1 M$_\\odot$ (see the grey dashed line in Figure \\ref{fig:IMFs}). This is also highly unlikely because the stellar density of the cluster suggests that it is T-association (see section~\\ref{sec:SD}), which does not have a significant number of high-mass stars. Therefore, even with the uncertainties, we can conclude that the age of the Cloud2-N cluster is less than 1 Myr.     \\subsection{Comparison with Other Embedded Clusters in the Far Outer Galaxy} \\label{subsec:SFEOG}   \\citet{Santos2000} discovered two distant embedded young clusters in far outer Galaxy ($R_g = 16.5$ kpc; $D = 10.2$ kpc) with near-infrared imaging of a CO cloud associated with an IRAS source. Their detection-limit ($m_K = 16.4$mag) corresponds to a 1 Myr $\\sim 1-2 M_\\odot$ star seen through 10 mag of visual extinction.  Snell, Carpenter, \\& Heyer (2002) has conducted a comprehensive NIR survey of embedded clusters in far outer Galaxy ($R_g = 13.5-17.3$ kpc) based on the FCRAO CO survey in the far outer Galaxy. They have found 11 embedded clusters with the detection limit of $m_{K'} = 17.5$ mag.  For the most distant cluster, this magnitude corresponds to a 1 Myr old 0.6 $M_\\odot$ star with no extinction, or a 1 Myr old 1.3 $M_\\odot$ star seen through 10 mag of visual extinction. These pioneering works revealed the existence of embedded clusters in the far outer Galaxy for the first time. Most importantly Snell et al.'s comprehensive work found that the star formation activity that is similar to that found throughout the Galaxy is ubiquitously present in the far outer Galaxy.   Our NIR imaging is the first deep NIR imaging of the embedded cluster in the far outer Galaxy with the detection limit of $m_K \\sim 20$, which is 3-4 mag deeper than the previous studies. Assuming that the age of the Cloud2-N cluster is 0.5 Myr old (section~\\ref{sec:Age}), the mass of the brightest star ($m_K \\sim 16$) is $\\sim 2.5 M_\\odot$ and the mass of the faintest star ($m_K \\sim 20)$ is $\\sim 0.12 M_\\odot$, taking into account A$_V$ = 6.4 mag and the disk color excess $\\Delta K_{disk} = 0.21$ mag (section~\\ref{sec:AvDISK}). Despite the largest distance ($R_g=19$ kpc; $D=12$ kpc) of the Cloud 2-N cluster among the embedded clusters in the far outer Galaxy, our deep imaging has reached to the limit close to the substellar mass range for the first time.  Because the model KLFs based on the typical IMFs reasonably fit the observed KLF including the peak near $m_k = 19$, we suggest that the IMF of the Cloud 2-N cluster can be approximated with the typical IMFs. An assessment to more accurate IMF requires an independent study of the age of the cluster based on spectroscopy of each cloud members.    \\citet{Snell2002} found $\\sim 25 - 95$ cluster members for the 11 embedded clusters in the far outer Galaxy with the detection limit of $M_K = 2.5 - 3.5$ mag ($m_{K'} = 17.5$ mag). \\citet{Santos2000} found $\\sim 30$ cluster members for a cluster in the far outer Galaxy with the detection limit of $M_K = 1.4$ mag ($m_K = 16.4$ mag). The numbers of stars with those detection limits are consistent with that in the near-by embedded clusters within 1 kpc from the solar system. We have detected 20 cluster members in Cloud2-N down to $M_K = 2.5$ mag and 40 cluster members down to $M_K = 3.5$ mag. Therefore, the properties of the Cloud 2-N cluster appear to be quite similar to the star clusters found by Snell et al and Santos et al.       \\subsection{Supernova Triggered Star Formation} \\label{SNTSF} \\citet{NK2006} suggested that the supernova remnant (SNR) HI shell \\citep[GSH138-01-94;][]{Stil2001} triggered the star formation activity in Cloud 2. The uniform extinction of the cluster members ($A_V \\sim 6$ mag) measured in section~\\ref{sec:AvDISK} is close to the total extinction of Cloud2-N that is estimated at $A_V \\sim 9$ mag from $^{13}$CO data \\citep{Digel1994}, suggesting that the formation of the Cloud2-N cluster occurs behind Cloud 2. The geometry of the SNR HI shell, Cloud 2-N, and Cloud2-N cluster is summarized in Fig.~\\ref{fig:GEO}. In section~\\ref{sec:Age}, the age of Cloud2-N cluster was estimated at $\\sim$ 1 Myr, which is much less than that of the SNR HI shell (4.3 Myr). These facts strongly suggest that the star formation activity in Cloud 2 was triggered by the SNR HI shell."
    },
    "0606/astro-ph0606509_arXiv.txt": {
        "abstract": "{}{ Most of current models of TeV blazars emission assume a Synchrotron Self-Compton mechanism where relativistic particles emit both synchrotron radiation and Inverse Compton photons. For sake of simplicity, these models usually consider only steady state emission. The spectral features are thus only related to the shape of the particle distribution, and do not depend on the timing of observations. }{ In this letter, we study the effect of, firstly, the lag between the beginning of the injection of the fresh particles and the trigger of the observation, and secondly, of a finite injection duration. We illustrate these effects considering an analytical time-dependent model of the synchrotron emission by a monoenergetic distribution of leptons. }{ We point out that the spectral shape can  be in fact very dependent on observational conditions if the particle injection term is time-dependent, particularly taking into account the effect of the time averaging procedure on the final shape of the SED. Consequences on the acceleration process are also discussed. }{} ",
        "introduction": "TeV blazars  belong to the radio-loud active galactic nuclei (AGN) and are characterized by non-thermal continuum spectrum, optical polarization, flat radio spectrum and strong variability in all frequency bands. The broad band of the SED which extends from radio up to the TeV range, consists typically in two bumps. The low energy one --- peaking in the $0.1-100\\rm\\ keV$ domain --- is commonly attributed to synchrotron emission of ultra-relativistic leptons. In the \\emph{Synchrotron-Self-Compton process} (SSC) the second component is attributed to the inverse Compton (IC) scattering of synchrotron photon field by the same population of leptons.  Emission of TeV photons provides evidence for the presence of particles at roughly the same energy if we suppose that they radiate in the IC Klein-Nishina regime \\citep{bg70}.  Therefore, it requires an extremely efficient particle acceleration mechanism at work in the close environment of supermassive black hole.  Extended observations of TeV blazars like Mrk\\ 501 ($z_s=0.034$, \\citealp{quinn96}) show that synchrotron component can be very variable in the UV/X-ray range. The simplest approach is the well-known homogeneous or one-zone modeling. It considers the emission of a relativistic moving blob of plasma filled by a tangled magnetic field, where ultra-relativistic particle are injected and can cool freely (via both synchrotron and inverse Compton scattering process). Generally, in these models, the acceleration mechanism is not consistently taken into account and the resulting \\emph{energy distribution function} (EDF) of injected particles is prescribed. If we supposed that the acceleration zone could be dissociated from the radiative one, then we must distinguish two different kinds of EDF : the \\emph{injected} one and the \\emph{cooled} one. In the usual one-zone model, the first is directly the result of the acceleration mechanism. The second ensues from the radiative cooling of particles (and possible re-acceleration mechanism). It is mathematically speaking the solution of the standard kinetic equation which the injected EDF is the main source term. The exact shape of the SED could be only understood in the framework of time dependent modeling, depending on the detail of the cooling and injection processes \\citep{krm98,cg99}. The majority of SSC models are steady-state: therefore the instantaneous SED is equal to the observed one. In this case, the authors use broken power-law EDF as the solution of the kinetic equation (characterized by the indices $\\alpha_{\\ell}$ and $\\alpha_u$). Then the subsequent synchrotron spectra have also a broken power-law shape\\footnote{% Only true in the case where particle energy dynamical range, \\textit{i.e.} the ratio $\\gamma_{\\rm max}/\\gamma_{\\rm min}$ is much larger than unity.  } $\\mathcal{F}_{\\nu}\\propto\\nu^{-\\beta_{\\ell,u}}$ with indices\\footnote{% At low frequency, this latter can not be steeper than the one produced by a monoenergetic EDF, given the lower limit $1/3$. } $\\beta_{\\ell,u}=(\\alpha_{\\ell,u}-1)/2 \\leqslant 1/3$ over suitable frequency ranges \\citep{bg70}. Spectral index variations can arise from the variation of $\\alpha_{\\ell,u}$ \\textit{i.e.} by varying the physical conditions of the emission zone (see \\textit{e.g.} \\citealp{kt04}).  We have recently investigated another solution, still in the framework of homogeneous modeling, but with a time dependent particle injection term \\citep[][\\null, hereafter paper I]{sauge04a}.  The kinetic equation is solved numerically taking into account \\emph{(i)} the possible in-situ pair reprocessing and  \\emph{(ii)} both synchrotron and IC cooling term. The source term of the equation is chosen as a quasi-monenergetic distribution of particles (or ``\\emph{pileup}''), injected during a finite time. Physically, the formation of such a distribution results from the combination of a stochastic heating via second order Fermi process and radiative cooling \\citep{hp91,s85}. Then the resulting cooled EDF is also partially a power-law, but with a constant index value $\\alpha=2$ over a dynamical range depending on the details of the cooling process.  In this case, the instantaneous synchrotron spectrum is a power-law of index $\\beta=1/2$ and no spectral index variation could be expected at this stage.  However, such a variation can be obtained by averaging out the instantaneous SED over the time and considering that the observation starting time does not necessarily coincide with the beginning of injection phase.  In this letter we study in more detail this effect. First of all, we derive a basic model in Sect. 2 in the idealized case of monoenergetic injected EDF subject only to synchrotron cooling. Then we derive the global shape of the time average synchrotron part of the SED as a function of the observational parameters. Finally, we illustrate our approach on {Beppo}SAX archival data sets of the object Mrk~501 in Sect. 3, before concluding. ",
        "conclusions": "We have exposed a simple time dependent mechanism in order to explain the spectral shape of the synchrotron spectrum of blazar considering the injection of a pure monoenergetic distribution of ultra-relativistic particles over a finite time range. This latter EDF arises from stochastic acceleration mechanism (Second order Fermi process). Spectral index and SED shape variations can be explained by the value of the starting observation time with respect to the beginning of the injection of fresh particles and from the variation of the acceleration condition even if the instantaneous synchrotron spectrum is universal with a constant spectral index equals to $1/2$. Because blazars are extremely variable sources, there are often observed in ``\\emph{Target-of-Opportunity}\" mode, where satellites observations are triggered by some other instruments after some delay leading naturally to $t_{\\rm obs}\\neq 0$. This mechanism leads to complex time dependent behaviors of the spectral shape ; we show that we can reproduce various spectral shapes from the single to broken power-law shape, with strong or soft cut-off. This model differs from the usual one in the context of the shock acceleration model, where the spectral variability arises from the physical conditions at the shock."
    },
    "0606/astro-ph0606353_arXiv.txt": {
        "abstract": "{Solar flares release a large fraction of their energy into non-thermal electrons, but it is not clear where and how. Bremsstrahlung X-rays are observed from the corona and chromosphere. } {We aim to characterize the acceleration process by the coronal source and its leakage toward the footpoints in the chromosphere. The relations between the sources reflect the geometry and constrict the configuration of the flare.} {We studied solar flares of GOES class larger than M1 with three or more hard X-ray sources observed simultaneously in the course of the flare. The events were observed with the X-ray satellite RHESSI from February 2002 until July 2005. We used imaging spectroscopy methods to determine the spectral evolution of each source in each event. The images of all of the five events show two sources visible only at high energies (footpoints) and one source only visible at low energies (coronal or looptop source, in two cases situated over the limb). } { We find soft-hard-soft behavior in both, coronal source and footpoints. The coronal source is nearly always softer than the footpoints. The footpoint spectra differ significantly only in one event out of five.} {The observations are consistent with acceleration in the coronal source and an intricate connection between the corona and chromosphere.} ",
        "introduction": "\\label{Introduction} The current understanding of solar flares leaves open fundamental questions such as: where is flare energy released, how are particles accelerated? A large part of the energy released in a solar flare is converted into energetic electrons emitting hard X-rays. Therefore, observations in X-ray wavelengths give quantitative measures of heating and particle acceleration in the flare. X-ray observations by Hoyng et al. (\\cite{Ho81}) showed hard X-ray (HXR) sources at both ends of a loop structure, commonly called footpoints. They are thick target bremsstrahlung emission produced by precipitating electrons, accelerated somewhere in the loop. Footpoints can also be seen in H$\\alpha$ and EUV (e.g. Gallagher et al. \\cite{Ga00}, Fletcher et al. \\cite{Fl04}), indicating the precipitation of flare particles and the reaction of the thermal plasma. In an event observed by \\textit{Yohkoh}, Masuda et al. (\\cite{Ma94}) first noted a third HXR source situated above the looptop (looptop or coronal source). Alexander \\& Metcalf (\\cite{Al97}) analyzed this event carefully, concluding that the loop top source can be best described by a thermal component and a non-thermal component which is harder than the footpoint spectrum. Petrosian et al. (\\cite{Pe02}) made an extended study of looptop sources and footpoints in \\textit{Yohkoh}-events. They find that the spectral index of the looptop source is softer than the footpoints on the average by about 1. The accuracy of their spectra however, was limited by the energy resolution of the \\textit{Yohkoh} detectors. \\\\  An important observation about the time behavior of the HXR flux has already been made in the late 1960s by Parks \\& Winckler (\\cite{Pa69}) and Kane \\& Anderson (\\cite{Ka70}). They found that the hardness of a spectrum changes in time and that there exists a correlation between the HXR flux and the hardness of the spectrum (soft-hard-soft or SHS). These observations were later confirmed by several authors, e.g. Benz (\\cite{Be77}); Brown \\& Loran (\\cite{Br85}); Lin \\& Schwartz (\\cite{Li86}); Fletcher \\& Hudson (\\cite{Fl02}); Hudson \\& F\\'arn\\'ik (\\cite{Hud02}). Beside the SHS pattern, a soft-hard-harder (SHH) pattern has also been observed in some events (Frost \\& Dennis \\cite{Fr71}; Cliver et al. \\cite{Cl86}; Kiplinger \\cite{Ki95}). A quantitative study of 24 solar flares observed by RHESSI on this subject has been made by Grigis \\& Benz (\\cite{Gr04}). They find that elementary flare bursts also show SHS. Battaglia et al. (\\cite{Ba05}) made a study of flares of different size, finding that events with smaller HXR flux are softer on the average and that the relation between HXR flux and spectral index at peak time of events of different size is the same as the one from several peaks of one event.  Is the SHS-behavior a feature of the acceleration mechanism as previously claimed? Or is it a transport effect produced by collisions or return currents? A further possibility could be a change in the dominating X-ray source from the coronal source (soft) to the footpoints (hard) and back to the coronal source again (soft). Thus, is the SHS-behavior nothing but a coronal-footpoint-coronal effect? The previous studies have been made using full sun spectra. To investigate the cause of the SHS, the spectra of each source must be analyzed separately. The Ramaty High Energy Solar Spectroscopic Imager (RHESSI, Lin et al. \\cite{Li02}) provides the possibility of making high resolution imaging spectroscopy at different locations on the sun. One can therefore study each source separately in events with several contemporaneous HXR-sources. The high energy resolution yields detailed spectra, allowing a reliable differentiation between thermal and non-thermal emission. Emslie et al. (\\cite{Em03}) made an analysis of a very large event with 4 HXR-sources observed by RHESSI. They find a coronal source with a strong thermal component and two (at times three) footpoints in regions with opposite magnetic polarity.  They report that the spectral indices of the footpoints differ notably and accredit this to collisional losses by different column densities in the loop connecting the footpoints to the coronal source.  The purpose of this work is a systematic study of the relation between coronal source and footpoints in time and spectra for several well observed events. The events were carefully selected, not necessarily the largest ones, but those with informative data concerning both, thermal and non-thermal source parameters. The RHESSI data has been searched for well separated, bright events without strong pileup, situated near the limb. We present here the results for the best observed events of the first 40 months since launch. ",
        "conclusions": "A selection of five RHESSI events with three concurrent X-ray sources (coronal source and two footpoints) has been studied regarding the spectral relations between the sources. All spectra can be fitted with a non-thermal component having a power-law photon distribution. Although no low-energy soft X-ray observations are available for comparison, we believe that they probably are examples for Masuda-type sources (Masuda \\cite{Ma94}), but our looptop sources are generally softer than the one found by Masuda. Therefore, such events are easier to detect with RHESSI than they were with \\textit{Yohkoh} and are more frequent than inferred previously. In addition, all coronal sources and some of the footpoints at times show a thermal component. The major results are:  \\begin{itemize} \\item All coronal sources evolve according to the same time evolution in spectral hardness. The higher the flux, the harder (smaller $\\gamma$) the non-thermal component. This soft-hard-soft pattern correlates with the non-thermal flux without measurable delay. Transport effects such as collisions or an induced electric field cannot cause SHS in the coronal source. \\item As the emission of the footpoints often dominates at 35~keV, it is not surprising that the pattern, previously reported for full sun observations, is also found in the footpoints of three out of 5 events. Imaging spectroscopy suggests that SHS is a feature of all sources, and thus possibly of the accelerator itself. \\item SHS in full sun observations cannot be explained by a change of the dominant source (softer coronal source present at all times plus a hard footpoint source with time-varying intensity). If SHS was caused by such an effect, neither source would display it individually. \\item The difference in spectral index between non-thermal coronal and footpoint emission is not 2, as would be the case if the difference was simply caused by thin and thick target bremsstrahlung, respectively. Smaller differences in $\\gamma$ may be explained by a an intermediate situation between the two extremes. The plasma of the coronal source could act as thick target for low energetic electrons and as thin target for higher electron energies. The cases with $\\gamma > 2$ require a filter effect in the propagation preferentially reducing the distribution at lower energies. Such a filter may be collisions or an electric field. \\item The pivot energy, characterizing the SHS-behavior of the non-thermal emission, is at the energy where the distribution of the thermal and non-thermal components balance in half of the cases. In the other half, the pivot energy is higher than this point. \\item The pivot energy at the footpoints is significantly lower in all cases ($16 -  23$ keV for the coronal source vs. $14 - 15$ keV for the footpoints). Such a difference suggests a filter acting during particle propagation to the footpoints, reducing lower energies more than higher energies. \\item In one out of 5 events the two footpoints have significantly different spectral indices, $\\Delta \\gamma = 0.33\\pm0.04$. The difference is constant during the event, although the spectral indices vary in time. Again, an energy filter during propagation seems to be at work, differing in one flare for the two legs of the loop. \\item The photon flux at energies below about 15 keV is dominated by thermal emission. Most of this emission originates from the coronal source. If its temperature correlates with the pivot energy it may hint at a physical significance of the pivot energy for the acceleration process, but needs further investigation. \\item As pointed out before from \\textit{Yohkoh/SXT} observations, the thermal emission from the coronal source often significantly exceeds the thermal emission of the footpoints, which is detectable in some events and at some times. \\end{itemize} This analysis has shown that the non-thermal X-ray emission in coronal sources cannot only be detected by RHESSI, but can also be studied in time. As the coronal source is directly related to flare energy release, this opens the possibility of further investigating the enigmatic acceleration process of electrons. The temporal and spectral relation of the coronal source to the footpoints suggests an intricate connection between corona and chromosphere. While a comprehensive interpretation of our results in terms of particle propagation and thermal conduction is beyond the scope of this study, the idea that flare energy release and particle acceleration are closely related to the coronal source is supported by our results."
    },
    "0606/astro-ph0606679_arXiv.txt": {
        "abstract": "The capture of electrons by the nucleus $^{7}$Be from the three-body initial state $p+e^{-}+^{7}{\\rm Be}$ in the continuum is studied. On the basis of the expansion of the three-body continuum wave function in a small parameter $\\epsilon\\approx\\,\\left({m_e}/{m_p}\\right)^{1/2}$ [$m_e$ ($m_p$) is the electron (proton) mass], the role of the protons on the electron capture is considered. The results are compared with the traditional treatment of the electron capture by the nucleus $^{7}$Be. For stars with the density and temperature like in the center of the Sun the studied mechanism can make non-negligible contribution to the capture rate. ",
        "introduction": "The process of the electron capture by the nucleus $^{7}$Be is important since it contributes to the low energy part of the spectrum of neutrinos radiated by the Sun. Besides, it is obvious that the balance of the disappearance channels of $^{7}$Be in the Sun regulates the amount of the nucleus $^{8}$B which is the source of the high energy solar neutrinos. This is the main reason why this process attracted considerable attention over many years \\cite{b1,b2,b5,b6,b7,b9,b14}.  Practically all the discussion so far of the electron capture in $^{7}$Be is concentrated on considering the electron wave function in the vicinity of the nucleus and on the screening effects on it; the study of this capture in the plasma was done in Ref.\\,\\cite{b14}.  In what follows we will estimate the role of the process that is usually not included in the standard theory of the $pp$ cycle in the Sun. Let us first note that in the standard theory of this cycle the destruction of the nucleus $^{7}$Be takes place in the following binary reactions \\begin{eqnarray} p\\,+\\,^{7}{\\rm Be}\\,&\\rightarrow&\\,^{8}{\\rm B}\\,+\\,\\gamma\\,,  \\label{br1} \\\\ e^-\\,+\\,^{7}{\\rm Be}\\,&\\rightarrow&\\,^{7}{\\rm Li}\\,+\\,\\nu\\,. \\label{br2} \\end{eqnarray} Since the nucleus $^{7}$Be participates in both processes, instead of the binary reactions (\\ref{br1}) and (\\ref{br2}) we consider the contribution to the electron capture rate from the three-particle initial state \\mbox{$p+e^{-}+^{7}$Be}. In this case the following reactions can take place \\begin{eqnarray} &\\nearrow&\\,^{7}{\\rm Li}\\,+\\,\\nu\\,+p\\,,    \\label{3a} \\\\ p\\,+\\,e^{-}\\,+\\,^{7}{\\rm Be}\\,&\\rightarrow&\\,^{8}{\\rm B}\\,+\\,\\gamma\\,+\\,e^-\\,, \\label{3b}  \\\\ &\\searrow&\\,^{8}{\\rm B}\\,+\\,e^-\\,.  \\label{3c} \\end{eqnarray} As it was shown in Ref.\\, \\cite{b2} the screening corrections for the electrons in the continuum are rather small. Therefore, we consider in the initial state the bare Coulomb interaction in all two-body subsystems $e^-+p$, $p+^{7}$Be, and $e^-+^{7}$Be. In this case one can immediately realize that there is a qualitative difference between the binary and ternary mechanisms\\footnote{We use the terms binary- and ternary reactions as synonyms for the reactions in the two- and three component systems.} of the electron capture. Indeed, if one starts from the three-body initial state, then the processes (\\ref{br1}) and (\\ref{br2}) should be interdependent because the wave function of three charged particles cannot be presented as a product of pair wave functions, as is required by the binary processes (\\ref{br1}) and (\\ref{br2})\\footnote{There is only one exception corresponding to the case when all three particles are at very large distances between themselves \\cite{b4}, which is not applicable to the electron capture.}. Here we apply the method and results obtained in Ref.\\, \\cite{b3} where an alternative to the Born-Oppenheimer approach has been suggested for the calculations of wave functions in the continuum for three charged particles. ",
        "conclusions": "The results of the calculations are presented in Table I and Figs.\\,\\ref{Fig2}-- \\ref{Fig4}. \\begin{figure}[h!] \\centerline{ \\epsfig{file=fig2f.eps} } \\vspace{0.4cm} \\caption{ The dependence of the mean value \\mbox{$<{\\varsigma}(R_0, T)>$} on the temperature $T$ and the value of the mean  distance $R_0$. Solid line, $R_0=0.1\\times 10^4$ fm; dashed line, $R_0=0.25\\times 10^4$ fm; dotted line, $R_0=0.5\\times 10^4$ fm; dashed and dotted line, $R_0=1.0\\times 10^4$ fm; dashed and double dotted line, $R_0=1.5\\times 10^4$ fm.}\\label{Fig2} \\end{figure} In Fig.\\,\\ref{Fig2}, the dependence of the mean value \\mbox{$<{\\varsigma}(R_0, T)>$} given in Eq.\\,(\\ref{BS}) on the temperature $T$ and the value of the mean distance $R_0$ is shown. A weak dependence of \\mbox{$<{\\varsigma}(R_0, T)>$} on the temperature means, as it follows from Eq.\\,(\\ref{Rb}), that the temperature dependence of the electron capture by $^{7}$Be is almost the same for the ternary and binary reactions. Such a behaviour can be understood from the fact that in both cases only rarely is all the kinetic energy carried by one particle, which is the electron. On the other hand, the dependence of \\mbox{$<{\\varsigma}(R_0, T)>$} on the value of $R_0$ shows that the contribution to the capture rate of the ternary reaction is presumably suppressed in stars  but it can be at the same level as the contribution to the capture rate for the binary reaction or even prevail over it at very high densities. This is natural because at short distances between the particles the factor of the larger effective charge acting on the electron will dominate. The same conclusion can be drawn from Fig.\\,\\ref{Fig3}. In this figure, the solid and dashed curves practically coincide. This again shows a very smooth dependence of \\mbox{$<{\\varsigma}(R_0, T)>$} on the temperature. Let us note that the values of the $R_0$, considered in Fig.\\,\\ref{Fig2}, correspond to rather dense stars. For example, the value of \\mbox{$R_0=10^4$ fm} corresponds to the proton density $\\rho_p=1673$ g/cm$^3$, which is about 11 times larger than in the center of the Sun. Let us further discuss the electron capture by $^{7}$Be solely in the Sun. \\begin{figure}[h!] \\centerline{ \\epsfig{file=fig3f.eps} } \\vspace{0.4cm} \\caption{ The dependence of the mean value \\mbox{$<{\\varsigma}(R_0, T)>$} on the values of the mean distance $R_0$. Solid curve, $kT=1.161$ keV, $\\beta=1/7$, and $Z=5$; dashed and dotted curve, $kT=1.161$ keV, $\\beta=4/7$, and $Z=6$; dashed curve, $kT=1.5$ keV, $\\beta=1/7$, and $Z=5$. The dashed and dotted curve corresponds to analogous calculations for the nuclei $^{4}$He.}\\label{Fig3} \\end{figure}  In Table I, we show the influence of protons on the electron capture in the Sun in more detail. For the Standard Solar Model, we choose Model SSMBP2004 \\cite{b12}. According to Fig.\\,6.1 of Ref.\\,\\cite{b13}, the maximal intensity of the electron capture by the nuclei $^{7}$Be takes place at the distance $R_s/R_\\odot\\approx 0.06$, where $R_\\odot$ is the radius of the Sun, and it drops to one half at $R_s/R_\\odot\\approx 0.03$ and $R_s/R_\\odot\\approx 0.1$. Using the data on the temperature, the density, and the fraction of the hydrogen in this area of the Sun, we obtain the mean values \\mbox{$<{\\varsigma}(R_0, T)>$} and \\mbox{$<{\\varsigma}_C(R_0, T)>$} presented in Table I. In the second column, we add the mean value calculated at $R_s/R_\\odot$ = 0.007 which is close to the very centre of the Sun. It can be seen that the change of the average quantities \\mbox{$<{\\varsigma}(R_0, T)>$} and \\mbox{$<{\\varsigma}_C(R_0, T)>$} is very smooth.  Table I. The mean values \\mbox{$<{\\varsigma}(R_0, T)>$} and \\mbox{$<{\\varsigma}_C(R_0, T)>$} for the electron capture by the nuclei $^{7}$Be in the Sun. \\begin{center} \\begin{tabular}{|l||c|c|c|c|}\\hline % $<{\\varsigma}(R_0, T)>$& 0.0991 & 0.0991 & 0.0965 & 0.0913 \\\\ $<{\\varsigma}_C(R_0, T)>$& 0.0718 & 0.0717 & 0.0696 & 0.0658 \\\\\\hline\\hline $R_s/R_\\odot$& 0.007 & 0.03 & 0.06 & 0.1\\\\ kT (keV) & 1.353 & 1.300 & 1.161 & 1.088 \\\\ $\\rho_p$ (g/cm$^{3}$)& 52.1 & 51.6 & 48.9 & 44.9 \\\\ $R_0 \\times 10^{-4}$ (fm)  & 3.179 & 3.188 & 3.250 & 3.340 \\\\\\hline \\end{tabular} \\end{center} As can be seen from the first row of Table I and from Fig.\\,\\ref{Fig4}, the contribution to the capture rate of the ternary reaction at the Sun is about 10 \\% of the binary one\\footnote{The next term in the expansion of Eq.\\,(\\ref{EXP}) is expected to change it only by $\\approx$ 2 \\%.}. This means that it should increase sensibly the burning out of the nuclei $^{7}$Be in comparison with the binary reaction, thus decreasing the concentration of the nuclei $^{8}$B that appear after the capture of protons by $^{7}$Be.  Comparing the first and the second rows of Table I shows a difference of 3 \\% between the values of \\mbox{$<{\\varsigma}(R_0, T)>$} and \\mbox{$<{\\varsigma}_C(R_0, T)>$}. This variation arises from the difference between the relativistic and nonrelativistic estimations of the electron wave function at zero distance for the binary reaction (\\ref{br2}).   In Fig.\\,\\ref{Fig4} we show the variation of \\mbox{$<{\\varsigma}(R_0, T)>$} for the reaction (\\ref{3a}) for larger intervals of $T$ and $R_0$. \\begin{figure}[h!] \\centerline{ \\epsfig{file=fig4f.eps} } \\vspace{0.4cm} \\caption{ The dependence of the mean value \\mbox{$<{\\varsigma}(R_0, T)>$} on the temperature $T$ and the value of the mean  distance $R_0$. Solid line, $R_0=3.25\\times 10^4$ fm; dashed line, $R_0=2.75\\times 10^4$ fm; dashed and dotted line, $R_0=3.75\\times 10^4$ fm.}\\label{Fig4} \\end{figure}  One can consider in analogy the influence of the nuclei $^{4}$He on the electron capture by $^{7}$Be in the Sun. In this case, $Z=6$, $\\beta=4/7$, and at the radius $R_s/R_\\odot=0.06$ the mean distance between the nuclei $^{4}$He is $R_0=5.34\\times 10^4$ fm. Then one obtains from Eq.\\,(\\ref{BS}) that \\mbox{$<{\\varsigma}(R_0, T)>$}=0.0036, which is about 27 times smaller that the analogous value of \\mbox{$<{\\varsigma}(R_0, T)>$} for the protons given in the fourth column of Table I. Evidently, this influence on the electron capture is negligible. This can also be seen from Fig.\\,\\ref{Fig3}.  The main conclusion following from our calculations is that the three-body process due to the presence of the proton in the vicinity of the nucleus $^{7}$Be results in the capture of the electron by an effective charge $Z=5$ instead of $Z=4$, which is qualitatively the new effect that cannot be simulated by introducing the Debye screening. This effect can increase the rate of the electron capture from the continuum by $^{7}$Be  in the Sun, which will reduce the concentration of the nuclei $^{8}$B that appear after the capture of protons by $^{7}$Be."
    },
    "0606/astro-ph0606165_arXiv.txt": {
        "abstract": "We used $V$ and $I$ CCD  photometry to search for low-surface brightness dwarf galaxies in the central ($<0.5$ h$^{-1}$ Mpc) region of the loose groups NGC 6868 (Telescopium) and NGC 5846, the compact group HCG 42 and the poor cluster IC 4765. We used the parameters given by the exponential profile fit (central surface brightness, scale length and limiting diameter) to identify 80 low surface-brightness dwarf candidates with magnitudes  $17 < V < 22$ mag ($-16.7 > M_{V} > -11.4$ at the assumed distances of the groups) and with colors $V-I<1.5$ mag at the limiting isophote of   25.8 V mag/arcsec$^2$. The selected galaxies have central surface  brightnesses $\\mu_{0}>22.5$ V  mag/arcsec$^{2}$, scale lengths $h>1.5$\\arcsec,  and diameters larger than 1.2 h$^{-1}$  kpc. Twenty of the eighty galaxies  are extended low surface-brightness galaxies that were detected only on smoothed  images, after masking all high surface brightness objects. The  completeness fraction in  the detection of the low surface-brightness dwarf galaxies is $\\sim80$\\% for $V\\la20$  and $22.5<\\mu_{0}<24.5$  V mag/arcsec$^{2}$, and below 50\\% at fainter magnitudes and central surface  brightnesses ($\\mu_{0}>24.5$ V mag/arcsec$^{2}$ and $V>20$ mag). In this last bin, the completeness increases to $\\sim 80$\\% when we search for galaxies in smoothed images instead.  The detected low surface-brightness dwarf galaxies are highly concentrated towards the center of the four groups in the inner 250 h$^{-1}$ kpc. At larger radii, the projected number density is similar to the value found in the control fields. The best fit power-law slope of the surface density  distribution is, on average, $\\beta \\sim -1.5$ ($R < 250$ h$^{-1}$ kpc), in  agreement with the values found for satellites dwarfs around isolated E/S0 galaxies and in X-ray groups. The distribution of the low surface-brightness dwarf galaxies in the $M_{V} - \\mu_{0}$ plane does not show a clear correlation, suggesting that  the correlation noted by other studies could be produced by selection  effects. The low surface-brightness dwarf galaxies follow a well defined  color-magnitude relation, extending for more than ten magnitudes (from bright ellipticals to faint dwarfs). A similar well defined color-magnitude relation from giants to dwarfs is known to be valid for galaxy clusters but it is the first time that it is demonstrated in the sparse environments of groups.  The spectroscopic follow-up shows that only 78 of the 412 galaxies with measured velocities in the fields of these groups are group members. We were able to obtain the radial velocity for 11 of the 80 low-surface brightness  dwarf galaxies. Five of these galaxies are group members while six are in the background of the groups. We infer, then, that more than 90\\% of the 80 low surface-brightness dwarf galaxies in our sample must be bonafide group members. In addition, new structures along the group's lines-of-sights were discovered. These new structures are groups and poor  clusters that extend to $\\sim 0.3$ in redshift space. ",
        "introduction": "Groups are the most common environment of galaxies in the Universe. Most galaxies are associated to these poor systems and, due to their low velocity dispersions, these systems are good laboratories to study the different physical processes that regulate galaxy formation and evolution.  Nearby galaxy groups are characterized by their small numbers of bright galaxies (typically 5--7 galaxies brighter than M$^{*}$). Usually redshifts are known only for the brightest members, and no redshift information is available for the galaxies at fainter magnitudes.  The faint-end of the luminosity distribution is the regime where  the dwarf galaxies dominate. Dwarf galaxies are exceptional laboratories for the study of many physical processes, such as structure formation, galaxy  evolution, star formation, and dark matter distribution. However, the low luminosity and surface brightness of these objects have limited most of the previous studies done so far. Because of their low-luminosities, spectroscopic observations of these objects are extremely difficult and time consuming.  Except for the Local Group and other nearby structures \\citep[e.g. M81, Sculptor, LeoI, ComaI:][1999]{jer98,jer00,tre02,bre98,flint01} the physical characterization of dwarf galaxies and its implication to galaxy formation and evolution is poorly constrained. In the last few years, the advent of a new generation of telescopes and new large-format CCD detectors opened a unique possibility to search and detect the low surface-brightness dwarf galaxy population of nearby systems. It became possible to obtain deep photometry for statistically significant samples of these galaxies in large areas of nearby groups and clusters \\citep[e.g.][]{car01,tre02,tre05}.  This is the second paper of a series dedicated to survey nearby groups with the aim of identifying the population of low surface brightness dwarf galaxies (LSBD) down to  M$_V =-13$ mag. The main goal is to identify the LSBD and to characterize and study the physical properties and the Galaxy Luminosity Function of the groups (the latter will be presented in  a later paper).   The program started with the observation of the Dorado group at  $cz\\sim1200$ km s$^{-1}$ \\citep{car01} and continued with the observation of  other four nearby groups with $1000<cz<4500$ km s$^{-1}$, described here. In this paper we present  the results of the photometric and spectroscopic study of the galaxy population   in NGC 6868, NGC 5846, HCG 42 and IC 4765.  The paper is organized as follows.  In Section 2 we describe the sample. The observations, the data reduction and the estimation of the photometric errors and completeness of the spectroscopic sample is presented in section 3. In section 4 we explain the criteria used to select the LSBD sample, the determination of the completeness fractions in the detection of Local Group-like dwarf  galaxies at the distance of the groups,  and the photometric analysis  of the LSBD galaxies detected. Section 5 is  devoted to analyse the spectroscopic survey, while in section 6 we present the main results. Finally, our conclusions are presented in Section 7.  For all calculations, the following cosmological parameters were assumed: $\\Omega_{M}=0.2$, $\\Omega_{\\Lambda}=0.8$ with a Hubble constant of H$_{0}=75$ km s$^{-1}$ Mpc$^{-1}$. Here $h$ is a dimensionless constant defined as the Hubble constant divided by 75 km s$^{-1}$ Mpc$^{-1}$, i.e. $h=H/75$. ",
        "conclusions": "Despite the fact that the low surface-brightness dwarf galaxies detected in the central region of groups have no membership determination, we have demonstrated the success of our method in discovering a population  of dwarf galaxies with low and very low surface brightnesses that seem to be physically associated to the groups (see section 6.3 and below). Using the parameters given by the exponential fit profile, we were able to characterize the properties of the  low surface-brightness dwarf galaxies in groups using an uniform and well defined criteria. At the 80\\% completeness limit, we were able to detected eighty low surface-brightness dwarf galaxies in the $r<0.5$ h$^{-1}$ Mpc  region of the four groups studied here. These galaxies have central surface  brigthnesses $\\mu_{0}\\la 25$ V mag/arcsec$^{2}$, scale lengths  $h>1.5$ arcsec, and absolute magnitudes $-16.5<M_{V}<-11.5$. The surface brightnesses and scale factors of the LSBD sample are similar to those seen in Local Group dwarf galaxies. Eleven of these galaxies have large sizes and  very low surface-brightness. Many of the large LSBD galaxies  were not  identified in the original images, using the selection criteria  described in section 4.1, but they were detected in the images after we  convolved  them with an exponential profile function to search for  large and diffuse objects. Some of these galaxies are similar in size  and brightness to the LSBD galaxies discovered in Virgo and Fornax  \\citep{imp88,bot91} and are apparently absent in the Local Group.  A possible limitation in our survey is the criteria used to select the sample of LSBD galaxies. The criteria used have been optimized to search for objects that are similar in magnitudes, sizes and surface brightnesses to those dwarf galaxies present in the  Local Group. Compact, high surface  brightness galaxies like M32 and very compact, high surface brightness, but faint, objects like UCD galaxies discovered in  Virgo \\citep{jones06} and Fornax clusters \\citep{drink98} are virtually  missing from our sample survey. UCD and M32-like galaxies are very compact and in general they are classified as stars by SExtractor. We did not detect such objects due  to the limitation imposed by the seeing and by the CLASS\\_STAR limit of  0.35 used to select the LSBD galaxies. Except for the Local group, M32-like galaxies are very rare and are nearly  absent in other environments. If these type of galaxies are present in  our sample, their numbers should be very small. In the case of UCD galaxies the  situation might be different. These galaxies are not seen in the Local  Group. However, as mentioned in section 6.1, a large number of these galaxies  were discovered in nearby clusters. These objects are  apparently absent in groups, although searches for UCD in poor  environments have just started \\citep{mieske06}. We cannot discard  that  a population of UCD galaxies could be present in our groups.  This is a pending task that requires deep, high resolution observations to search for these types of objects.  Monte Carlo simulations performed to derive the completeness fractions of LSBD detection on real images were also used to estimate the photometric  errors. At $\\mu_{0}\\sim 25.5$ mag/arcsec$^{2}$ (our limit in central surface  brightness), the completeness in the detection is $\\sim 80$\\%, with an error  in the magnitude and in the central surface brightness $<0.3$ mag and $<0.5$  V mag/arcsec$^{2}$ respectively. One important result given by the simulation is our ability to determine correctly the scale factors of the galaxies. The results show that the scale factors for galaxies with $22.5<\\mu_{0}<24.5$  mag/arcsec$^{2}$ and  $V<19.0$ mag are overestimated in $\\sim30$\\%. At fainter  magnitudes and central surface brightnesses (the last bins in our simulation),  the overestimation is $\\sim 60$\\%. This limitation in estimating correctly the  scale factor is stronger at fainter magnitudes and could affect the correct determination of the faint-end of the Galaxy Luminosity Function by adding to the sample an undesired population of background galaxies.  In clusters (e.g. Coma and Fornax) there exists a tight  color-magnitude relation for the early-type galaxies that extends to the  dwarf regime. We demonstrate, as far as we know for the first time, that this relation is also present in groups (see Fig. \\ref{fig11}), extending  the relation in more than 10 magnitudes (from brigth early-type to faint dwarf spheroidal galaxies). We also found that the LSBD galaxies do not show a clear  correlation between magnitude and surface brightness, in contrast  of what is reported in other  recent studies, perhaps because the apparent  observed correlation seen  in other environments could be produced by selection effects. As we discussed in section 6.1, extended, very low-surface brightness galaxies could be absent from most of the surveys in groups and  cluster due to a limitation imposed by the sky brightness. In fact, we were  able to detect a number of large very low surface-brightness galaxies only  after we convolved the images with an exponential profile function that  enhanced the large, low surface-brightnes  features in the images.  The LSBD galaxies are clearly concentrated around the brightest group members of the  groups, in a scale of  $\\sim 250$ h$^{-1}$ kpc  (see Fig. \\ref{fig12}). At larger radius, the number density is similar  to the values found in our control fields. A similar correlation is  seen in other low-density environments \\citep{vad91,car01}, where bright early-type galaxies  are predominant. The central concentration detected suggests that these galaxies are true dwarf galaxies of the groups. The slopes obtained from the fit to the number  density profile  of LSBD galaxies agree with the values derived for satellites dwarfs of isolated bright galaxies \\citep{vad91,chen05}. Furthermore, the distribution of satellites around the centers of the groups are similar  in extent to that seen in some fossil groups \\citep[e.g. NGC 1132][]{mul99}.  Only 78 ($\\sim 19$\\%) of the 412 galaxies with radial velocities were confirmed as members of the groups. In addition, several new background structures were identified along the  line-of-sight of the groups. In the direction of HCG 42 we identified  one  cluster and one group of galaxies (at $z=0.18$ and $z=0.25$).  In the direction  of IC 4765 we identified one cluster ($z\\sim0.04$) and towards NGC 6868, only  one poor group ($z\\sim0.016$) was identified.  We have attempted to measure the radial velocities of the majority of the 80 LSBD galaxies of our list, but only eleven spectra had enough S/N for a realiable determination  of their radial velocities. Only five LSBD galaxies turned out to be members of the groups. Of the six background galaxies, two are in the direction of NGC 6868 group and they are genuine dwarf galaxies which belong to a background structure located only 2000 km s$^{-1}$ of NGC 6868. For the  remaining four galaxies, two of them have colors $V-I>1.3$, two are blue and all show emission lines. It is not surprising that only few galaxies  of the LSBD catalog, for which we obtained spectroscopy, are members of the groups, given that $\\sim 75$\\% of our sample have morphological types dE/dSph and therefore could not be observed due to the low S/N absorption  spectra they present. About 25\\% of the LSBD sample are dwarf galaxies that show an irregular morphology and there is indication that they are forming stars. The four background emission-lines galaxies belong to  this sub-sample and they comprise only the 30\\% of the total population  of LSBD galaxies  classified as dIrr. Hence, we estimated that only a small fraction of our sample ($\\sim 10$\\% or less) could be contaminated with background galaxies.  The dwarf galaxies ($M_{V}>-17.5)$ in the groups studied follow the same morphology-radius relation seen in much denser environments (e.g. Virgo).  The fraction of red ($V-I>0.95$) dwarfs (dE, dE,N, dSph) decrease towards the  outskirts of the groups, while the fraction of blue  ($V-I<0.95$), star  forming dwarfs (dIrr) increase with cluster-centric distance. The same  trend is found for low-mass systems in SDSS galaxy groups analyzed by   \\citet{weinmann06}.  In this work we presented a searche for low surface-brightness dwarf galaxies in poor environments at distances $>$10 h$^{-1}$ Mpc. We showed here that the technique  used to search for LSBD galaxies is reliable and robust if we can derive correctly all errors associated to the detection and photometry. A good knowledge of the errors and the selection effects may give us the possibility to construct a complete, well defined sample of low-luminosity, low surface-brightness galaxies beyond the Local Group, and characterize their properties and intrinsic nature which is an important clue for testing the models of galaxy formation, in particular, at  low-mass levels."
    },
    "0606/astro-ph0606486_arXiv.txt": {
        "abstract": "s{ The \\HESS\\  Imaging Atmospheric Cherenkov Telescope Array is currently the most sensitive instrument for Very High Energy (VHE) $\\gamma$-ray observations in the energy range of about 0.1-10~TeV. During more than two years of operation with the complete 4-telescope array, many galactic and extragalactic VHE $\\gamma$-ray sources have been discovered. With its superior sensitivity and its large field-of-view camera, \\HESS\\  is particularly suited for surveys and detailed studies of extended sources. A selection of recent \\HESS\\  results is presented in this proceeding.  } ",
        "introduction": " ",
        "conclusions": "The \\HESS\\  experiment has been fully operated since December 2003. For the past two years it has run with full performance and produced numerous exciting new results in VHE $\\gamma$-ray observations. \\HESS\\  showes for the first time that the latest generation of ground-based Cherenkov telescope arrays has reached a sensitivity where real $\\gamma$-ray astronomy is feasible. The sensitivity of \\HESS, combined with its good angular resolution, provides an excellent handle to study and understand the physics processes that take place in galactic and extragalactic sources of VHE $\\gamma$-rays.\\\\ Starting in 2008, H.E.S.S. II will explore the $\\gamma$-ray sky with even higher sensitivity and at least two-times lower energy threshold. This will provide a comfortable overlap in energy with the upcoming GLAST \\cite{GLAST} satellite experiment."
    },
    "0606/astro-ph0606603_arXiv.txt": {
        "abstract": "The evolution of central stars of planetary nebulae can proceed in several distinct ways, either leading to H-deficiency or to H-normal surface composition. Several new simulations of the evolution channels that lead to H-deficiency are now available, mainly the born-again scenarios that are triggered by a He-shell flash during the hot pre-white dwarf evolution phase. A realistic AGB progenitor evolution is important for correct HRD tracks, that allow mass determinations. New hydrodynamic simulations of He-shell flash convection including cases with H-ingestion are now performed, and allow a determination of the convective extra-mixing efficiency. This has direct consequences for the intershell abundance distribution of AGB stars that can be observed in the H-deficient CSPN. ",
        "introduction": "The evolution of the central stars of planetary nebulae (CSPN) is closely connected to the evolution of the progenitor phase, the Asymptotic Giant Branch (AGB) stars, and has observable effects for the CSPN progeny, the WD stars. The AGB evolution phase is the result of previous main-sequence and horizontal-branch evolution of low- and intermediate mass stars. An updated classification of mass ranges involving AGB stars has been recently suggested by \\citet[][Fig.\\, 2]{herwig:04c}. Accordingly, low-mass stars are those below an initial mass of $\\approx 1.8\\msun$ that ignite a He-core flash at the tip of the RGB, while intermediate mass stars are those initially more massive than $\\approx 1.8\\msun$ but less massive than required to explode as a core-collapse supernova. The dividing mass at $\\approx 10\\msun$ separates those super-AGB stars that form ONe-white dwarfs and those that ignite as core collapse supernova \\citep[e.g.][]{garcia-berro:94}. However, quantitatively the fraction of AGB stars that explode as core-collapse supernova is still very uncertain, largely due to uncertainties in super-AGB mass loss, and the physics of mixing and burning, e.g. during dredge-up and hot-bottom burning (see below). For AGB stars it is useful to use the limiting mass for hot-bottom burning (HBB) as a dividing line for low-mass and massive AGB stars. Hot-bottom burning transforms dredged up C into N and prevents the formation of C-rich composition \\citep[e.g.][]{scalo:75,lattanzio:92b}. HBB has therefore an important influence on the evolution of AGB star, and the composition of the ejecta that later form the planetary nebula. AGB stars that are massive enough to ignite carbon burning are the super-AGB stars.  Conceptually, the evolution of AGB stars is well understood, at least at solar and modest metal deficiency, like in the Magellanic Clouds. The review by \\citet{iben:83b} describes the basic properties of AGB stars: thermal pulses, the necessity of the third dredge-up, the effect of mass loss and the potential for important nucleosynthesis. More recent reviews by \\citet{lattanzio:97} and \\citet{bloecker:99b} include accounts of hot-bottom burning, and an improved quantitative description of AGB evolution. Dedicated reviews were written on mass loss \\citep{wilson:00}, the \\sprn\\ \\citep{busso:99,meyer:94} and the cool post-AGB stars that allow important observations concerning the \\sprn\\ \\citep{vanwinckel:03}.  \\begin{figure} \\begin{center} \\includegraphics[width=10.5cm]{herwigf1.ps}  % \\caption{ Comparison of three different evolution channels in the $g$--\\teff--plane for a $0.604\\msun$ post-AGB star that had a main-sequence mass of $2\\msun$ and evolved through about a dozen thermal pulses with dredge-up. This figure originally appeared in the Publications of the Astronomical Society of the Pacific \\citep{werner:06}.  Copyright 2005, Astronomical Society of the Pacific; reproduced with permission of the Editors. }\\label{fig:fig_gteff_herwig} \\end{center} \\end{figure} Unfortunately, major unresolved issues remain, that severely limit the predictive quality of current AGB modeling \\citep{herwig:04c}. The issues concern the mass-loss history, the nature and importance of hydrodynamic mixing, including the physics of mixing at the bottom of the envelope convection zone that seems to be critical for the third dredge-up, and for formation of the assumed \\cdr-pocket. Another problem is the rather poorly known effect of rotation and magnetic fields on the evolution, nucleosynthesis and formation of asymmetries of bi-polar structure in proto-planetary nebulae. All these problems are amplified for stars at very low metallicity, partly because less observational checks are available.  The evolution of CSPN (described in \\kap{sec:mod-cspn}) depends sensitively on the progenitor AGB evolution. In particular, the observed abundances of the H-deficient CSPN are linked to mixing processes in the He-shell flash of the AGB phase, and some new results in this area are presented in \\kap{sec:hydro}. ",
        "conclusions": "Among the CSPN channels, the evolution leading to H-deficient surface composition is of great importance for constraining nucleosynthesis and mixing in the AGB progenitors. A clear correlation exists, e.g. between the mixing efficiency at the bottom of the He-shell flash convection and the O abundance in the intershell that is subsequently observed in H-deficient CSPN. AGB calculations with parameterized overshoot calculations by \\citet{herwig:99a} showed that more convective extra-mixing leads to higher O abundance, and that an exponential decay parameter for extra-mixing (see \\citet{herwig:99a} for the definition of) $f \\approx 0.01$ would be consistent with CSPN observed O abundances in the range $8 - 20\\%$ by mass. Our preliminary analysis of hydrodynamic simulations of the He-shell flash convection zone quantitatively agrees with this level of convective extra-mixing. We regard this agreement of hydrodynamic simulation and astrophysical observation as a successful validation of the computational simulations, and feel encouraged to continue this work towards more complex simulations of convective-reactive mixing."
    },
    "0606/astro-ph0606435_arXiv.txt": {
        "abstract": "We analyze the processes relevant for star formation in a model with dark matter in the form of sterile neutrinos.   Sterile neutrino decays produce an x-ray background radiation that has a two-fold effect on the collapsing clouds of hydrogen.  First, the x-rays ionize the gas and cause an increase in the fraction of molecular hydrogen, which makes it easier for the gas to cool and to form stars.  Second, the same x-rays deposit a certain amount of heat, which could, in principle, thwart the cooling of gas.  We find that, in all the cases we have examined, the overall effect of sterile dark matter is to facilitate the cooling of gas.  Hence, we conclude that dark matter in the form of sterile neutrinos can help the early collapse of gas clouds and the subsequent star formation. ",
        "introduction": "\\label{sec_introduction}  Numerous observations, from galaxy rotation curves to gravitational lensing, to cosmic microwave background radiation, all point to the existence of dark matter which is not made of ordinary atoms, but, rather, of some new, yet undiscovered particles.  In particular, the abundance of dark matter has been precisely determined recently by the WMAP~\\citep{spergel,spergel06}.  However, the nature of dark-matter particles remains unknown.  One attractive candidate is a sterile neutrino with mass of several keV and a small mixing with the ordinary neutrinos.  If such a particle exists, it could be produced in the early universe with the right abundance to be dark matter either from neutrino oscillations~\\citep{Dodelson,shi,abaz01a,abaz01b,abaz05,Dolgov02}, or from some other processes, for example from inflaton decay~\\citep{Shaposhnikov:2006xi}. The same particle can explain the observed velocities of pulsars because its emission from a cooling neutron star would be anisotropic, hence providing the neutron star with a recoil velocity of a sufficient magnitude \\citep{Kusenko97,b11,Kusenko98,fuller,barkovich,2005hep.ph....3113B,Kusenko04}. In addition, the neutrino kick can enhance the convection during the first second of the supernova explosion, which can increase the energy of the supernova shock and bring the supernova calculations in better agreement with the observations \\citep{fryer}. Sterile neutrinos can also help the formation of supermassive black holes in the early universe \\citep{mun,bb}.  Although the dark-matter sterile neutrinos have interactions that are too weak to be discovered in the laboratory,  some other sterile neutrinos may exist. Most models of neutrino masses introduce sterile (or right-handed) neutrinos to generate masses of the ordinary neutrinos via the seesaw mechanism~\\citep{seesaw1,seesaw2,seesaw3}.  The name {\\em sterile} was coined by~\\cite{pontecorvo}.  Many seesaw models assume that sterile neutrinos have very large masses, which makes them unobservable. However, a number of theoretical models predict the existence of sterile neutrinos at or below the electroweak scale~\\citep{smirnov,shrock,asakaa,deGouvea}. There is an indication that some light sterile mass eigenstates may cause the neutrino flavor transformation seen in the Los Alamos Liquid Scintillator Neutrino Detector (LSND) experiment \\citep{b40,sorel,deGouvea}. The results of LSND are being tested by the MiniBooNE experiment. One can expect that there are at least three light sterile neutrinos \\citep{asakab}, in which case the neutrino oscillations can explain the baryon asymmetry of the universe \\citep{akhmedov,asakaa}.   Although dark-matter sterile neutrinos are stable on cosmological time scales, they nevertheless decay. The mass matrix is assumed to have the seesaw form with the Majorana mass term of the order of several keV and with a much smaller Dirac mass term.  The diagonalization of such a mass matrix yields the light mostly active Majorana mass eigenstates and the heavy sterile Majorana eigenstates, which can decay into the light ones.  The dominant decay mode, into three light neutrinos, is \"invisible\" because the daughter neutrinos are beyond the detection capabilities of today's experiments.  The most prominent \"visible\" mode is decay into one active neutrino and one photon, $\\nu_s \\rightarrow \\nu_a \\gamma$.  Assuming two-neutrino mixing for simplicity, one can express the inverse width of such a decay as \\begin{equation} \\tau \\equiv \\Gamma_{\\nu_s \\rightarrow \\nu_a \\gamma}^{-1}= 1.3 \\times 10^{26} \\rm{ s } \\left(\\frac{7\\ \\rm keV}{m_{s} }\\right)^{5} \\left(\\frac{0.8 \\times 10^{-9}}{\\sin^{2} \\theta}\\right), \\label{meanlife} \\end{equation} where $m_s$ is the mass and $\\theta $ is the mixing angle.  For mass units we set $c=1$.  Since this is a two-body decay, the photon energy is half the mass of the sterile neutrino.  The monochromatic line from dark matter decays can, in principle, be observed by x-ray telescopes.  No such observation has been reported, and some important limits have been derived on the allowed masses and mixing angles~\\citep{abaz01a,abaz01b, Boyarsky,2006hep.ph....1098B,2006astro.ph..3368B,2006astro.ph..3660B, Riemer-Sorensen:2006fh,2006PhRvD..73f3513A,Watson:2006qb}.   These constraints are based on different astrophysical objects, from Virgo and Coma clusters, to Large Magellanic Clouds, to the Milky Way halo and its components, etc.  There are different uncertainties in modeling the dark matter populations in these objects. Different groups have also used very different methodologies in deriving these bounds: from a conservative assumption that the dark-matter line should not exceed the signal, to more ambitious approaches that involved modeling the signal or merely fitting it with a smooth curve and requiring that the line-shaped addition not affect the quality of the fit.  The x-ray limits are usually shown based on the assumption that the sterile neutrinos constitute the entire dark matter ($\\Omega_s=\\Omega_{dm}$, where $s$ and $dm$ refer to the sterile neutrinos and the dark matter, respectively). However, the x-ray observations exclude certain masses and mixing angles even if the sterile neutrinos make up only a fraction of dark matter.  The production via mixing \\citep{Dodelson} provides the lowest possible abundance of sterile neutrinos. In Fig.~\\ref{fig:models} we show both the bounds based on the assumption $\\Omega_s=\\Omega_{dm}$ and the model-independent exclusion region~\\citep{Kusenko06} based on the production only via the \\cite{Dodelson} mechanism.  \\cite{Kusenko06} used the analytical fit to the numerical calculation of \\cite{abaz05}, and re-scaled the x-ray flux in proportion to the amount of relic neutrinos produced for a given mass and mixing angle.   A  different set of bounds comes from the observations of the Lyman-alpha forest~\\citep{viel1,viel2,seljak}, which limit the sterile neutrino mass from below. Based on the high-redshift data from SDSS and some modeling of gas dynamics, one can set a limit as strong as 14~keV~\\citep{seljak}. However, the high-redshift data have systematic errors that are not well understood~\\citep{jena}, and more conservative approaches, based on the relatively low-redshift data, have led to some less stringent bounds~\\citep{viel1}.  More recently, \\cite{viel2} have reanalyzed the high-redshift data and arrived at the bound $m_s>10$~keV.  The mass bound as quoted depends on the production mechanism in the early universe.  The Lyman-alpha observations constrain the free-streaming lengths of dark matter particles, not their masses.  For each cosmological production mechanism, the relation between the free-streaming length and the mass is different~\\citep{Hansen_et_al}.  For example, the bound $m_s >10$~keV~\\citep{viel2} applies to the production model due to~\\cite{Dodelson}.  Even within a given cosmological scenario, there are uncertainties in the production rates of neutrinos for any given mass and mixing angle~\\citep{Asaka:2006rw}.  These uncertainties may further affect the interpretation of the Lyman-alpha bounds in terms of the sterile neutrino mass. The x-ray bounds depend on a combination of mass and mixing angle as in eq.~(\\ref{meanlife}), regardless of one's cosmological assumptions, as long as all the dark matter is comprised of sterile neutrinos. However, the x-ray bounds are subject to uncertainties in modeling the dark matter populations of different objects.  If the sterile neutrinos make up only a part of dark matter, none of the above bounds apply.  Also, if inflation ended with a low reheating temperature, the bounds are significantly weaker~\\citep{b22}.  It should also be mentioned that the Lyman-alpha bounds appear to contradict the observations of dwarf  spheroidal galaxies~\\citep{Wilkinson,gilmore_06,strigari}, which  suggest that dark matter is warm and which would favor the 1--5~keV mass range for sterile neutrinos. There are several inconsistencies between the predictions of N-body simulations of cold dark matter (CDM) and the observations \\citep{willman,bode,Peebles,Dalcanton,bosch,zentner,abaz05b,Dolgov01, b38,b39,halo_shape,halo_shape2,halo_shape_cdm}. Each of these problems may find a separate independent solution.   Perhaps, a better understanding of CDM on small scales will resolve these discrepancies. It is true, however, that warm dark matter in the form of sterile neutrinos is free from all these small-scale problems altogether, while on large scales WDM fits the data as well as CDM.   Several independent considerations, from the pulsar kicks to small-scale behavior of dark matter, favor sterile neutrinos as the dark matter candidate. It is important, therefore, to examine the effects of sterile-neutrino dark matter on star formation and reionization of the universe. \\citet{yoshida03} have attempted to study such effects for some kind of warm dark matter, but they have used a linear matter power spectrum and assumed that the dark matter particles have a thermal distribution, which is not the case for sterile neutrinos \\citep{abaz05}. Even more importantly, they have not considered the effects of sterile neutrino decays and the resulting x-ray background. It was shown by \\citet{biermann06} that the x-rays produced by the decays of relic sterile neutrinos can speed up the formation of molecular hydrogen, which can help the cooling of the gas and the star formation, which can lead, in turn, to a reionization of the universe at the redshift consistent with the WMAP results.  More specifically, the WMAP (three years) measurement of the reionization redshift $z_r=10.9^{+2.7}_{-2.3}$ \\citep{spergel06} has posed a challenge to theories of star formation. On the one hand, stars have to form early enough to reionize gas at redshift~11. On the other hand, the spectra of bright distant quasars imply that reionization must be completed by redshift~6.  Stars form in clouds of hydrogen which collapse at different times, depending on their sizes: the small clouds collapse first, while the large ones collapse last.  If the big clouds must collapse by redshift~6, then the small halos must  undergo the collapse much earlier.  It appears that the star formation in these small halos would have occurred at high redshift, when the gas density was very high, and it would have resulted in an excessive Thompson optical depth \\citep{haiman_bryan}. To be consistent with WMAP, the efficiency for the production of ionizing photons in minihalos must have been at least an order of magnitude lower than expected~\\citep{haiman_bryan}. One solution is to suppress the star formation rate in small halos by some dynamical feedback mechanism. The suppression required is by at least an order of magnitude.  Alternatively, one can consider warm dark matter, in which case the number of small clouds is strongly suppressed.  However, for some WDM candidates (for example, the gravitino), the lack of power on small scales in the spectrum of density perturbations leads to an unacceptable delay in the onset of star formation \\citep{yoshida03}. This is not the case for sterile neutrinos, which decay slowly during the \"dark ages\" preceding the birth of stars.  The decay photons could ionize gas, and, although this ionization is too weak to affect the WMAP measurements directly~\\citep{mapelli1,mapelli2}, the increase in the ionization fraction could cause a rapid production of molecular hydrogen at redshifts as high as 100~\\citep{biermann06}. Molecular hydrogen is the essential cooling agent that allows the gas to cool and collapse to form the first stars.  \\citet{biermann06} have not considered the effects of gas heating associated with the decays, which may, in some cases,  thwart the effects of ionization.   The increase in the fraction of molecular hydrogen can help cooling the gas, but if the heating from sterile neutrino decays is significant, it can overcome the radiative cooling and stymie the collapse and star formation.   In this paper we will examine thermal evolution of gas clouds in detail, taking into account both effects of the sterile neutrino decays, namely the ionization and the heating of gas.  The paper is organized as follows. In Section \\ref{sec_H2} we will discuss the role of sterile neutrino decays and the production of molecular hydrogen in the cooling of the primordial gas clouds. Next, in Sections 3-6 we will briefly describe the top-hat overdensity simulation. The results are shown in Section~\\ref{sec_results}. ",
        "conclusions": "\\label{sec_conclusion}  We have performed a detailed analysis of the cooling and collapse of primordial gas in the model with warm dark matter, taking into account both the increase in the fraction of molecular hydrogen~\\citep{biermann06} and the heating due to the sterile neutrino decay.  As expected, the effect on the largest gas clouds is negligible, but the smaller clouds are, in fact affected.  We have performed the analysis for some benchmark cases which arise in  realistic scenarios.  For the largest clouds, the additional molecular hydrogen made no difference.  For smaller clouds, especially for those with masses $10^5-10^6~M_\\odot$, the increase in the x-ray background made the successful collapse possible in cases where it could not occur in the absence of sterile neutrino decays."
    },
    "0606/physics0606122_arXiv.txt": {
        "abstract": "We consider the magnetorotational instability (MRI) of a hydrodynamically stable Taylor-Couette flow with a helical external magnetic field in the inductionless approximation defined by a zero magnetic Prandtl number ($\\Pm=0)$. This leads to a considerable simplification of the problem eventually containing only hydrodynamic variables. First, we point out that the energy of any perturbation growing in the presence of magnetic field has to grow faster without the field. This is a paradox because the base flow is stable without the magnetic while it is unstable in the presence of a helical magnetic field without being modified by the latter as it has been found recently by Hollerbach and R\\\"{u}diger {[}Phys. Rev. Lett. \\textbf{95}, 124501 (2005){]}. We revisit this problem by using a Chebyshev collocation method to calculate the eigenvalue spectrum of the linearized problem. In this way, we confirm that MRI with helical magnetic field indeed works in the inductionless limit where the destabilization effect appears as an effective shift of the Rayleigh line. Second, we integrate the linearized equations in time to study the transient behavior of small amplitude perturbations, thus showing that the energy arguments are correct as well. However, there is no real contradiction between both facts. The linear stability theory predicts the asymptotic development of an arbitrary small-amplitude perturbation, while the energy stability theory yields the instant growth rate of any particular perturbation, but it does not account for the evolution of this perturbation. Thus, although switching off the magnetic field instantly increases the energy growth rate, in the same time the critical perturbation ceases to be an eigenmode without the magnetic field. Consequently, this perturbation is transformed with time and so looses its ability to extract energy from the base flow necessary for the growth. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606059_arXiv.txt": {
        "abstract": "We report on the detection of localized off-center emission at 1--4~AU in the circumstellar environment of the young stellar object AB~Aurigae. We used closure phase measurements in the near-infrared made at the long baseline interferometer IOTA, the first obtained on a young stellar object using this technique. When probing sub~AU scales, all closure phases are close to zero degrees, as expected given the previously-determined size of the AB~Aurigae inner dust disk. However, a clear closure phase signal of $-3.5 \\pm 0.5 \\degr$ is detected on one triangle containing relatively short baselines, requiring a high degree of non-point symmetry from emission at larger (AU-sized) scales in the disk.  We have not identified any alternative explanation for these closure phase results and demonstrate that a ``disk hot spot'' model can fit our data.  We speculate that such asymmetric near-infrared emission detected might arise as a result of localized viscous heating due to a gravitational instability in the AB Aurigae disk, or to the presence of a close stellar companion or accreting sub-stellar object. ",
        "introduction": "\\label{intro_sec}  AB~Aurigae (A0~Ve, $d=144$~pc, V~=~7.1, H~=~5.9) is often referred to as the prototype of the Herbig~Ae/Be (HAeBe) class of pre-main sequence stars of intermediate mass. Like their solar-type analogs, the T~Tauri objects, HAeBe objects are known to be surrounded by pre-planetary disks of gas and dust.  AB~Aur was one of the first young stellar objects (YSO) to be spatially resolved in the near-infrared (NIR) \\citep{rmg1999}, using the technique of long baseline optical interferometry. These and subsequent observations of a relatively large sample of HAeBe (and T~Tauri) objects found that essentially all objects have characteristic NIR sizes much larger than expected from previous-generation disk models that were tuned to fit the spectro-photometric data alone \\citep{tuthill1999, rla2000, rmg2001, rla2002, eisner2004, rla2005, eisner2005, jdm2005}.  These results have prompted in part a significant revision of disk models, whereby previously ignored detailed physics of the disk inner edge are now incorporated with some success \\citep{natta2001, dullemond2001, jdm2002, muzerolle2003,isella2005}. For a recent review of these developments the reader is also referred to \\citet{rmg2006}.  A common characteristic of all the interferometer studies mentioned is that they were based on single-baseline data, and therefore depended on the interpretation of only visibility amplitudes and with very sparse spatial frequency coverage.  One of the unexplored frontiers in optical long baseline interferometry, particularly in the area of YSO science, is the measurement and interpretation of not only visibility amplitudes, but also closure phases (CP), obtained at arrays containing three or more telescopes. Contrary to the visiblity amplitudes or phases measured for baseline pairs, the CP is uncorrupted by atmospheric turbulence \\citep{jennison1958,monnier2003}, and can be calibrated to higher precision. The CP can have values different from 0 or $180 \\degr$ only for a non-point-symmetric sky brightness, and is therefore a highly effective probe of this type of morphology.  Moreover, CP and visibility amplitude data may be used together to provide powerful constraints to parametric models describing relatively complex morphologies. In principle, given sufficient spatial frequency coverage, direct image reconstruction becomes possible, as is commonly done at radio wavelengths \\citep{readhead1980,cornwell1981}.  The new models of HAe objects, and of AB~Aur in particular, successfully reproduce the spectral features and broadband spectral energy distribution (SED) with a model in which no or little optical depth exists between the central star and the dust disk inner edge, such that the frontal stellar heating forms a hot wall of emission from which essentially all the NIR (disk) flux arises.  This model is also successful at reproducing the characteristic NIR sizes measured interferometrically.  Moreover, the outer regions of circumstellar disks are believed to have scale heights that increase with radius \\citep[flaring,][]{kenyon1987}. Thus, non-point-symmetric brightness distributions and non-zero CPs are expected for star-disk systems such as AB~Aur if they have some inclination to the observer's line of sight, as the cooler foreground disk regions occult part of the hotter inner disk regions \\citep[see also][]{tuthill1999, malbet2001, isella2005}. The original goal of the work presented here, was to search for such {\\it skewed} emission coming from the inner dust disk; and in a forthcoming publication \\citep{jdm2006} we discuss in detail the results of our survey of 14 YSOs in the context of hot inner-dust wall models. In this Letter, we report instead on the unexpected detection of a strong CP signal that originates further out in the AB~Aur disk. ",
        "conclusions": "The principal result of this Letter is the clear detection of an unexpected localized asymmetry in the circumstellar environment of AB~Aur at spatial scales of 1--4~AU; this is an {\\it intermediate} scale between the the inner wall containing the hottest dust and the outer disk.  We first considered that the physical origin of this emission could be H-band scattering off the disk atmosphere or ``halo''.  The required asymmetry could easily arise from a small photocenter shift due to e.g., forward scattering in combination with a small disk inclination. However, this emission would arise from large angular scales (up to $\\sim 1 \\arcsec$) and would be completely resolved by even the shortest IOTA baselines, and therefore could not result in a CP signal.  Indeed, the model fitting described above shows that the offset Gaussian must be relatively compact ($\\lesssim 12$~mas or 1.7~AU), and it is difficult to reconcile the required amount of flux with scattering over such a small area.  Due to the low inclination known for the AB~Aur disk, we also do not favor outer disk occultation or opacity effects.  A more likely explanation, is that the asymmetry is due to compact thermal emission. The presence and persistence of such a local inhomogeneity in the AB~Aur disk at these spatial scales would however be surprising, given that the dynamical time-scales are very short (1--5~years) compared to the age of the system ($\\sim 1-4$~Myr).  A possible explanation is that the disk is not steady state, and the feature corresponds to a gravitational instability, possibly associated with an inner extension of the spiral density waves seen in near-IR coronographic and millimiter interferometry images \\citep{fukagawa2004,pietu2005}.  The existence of a disk instability at few~AU radii and for a disk of this age clearly has important implications for the formation of planets in these disks.  Alternatively, we could be detecting a companion embedded in the AB~Aur disk.\\footnote{We note that \\citet{baines2006} report ``spectro-astrometric'' evidence for binarity in AB~Aur (and several other HAeBe objects). However, they infer large separations of $0.5 - 3.0 \\arcsec$; therefore this putative companion can not be related to the CP signatures modelled here.} A stellar companion at the few--AU separation we infer is however unlikely, given that a substantial circumstellar disk has in fact survived around AB~Aur.  The intriguing possibility therefore arises that the companion is instead a proto-planetary object forming in the disk.  Similar claims (but based on visibility amplitudes rather than the more robust CP) for a disk thermal instability or close companion have been recently made for the active-disk system FU~Orionis \\citep{malbet2005}. We also note the possible connection between the extended structure detected here, and the intriguing compact halo (45~mas FWHM) seen in the T~Tauri object DG~Tau in lunar occultation observations by \\citep{leinert1991}. As discussed by \\citet{jdm2006}, disk ``halos'' appear to be common (although not the norm) in young star~$+$~disk systems, and deserve new observational and modelling scrutiny.  Our hypothesis is readily testable. If the non-zero CP arises in a ``disk hot spot,'' variability in the CP signal over 1--5~years should reveal the orbital motion that could be followed using short or long baselines from the VLTI AMBER \\citep{malbet2004} or MIDI \\citep{leinert2003} instruments, respectively.  The angular scales involved also make this object a prime target for direct imaging using high resolution techniques on large single-aperture telescopes."
    },
    "0606/astro-ph0606573_arXiv.txt": {
        "abstract": "We present abundance measurements for two super Lyman Limit systems (SLLS; quasar absorption line systems with $10^{19} \\cm{-2} < \\mnhi < 10^{20.3} \\cm{-2}$) selected from a set of metal-strong absorbers in the Sloan Digital Sky Survey quasar database.  After applying estimate corrections for photoionization effects, we derive gas-phase metallicities of $\\lbrack$M/H$\\rbrack$=$+0.7 \\pm 0.2$\\,dex for the SLLS at $z=1.7749$ toward \\qlla\\ and $\\lbrack$M/H$\\rbrack$=$+0.05 \\pm 0.1$\\,dex for the SLLS at $z=1.7678$ toward \\qllb. The former exhibits among the highest gas metallicity of any astrophysical environment and its total metal surface density exceeds that of nearly every known damped \\lya\\ system. The properties of these absorbers -- high metallicity and large velocity width ($\\delv > 300 \\mkms$) -- resemble those of gas observed in absorption in the spectra of bright, star-forming galaxies at high redshift. We discuss the metal mass density of the SLLS based on these observations and our ongoing SLLS survey and argue that a conservative estimate to the total metal budget at $z=2$ is greater than $15\\%$ of the total, suggesting that the metal-rich LLS may represent the dominant metal reservoir in the young universe. ",
        "introduction": "With the recent successes of high redshift galaxy surveys, the star formation history of the young universe is revealing itself \\citep{madau96,goods04}.  Although debate continues on various aspects of the measurements (e.g.\\ dust extinction, sample selection, sample variance), it is evident that the star formation rate at $z>2$ was substantially higher than the current epoch. Accordingly, the massive stars which light up these galaxies must have produced copious metals and redistributed them via supernovae to the galaxy and surrounding medium.  A valuable test of this picture, therefore, is to obtain an accurate census of metals at $z \\approx 2$ and compare against the amount predicted by the integrated the star formation history \\citep{pagel02,pettini04,bouche05I}. In the following, we will adopt a metal mass density at $z=2$ of $\\Omega_m^{Total} = 3 \\sci{-5} h_{70}^{-2}$ based on the work of \\cite{bouche06II}.  Prior to these galaxy surveys, analysis of heavy metals in quasar absorption line (QAL) systems -- the gas within and in between high $z$ galaxies -- demonstrated that star formation was ubiquitous at $z>2$ \\citep{tytler95,wolfe94,pettini94}. Metals are present in excess of the primordial value throughout much of the universe \\citep{schaye03,simcoe04,bergeron05} and a direct census of the intergalactic medium (IGM) adds up to $5-15\\%$ of \\omt \\citep{pettini04,schaye03}. Regarding the interstellar medium of galaxies, the damped \\lya\\ systems (QAL systems with $\\mnhi \\geq 2\\sci{20} \\cm{-2}$) offer the most direct means of measurement \\citep{pettini94,pro_mtl03}. Surveys of these absorption systems have demonstrated that the metals associated with \\ion{H}{1} gas in DLAs contribute as comparable a fraction as the IGM probed through the \\lya\\ forest. The metals in stars can be estimated from the luminosity functions and metallicity estimates of galaxy surveys; \\cite{bouche06II} estimates this contributiotn to be $20-50\\%$ at $z=2$. Altogether the census of metals in the IGM, the ISM probed by DLA, and stars falls short by as much as 70\\% of \\omt.  At present, there are several suggested explanations for this discrepancy. These include: (1) a significant portion of the metals are within dusty neutral gas (possibly molecular) which obscures any background quasar avoiding detection \\citep{fall93,vladilo05}, (2) a significant fraction of metals are sequestered within a hot, diffuse, collisionally ionized gas, which is difficult to probe with rest-frame ultraviolet transitions \\citep{ferrara05}, (3) metals (primarily O~VI) are locked in warm--hot gas \\citep{simcoe02}, or (4) metals are located in high metallicity ``feedback'' systems \\citep{simcoe06}. All of these phases are certain to contribute, some to a larger extent than others, and it is possible that a complete solution will include significant contributions from each.  In this paper, we wish to highlight another reservoir of metals: highly photoionized, yet optically thick absorption systems. Specifically, we wish to examine the possible contribution from the Lyman Limit systems (LLS; absorbers with $\\mnhi > 10^{17.2} \\cm{-2}$ and restricted here to have $\\mnhi < 10^{20.3} \\cm{-2}$), the heretofore neglected sibling of the QAL systems. While the intergalactic medium examines gas arising in regions with overdensity $\\delta \\equiv \\rho/\\bar\\rho < 10$ and the damped \\lya\\ systems describe gas in overdense ($\\delta > 200$) galactic regions,  the LLS are likely to represent the interface between dense galactic and tenuous intergalactic gas.  It is reasonable to speculate that these absorbers probe gas related to galactic feedback processes and could therefore be a significant metal reservoir. The IGM and DLA systems have been surveyed extensively for metals \\citep{schaye03,simcoe04,pro_mtl03} but analysis of the LLS has been limited to a few systems \\citep{prochaska99,pb99} and a small survey of super Lyman limit systems (SLLS), systems with $\\mnhi = 10^{19} - 10^{20.3} \\cm{-2}$ \\citep{mirka03,peroux_slls03,peroux05}. Because even the SLLS outnumber DLA by approximately a factor of four, the LLS could easily contribute a significantly larger metal mass than the DLA.  To make a direct comparison, one must consider the frequency distribution, ionization state, and mean metallicity of the LLS, all of which are poorly constrained. In this Letter, we will consider the current constraints and highlight the prospective importance of the SLLS. \\cite{peroux06} have recently presented measurements on a low redshift SLLS which shows a super-solar metallicity\\footnote{One notes that \\cite{rtn06} measured this absorber to have $\\log \\mnhi = 20.5$ and therefore identified it as a DLA system.  Our analysis of the HST/STIS spectra gives $\\log \\mnhi = 20.3 \\pm 0.2$.} and the authors proposed that similar absorbers at high redshift could contain an important fraction of the metal budget.  Here, we report on two super-solar SLLS at $z_{abs}=1.8$ drawn from our survey of metal-strong DLA candidates \\citep{herbert06} and an on-going survey of LLS. The more extreme of the two systems exhibits greater than 5 times solar metallicity and has a metal surface density which matches all of the DLA with metallicity measurements at $z=1.6$ to 2.2. We will argue that the LLS are likely to account for at least $15\\%$ of the metal budget at high redshift and possibly the remainder of `missing' metals.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=3.5in]{f1.eps} \\end{center} \\caption{\\lya\\ profiles of the SLLS at (a) $z=1.775$ toward \\qlla\\ and (b) $z=1.768$ toward \\qllb.  The dotted curves indicate our estimate of the quasar continuum convolved with the blaze function of the HIRES spectrometer.  The dashed vertical lines indicate the redshift of `clouds' included in our best-fit solution (solid red line). } \\label{fig:lya} \\end{figure} ",
        "conclusions": "Before commenting on the implications for metals in the young universe, let us consider the physical origin of this gas. It is notable that both absorbers exhibit relatively wide absorption line profiles revealing a velocity field of several hundred \\kms. While this is partly the effect of selection bias (large velocity width yields larger EW), the MSDLA candidate sample \\citep{herbert06} focused primarily on \\ion{Si}{2}~1808 and \\ion{Zn}{2}~2026 which are optically thin in these systems. We suggest that the kinematics are indicative of feedback processes correlated with star formation.  Indeed, these absorbers may represent sightlines which pass through the gas observed in absorption in the spectra of bright Lyman break galaxies \\citep{steidel99,prs02}. The systems also exhibit very large \\ion{C}{4} column density. Current measurements of the frequency distribution of \\ion{C}{4} column densities show a dependence $N_{\\rm CIV}^{-\\alpha}$ with $\\alpha < 2$ to \\ion{C}{4} column densities of $\\mnciv = 10^{15} \\cm{-2}$ \\citep{songaila05}. Therefore, the mass density of C$^{+3}$ ions is dominated by the largest column density absorbers.  We suspect that LLS like the ones presented here contain the majority of C$^{+3}$ ions at all redshifts. Again, this is an assertion we will test through a large sample of LLS observations. It would be valuable to compare the observations of the super-solar SLLS with models of outflows from star-forming galaxies.  The detection of two super-solar SLLS underscores the prospect for LLS to contain a substantial fraction of the metal mass in the high $z$ universe.  Consider the results for \\slla\\ alone. The total metal column density of the gas is given by $\\log N_M = \\log N_H + \\log(M/H)$ where $N_H$ is the total hydrogen column density and M/H is the mean metallicity of the gas. Taking $x = 0.9$ ($N_H = 10 \\mnhi$) and five times solar abundance, we find that the metal column density is $50\\%$ larger than any DLA at $z \\approx 2$. In fact, the $N_M$ value roughly matches the total of the $\\approx 20$ DLA with metallicity measurements at $z=1.6$ to 2.2 \\citep{pro_mtl03}. Because SLLS outnumber DLA by $\\approx 4$ times \\citep{omeara06}, even if systems like \\slla\\ represent only 1\\% of the SLLS population, their contribution would match the entire DLA population.   This remains true even if the other 99\\% of SLLS were primordial!  To better illustrate the contribution of LLS (specifically SLLS) on the cosmological metal budget, consider the following calculation. First, define the mass density of metals in the DLA relative to the critical density  \\begin{equation} \\Omega_m^{DLA} = \\Omega_{HI}^{DLA} \\bar Z_{DLA} \\end{equation} where $\\bar Z_{DLA}$ is the mean metallicity in mass units (i.e.\\ $Z_\\odot = 0.022$ is the Solar metallicity). Adopting the results from \\cite{phw05} $\\Omega_{HI}^{DLA} = 0.5\\sci{-3}$ and \\cite{pro_mtl03} $\\bar Z_{DLA} = 0.0022$, we derive $\\Omega_m^{DLA} = 1.1 \\sci{-6}$, i.e.\\ a few percent of \\omt.  Consider the metal mass density of the SLLS, $\\Omega_m^{SLLS}$, motivated by observational considerations: \\begin{equation} \\Omega_m^{SLLS} = \\int \\frac{1}{1-x} \\mnhi f(\\mnhi) \\bar Z_{SLLS} \\; d\\mnhi \\label{eqn:mlls} \\end{equation} where $x$ is the ionization fraction of the LLS (likely a function of \\nhi), $f$ is the frequency distribution, and $\\bar Z_{LLS}$ is the mean metallicity of the gas. Currently, all of the expressions in the integrand of Equation~\\ref{eqn:mlls} are poorly constrained and so we will proceed conservatively. First, we adopt our new measurement of $\\mfnhi$ for the SLLS at $z \\approx 2.5$ \\citep{omeara06}: $\\mfnhi = 10^{7.17} \\mnhi^{-1.4}$. Second, we will adopt a value to the mean metallicity based on the \\cite{peroux_slls03} survey ($\\bar Z_{SLLS} = 0.1 Z_\\odot$) and assume the value is independent of \\nhi\\ value. In light of the results presented in this paper, we consider this to be a conservative lower limit to $\\bar Z_{SLLS}$. Even though super-solar SLLS are very rare they will contribute significantly to the mean if they represent $>1\\%$ of the sample.   The most uncertain quantity in this calculation is the ionization fraction which undoubtedly varies with \\nhi\\ column density (and possibly metallicity).  Photoionization calculations estimate that $x \\approx 0$ at the DLA threshold and increases to $\\approx 0.9$ at $\\mnhi = 10^{19} \\cm{-2}$ with strongest dependence on the gas density and intensity of the radiation field \\citep[e.g.][]{vladilo01,pro_ion02}. \\cite{peroux_slls03} and \\cite{mirka03} stressed that ionization corrections are small for their SLLS sample and gave the misleading impression that the gas is predominantly neutral.  Indeed, this is not the case (Dessauges-Zavadsky, priv.\\ comm.); we also find that the SLLS can be significantly ionized (Prochaska 1999; Prochter et al.\\, in prep.). For the following, we will parameterize $x$ with an empirical motivation: $x = C_x (1 - \\mnhi/10^{20.3})$ where $C_x$ can be varied to reduce/increase the degree of photoionization.  Evaluating Equation~\\ref{eqn:mlls}, we derive $\\Omega_m^{SLLS} = 5\\sci{-6} = 0.15 \\momt$ for $C_x=1$ and $2\\sci{-6}$ for $C_x = 0.5$.  These values are significantly in excess of the metal mass density implied by the damped \\lya\\ systems and the IGM. We believe this calculation is reasonably conservative because (i) we have neglected the contribution of LLS with $\\mnhi < 10^{19} \\cm{-2}$ and (ii) the mean metallicity of the SLLS could easily be 1/3 or 1/2 solar. Therefore, the LLS may contain the majority of metals at high $z$. In future papers, we will derive empirical measurements of the ionization fraction and metallicity of a large sample of LLS to directly determine the LLS contribution."
    },
    "0606/astro-ph0606090_arXiv.txt": {
        "abstract": "We study the reionization history of the Universe in cosmological models with non-Gaussian density fluctuations, taking them to have a renormalized $\\chi^2$ probability distribution function parametrized by the number of degrees of freedom, $\\nu$.  We compute the ionization history using a simple semi-analytical model, considering various possibilities for the astrophysics of reionization. In all our models we require that reionization is completed prior to $z=6$, as required by the measurement of the Gunn--Peterson optical depth from the spectra of high-redshift quasars. We confirm previous results demonstrating that such a non-Gaussian distribution leads to a slower reionization as compared to the Gaussian case. We further show that the recent WMAP three-year measurement of the optical depth due to electron scattering, $\\tau=0.09 \\pm 0.03$, weakly constrains the allowed deviations from Gaussianity on the small scales relevant to reionization if a constant spectral index is assumed. We also confirm the need for a significant suppression of star formation in mini-halos, which increases dramatically as we decrease $\\nu$. ",
        "introduction": "The reionization of the Universe is a direct result of the formation of luminous sources which produce the photons responsible (see Barkana \\& Loeb 2001 for a review). Given that the first stars appear in low-mass halos formed at high redshift, the reionization history of the Universe is expected to be a powerful probe of the amplitude and nature of density fluctuations on small scales.  The detection of Gunn--Peterson troughs (Gunn \\& Peterson 1965) in the absorption spectra of distant quasars suggests a late reionization for the Universe, at a redshift $z \\sim 6$ (Becker et al.~2001; Fan et al.~2003; White et al.~2003; Fan et al.~2005; Gnedin \\& Fan 2006). The indications of a high optical depth from the first-year WMAP results (Kogut et al.~2003; Spergel et al.~2003), implying early reionization, led to a flurry of papers on quite complex ionization scenarios (e.g.~Haiman \\& Holder 2003; Hui \\& Haiman 2003; Ciardi, Ferrara \\& White 2003; Cen 2003; Chiu, Fan \\& Ostriker 2003; Melchiorri et al.~2005). Simplicity has now largely been restored by the recent more precise third-year WMAP results (Page et al.~2006; Spergel et al.~2006), which indicate a much smaller optical depth, $\\tau = 0.09 \\pm 0.03$. This is consistent both with the quasar results and with a much simpler reionization history with a fast transition from a neutral to a ionized universe (see for example Choudhury \\& Ferrara 2006).  Recent results by Haiman \\& Bryan (2006) suggest that this can only be achieved if star formation in high-redshift mini-halos has been significantly suppressed (see also Wyithe \\& Loeb 2006).  In the present paper we perform a complementary study, investigating the dependence of the reionization history of the Universe on the nature of the cosmological perturbations, in the context of cosmologies permitted by the WMAP three-year results.  We are motivated by the suggestions that a non-Gaussian contribution, for example from cosmic defects, might have been able to reconcile the high optical depth suggested by the WMAP first-year results and a low redshift of (complete) reionization implied by the quasar data (Avelino \\& Liddle 2004; Chen et al.~2003). With the new simpler ionization models, non-Gaussian models should be better constrained.  We compute the reionization history of the Universe using a simple semi-analytical model similar to that used by Haiman \\& Bryan (2006) (based on Haiman \\& Holder 2003). We consider various possibilities for the astrophysics of reionization, subject to the constraint that reionization is completed prior to $z=6$ as required by the measurement of the Gunn--Peterson optical depth from the spectra of bright quasars at high redshift. ",
        "conclusions": "We have investigated the possible impact of non-Gaussian density perturbations on the reionization history of the Universe in the context of cosmological models with non-Gaussian density fluctuations with a renormalized $\\chi^2$ probability distribution function parametrized by the number, $\\nu$, of degrees of freedom. We assumed a constant spectral index so that the amplitude of the density perturbations on the small scales relevant to reionization could be calibrated using the three-year WMAP results. We have confirmed the results of Haiman \\& Bryan (2006) suggesting a significant suppression of star formation in mini-halos.  We have also shown that the reionization history of the Universe is able to constrain the non-Gaussian nature of the density perturbations. However, we have seen that reionization alone is not yet able to rule out any value of $\\nu$ (although very small values of $\\nu$ appear to be disfavoured). Given that large deviations from a Gaussian probability distribution function on large cosmological scales are already excluded by the WMAP data, we conclude that unless the probability distribution function has a strong scale dependence the cosmic reionization constraints are not competitive with those coming from the cosmic microwave background anisotropies. Still, it is important to bear in mind that the scales probed by reionization are much smaller than those probed by the WMAP (or even Planck), and consequently the consistency is reassuring, in particular taking into account recent claims (Mathis, Diego \\& Silk 2004) of hints for non-Gaussianity on scales much larger than those relevant for reionization.  We should also emphasize that our assumption of a constant spectral index excludes an important class of models where the density field is the sum of Gaussian and non-Gaussian contributions with different power spectra. On large cosmological scales the non-Gaussian component is severely constrained and only very small deviations from Gaussianity are allowed. However, if the power spectrum of the Gaussian contribution is steeper than that of the non-Gaussian one, as happens in hybrid models with both inflationary and defect perturbations, the non-Gaussian part may be the dominant component on small scales while being completely negligible on large cosmological scales.  Avelino \\& Liddle (2004) have recently shown that the reionization history of the Universe leads to stringent constraints on the energy scale of defects (see also Olum \\& Vilenkin 2006). However, since in that case the redshift at which the Universe becomes fully ionized no longer fixes the efficiencies of type I halos, the analysis is more dependent on the astrophysics of reionization. We shall return to this issue in a forthcoming publication."
    },
    "0606/gr-qc0606076.txt": {
        "abstract": "{ Using some supernovae and CMB data, we constrain the Cardassian, Randall-Sundrum, and Dvali-Gabadadze-Porrati brane-inspired cosmological models. We show that a transient acceleration and an early loitering period are usually excluded by the data. Moreover, the three models are equivalent to some usual quintessence/ghost dark energy models defined by a barotropic index $\\gamma_\\phi$ depending on the redshift. We calculate this index for each model and show that they mimic a universe close to a $\\Lambda CDM$ model today. } %------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------% ",
        "introduction": "\\label{s0} The notion of extradimensions (see \\cite{AbeMar00} for a popularised review) takes root in the search for a unified theory of fundamental interactions. In the twenties, \\cite{Kal21} tried to unify electromagnetism with Einstein's general relativity (GR) thanks to a geometrical scheme involving a compactified extradimension. However, such a solution does not work with weak and strong interactions. Interest in extradimensions returned in the seventies and the eighties with the development of string theory and supersymmetry, here most symmetric forms occur for a ten or eleven dimensional universe. In particular, in the mid eighties, Michael Green and John Schwartz discovered the superstring theory that is able to incorporate quantum mechanics without any problem of infinity, provided we live in a ten-dimensional universe, that is, our four-dimensional universe with six compactified extradimensions.  For all these theories, it is assumed that the four fundamental interactions are everywhere in the ten-dimensional spacetime. A breakthrough came in the late nineties when one wonder if gravitational interaction could be the only one able to travel in the ten dimensions while the other interactions would stay confined in our four-dimensional brane world. One of the interesting things about the brane idea is that it could explain why gravity is so weak with respect to the other interactions. For instance, it could be diluted in the extradimensions as explained by \\cite{ArkDimDva98}. Another interest of branes could be to explain why the cosmological constant $\\Lambda$ is so small. Randall and Sundrum looked for branes in which the bulk dimensions are strongly curved or warped but not necessarily compactified. Then, if they are warped in the right way, it is possible to explain the smallness of $\\Lambda$.  As a result, branes offer many new and exciting possibilities. Some of them, related to the universe expansion, have been qualified \"cosmological surprise\" by \\cite{Sah05}. For instance, some brane theories are able to produce a so-called transient acceleration, i.e. an acceleration of the expansion that is not eternal. Such behaviour would allow some string theories problems related to the S-Matrix to be avoided. Another cosmological surprise is the existence of some early loitering periods as defined by \\cite{SahSht05}, that is some early periods of time during which the Hubble function grows more slowly than the Hubble function of the $\\Lambda CDM$ model. Thus \\cite{SahSht05} mathematically define the loitering as the existence of a minimum value of the ratio $H(z)/H_{\\Lambda CDM}$. As a result, the age of a loitering universe is more than of a $\\Lambda CDM$ model. It also increases the growth rate of density inhomogeneities with respect to the $\\Lambda CDM$ model. Hence, an early epoch of loitering should help to boost the formation of high-redshift, gravitationally bound systems as noted by \\cite{Sah05}.\\\\ The aim of this paper is to constrain the free parameters of three brane theories, the so-called Cardassian, Randall-Sundrum, and Dvali-Gabadadze-Porrati models (DGP). We use some supernovae and CMB data. We show that a transient acceleration and an early loitering period are excluded by the observations most of the time. Moreover, we relate each of these models to some of the usual quintessence or ghost dark energy theories; so we begin by introducing each of the three brane theories.\\\\ The Cardassian model was introduced for the first time by \\cite{ChuFre00}. The authors consider a $3+1$ brane located at the $Z_2$ symmetry fixed plane of a $Z_2$ symmetric five-dimensional spacetime. With a suitable (and non unique) choice of the bulk energy momentum tensor, they found that on the brane, the Hubble function is proportional to a power of the matter density. The cosmological consequences of the expansion acceleration were analysed by \\cite{FreLew02}. A fluid interpretation of the Cardassian model was also developed by \\cite{GonFre03} as a description of the matter with some interaction terms which would give birth to an effective negative pressure able to drive the expansion acceleration.\\\\ There are two types of Randall-Sundrum models\\citep{RanSun99A}. In the Randall-Sundrum type $I$ model\\citep{Gum03}, two $3+1$ branes are embedded in a five-dimensional anti-de-Sitter ($AdS_5$) bulk with a negative cosmological constant. The extra dimension has $Z_2$ symmetry of about $y=L$ (where the brane stands with a negative tension where we live) and $y=0$ (where the other brane stands with a positive tension). We are interested by the Randall-Sundrum type $II$ model where the brane in $y=L$ is sent to infinity, whereas our brane is in $y=0$. It has a positive tension; otherwise, the Newton constant is negative.\\\\ Last, the Dvali-Gabadadze-Porrati (DGP) model\\citep{DvaGabPor00} rests upon a $3+1$ brane in a five-dimensional Minkowski space with an infinite-size extra dimension. There is neither a bulk cosmological constant nor any brane tension.\\\\ The plan of the paper is the following. In Sect. \\ref{s1}, we define the geometry and the matter content of our brane universe. In Sect. \\ref{s2}-\\ref{s4}, we use some observational data to constrain the Cardassian, Randall-Sundrum, and Dvali-Gabadadze-Porrati brane models. We conclude in Sect. \\ref{s5}. %------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------% ",
        "conclusions": "\\label{s5} The parameter values corresponding to the lowest $\\chi^2$ for each brane theory in agreement with and reasonably constrained by both the supernovae and CMB data are given in Tab. \\ref{tabChi2A}-\\ref{tabChi2B}. These are the models we are going to discuss now.\\\\ \\\\ The Cardassian theory with only CDM is equivalent to GR with dark energy $\\rho_\\phi$ with a constant eos very close to the one of a cosmological constant $p_\\phi/\\rho_\\phi=-1$: for a flat universe, the barotropic index $\\gamma_\\phi=p_\\phi/\\rho_\\phi+1=-0.01$ whereas with curvature, $\\gamma_\\phi=0.02$. Mathematically, there is no transient acceleration. A late loitering period agrees with the data. Adding dark energy does not allow us to fit the data better, but in this case the $\\nu$ parameter may take some positive values that otherwise disagree with an accelerated expansion.\\\\ The Randall-Sundrum theory illustrates how the conclusions got with the data strongly depends on the assumption we made on the cosmological parameters.\\\\ First, we assume that $\\Omega_{d_0}= 0\\pm 0.1$. The Randall-Sundrum theory is equivalent to the Cardassian theory or to GR, both with dark energy having a varying eos. In this last case, the barotropic index of dark energy mimicking the Randall - Sundrum theory with only CDM is represented in Fig. \\ref{DarkRS1}. In the past, the barotropic index tends to a stiff fluid, and in the future to a cosmological constant. A transcient acceleration disagrees with the data (at the border of the $2\\sigma$ confidence contours) but an early loitering period may occur.\\\\ \\begin{figure*}[h] \\centering \\includegraphics[width=12cm]{4154fi33.eps} \\caption{\\scriptsize{\\label{DarkRS1}The Randall - Sundrum theory is equivalent to GR with dark energy. These curves represent the barotropic index of dark energy for a Randall - Sundrum model with only CDM and for several intervals of the redshift when one assumes $\\Omega_{d_0}= 0\\pm 0.1$.}} \\end{figure*} If we decide to change the prior and take, for instance, $\\Omega_{m_0}=0.30\\pm 0.1$, something new happens. The best $\\chi^2$ is got when $\\Omega_{m_0}=0.29$, $\\Omega_{\\Lambda_0}=0.78$, and $\\Omega_{d_0}=-0.09$. A loitering period is always in agreement with the data, and a transient acceleration always at the border of the $2\\sigma$ confidence contours. But this dark energy eos mimicking the Randall-Sundrum model and shown on the Fig. \\label{DarkRS} presents some very interesting properties. In the past, the barotropic index tends to a stiff fluid. Later, it mimics quintessence dark energy and then ghost dark energy: dark energy is thus able to cross the line $\\gamma_\\phi=0$. In the far future, it will tend to a cosmological constant. This crossing of the line $\\gamma_\\phi=0$ has also been noticed in \\cite{AreKos05}. Sometimes, one wonders what the physical meaning of ghost dark energy is that violates the weak energy condition or what it means when quintessence dark energy becomes ghost. Here, we see that these questions could be physically meaningless if we try to answer in the framework of dark energy. However, they become physically meaningful if we consider that quintessence/ghost dark energy is just a mathematical representation of a physically well justified theory such as the Randall-Sundrum theory.\\\\ \\begin{figure*}[h] \\centering \\includegraphics[width=12cm]{4154fi34.eps} \\caption{\\scriptsize{\\label{DarkRS}The Randall - Sundrum theory is equivalent to GR with dark energy. These curves represent the barotropic index of this dark energy for a Randall - Sundrum model with only CDM and for several intervals of the redshift when one assumes $\\Omega_{m_0}=0.30\\pm 0.1$.}} \\end{figure*} What about the DGP theory? It is also equivalent to GR with dark energy. The barotropic index corresponding to the DGP theory with only CDM or with dark energy on the brane is represented in Fig. \\ref{DarkDGP}. In the first case, dark energy is quintessence. It tends to a cosmological constant in the future and to the special value $\\gamma_\\phi=0.5$ in the past. In the second case, dark energy is a ghost and tends to a cosmological constant in the past and to $\\Gamma$ in the future.\\\\ \\begin{figure*} \\centering \\includegraphics[width=12cm]{4154fi35.eps} \\caption{\\scriptsize{\\label{DarkDGP}The DGP theory is equivalent to GR with dark energy. These curves represent the barotropic index of this dark energy when one considers the RS with only CDM (left graph) or with dark energy on the brane (right graph).}} \\end{figure*} For the DGP theory with only CDM, a transient acceleration is mathematically impossible and a late loitering period ($z_l<1.5$) is possible. With dark energy, a transient acceleration is mathematically impossible and an early loitering period seems to agree with the data if $\\Gamma=O(10^{-1})$. Note that the DGP theory with minus sign and only CDM, disagrees with the data if one adds dark energy. Then a transient acceleration is still mathematically impossible and an early loitering period agrees with the data for positive values of $\\Gamma$.\\\\ \\begin{table*} \\begin{center} \\caption[]{Values corresponding to the smallest $\\chi^2$ for each brane theories in agreement with and reasonably constrained by both the supernovae and CMB data.} \\begin{tabular}{lllllll} \\hline Model & prior & $\\chi^2_{DOF}$ & $\\Omega_{m_0}$ & $\\Omega_{k_0}$ & $\\Omega_{e_0}$ & Others\\\\ \\hline Card. 1 & $\\Omega_{m_0}=0.27\\pm0.05$ & 1.16 & $0.29$ & - & - & $\\nu=-1.01$\\\\ 2 & $\\Omega_{m_0}=0.27\\pm0.05$ & 1.17 & $0.30$ & $-0.02$ & - & $\\nu=-0.98$\\\\ & and $\\Omega_{k_0}=-0.02\\pm0.05$ &  &  &  &  & \\\\ \\hline \\hline RS      & $\\Omega_{d_0}=0\\pm0.1$ &  & $0.15$ & - & $0.8$ & $\\Omega_{d_0}=-0.008$ \\\\ \\hline \\hline DGP 1   & $+$ sign, $\\Omega_{m_0}=0.27\\pm0.05$ & $1.17$ & $0.23$ & $-0.04$ & - & -\\\\ & and $\\Omega_{k_0}=-0.02\\pm0.05$ &  &  &  &  & \\\\ 2   & $+$ sign, $\\Omega_{m_0}=0.27\\pm0.05$ & $1.14$ & $0.27$ & - & $0.17$ & $\\Gamma=-12.83$\\\\ \\hline \\end{tabular} \\label{tabChi2A} \\end{center} \\end{table*}  \\begin{table*} \\begin{center} \\caption[]{Values corresponding to the smallest $\\chi^2$ for each brane theories in agreement with and reasonably constrained by both the supernovae and CMB data.} \\begin{tabular}{llllllll} \\hline Model & $H_0$ & Age & $z_a$ & $z_d$ & Trans. & Loit.\\\\ \\hline Card. 1 & $64.6$ & $14.64$ & $0.69$ & $0.34$ & No & Yes($z_l<1$) \\\\ 2 & $64.6$ & $14.63$ & $0.69$ & $0.32$ & No & Yes($z_l<1$) \\\\ \\hline \\hline RS     &  $64.7$ & $11.4$ & $0.44$ & $0.34<z<1.17$ & No & Yes (early-time) \\\\ &   \t\t & \t\t  & \t   & \t           &    & if $\\Omega_{\\Lambda_0}<0.8$ \\\\ \\hline \\hline DGP 1  & $64.86$ & $14.96$ & $0.76$ & $0.68$ & No & Yes($z_l<1.5$)\\\\ 2  & $67.75$ & $14.13$ & $0.32$ & $0.33$ & No & Yes(early-times)\\\\ &   \t\t & \t\t   & \t    & \t     &    & if $\\Gamma=O(10^{-1})$\\\\ \\hline \\end{tabular}  \\label{tabChi2B} \\end{center} \\end{table*} Hence, a transient acceleration and an early loitering period - two interesting properties of branes - usually disagree with the data. In particular, the loitering period usually occurs at late time. All the theories shown in Tab. \\ref{tabChi2A}-\\ref{tabChi2B} predict a Hubble constant around $65$ (around $68$ for the DGP model with curvature) and an age for the universe around $15$ billions years (around $11.4$ for the Randall-Sundrum model).\\\\ Moreover, all the brane models we have considered are equivalent to GR with dark energy defined by a barotropic index $\\gamma_\\phi$. Each time one takes the supernovae and CMB data into account, this dark energy is close to the form of a cosmological constant, i.e. $\\gamma_\\phi=0$. Hence for the Cardassian model, $\\gamma_\\phi\\approx 10^{-2}$. For the Randall-Sundrum and DGP models, their barotropic index varies but is also close to $0$ at present time. Indeed, their barotropic index tends to some constants in the past and in the future, these constants being related by a fast transition period and the constant value $\\gamma_\\phi=0$ reached in the future (past) for the Randall-Sundrum (DGP) model. What is striking is that the data predict we should live just during the short transition period. What does it mean? Probably that the value of $\\gamma_\\phi$ is so close to $0$ today that we are not able to measure it precisely enough to prove that we are in the asymptotical period when $\\gamma_\\phi\\rightarrow 0$ instead of the transition period where $\\gamma_\\phi$ may be slightly different from zero (in the same way that if the universe is nearby flat, we can prove it only if the measures are precise enough to measure a very very sligth curvature). This is probably why the data lead us to believe that we are living during a transition period when $\\gamma_\\phi$ is close to zero and not yet in an asymptotical epoch when $\\gamma_\\phi\\rightarrow 0$.\\\\\\\\ To conclude, if our universe is described by one of the three brane-inspired models studied in this paper, today it is probably very similar to a $\\Lambda CDM$ universe. But how similar? The supernovae data are not precise enough to answer (and consequently could fake a universe undergoing a transition to a $\\Lambda CDM$ model instead of being already asymptotically enter into such a state) but it could be the case of structures formation data (as SDSS or 2dFGRS) as recently shown in \\cite{AmaElgMul04}. \\appendix %------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------------%"
    },
    "0606/hep-th0606090_arXiv.txt": {
        "abstract": " ",
        "introduction": "According to our present understanding of cosmology the history of the universe has been characterized by an incessant slowing down of the Hubble rate $H = \\dot{a}/ a$. Even the ``accelerated'' expansion ($\\ddot{a} > 0$) today and during inflation at best corresponds to $H = \\hbox{\\rm constant}$, and therefore can only halt the decrease of $H$. Intuitively, the fact that $H$ cannot increase is a manifestation of the fact that gravity is generally attractive, and in some sense never ``too repulsive.'' More formally, $\\dot{H} \\le 0$ can be viewed as a consequence of the null energy condition (NEC), which requires that for all null vectors $n^\\mu$ the matter stress-energy tensor must satisfies $T_{\\mu\\nu} n^\\mu n^\\nu \\ge 0$. In a Friedmann-Robertson-Walker (FRW) spacetime this condition reduces to $\\rho + p \\ge 0$, which for a spatially flat universe directly implies $\\dot H \\le 0$.  But if the NEC is satisfied and $H$ is always decreasing, why is the universe expanding in the first place? The conventional view is that $H$ increases going backward in time until we reach an era where $H \\sim \\MP$ (the ``big bang'') where quantum gravity effects become important. At this scale general relativity breaks down as an effective theory, so in this view the mystery of the origin of the expansion of the universe is inseparable from the issue of the UV completion of gravity. On the other hand, if the NEC is violated it is possible that the history of the universe is drastically modified without quantum gravity effects becoming important. For example, the present expansion might be due to a smooth ``bounce'' from an earlier contracting phase. An even more interesting possibility, which is just a limiting case of a bouncing cosmology, is a universe that asymptotically in the past is Minkowski space and as time goes on it simply starts expanding connecting smoothly to the subsequent familiar FRW cosmology and thus resolving the Big Bang singularity. These possibilities are discussed more fully below.   Another motivation to consider violations of the NEC come from observations of the current accelerated expansion of the universe. These are conventionally summarized by the equation of state parameter $w = p/\\rho$. A cosmological constant has $w = -1$, but present data allow $w \\gsim -1.2$ \\cite{Spergel:2006hy,Seljak:2006bg}. The NEC implies $w \\ge -1$, and it is natural to ask whether $w < -1$ is a viable possibility.   One simple way to violate the NEC is to add fields with wrong-sign kinetic terms. Scalar excitations  have energy unbounded from below (they are ``ghosts''), and are therefore subject to catastrophic instabilities at all scales. In particular the vacuum is unstable to decay into ghosts and gravitons with an infinite rate in any theory that is Lorentz invariant in the UV \\cite{riccardo,ghostconstraints2}. This kind of model can nonetheless be made compatible with observation by coupling it only to gravity and postulating that unknown UV physics breaks Lorentz invariance and cuts off the divergence for short wavelenths \\cite{riccardo,ghostconstraints2}. In this approach, the stability of the universe depends on unknown UV physics as the leading instability is at short distances\\footnote{See \\cite{Holdom:2004yx} for an explicit realization of a Lorentz violating cut-off of a ghost instability.}.  The NEC is believed to hold for all well-behaved systems without instabilities. Other energy conditions can be violated by a sensible system simply by adding a suitable (positive or negative) cosmological constant. The NEC cannot, since it is saturated by a cosmological constant. Indeed it has been shown in Ref.~\\cite{fluids} for a very broad class of models that whenever the stress-energy tensor violates the NEC the system has catastrophic instabilities, either in the form of ghosts or in the form of exponentially growing modes with arbitrarily short wavelengths (``tachyons'').% \\footnote{The only exceptions are certain anisotropic systems which are stable but admit superluminal excitations \\cite{fluids}. However such exceptions are not relevant for cosmology, where one is only interested in isotropic systems.} Both types of instability occur for all wavelengths down to the UV cutoff of the theory. Therefore, unlike the Jeans instability, these instabilities cannot be damped by Hubble friction, which is only effective at frequencies smaller than $H$.  The results of Ref.~\\cite{fluids} are derived at the two-derivative level in a systematic low-energy expansion. Higher derivative corrections are suppressed by some scale $M$, and in general they can be safely ignored at frequencies and momenta smaller than $M$. However this is not the case if the two-derivative action is degenerate, and some modes propagate only thanks to higher derivative terms. This is precisely what happens in the ghost condensate \\cite{Arkani-Hamed:2003uy} and in more general theories of massive gravity \\cite{sergio}. For example, scalar fluctuations around the ghost condensate have dispersion relation $\\omega^2 = \\vec{k}^4 / M^2$. Theories of modification of gravity can be thought as non-trivial scalar backgrounds with a stress-energy tensor equal to the one of a cosmological constant \\cite{fluids}. Therefore they lie on the threshold of violating the NEC. This explains the origin of the degenerate dispersion relations and tell us that these theories are a first step towards violating the NEC.  In this paper we use the ghost condensate as a starting point to construct a consistent and technically natural low-energy effective field theory that violates the NEC without any instabilities. The simplest example is a ghost condensate $\\phi$ rolling {\\em up} a linear potential, which gives rise to  superacceleration ($\\dot{H} > 0$). The late-time solution for $\\phi$ climbs up the potential indefinitely, making the universe expand with an indefinitely growing Hubble rate, $H \\propto \\sqrt{t}$. This solution is an attractor, in the sense that trajectories flow to a common late time behavior, like in conventional slow-roll inflation. An interesting feature of this model is that we cannot violate the NEC by an arbitrary amount, in the sense that we must have \\be \\label{eq:thebound} \\dot{H} \\lsim H^2 \\ee to avoid instabilities. This comes from an interplay between gradient and Jeans instabilities. As discussed in Ref.~\\cite{fluids}, violating the NEC gives rise to a dispersion relation of the form $\\omega^2 \\sim -k^2$, which has a gradient instability. This instability is cured at short distances by the higher-derivative term giving a dispersion relation schematically of the form $\\omega^2 \\sim -k^2 + k^4$, which is unstable only for long times $\\omega < \\omega_{\\rm grad}$. However, the $k^4$ term increases the mixing with gravity, leading to a Jeans instability similar to that for ordinary matter, so the system is unstable for $\\omega < \\omega_{\\rm Jeans}$. Remarkably, we find the model-independent relation between the instabilities \\be \\label{riccardo} \\omega_{\\rm grad} \\,  \\omega_{\\rm Jeans} \\sim \\dot{H} \\;. \\ee Hubble friction will damp all instabilities only if both $\\omega_{\\rm grad}$ and  $\\omega_{\\rm Jeans}$ are smaller than $H$, which leads to Eq.~(\\ref{eq:thebound}). We emphasize that if this inequality is violated, the instability does not extend to arbitrarily short distances and time scales, and is under control of the effective theory. These instabilities are therefore much milder than ghost or tachyon instabilities, and may in fact have interesting cosmological consequences. (For example, the conventional Jeans instability gives rise to structure formation in the universe.)  In fact, these features are much more general than this particular model. We demonstrate this by performing a general analysis of scalar fluctuations about an arbitrary FRW cosmological history. The analysis can be viewed most simply as the general case of an expansion driven by a rolling scalar. However, it also applies to the fluctuations of the Goldstone mode of time translations in the case where the expansion is driven by a mixture of fluids and rolling scalar fields, and so the analysis is quite general. We find that the qualitative features of the ghost condensate in a linear potential are completely generic. In particular, whenever $\\dot H$ is positive the system is unstable at the two-derivative level, and the same interplay between gradient and Jeans instabilities leads to Eq.~(\\ref{riccardo}), so we must have $\\dot{H} \\lsim H^2$ in order for instabilities to be completely absent.  We then consider a number of cosmological applications of models based on ghost condensation that violate the NEC without instabilities. We first discuss the possibility that the present expansion of the universe is superaccelerating ($\\dot{H} > 0$), corresponding to $w < -1$. The model consists of the ghost condensate rolling up a potential. In such a model $H$ can in principle increase smoothly forever, although at some point $H$ becomes so large that the model exits the regime of validity of the effective field theory. To get a measurable violation of the NEC we must be close to saturate the inequality (\\ref{eq:thebound}), suggesting that there may be long wavelength growing modes in the dark energy sector that just started evolving. These may provide an additional signal for this class of models.  We then construct a model where the universe starts from Minkowski space in the asymptotic past. The model consists of a ghost condensate with $\\phi=t$ that rolls up a potential $V \\sim \\phi^{-2}$. For large negative $t$, this model gives $H \\sim |t|^{-1}$, so the universe goes from a zero curvature, zero energy state to higher curvatures and energies, and this can then be smoothly connected to a standard FRW expansion.  Next, we consider a model which gives rise to a cyclic universe that is always expanding. While the scale factor $a(t)$ steadily grows larger and larger, the Hubble parameter and the energy density are periodic functions of time. In this model, a ghost condensate travels along a periodic potential, giving rise to a phase of superacceleration followed by reheating, a radiation-dominated and then matter-dominated phase, followed again by superacceleration, and so on. In such a scenario, the present accelerated phase is the beginning of primordial inflation!  Finally, we consider the possibility that the universe smoothly bounces from a contracting to an expanding phase. Such a possibility has been considered previously in the literature (for a review, see Ref.~\\cite{justin}), but in previous treatments the bounce was not under theoretical control. We construct a completely smooth bouncing solution with no instability. Note that since $H \\to 0$ at the bounce, while $\\dot{H} \\ne 0$, there is always an unstable mode violating Eq.~(\\ref{eq:thebound}). However, it is easy to arrange that instabilities do not have time to grow much during the bounce, so that this phase is under theoretical control without catastrophic instabilities.   We conclude that it is possible to have physical systems that violate the NEC without instabilities or other pathologies, and that this opens up a large number of interesting new possibilities for cosmology. The illustrative examples we consider are intended only as toy models; we leave the construction of realistic models and their phenomenology to future work. ",
        "conclusions": "} In this paper we have shown that at the level of effective field theory there is no obstruction in constructing systems that violate the null energy condition (NEC). In particular, the instabilities that plague a very broad class of NEC-violating systems \\cite{fluids} can be completely avoided. We use this to construct cosmological models with $\\dot{H} > 0$ that do not have any instabilities.  These results depend crucially on the interplay between two potential types of instabilities. The first are gradient instabilities that occur very generally whenever the NEC is violated. These can be avoided due to higher-derivative terms in the effective theory. However, increasing the higher-derivative terms increases the effect of another potential instability, the Jeans instability that arises from the attractive nature of gravity. The Jeans instability can be avoided if the time scale for the instability is longer than the Hubble time. Demanding that both of these instabilities are absent requires \\be \\label{conditionconclusion} \\dot{H} \\lsim H^2 \\; . \\ee This means that we cannot get an arbitrarily large violation of the NEC. These results are derived both in a specific model (a deformation of ghost condensation \\cite{Arkani-Hamed:2003uy} deformed by the addition of a potential \\cite{Arkani-Hamed:2003uz,Senatore:2004rj}) and in a very general effective field theory analysis of adiabatic scalar fluctuations about an arbitrary FRW background.  We emphasize that if Eq.~(\\ref{conditionconclusion}) is violated the effective field theory does not break down, but it contains long wavelength exponentially growing modes. These might be interesting for cosmology: for example, if the present acceleration of the universe has $w = p /\\rho < -1$, then it becomes possible that they might have shown up around our present epoch.  We constructed a number of explicit examples to show how this can lead to interesting alternative cosmological scenarios. We constructed a simple explicit model that gives $w < -1$ today. We also presented a model in which the universe starts from Minkowski space in the distant past, giving a possible alternative to standard inflation for solving the horizon problem. Putting these together, we consider an eternally expanding cyclic model in which the present accelerating phase is the beginning of inflation in the next cycle. This opens up the intriguing possibility that measurements today can give information about the inflationary phase that gave rise to the structure we observe today. Finally, we constructed a model which has a smooth ``bounce'' connecting a contracting phase to the present expanding phase. Although bouncing cosmologies have been previously considered in the literature, we believe ours is the first example in which the bounce is under theoretical control. These illustrate some of the phenomenologically interesting possibilities opened up by violation of the NEC. We leave detailed analysis of these ideas to future work.  We close with some theoretical questions. It would be interesting to understand whether the requirement $\\dot H \\ll H^2$ to have a parametric suppression of instabilities is completely generic or not. Models involving more degrees of freedom could be completely stable even when the NEC is strongly violated.  While the results of this paper are under control in a systematic effective field theory expansion, it is natural to ask whether this particular effective field theory can be embedded in a UV complete theory such as renormalizable quantum field theory or string theory. In particular, the somewhat exotic features of the ghost condensate (e.g. Lorentz breaking and complicated nonlinear dynamics \\cite{Arkani-Hamed:2003uy}) may make one suspicious that the ghost condensate lies in the ``swampland'' of effective field theories without UV completions \\cite{Vafa:2005ui, Arkani-Hamed:2006dz}. In fact, in Ref.~\\cite{adams} it was shown that a large class of consistent effective field theories (including some interesting modifications of gravity) have no Lorentz invariant UV completion, essentially because they allow configurations in which signals travel faster than light \\cite{adams}. The present model does not suffer from this problem, since the excitations of the ghost condensate have a dispersion relation $\\omega^2 \\simeq k^4/M^2$, and therefore travel much \\emph{slower} than light. However, it is true that there is currently no known UV completion of the model; for an interesting recent attempt, see Ref.~\\cite{ian}.  Despite these open questions about the full consistency of the model at all energy scales, we believe that the phenomenological possibilities opened up by violation of the NEC are well worth further exploration."
    },
    "0606/astro-ph0606528_arXiv.txt": {
        "abstract": "We present a new technique for choosing spatial regions for X-ray spectroscopy, called ``contour binning''. The method chooses regions by following contours on a smoothed image of the object. In addition we re-explore a simple method for adaptively smoothing X-ray images according to the local count rate, we term ``accumulative smoothing'', which is a generalisation of the method used by \\textsc{fadapt}. The algorithms are tested by applying them to a simulated cluster data set.  We illustrate the techniques by using them on a 50~ks \\emph{Chandra} observation of the Cassiopeia A supernova remnant. Generated maps of the object showing abundances in eight different elements, absorbing column density, temperature, ionisation timescale and velocity are presented.  Tests show that contour binning reproduces surface brightness considerably better than other methods. It is particularly suited to objects with detailed spatial structure such as supernova remnants and the cores of galaxy clusters, producing aesthetically pleasing results. ",
        "introduction": "Many X-ray telescopes such as \\emph{XMM-Newton}, \\emph{Chandra} and \\emph{ROSAT} contain detectors which capture individual photons, recording their energy and position on the detector. Therefore unlike conventional optical observing techniques, we simultaneously collect imaging and spectroscopic data for each part of the object. In addition the \\emph{Chandra X-ray Observatory} has very high spatial resolution ($\\sim 1$~arcsec), allowing the properties of the emitting object to be studied in unprecedented detail.  With this spectroscopic and imaging information we can select ``events'' from the observation corresponding to a particular part of an object with a ``region filter''. From these events a spectrum can be built-up. An X-ray spectral package such as \\textsc{xspec} (Arnaud 1996) can then be used to fit a physical model to the spectrum. Conventionally simple geometric shapes such as annuli, sectors, boxes or ellipses are used to define a region filter. Using sectors, for example, and assuming spherical or elliptical symmetry, one can account for projection in a cluster of galaxies. However most extended objects are not symmetric when observed in detail (for example, the Perseus cluster [Fabian et al 2000; Sanders et al 2004], the Cassiopeia-A supernova remnant [Hwang et al 2004], and Abell 2052 [Blanton, Sarazin \\& McNamara 2003]).  Given the morphological diversity of extended X-ray sources it is important to have techniques which allow us to analyse the spectral variation of a source over its extent. We first investigated this problem when looking for cool gas in a sample of X-ray clusters using \\emph{ROSAT} PSPC archival data (Sanders, Fabian and Allen 2000). We devised a technique which used adaptively smoothed maps (Ebeling et al, in preparation) to define contours in surface brightness. The ratio of the number of counts in different energy bands between each contour was used to define an X-ray colour. By using a grid of models, the absorption and temperature of the gas could be estimated between the contours.  We approached this problem again with the advent of data from \\emph{Chandra} with its high spatial resolution. We created an algorithm called ``adaptive binning'' (Sanders \\& Fabian 2001) which used the uncertainty on the number of counts or the error on the ratio of counts in different bands to define the size of binning region used. The process was simple: pass over the image, copying those pixels which have a small enough uncertainty on the number of counts or colour to an output image. Bin up the remaining pixels by a factor of two. Repeat until all the pixels have been binned. On the final pass we bin any pixels which are not yet binned. This simple approach works well and was used by us and other authors on data from several clusters of galaxies (e.g.  Centaurus -- Sanders \\& Fabian 2002; Perseus -- Fabian et al 2000; Abell~4059 --- Choi et al 2004).  The disadvantage of this approach is that the binning scale varies by factors of two. It is very noticeable where the scale changes, and some regions are overbinned. Therefore we started using the ``bin accretion'' algorithm of Cappellari \\& Copin (2003). The algorithm adds pixels to a bin until a signal-to-noise threshold is reached. After all the pixels have been accreted, it uses Voronoi tessellation to make tessellated regions based on the weighted position of the original bins. This technique has the advantage of creating bins which are compact, varying in size smoothly with the surface brightness, and also provides bins with similar signal-to-noise ratios.  We applied the method to X-ray observations of the Perseus cluster (Sanders et al 2004) and Abell~2199 (Johnstone et al 2002).  Rather than use X-ray colours, we extracted spectra for each of the regions and used spectral fitting to derive, for example, temperature and abundance maps. Recently Diehl \\& Statler (2006) have generalised this algorithm to allow for data whose signal-to-noise does not add in quadrature.  The motivation for further work in this area is that the methods above do not use the surface brightness distribution to change the shape of each bin.  Physical parameters (e.g. density, temperature and abundance) usually change in the direction of surface brightness changes.  The method we describe here uses the surface brightness to define bins which cover regions of similar brightness.  Other methods have been presented for mapping the parameters of the intracluster medium. These included wavelet techniques (Bourdin et al 2004) and monte carlo methods (Peterson, Jernigan \\& Kahn 2004). The advantage of binning techniques is that they provide errors on individual spectral fit parameters, or colours, from a particular part of the sky. The individual measurements made using binning methods are independent, making it easy to easily measure the significance of individual spatial features.  The techniques presented in this paper have already been applied to a number of \\emph{Chandra} observations of clusters, including a deep observation of the complex structure of the Centaurus cluster (Fabian et al 2005), the possible detection of nonthermal radiation and the identification of a high metal shell likely to be associated with a fossil radio bubble in the Perseus cluster (Sanders et al 2005), a sample of moderate redshift clusters (Bauer et al 2005), and a 900~ks observation of the Perseus cluster (Fabian et al 2006), finding little evidence for temperature changes associated with shock-like features, and producing evidence of a substantial reservoir of cool X-ray emitting material.  We first present a simple smoothing method (``accumulative smoothing''), and then present the binning method based on the smoothed image (``contour binning''). ",
        "conclusions": "We have shown that the accumulative smoothing method, or that used by \\textsc{fadapt}, is a robust way to estimate the real surface brightness distribution of an extended object. We have discussed the method in the cases of varying exposures, and including backgrounds.  We have described a new method for choosing regions for spectral extraction, or colour analysis, called contour binning. We have shown that method reliably creates bins which follow the surface brightness. We demonstrate the method using a simulated cluster dataset, and apply it to a \\emph{Chandra} observation of the supernova remnant Cassiopeia A.  The algorithm is ideal where spectral changes are associated with changes in surface brightness, as is often the case in X-ray astronomy. The results on objects with fine spatial detail are very closely matched to the object, aesthetically pleasing, and easy to interpret, compared to other methods (as exemplified by Fig.~\\ref{fig:cassa_binned} and Fig.~\\ref{fig:cassa_deltas}). As the algorithm follows the surface brightness it naturally bins regions which are likely to be physically associated. It avoids mixing spectra together for physically disassociated regions."
    },
    "0606/astro-ph0606002_arXiv.txt": {
        "abstract": "We stacked the X-ray data from the \\rosat\\ All Sky Survey for over 4,000 clusters selected from the 2MASS catalog and divided into five richness classes.  We detected excess X-ray emission over background at the center of the stacked images in all five richness bins.  The interrelationships between the mass, X-ray temperature and X-ray luminosity of the stacked clusters agree well with those derived from catalogs of X-ray clusters.  Poisson variance in the number of galaxies occupying halos of a given mass leads to significant differences between the average richness at fixed mass and the average mass at fixed richness that we can model relatively easily using a simple model of the halo occupation distribution.  These statistical effects probably explain recent results in which optically-selected clusters lie on the same X-ray luminosity-temperature relations as local clusters but have lower optical richnesses than observed for local clusters with the same X-ray properties.  When we further binned the clusters by redshift, we did not find significant redshift-dependent biases in the sense that the X-ray luminosities for massive clusters of fixed optical richness show little dependence on redshift beyond that expected from the effects of Poisson fluctuations. Our results demonstrate that stacking of RASS data from optically selected clusters can be a powerful test for biases in cluster selection algorithms. ",
        "introduction": "The primary objective of many new extragalactic surveys is to determine the equation of state of the dark energy.  One approach is to determine the evolution of the cluster mass function with redshift, which depends on dark energy through the growth factor and the volume element (e.g., Haiman, Mohr \\& Holder 2001; Huterer \\& Turner 2001; Podariu \\& Ratra 2001; Kneissl et al. 2001; Newman et al. 2002; Levine, Schulz, \\& White 2002; Majumdar \\& Mohr 2003; Hu 2003).  Clusters are identified using some observable proxy for their mass over a broad range of redshifts and the proxy must then be calibrated to compare the results to the theoretically predicted mass functions.  Possible proxies and identification methods are the optical richness of the cluster (e.g., Abell 1958; Huchra \\& Geller 1982; Dalton et al. 1992; Postman et al. 1996, 2002; Zaritsky et al. 1997; Gonzalez et al. 2001; Bahcall et al. 2003; Gal et al. 2003; Gladders \\& Yee 2005), the X-ray luminosity or temperature (e.g., Gioia et al. 1990; Henry et al. 1992; Rosati et al. 1995,1998; Ebeling et al. 1998, 2000; B\\\"ohringer et al. 2000, 2004; Bauer et al. 2002; Giacconi et al. 2002), the Sunyaev-Zeldovich decrement (e.g., Carlstrom et al. 2000; LaRoque et al. 2003) and the weak lensing shear (e.g., Wittman et al. 2001; Dahle et al. 2002; Schirmer et al. 2003).   Examples of large scale surveys planning on using these methods are the Dark Energy Survey (DES\\footnote{http://www.darkenergysurvey.org/} using optical richness, Sunyaev-Zeldovich and weak lensing), the Large Synoptic Survey Telescope (LSST\\footnote{http://www.lsst.org/} optical richness and weak lensing) and the Supernova / Acceleration Probe (SNAP\\footnote{http://snap.lbl.gov/} optical richness and weak lensing).  This is in addition to smaller scale surveys based on optical richness such as the SDSS cluster surveys (Goto et al. 2002; Bahcall et al. 2003)  and the Red Cluster Sequence Survey (RCS, Gladders \\& Yee 2005).  A persistent worry about optically selected cluster samples is that chance projections of foreground and background galaxies significantly bias the resulting catalogs.  The problem can range from false positives, detections of non-existent clusters, to a richness bias in which chance projections lead to overestimates of the cluster richness.  The problems should be more severe for higher redshift and lower richness clusters because the effects of a chance projection increase as the number of detectable galaxies in the cluster diminishes.  One approach to checking the extent to which a cluster finding algorithm is affected by these issues is to generate mock galaxy catalogs matching the actual data as closely as possible, find clusters in the mock catalogs and then compare the output and input catalogs (e.g., Kochanek et al. 2003; Miller et al. 2005).  In general these comparisons have found only modest biases.  A second approach is to compare the mass estimates based on the optical richness to those from another method.  At present this has meant examining either the X-ray properties (e.g., Bahcall 1977; Donahue et al. 2002; Kochanek et al. 2003; Popesso et al. 2004) or the weak lensing masses (e.g., SDSS/RCS) of the optically-selected clusters.  A basic problem of most existing tests of optically-selected cluster catalogs using X-ray data is that the comparisons are made to catalogs of X-ray clusters rather than through X-ray observations of the optically-selected clusters.  While this provides a simple means of calibrating the relationship between optical and X-ray properties, it is not an optimal approach to searching for biases in the optical selection methods.  The problems could be more severe because several recent studies (e.g., Bower et al.\\ 1997; Lubin, Mulchaey \\& Postman 2004; Gilbank et al. 2004) have found offsets between the X-ray temperature-optical richness relations of local and higher redshift optically-selected clusters even though the clusters lie on the same X-ray temperature-luminosity relations.  In this paper we address these questions by measuring the mean X-ray properties of a sample of galaxy clusters selected from the 2MASS survey (Kochanek et al. 2003, 2006 in preparation) by averaging (``stacking\") the \\rosat\\ All-Sky Survey (RASS, Voges et al. 1999) data for the clusters as a function of richness and redshift. We outline the cluster catalog and our procedures for analyzing the RASS data in \\S2.  In \\S3 we discuss the theoretical differences between measuring the mean properties of clusters at fixed richness rather than fixed mass.  In \\S4 we present our results for the correlations between the optical richness, X-ray properties and cluster mass estimates, as well as the dependence of the results on cluster redshift. We summarize our results in \\S5.  We assume that $H_0 = 70~\\rm{km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\Lambda}= 0.7$ throughout the paper.  \\begin{figure} \\epsscale{1} \\plotone{f1.eps} \\caption{RASS images of four clusters over-plotted with probable cluster members (squares, $> 50\\%$ likelihood of membership as determined from the cluster finding algorithm in Kochanek et al. 2003) and field galaxies (X symbols, $< 50\\%$ likelihood) along with the error ellipse for the cluster centroids provided by optical cluster detection algorithm. Panel A shows the typical cluster, where the optical position is roughly consistent with the X-ray position.  Panel B shows an example where the existence of a cluster is correctly inferred in the optical catalog, but the position is seriously in error.  Panels C and D show two examples where the optical cluster is confused by the presence of multiple (or merging) clusters.\\label{fig:oxplot}} \\end{figure}  \\begin{figure} \\epsscale{1} \\plotone{f2.eps} \\caption{The stacked \\rosat\\ images of 2MASS clusters in richness bins 0--4 and at random positions. The images are 2.0, 2.0, 1.5, 1.0, 1.0, and 1.0 \\h Mpc in radius for richness bins 0--4 and the random positions.  The contours are drawn at levels of 1--4~$\\sigma$ above the background. Excess X-ray emission is clearly detected in all richness bins, whereas the stacked image using random positions does not show any significant emission at the center. \\label{fig:img}} \\end{figure} ",
        "conclusions": "Optically-selected cluster catalogs have always been suspect because of concerns about false detections and redshift-dependent biases created by chance alignments of galaxies.  Current attempts to understand the biases of optically-selected catalogs have focused either on testing the algorithms in mock galaxy catalogs (Kochanek et al. 2003; Miller et al. 2005) or by using weak lensing measurements (Sheldon et al. 2001). Unfortunately, there has never been a complete X-ray survey of a large optically-selected cluster catalog to provide an independent test of these concerns. In this study we have used a stacking analysis of the \\rosat\\ All-Sky Survey data to determine the average X-ray properties of a sample of 4333 clusters found using the matched filter algorithm of Kochanek et al. (2003) applied to the 2MASS galaxy survey with a magnitude limit of K$<13.25$~mag.  After dividing the clusters into bins of optical richness, we could successfully measure the surface brightness profiles, bolometric X-ray luminosities, X-ray temperatures, masses and gas entropies for the averaged clusters, finding results that are generally consistent with those from analyzing individual clusters.  The one exception is that we tend to find larger core radii for $\\beta$-model fits to the surface brightness presumably because of additional smoothing created by errors in the optical positions.  The stacked clusters cleanly follow the correlations observed for samples of individual clusters between X-ray luminosity, temperature and mass.  However, in order to interpret the correlations between the optical richness and the X-ray luminosity, temperature and mass, it is necessary to model the effects of Poisson fluctuations of the number of galaxies in a cluster because in the presence of these fluctuations averaging at fixed richness is very different from averaging at fixed mass. These effects matter for high richness clusters observed at redshifts where they will contain few detectable galaxies, and for all low richness clusters.  A very simple Poisson model for the effects of these fluctuations naturally reproduces our results.  We note that Stanek et al. (2006) have recently demonstrated similar effects arising from the scatter in X-ray luminosity at fixed cluster mass.  We also examined the problem of redshift-dependent biases in optical catalogs by examining the X-ray properties of clusters of fixed richness as a function of redshift.  The low redshifts of the 2MASS clusters ($z<0.1$) means that we need worry little about genuine cosmological evolution.  Our statistical model predicts that high richness clusters will show a slowly declining X-ray luminosity, while low richness clusters will show a rising X-ray luminosity, and this agrees with our measurements. Essentially, Malmquist biases created by the much more abundant low-mass clusters become steadily more important for the high mass clusters as the numbers of detectable member galaxies decline, while the reverse occurs for the low-mass clusters.  In general, however, samples restricted to clusters containing at least 3 detectable galaxies (already a floor required to minimize the abundance of false positives) and to the higher mass clusters rather than groups ($T \\gs 2$~keV, $M \\gs 10^{14}M_\\odot$, $L_X \\gs 10^{43}$~ergs/sec) have corrections due to the effects of Poisson fluctuations that are relatively simple to model and depend little on the model for the cutoff between groups and galaxies.  Models for the lower mass groups are more challenging, particularly if variations with redshift must be included.  We see no evidence for large systematic errors created by chance superpositions.  Despite the shallowness of the RASS, our approach (see also Bartelmann \\& White 2003) can be applied to analyses of cluster samples with much higher mean redshifts than the 2MASS sample, where even the richest clusters are tracked only to $z \\ls 0.1$.  Modulo redshift scalings for K-corrections and between cosmological distances, the signal-to-noise ratio of a stacking analysis scales as the square root of the product of the survey area and depth.  Thus, an analysis of clusters in 10000~deg$^2$ of the SDSS survey (25\\% of the area of 2MASS), which can detect rich clusters with $z\\ls 0.5$ rather than $z\\ls 0.1$ (5 times the depth), should have a signal-to-noise ratio comparable to our present analysis for the measurement of the mean X-ray properties of the clusters.  Analyses of rich clusters over 1000~deg$^2$ would have a signal-to-noise ratio comparable to our analysis of clusters with $\\sim 1/3$ the richness.  This assumes that the analysis remains limited by statistical errors rather than being dominated by systematic problems such as contamination by faint point sources or difficulty in controlling the background."
    },
    "0606/astro-ph0606144_arXiv.txt": {
        "abstract": "We present a two-dimensional kinematic analysis for a sample of simulated binary disc merger remnants with mass ratios 1:1 and 3:1. For the progenitor discs we used pure stellar models as well as models with 10\\% of their mass in gas. A multitude of phenomena also observed in real galaxies are found in the simulations. These include misaligned rotation, embedded discs, gas rings, counter-rotating cores and kinematic misaligned discs. Using the 2D maps we illustrate projection effects and the change in properties of a merger remnant when gas is included in the merger. We find that kinematic peculiar subsystems are preferably formed in equal mass mergers. Equal-mass collisionless remnants  can show almost no rotation, regular rotation or strong kinematic misalignment. The inclusion of gas makes the remnants appear more round(1:1) and axisymmetric(3:1). Counter-Rotating Cores (CRCs) are almost exclusively formed in equal-mass mergers with a dissipational component. 3:1 remnants show a much more regular structure. We quantify these properties by applying the kinemetric methods recently developed by Krajnovi\\'c et al. This work will help to understand observations of elliptical galaxies with integral field spectrographs, like SAURON. \\\\ ",
        "introduction": "N-body simulations of mergers of disk galaxies have been very succesful in reproducing many observational features of low and intermediate luminosity elliptical galaxies (\\citealp{1972ApJ...178..623T}; \\citealp{1992ApJ...393..484B}; \\citealp{1992ApJ...400..460H}), while the most massive ellipticals are probably a result of collisionless early-type mergers or mixed mergers (\\citealp{2003ApJ...597L.117K}; \\citealp{2006ApJ...636L..81N}). However, the observational information obtained from  galaxies has been greatly enhanced with the advent of Integral Field Unit (IFU) instruments such as SAURON \\citep{2001MNRAS.326...23B}, OASIS and GMOS. While mostly long-slit data was used to constrain the kinematic properties of early-type galaxies until recently, 2-dimensional data contains more information, and makes e.g. the detection of a second distinct kinematic component, which might be aligned with an angle not covered by a slit in a conventional observation, much easier or even possible at all.  The SAURON collaboration presented the first representative sample of 2-dimensional observations of Elliptical and S0 galaxies (\\citealp{2004MNRAS.352..721E}, henceforth EM04). They presented 2D maps of line-of-sight velocities, dispersions and the third and fourth Gauss-Hermite moments. These maps showed a multitude of features such as kinematically decoupled components (KDCs, see also \\citealp{2001ApJ...548L..33D}), central velocity dispersion dips, counter-rotating disks  and many others. Dynamical modelling of the kinematic data revealed that galaxies with significant rotation are more anisotropic than slow-rotating ellipticals \\citep{2005nngu.confE...5C}. An interesting new correlation was found that the slow-rotating ellipticals host large KDCs which have stellar populations of the same age as the main body, in contrast to fast-rotating ellipticals which have much smaller KDCs and are younger on average \\citep{2006astro.ph..9452M}. This correlation must contain some important information on the formation history of early-type galaxies.  Numerical studies of merging galaxies in the past showed that many peculiar kinematical features occur naturally in merger remnants, especially when a dissipational component was present in the progenitor galaxies. The dissipative component which can cool falls to the center and  changes the morpholgy of the remnant drastically by re-arranging its orbital content \\citep{1996ApJ...471..115B}. Disk like counter-rotating cores can also be formed from infalling gas, which has a spin vector with the opposite sign than the main rotating body \\citep{1991Natur.354..210H}, however also dissipationless scenarios have been proposed for counter-rotating stellar populations in early-type galaxies (\\citealp{1998ApJ...505L.109B};\\citealp{1990ApJ...361..381B}; \\citealp{2000MNRAS.316..315B}). \\citet{2002MNRAS.333..481B} showed that the gaseous material which does not fall into the center can form large scale gas discs which can be warped or form bars. Polar-ring like features have been succesfully produced by several authors either in binary mergers \\citep{1998ApJ...499..635B}, tidal accretion events \\citep{2003A&A...401..817B} or in a cosmological context \\citep{2006ApJ...636L..25M}.  The aim of this work will be to compare 2-dimensional observations of a representative sample of 1:1 and 3:1 disc-disc merger remnants to the observations of the SAURON galaxy sample (EM04). The only similar study of this type was carried out by \\citet{2000MNRAS.316..315B}, who identified peculiar  2-dimensional kinematic features in collisionless merger models, such as as an orthogonally decoupled core, counter-rotating populations and misaligned rotation. Their sample contained equal-mass mergers as well as mergers with mass ratios of 3:1 and they found that equal-mass mergers produce a larger variety of kinematical features. But as pointed out before, collisionless mergers are unlikely to be the only formation channel for early-type galaxies and it is unclear how the gas will influence 2D maps of such remnants. Gas infall also will vary from merging geometry to merging geometry as the graviational torques on the gas will vary \\citep{2002MNRAS.333..481B}.  We will examine the merger sample of \\citet{2006MNRAS.tmp..996N}, henceforth NJB06, which includes for every merging geometry a collisionless merger and a merger where each progenitor galaxy had 10\\% of its disc mass in gas. In total we study a sample of 96 mergers which should hopefully give us a similar variety of kinematic features than found in the SAURON sample. As we can compare the mock obsevations of each collisionless remnant with its gaseous counterpart, we can easily assess which kinematical feature was already present in the dissipationless remnant and which was caused by the influence of the gas..  Although 2-dimensional features can many times be intuitively grasped by visual inspection of the maps, this is hardly a quantifiable way to compare to simulations. Recently \\citet{2006MNRAS.366..787K}, henceforth K06, developed a method to describe the properties of 2-dimensional maps of arbitrary moments of the line-of-sight-velocity distribution(LOSVD). They termed this method {\\it kinemetry} which works much in analogy to the common photometric fitting of isophotes. With this method one can determine, for example, in an easy way position angle, amplitude and shape of the rotation of a galaxy, or merger remnant,  and identify kinematic subsystems more easily. The shape of the 2-dimensional maps are of course viewing angle dependant, and we try to take this into account by analysing the projections along the three principle axes. However, we are not carrying out a comprehensive statistical study in this work. We will use kinemetry throughout this work to quantify the kinematic features which arise through merging of disc galaxies.  In Section \\ref{sec:simu} we introduce the simulation sample. We explain the observational method to obtain 2D maps of various moments of the LOSVD in Section \\ref{sec:2dana}. The maps and their kinemetric analysis are presented in Section \\ref{sec:results}. We compare these results to 2D-kinematics of observed galaxies in Section \\ref{sec:observ} and summarize our findings in Section \\ref{sec:2dconcl}. Velocity maps and kinemetry of the whole sample are shown in the Appendix. ",
        "conclusions": "\\label{sec:2dconcl}  We presented 2D maps of various moments of the LOSVDs of a large sample of 1:1 and 3:1 disc merger remnants with and without gas. Every remnant was resimulated with a dissipational component containing 10\\% of the luminous mass, allowing us to assess the influence of gas on the 2D fields. Additionally we performed a kinemetric analysis using the method devised by K06 to quantify properties such as kinematic position angle or deviations from regular rotation.  The difference between equal-mass and an unequal-mass merger remnants is not only seen in photometric and global kinematic properties \\citep{1999ApJ...523L.133N}, but also in the 2D kinematics \\citep{2000MNRAS.316..315B}, which exhibits very different features. 1:1 mergers can lead to a merger remnant with low rotation. Also, orbit classes rotating around the major axis of the remnant can only be populated in significant numbers in 1:1 mergers, leading to kinematic misalignment and twists, while 3:1 remnants rotate much faster and have almost no kinematic twists. All rotating collisionless mergers fail to reproduce the observed $v-h_3$ anti-correlation, because the population of box orbits in the center is not inhibited (NJB06).  The presence of a dissipative component of only 10\\% of the luminous disc mass in the progenitors can lead to a considerable change in the properties of 1:1 remnants. They are more round, kinematic misalignment is reduced, counter-rotating cores are forming and kinematic twist are more pronounced. Rare but observed features like polar gas rings form also in equal mass mergers. The effect of gas on 3:1 remnants is less dramatic. KTs, CRCs and polar rings are not formed or at least must be very rare. A noticeable change is that the opening angles of the iso-velocity contours become smaller. As shown in NJB06, the $v-h_3$ anti-correlation is now much better reproduced, especially if the gas component is included in the analysis.  Merger remnants show a variety of velocity dispersion features. Significant central velocity dispersion dips are caused by infalling gas. The central star component in the gaseous runs, however, is heated. Edge on gas discs can exhibit lower velocity dispersion (as found in 3:1 remnants) or edge-on gas rings can produce double-peaked velocity dispersion maps (as found in 1:1 remnants). Gaseous structures  can imprint a ring-like  depression in $\\sigma$ even on the stellar component, which is also visible in $h_4$.  2D analysis of LOSVDs of N-Body simulations of galaxy formation  are an additional tool to connect the formation history of galaxies with their observable kinematic fine structure. We defer a full statistical analysis to future work, with a larger sample of galaxies mergers, which will also take star formation into account.  The simulations and analysis presented her are a step forward in understanding the possible formation of ellipticals by mergers of discs. We focus on the influence of a small dissipative component, which is important to understand its impact on the dynamical properties of the remnants. However, additional physical processes like star-formation and feedback will change the detailed properties especially at the center of the remnants (\\citealp{1996ApJ...464..641M};\\citealp{2005ApJ...620L..79S}) and definitely have to be considered for simulations with higher gas fractions (\\citealp{2004ApJ...606...32R};\\citealp{2005ApJ...622L...9S}). We believe that the global effects of dissipation even with star formation will be very similar to the results presented here. However, larger gas fractions might result in more massive stellar discs in the remnant. From a kinematic point of view disc merger remnants appear very similar to observed ellipticals, although questions regarding the age and metalicity of the stellar populations have to be addressed in the future. Analytical models of disc formation \\citep{2006MNRAS.366..899N} as well as semi-analytical modeling \\citep{2005astro.ph..9375K} in combination with merger simulations \\citep{2006ApJ...636L..81N} can be used to place further constraints on the disc merger hypothesis."
    },
    "0606/astro-ph0606658_arXiv.txt": {
        "abstract": "Using Particle-In-Cell simulations i.e. in the kinetic plasma description the discovery of a new mechanism of parallel electric field generation was recently reported. Here we show that the electric field generation parallel to the uniform unperturbed magnetic field  can be obtained in a much simpler framework using the ideal magnetohydrodynamics (MHD) description. In ideal MHD the electric field parallel to the uniform unperturbed magnetic field  appears due to fast magnetosonic waves which are generated by the interaction of weakly non-linear Alfv\\'en waves with the transverse density inhomogeneity. In the context of the coronal heating problem  a new {\\it two stage mechanism} of plasma heating is presented by putting emphasis, first, on the generation of parallel electric fields within an {\\it ideal MHD} description directly,  rather than focusing on the enhanced dissipation mechanisms of the Alfv\\'en waves and, second, dissipation of these parallel electric fields via {\\it kinetic} effects.  It is shown that for a single Alfv\\'en wave harmonic  with frequency $\\nu = 7$ Hz, and longitudinal wavelength $\\lambda_A = 0.63$ Mm for a putative  Alfv\\'en speed of 4328 km s$^{-1}$, the generated parallel electric field could account for 10\\% of the necessary coronal heating requirement. We conjecture that  wide spectrum (10$^{-4}-10^3$ Hz) Alfv\\'en waves, based on the observationally constrained spectrum, could provide the necessary coronal heating requirement. By comparing MHD versus kinetic results we also show that there is a clear indication of the {\\it anomalous resistivity} which is 100s of times greater than the classical Braginskii value. ",
        "introduction": "The coronal heating problem, the puzzle of what maintains the solar corona 200 times hotter than the photosphere, is one of the main outstanding questions in solar physics. A significant amount of work has been done in the context of heating of open magnetic structures in the solar corona (e.g. phase mixing, one of the possible mechanisms of the heating). Historically all phase mixing studies have been performed in the Magnetohydrodynamic (MHD) approximation, however, since the transverse scales in the Alfv\\'en wave collapse progressively to zero, the MHD approximation is inevitably violated. Thus, \\citet{tss05a,tss05b} studied the phase mixing effect in the kinetic regime, using Particle-In-Cell simulations, i.e. beyond a MHD approximation, where a new mechanism for the acceleration of electrons due to the generation of a parallel electric field in the solar coronal context was discovered. This mechanism has important implications for various space and laboratory plasmas, e.g. the coronal heating problem and acceleration of the solar wind. It turns out that in the magnetospheric context, a similar parallel electric field generation mechanism in the transversely inhomogeneous plasmas has been previously reported \\citep{glm04,glq99}. See also \\citet{mgl06} and references therein.  After the comment paper by \\citet{mgl06} we came to the realisation that the electron acceleration seen in both series of works \\citep{tss05a,tss05b,glm04,glq99} is a non-resonant wave-particle interaction effect. In works by \\citet{tss05a,tss05b} the electron thermal speed was $v_{th,e}=0.1c$ while the Alfv\\'en speed in the strongest density gradient regions was $v_A=0.16c$; this unfortunate coincidence led us to the conclusion that the electron acceleration by parallel electric fields was affected by the Landau resonance with the phase-mixed Alfv\\'en wave. In works by \\citet{glm04,glq99} the electron thermal speed was $v_{th,e}=0.1c$ while the Alfv\\'en speed was $v_A=0.4c$ because they considered a more strongly magnetised plasma applicable to Earth magnetospheric conditions. However, the interaction of the Alfv\\'en wave with a transverse density plasma inhomogeneity when the Landau resonance condition $\\omega=k_\\parallel c_A(x)$ is met can be quite important for the electron acceleration \\citep{chaston2000,hs76}. \\citet{chaston2000} assert that the electron acceleration observed in density cavities in aurorae can be explained by the Landau resonance of the cold ionospheric electrons with the Alfv\\'en wave. \\citet{hs76} also established that at the resonance Alfv\\'en wave fully converts into the kinetic  Alfv\\'en wave with the perpendicular wavelength comparable to the ion gyro-radius. We can then conjecture that {\\it because of  kinetic Alfv\\'en wave front stretching} (due to phase mixing, i.e. due to the differences in local Alfv\\'en speed), this perpendicular component {\\it gradually realigns} with the ambient magnetic field and hence creates the time varying parallel electric field component. This points to the importance of the Landau resonance for electron acceleration when the resonance condition is met. But as witnessed from works of  \\citet{glm04,glq99}, even when the resonance condition is not met electron acceleration is still possible.   \\begin{figure*} \\centering \\epsfig{file=lk_vx_2.eps,width=6cm} \\epsfig{file=lk_ez_2.eps,width=6cm} \\epsfig{file=lk_vx_20.eps,width=6cm} \\epsfig{file=lk_ez_20.eps,width=6cm} \\caption{Snapshots of $V_x$ and $E_z$ at $t=2$ and 20 for the case of $k=10$, $\\nu = 7$ Hz, $\\lambda_A = 0.63$ Mm.} \\end{figure*}   There were three main stages that lead to the formulation of the present model.  (i) The realisation of the parallel electric field generation (and particle acceleration) being a {\\it non-resonant} wave-particle interaction effect lead us to the question: could such  parallel electric fields be generated in a MHD approximation?  (ii) Next we realised that if one considers {\\it non-linear} generation of the fast magnetosonic waves in the transversely inhomogeneous plasma, then $\\vec E = - (\\vec V \\times \\vec B)/c$ contains a non-zero component parallel to the ambient magnetic field $E_z = - (V_x B_y - V_y B_x)/c$.  (iii) From previous studies \\citep{bank00,tan01} we knew that the fast magnetosonic waves ($V_x$ and $B_x$) did not grow to a substantial fraction of the Alfv\\'en wave amplitude. However after reproducing the old parameter regime (k=1, i.e a frequency of 0.7 Hz), the case of k=10, i.e a frequency of 7 Hz was considered, which showed that  fast magnetosonic waves and in turn a parallel electric field were more efficiently generated. ",
        "conclusions": ""
    },
    "0606/astro-ph0606372_arXiv.txt": {
        "abstract": "We present two-dimensional spectroscopy covering the rest-frame wavelengths of strong optical emission lines in six luminous submillimetre galaxies at $z=1.3$--2.5.  Using this near-infrared integral field spectroscopy together with {\\it Hubble Space Telescope} ACS and NICMOS imaging we map the dynamics and morphologies of these systems on scales from 4--11\\,kpc.  Four of the systems show multiple components in their spatially-resolved spectra with average velocity offsets of $\\sim$\\,180$\\kms$ across 8\\,kpc in projection.  From the ensemble properties of eight galaxies, from our survey and the literature, we estimate the typical dynamical masses of bright submillimetre galaxies as 5$\\pm$3$\\times$10$^{11}\\Msol$. This is similar to recent estimates of their stellar masses -- suggesting that the dynamics of the central regions of these galaxies are baryon dominated, with a substantial fraction of those baryons in stars by the epoch of observation.  Combining our dynamical mass estimates with stellar luminosities for this population we investigate whether submillimetre galaxies can evolve onto the Faber-Jackson relation for local ellipticals.  Adopting a typical lifetime of $\\tau_{\\rm burst}\\sim 300$\\,Myr for the submillimetre-luminous phase -- using the latest estimates of gas masses, star formation rates and AGN contribution to the bolometric luminosities -- we find that the stellar populations of submillimetre galaxies should fade to place them on the Faber-Jackson relation, at M$_K\\sim -25.1$.  Furthermore, using the same starburst lifetime we correct the observed space density of submillimetre galaxies for the duty cycle to derive a volume density of the progenitors of $\\sim1\\times10^{-4}$\\,Mpc$^{-3}$.  This is consistent with the space density of local luminous early-type galaxies with M$_K\\sim -25.1$, indicating that submillimetre galaxies can evolve onto the scaling relations observed for local early-type galaxies, and the observed population at $z\\sim2$ is then sufficient to account for the formation of the whole population of $\\gsim3 L^*_K$ ellipticals seen at $z\\sim 0$. ",
        "introduction": "Deep optical and near-infrared imaging with the superlative resolution of the NICMOS and ACS cameras on-board {\\it Hubble Space Telescope} ({\\it HST}) has made it possible to study the morphologies and colours of high-redshift far-infrared luminous galaxies identified in the sub-millimetre (submm) waveband in unprecedented levels of detail \\citep{Smail98,Smail02,Chapman03b,Smail04,Pope05,Almaini05}.  By combining this data with spectroscopically confirmed redshifts \\citep{Chapman03a,Chapman05a}, we are beginning to understand the processes which triggers the immense bolometric luminosity output in these galaxies, allowing us to address issues such as their true contribution to the star-formation rate history of the Universe \\citep[e.g.,][]{Chapman05a}.  Submm galaxies (SMGs) are amongst the most bolometrically luminous galaxies in the Universe (with $L_{\\rm bol}\\sim$10$^{13}\\Lsol$ and a median redshift of $<z>\\sim$2.2). However, in contrast to quasars at the same early epochs, their bolometric output appears to be dominated by star-formation rather than AGN activity \\citep{Swinbank04,Alexander05a}.  It is clearly important to establish firmly the masses for these galaxies in order to understanding the rapid evolution of the submm population.  Such diagnostics will allow us to determine how they relate to present-day galaxies and in particular to test whether they represent the formation phase of the most massive elliptical galaxies at the present-day, as many suspect \\citep{Lilly99}.  To probe the dynamical structures of these frequently complex systems requires a reliable separation of the spatial and spectral information. One particularly powerful approach to achieve this is using millimetric CO emission line maps to trace the distribution and kinematics of dense gas within the galaxies \\citep{Frayer99,Frayer03,Neri03,Greve05} especially at high resolution \\citep{Genzel03,Downes03,Tacconi06}. Unfortunately, prior to the commissioning of the Atacama Large Millimetre Array, this approach remains observationally very demanding, typically requiring 30 hours per source with the IRAM interferometer and can only be applied to the most massive and gas-rich SMGs.  An alternative tool which is less demanding in terms of telescope time and can provide spatially-resolved dynamics of SMGs with sufficient spatial and velocity resolution to study their internal structure is near-infrared integral field spectroscopy targeting the rest-frame optical emission lines such as [O{\\sc iii}]$\\lambda$5007 and H$\\alpha\\lambda$6563 \\citep[e.g.\\,][]{Tecza04,Swinbank05b}.  The spectroscopic maps so produced allow us trace the dynamical and structural properties of SMGs on scales of a few kpc, pin-point the sites of active star formation and identify non-thermal emission from active galactic nuclei (AGN) in components within these systems. However, they are also subject to potential biases due to extreme extinction within some regions of the SMGs and from outflows in the emission-line gas which is being observed \\citep{Bower04,Wilman05}. While the star formation rates estimated from the H$\\alpha$ are typically a factor of 10--100\\,$\\times$ less than that predicted from the far-infrared emission (suggesting to 1--4 magnitudes of dust absorption; q.v. \\citealt{Smail04}), it is clear that the H$\\alpha$ emission provides a more secure probe of the star-formation rates and dynamics than either the Ly$\\alpha$ emission or rest-frame UV continuum and can thus provide a reliable and observationally efficient probe of the activity and dynamics of this important population.  Using the UIST near-infrared Integral Field Unit (IFU) at UKIRT and the GNIRS IFU at Gemini-South we have studied the rest-frame optical emission line properties of six submillimetre galaxies at $z=1.3$--2.5. These observations, together with high-resolution imaging from {\\it HST} allow us to probe the physical cause of the rapid evolution of submillimetre galaxies, enabling us to test whether these far-infrared luminous galaxies comprise merging systems which are likely to be the progenitors of local massive ellipticals, or whether instead they are simply high-luminosity episodes in the history of more mundane galaxies. Our paper is laid out as follows: in \\S\\ref{sec:obs_red} we present the observations and their reduction.  In \\S\\ref{sec:analysis} we discuss the analysis of these data, while in \\S\\ref{sec:discussion} and \\S\\ref{sec:conc} we present our discussion and conclusions respectively.  We use a cosmology with $H_{0}=70\\kms$, $\\Omega_{M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ in which 1$''$ corresponds to 8.2\\,kpc at $z=2.5$ and 8.5\\,kpc at $z=1.5$.   \\begin{figure*}  \\begin{minipage}{7.5in} \\centerline{\\psfig{figure=f1.ps,width=6.5in,angle=0}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent{\\bf Figure~1.} {\\it Left:} True colour {\\it HST} ACS/NICMOS $B_{450}I_{814}H_{160}$-band image of CFRS\\,03.15 at $z$=1.4076 showing the complex morphology (the colour image is displayed in log scale to emphasise the morphology).  {\\it Middle:} {\\it HST} NICMOS $H_{160}$-band image of CFRS\\,03.15 with the H$\\alpha$ emission line map from the GNIRS IFU observations overlaid as contours, showing that the bright knot located $\\sim$1.3$''$ ($\\sim$11\\,kpc) to the South-West is associated with the system.  The solid bar represents the FWHM of the seeing disk for the IFU observations.  {\\it Right:} Spectra from the two components seen in the {\\it HST} imaging of CFRS\\,03.15 from the Keck/NIRSPEC observation.  The top spectrum is offset in flux for clarity.  From our spectroscopy, the compact component to the South-West is blue-shifted by 90$\\pm$20$\\kms$ in projection from the nucleus, suggesting they are two components whose interaction has promoted the strong starburst and disturbed morphology.  It is interesting to note that both spectra have strong [N{\\sc iii}]/H$\\alpha$ emission line flux ratios which may indicate that both components host active AGN.  }} \\end{minipage}  \\begin{minipage}{7.5in} \\centerline{ \\psfig{figure=f2.ps,width=4.0in,angle=90} \\psfig{figure=f3.ps,angle=0,width=2.3in}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent{\\bf Figure~2.}  {\\it Left:} A true colour image of N2\\,850.7 comprising {\\it HST} ACS $I_{814}$-band and NICMOS $H_{160}$-band images illustrating the complex morphology of this system.  {\\it Middle:} NICMOS $H_{160}$-band image with H$\\alpha$ emission line intensity overlaid as contours.  Each panel is 4\\,arcsec square (34\\,kpc at $z=1.49$) and has East left and North top.  The bar shown on the NICMOS image indicates the 0.6$''$ resolution of the IFU observations.  The system comprises at least three components, marked as A, B and C (and there is also tentative evidence for weak H$\\alpha$ emission from the source $\\sim1''$ to the West, suggesting it may be associated with the system).  The continuum colours of C are bluer than the other components.  The intensity of the H$\\alpha$ emission traces the $H_{160}$-band morphology and it is interesting to note that the star formation appears distributed across the SMG. {\\it Right:} Spectra covering the redshifted H$\\alpha$ emission in the three components of N2\\,850.7 from the UIST IFU observations. The spectra have been offset in flux scale for clarity.  The lower panel shows a sky spectrum.  By centroiding the emission lines, we find velocity offsets between components A \\& B and A \\& C of $215\\pm80\\kms$ and $135\\pm90\\kms$ in projection respectively.  }} \\end{minipage} \\vspace{0.5cm}  \\begin{minipage}{7.5in} \\centerline{ \\psfig{figure=f4.ps,width=3.5in,angle=90} \\psfig{figure=f5.ps,angle=90,width=3.9in}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent{\\bf Figure~3.}  {\\it Left:} {\\it HST} $I_{814}$-band image of N2\\,1200.18 with the [O{\\sc iii}]$\\lambda$5007 and H$\\alpha$ emission line intensity maps from the UIST IFU observations overlaid as contours. Both panels are displayed on a $\\log$ scale to emphasis the morphology.  The {\\it HST} image shows a a central bright, compact component and a lower-surface brightness extension to the North.  The [O{\\sc iii}]$\\lambda$5007 and H$\\alpha$ emission also have different morphologies, with the strong, compact [O{\\sc iii}]$\\lambda$5007 probably arising from (unresolved) AGN activity in the bright central component, whilst the H$\\alpha$ morphology appears resolved and extends across the whole system.  {\\it Right:} Spectra from the central component and northern extension of N2\\,1200.18 around the [O{\\sc iii}]$\\lambda$5007 and H$\\alpha$ emission lines from the UIST IFU observations.  The spectrum of the central source has an [O{\\sc iii}]/H$\\alpha$ emission line flux ratio of 2.0$\\pm$0.5 which, combined with the unresolved [O{\\sc iii}] emission suggests AGN activity.  The northern component has a velocity offset of 250$\\pm$80$\\kms$ from the nucleus.  This northern component has star-burst (rather than AGN) characteristics, with a lower [O{\\sc iii}]/H$\\alpha$ and [N{\\sc ii}]/H$\\alpha$ emission line flux ratios.  {\\it Bottom:} The OH airglow emission from the sky (continuous line), as well as the Mauna Kea atmospheric transmission (dotted line) are also shown in the lower panel.  }} \\end{minipage}  \\end{figure*}  \\begin{figure*} \\begin{minipage}{7.5in} \\centerline{\\psfig{figure=f6.ps,width=6in,angle=90}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent{\\bf Figure~4.}  {\\it Left:} {\\it HST} I$_{814}$-band image of N2\\,1200.90.  The galaxy has a compact morphology ($\\ls$0.2$''$ (1.6\\,kpc) FWHM) in this continuum imaging.  {\\it Right:} Spectrum of N2\\,1200.90 from the UIST IFU observations showing H$\\alpha$ at a redshift of $z=2.198$. The rest-frame UV spectrum from Chapman et al.\\ (2006) shows clear signs of AGN activity, yet the apparently low [N{\\sc ii}]/H$\\alpha$ emission line ratio from our near-infrared spectroscopy is more suggestive of star-formation.  However, we note that the H$\\alpha$ emission line width of 500$\\pm$100$\\kms$ may indicate (non-thermal) AGN activity.  }} \\end{minipage} \\vspace{0.6cm}   \\begin{minipage}{7.5in} \\centerline{ \\psfig{figure=f7.ps,width=5.2in,angle=90} \\psfig{figure=f8.ps,width=2.4in,angle=90}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent{\\bf Figure~5.}  Images of \\b59 generated from the UIST IFU observations.  {\\it Left:} $K$-band continuum image (collapsed from the datacube between 2.0 and 2.4$\\mu$m) with contours overlaid from the $H$-band (collapsed between 1.5 to 1.8$\\mu$m). {\\it Middle:} Continuum subtracted H$\\alpha$ image. {\\it Right:} $K$-band image with contours from the continuum subtracted, H$\\alpha$ narrow-band image overlaid.  The ``X'' in each panel marks the position of the center of the $H$-band light distribution.  In the final panel we show the collapsed, one-dimensional spectrum around the redshifted H$\\alpha$ emission from \\b59.  The H$\\alpha$ redshift of $z$=1.328$\\pm$0.003 is in excellent agreement with the CO(5-4) emission, and the H$\\alpha$ line width and emission line flux are also comparable to those given in Borys et al.\\ (2006).  }} \\end{minipage} \\vspace{0.6cm}  \\begin{minipage}{7.5in} \\centerline{ \\psfig{figure=f9.ps,width=3.0in,angle=90} \\psfig{figure=f10.ps,width=2.8in,angle=90}} \\end{minipage} \\begin{minipage}{7.5in} {\\addtolength{\\baselineskip}{-3pt} {\\scriptsize \\noindent {\\bf Figure~6.}  {\\it Left:} $HST$ $I_{814}$-band image of SA\\,22.96 with the H$\\alpha$ emission from the UIST IFU observations overlaid as contours.  % {\\it Right:} Collapsed spectra around the H$\\alpha$ emission from the Western component ($A$) from the UIST IFU observations of SSA22.96. We identify $A$ as the submm source at redshift $z$=2.517 through identification of strong H$\\alpha$ and [N{\\sc ii}] emission at the same redshift as derived from the rest-frame UV spectroscopy in Chapman et al.\\ (2005).  We note that there is tentative evidence for weak H$\\alpha$ emission from the structure $\\sim1.5''$ to the East, suggesting it may be associated with the system.  }} \\end{minipage} \\vspace{0.6cm}  \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  In this paper we have studied the rest-frame optical emission line structures and dynamics of six high redshift, luminous submm galaxies. Including two sources from the literature with similar observations, we find that of the eight submillimetre galaxies with resolved spectroscopy, at least five appear to comprise two or more dynamical sub-components.  The average velocity offsets between these components is $\\sim180\\kms$ across a projected spatial scale of 8\\,kpc.  The obvious merging/interacting nature of these systems suggests there are analogous to the multi-component nature of the, typically less bolometrically luminous and slightly more compact, ULIRGs in the local Universe.  Since our IFU observations allow us to disentangle AGN and star-burst-like components (e.g.\\ from [N{\\sc ii}]/H$\\alpha$ flux ratios), it is perhaps significant that two of the eight submm galaxies shows possible signs of AGN activity in two spatially-resolved components on 1--2$''$ (10--20\\,kpc) scales.  Deep X-ray surveys of the {\\it Chandra} Deep Field North have produced similar results: \\citet{Alexander03a} report that two of seven submm galaxies in this region are individually associated with pairs of X-ray sources on similar spatial scales (approximately one galactic diameter).  Since the chance of a false association is $\\ls1$\\%, it may be that we are witnessing the interaction or merging of AGNs in these sources (a low redshift example of this binary AGN activity is seen in the ULIRG NGC\\,6240; \\citealt{Komossa03}).  The double AGN in some of these systems, suggests mergers may be responsible for disturbing the gas orbits in both components, hence fuelling subsequent AGN growth which currently lags behind a significant pre-existing stellar population \\citep{Alexander03a,Smail03a}.  By combining the spatially-resolved kinematic information from IFU, CO and longslit observations and assuming the velocity offsets we measure are eventually transformed into velocity dispersions in a pressure-supported galaxy, we derive dynamical masses of 5$\\pm$3$\\times$10$^{11}\\Msol$ for typical luminous submillimetre galaxies.  This is similar to the dynamical masses found using H$\\alpha$ line widths \\citep{Swinbank04} and resolved CO spectroscopy \\citep{Greve05,Tacconi06} and a factor of $\\sim 2\\times$ greater than the recent stellar mass estimates from \\citet{Borys05}, suggesting that the central regions of these galaxies are baryon dominated.  We combine our dynamical measurement with estimates of the stellar luminosities for this population from \\citet{Borys05}. We use these to constrain their likely evolution and so compare the properties expected for the descendents of submillimetre galaxies with $z=0$ galaxy populations.  We update the gas depletion time scales estimated by \\citet{Greve05} using recent results on the gas masses, bolometric luminosities and AGN contributions for this population, deriving typical gas depletion times of $\\sim 150$\\,Myr.  Adopting a typical burst duration of twice this age, $\\sim 300$\\,Myr, we predict the evolution of their stellar populations and compare these to the Faber-Jackson relation for local ellipticals.  For a $\\sim 300$\\,Myr burst with a Kroupa IMF, we find that the expected fading would result in the descendents of the SMGs lying within $\\sim 0.2$\\,mags of the FJ relation at their measured velocity dispersion.  The same estimate of the lifetime can be used to correct the apparent space density of SMGs (assuming a galaxy only goes through one SMG phase in its lifetime) to derive a space density for the descendents of this population of $\\sim 1.3\\times10^{-4}$Mpc$^{-3}$ for a population with $z=0$ absolute luminosities of M$_K\\sim -25.1$. Hence if SMGs have burst durations of $\\sim 300$\\,Myr, as indicated by gas consumption timescales, then their predicted fading places them close to the present-day FJ relation.  Moreover, the observed population at $z\\sim2$ is very comparable to that of luminous ellipticals with M$_{K}\\lsim -25.1$ ($\\sim 3 L^{*}_{K}$) at the present day.  Our observations show the power of using integral field spectroscopy around the rest-frame optical emission lines to probe the dynamics and power sources of these galaxies on $\\gs$5\\,kpc scales.  With forthcoming adaptive optics integral field spectroscopy on 8- and 10-m telescopes (e.g.\\ OSIRIS on Keck, NIFS on Gemini and SINFONI on the VLT), we will soon be able to probe the dynamics, metallicities and star-formation properties on $\\ls$0.1$''$ (sub-kpc at $z=2$) scales. Furthermore, the increased light grasp of large aperture telescopes will allow us to spatially resolve and map the absorption line spectra of these galaxies (including important diagnostic lines such as Ca{\\sc ii} H\\&K, CH G-band, H$\\delta$ and MgB), allowing us to accurately probe the stellar contents and metallicity gradients, and so investigate the stellar ages and populations produced in these extreme systems.  Such observations will allow us to probe further the timescales involved in the encounters as well as how and when these galaxies form most of their stars and thus the precise relation to similarly massive galaxies at the present day: luminous ellipticals."
    },
    "0606/astro-ph0606414_arXiv.txt": {
        "abstract": "About 2 million seconds of \\chandra\\ observing time have been devoted to the Galactic center (GC), including large-scale surveys and deep pointings. These observations have led to the detection of about 4000 discrete X-ray sources and the mapping of diffuse X-ray emission in various energy bands. In this review, I first summarize general results from recent studies and then present close-up views of the three massive star clusters (Arches, Quintuplet, and GC) and their interplay with the Galactic nuclear environment. ",
        "introduction": "While Sgr A* itself is only weakly active at present, much of the high-energy activity in the GC is initiated apparently by the three young massive stellar clusters, located  within 50 pc radius of the super-massive black hole Sgr A* (Fig.~\\ref{f:global}): Arches [with an age of $(2-3) \\times 10^6$ yrs], Quintuplet [$(3-6) \\times 10^6$ yrs], and GC [$(3-7) \\times 10^6$ yrs] (e.g., Figer et al. 2004; Stolte et al. 2002, 2005; Genzel et al. 2003). Massive stars themselves can be moderately bright X-ray sources (e.g., colliding stellar wind binaries). Such stars also release large amounts of mechanical energy in form of fast stellar winds and supernovae, heating and shaping the surrounding interstellar medium (ISM) and affecting the accretion of the black hole. Furthermore, stellar end-products of massive stars (neutron stars and black holes) can also be strong X-ray sources. X-ray observations are thus a powerful tool for probing such high-energy phenomena and processes.  The GC is a region where the spatial resolution of X-ray observations matters. This is why \\chandra\\ has invested heavily on the GC, conducting repeated large-scale raster surveys [existing 360 ks (Wang et al. 2002a) and upcoming 600 ks (PI: Muno)] and numerous deep pointed observations: about 1 Ms on Sgr A (Baganoff et al. 2003), 100 ks each on Sgr B (Takagi et al. 2002), Sgr C (PI: Murakami), and Arches (Wang et al. 2006a), as well as 50 ks on Radio Arc (Yusef-Zadeh et al. 2002a). The main instrument used in these \\chandra\\ observations is the ACIS-I, which covers a field of $17^\\prime \\times 17^\\prime$ and an energy range of 0.5-10 keV. But  below $\\sim 2$ keV, X-rays from the GC are heavily absorbed by the ISM. The spatial resolution ranges from $\\sim 1^{\\prime\\prime}$ on-axis to $\\sim 10^{\\prime\\prime}$ at the outer boundaries of the field.  While much of the data analysis is still ongoing, here I will first summarize some of the general results on detected discrete sources and diffuse X-ray emission and will then focus on the three massive star clusters, highlighting various high-energy phenomena and processes involved. ",
        "conclusions": ""
    },
    "0606/astro-ph0606622_arXiv.txt": {
        "abstract": "{}{% An interesting question of contemporary cosmology concerns the relation between the spatial distribution of galaxies and dark matter, which is thought to be the driving force behind the structure formation in the Universe.  In this paper, we measure this relation, parameterised by the linear stochastic bias parameters, for a range of spatial scales using the data of the Garching-Bonn Deep Survey (GaBoDS).}{% The weak gravitational lensing effect is used to infer matter density fluctuations within the field-of-view of the survey fields. This information is employed for a statistical comparison of the galaxy distribution to the total matter distribution.  The result of this comparison is expressed by means of the linear bias factor $b$, the ratio of density fluctuations, and the correlation factor $r$ between density fluctuations.  The total galaxy sample is divided into three sub-samples using $R$-band magnitudes and the weak lensing analysis is applied separately for each sub-sample.  Together with the photometric redshifts from the related COMBO-17 survey we estimate the typical mean redshifts of these samples with \\mbox{$\\bar{z}=0.35, 0.47, 0.61$}, respectively.}{% Using a flat $\\Lambda\\rm CDM$ model with \\mbox{$\\Omega_{\\rm m}=0.3, \\Omega_\\Lambda=0.7$} as fiducial cosmology, we obtain values for the galaxy bias on scales between \\mbox{$1^\\prime\\le\\theta_{\\rm ap}\\le20^\\prime$}. At $10^\\prime$, the median redshifts of the samples correspond roughly to a typical comoving scale of \\mbox{$3,5,7~h^{-1}\\rm Mpc$} with \\mbox{$h=0.7$}, respectively. We find evidence for a scale-dependence of $b$.  Averaging the measurements of the bias over the range \\mbox{$2^\\prime\\le \\theta_{\\rm ap}\\le 19^\\prime$} yields $\\bar{b}=0.81\\pm0.11, 0.79\\pm0.11, 0.81\\pm0.11$ ($1\\sigma$), respectively.  Galaxies are thus less clustered than the total matter on that particular range of scales (anti-biased). As for the correlation factor $r$ we see no scale-dependence within the statistical uncertainties; the average over the same range is $\\bar{r}=0.61\\pm0.16, 0.64\\pm0.18, 0.58\\pm0.19$ ($1\\sigma$), respectively.  This implies a possible decorrelation between galaxy and dark matter distribution. An evolution of galaxy bias with redshift is not found, the upper limits are: $\\Delta b\\lesssim0.2$ and $\\Delta r\\lesssim 0.4(1\\sigma)$.}{} ",
        "introduction": "In comparison to the total mass in the Universe, galaxies take -- considering their mass -- only a minor part in the big picture of structure formation.  \\stressa{They formed from the baryonic component, embedded in the fluctuations of the dark matter density field, whose total mean density is very much lower than that of dark matter.}  Due to their relatively easy observability, it would be very convenient if galaxies were perfect tracers -- unbiased tracers -- of the total mass distribution; all statistical properties of the mass structure could then be derived from galaxy catalogues.  Indeed, it is rather unlikely that galaxies are unbiased tracers, because the laws determining the galaxy distribution are very complex and highly non-linear.  The primordial gas from which they form requires special conditions to be able to cool and fragment into galaxies (White~\\&~Frenk 1991; White~\\&~Rees 1978). Due to shock heating of the baryons and energy feedback between galaxies and the intergalactic medium, the properties of the gas feeding galaxy formation gradually changed with time.  Furthermore, galaxies interact with each other or with the baryonic intergalactic medium, merge or get accreted into more massive galaxies (Lacey~\\&~Cole 1993).  These mechanisms probably produced the large diversity in galaxy masses, colours, morphologies and chemistry we observe today.  Based on our current knowledge it would be very surprising if this complexity would eventually result in a simple, linear, one-to-one relationship between the galaxy density and total matter density, making galaxies unbiased tracers.  But by studying this dark matter-galaxy relationship we can learn more about galaxies.  Observing the relation between the invisible dark matter field and the galaxies is a particularly tough problem.  However, with gravitational lensing at hand, we now have a technique to directly unravel this relationship. The importance of ``cosmic shear'' as a tool for cosmology was proposed in the early 1990s by Blandford et al. (1991), Miralda-Escud\\'e (1991) and Kaiser (1992). Since these pioneering days of gravitational lensing the techniques and surveys have been refined to make valuable contributions to the ongoing research on structure formation on cosmological scales. In particular, the investigation of the relation between galaxy and dark matter distribution using the weak gravitational lensing effect has become almost standard (Seljak et al.  2005; Kleinheinrich et al.  2005; Sheldon et al. 2004; Hoekstra et al.  2003; Guzik~\\&~Seljak 2001; McKay et al. 2001; Fisher et al. 2000; Hudson et al. 1998; Brainerd et al. 1996). This paper will focus on the lensing technique as well.  \\subsection{Quantifying galaxy bias}  From the point of view of statistics, quantifying galaxy bias leads to the question how one can parametrise differences in the statistical properties -- not the obvious differences between two particular realisations -- of two random fields. Both the distribution of galaxies and the distribution of dark matter are thought to be realisations of statistically homogeneous and isotropic random fields. Commonly, one uses a parametric way to describe the biasing between two (random) density fields, for instance galaxy and matter distribution or the distributions of two different galaxy populations. Biasing between two density fields, say $\\rho_{\\rm g}$ and $\\rho_{\\rm m}$, can in general be quantified using the joint probability distribution function (PDF) $P\\left(\\delta_{\\rm g},\\delta_{\\rm m}\\right)$ of the density contrasts (density fluctuations) \\begin{equation} \\delta_{\\rm g}\\equiv\\frac{\\rho_{\\rm g}}{\\ave{\\rho_{\\rm g}}}-1~~;~~ \\delta_{\\rm m}\\equiv\\frac{\\rho_{\\rm m}}{\\ave{\\rho_{\\rm m}}}-1 \\end{equation} at the same point in the density fields at some redshift (local Eulerian bias). The probability of finding density contrasts of $\\delta_{\\rm g}$ and $\\delta_{\\rm m}$ within an interval of $\\d\\delta_{\\rm g}$ and $\\d\\delta_{\\rm m}$, respectively, equals $P(\\delta_{\\rm g},\\delta_{\\rm m})\\d\\delta_{\\rm m}\\d\\delta_{\\rm g}$. The density contrasts are smoothed to a certain scale, $R$, before investigating their PDF.  Often the special type of smoothing kernel is set by the method that is used to determine the bias.  To lowest, second order (first order moments vanish by definition, $\\ave{\\delta_{\\rm g}}=\\ave{\\delta_{\\rm m}}=0$) the only relevant parameters for biasing are the scale-dependent parameters of the ``linear stochastic bias'' \\begin{equation}\\label{biascoeff} b(R)=\\sqrt{\\frac{\\ave{\\delta_{\\rm g}^2}}{\\ave{\\delta_{\\rm m}^2}}} ~~;~~ r(R)=\\frac{\\ave{\\delta_{\\rm g}\\delta_{\\rm m}}}{\\sqrt{\\ave{\\delta_{\\rm g}^2}\\ave{\\delta_{\\rm m}^2}}} \\; , \\end{equation} which differ from unity in the case of two biased fields. These two parameters provide a complete description for the bias between $\\delta_{\\rm g}$ and $\\delta_{\\rm m}$ for Gaussian fields only and therefore are a full description only on large smoothing scales where linear perturbation theory applies (on ``linear scales''). The linear bias factor $b$ is a measure for the difference in clustering strength. The correlation coefficient $r$, on the other hand, measures partly the stochasticity in the relation between the density contrasts, for instance how well minima and maxima of the density fields coincide. Only partly, because the correlation coefficient is also sensitive to the non-linearity in this relation, which becomes relevant on smaller smoothing scales where the fields are usually no longer Gaussian. In fact, they cannot be Gaussian since \\stressa{$\\delta\\ge-1$} by definition. Despite this degeneracy, the linear stochastic bias parameters are clearly defined on non-linear scales; it is just their interpretation which is no longer straightforward. To disentangle a non-linear but deterministic relation between two density contrasts from a stochastic one requires higher-order statistics, like for example in the context of the ``non-linear stochastic bias'' scheme (Dekel~\\&~Lahav 1999).  \\subsection{Galaxy bias in observations}  Observationally, galaxy bias can be derived from the one-dimensional PDF of galaxies, $P(\\delta_{\\rm g})$, (Marinoni et al. 2005; Sigad et al.  2000), redshift space-distortions (Pen 1998; Kaiser 1987), weak gravitational lensing (Seljak et al. 2005; Sheldon et al. 2004; Pen et al. 2003; Hoekstra et al.  2002; Wilson et al. 2001; van Waerbeke 1998; Schneider 1998) and counts-in-cells statistics (Conway et al. 2005; Tegmark~\\&~Bromley 1999; Efstathiou et al. 1990). Additionally, the large-scale flow of galaxies can be used to make a POTENT reconstruction of the total mass field on large scales which can be compared to the galaxy distribution (Sigad et al. 1998; Dekel et al. 1993). Gravitational lensing (Schneider et al. 2006; van Waerbeke~\\&~Mellier 2003; Bartelmann~\\&~Schneider 2001) provides a promising new method in this respect because it allows for the first time to map the total matter content (mainly dark matter) independent from the galaxy distribution. The work of this paper is based on this technique.  A brief overview of the current status of the observational results is given in the following. Note that the given conclusions to some extent depend on the assumed cosmological model. We quote only the conclusions for the concordance $\\Lambda\\rm CDM$ model (cf.  Tegmark et al. 2004). In the local universe, $L_\\star$ galaxies are almost unbiased tracers on linear scales of about $8h^{-1}\\rm Mpc$ and larger (Seljak et al. 2005; Verde et al.  2002; Lahav et al. 2002; Loveday et al.  1996).  However, this is probably not true on smaller scales. A comparison of the theoretical dark matter clustering -- which is constrained by the cosmic microwave background anisotropies, gravitational lensing and the Lyman-$\\alpha$ forest (Tegmark et al. 2004) -- and the observable galaxy clustering suggests that on smaller scales $\\sim1h^{-1}\\rm Mpc$ galaxies are less clustered than the dark matter (``anti-biased'') becoming positively biased, $b>1$, on even smaller scales below $\\sim0.1h^{-1}\\rm Mpc$. Hoekstra et al.  (2001, 2002) use in their work on the VIRMOS-DESCART survey (van Waerbeke et al.  2001) and RCS (Gladders~\\&~Yee 2001) weak gravitational lensing to measure the linear stochastic bias for galaxies with a median redshift of $z_{\\rm m}=0.35$, covering a range from $0.1h^{-1}{\\rm Mpc}$ to $6.3h^{-1}{\\rm Mpc}$.  They claim to have observed such a dip in the linear bias factor.  Also based on gravitational lensing there is evidence that the ratio $b/r$ stays close to unity from sub-megaparsec scales up to \\mbox{$\\sim8h^{-1}\\rm Mpc$} (Sheldon et al.  2004; Hoekstra et al.  2002; Guzik~\\&~Seljak 2001; Fisher et al. 2000), thus from non-linear to linear scales. The analysis of the bispectrum of the galaxy clustering in the 2dFGRS (Colless et al. 2001) led Verde et al. (2002) to the conclusion that on scales between $5h^{-1}\\rm Mpc$ and $30h^{-1}\\rm Mpc$ the biasing relation between dark matter and galaxies is essentially linear (see also Lahav et al.  2002). The same conclusion was drawn several years earlier by Gazta\\'{n}aga \\& Frieman (1994) based on the APM survey (Maddox et al. 1990).  However, recently the work of Kayo et al.  (2004) has questioned a strict linear relation on scales \\mbox{$\\lesssim10h^{-1}\\rm Mpc$} by studying the three-point correlation of galaxy clustering as a function of morphology, colour and luminosity, this time in the SDSS (York et al. 2000).  Subdividing the galaxies into various subsets gives a more detailed picture of galaxy biasing. At low redshift, clustering is a function of morphological type, spectral type, colour and luminosity (e.g. Madgwick et al. 2003;  Zehavi et al.  2002; Norberg et al.  2001; Benoist et al.  1996; Tucker et al. 1997; Loveday et al.  1995; Davis~\\&~Geller 1976).  Late-type, blue, spiral or star forming galaxies are less clustered than early-type, elliptical or red galaxies with a relative linear bias factor of about \\mbox{$b_{\\rm red}/b_{\\rm blue}\\approx1.4$} on scales of roughly $8\\,h^{-1}\\rm Mpc$ (e.g.  Wild et al.  2005; Conway et al.  2005; Zehavi et al. 2002; Norberg et al. 2002; Baker et al. 1998).  On large scales, the relative biasing between red and blue galaxies does not seem to be well described by a simple linear biasing function which according to Wild et al. (2005) (see also Conway et al.  2005) is ruled out with high significance using counts-in-cells statistics in redshift space. Wild et al. observe a scale-dependent non-linear bias between red and blue galaxies with a dominant stochasticity component for typical physical scales of about $7h^{-1}\\rm Mpc$ up to $31h^{-1}\\rm Mpc$. Relative bias between red and blue galaxies is therefore both non-linear and stochastic. It is therefore also expected that at least for some galaxy populations the bias with respect to the dark matter distribution is non-linear and stochastic as suggested by simulations (Yoshikawa et al. 2001). This, however, has not been measured directly so far.  Observational evidence for the relation between red and blue galaxies being non-deterministic was already given some years ago by the work of Tegmark~\\&~Bromley (1999), which was based on the Las Campas Redshift Survey (Shectman et al.  1996), and by Blanton (2000). Moreover, galaxy bias seems to be a function of redshift (Marinoni et al.  2005; Magliocchetti et al.  2000) which is expected both from simulations (Weinberg et al. 2004 and references therein) and analytical models (Tegmark~\\&~Peebles 1998; Mo~\\&~White 1996; Fry 1996).  In this paper, we apply the method of Hoekstra et al. (2002; Sect. \\ref{methodsection}) to the Garching-Bonn Deep Survey (Sect. \\ref{gabodssection}) to obtain the linear stochastic bias coefficients $b$ and $r$ of three galaxy subsets binned by their apparent $R$-band magnitude. The selection of the galaxy samples is outlined in Sect. \\ref{gabodssection}.  After presenting our results in Sect. \\ref{resultsection} we close with a discussion and conclusions in Sect.  \\ref{finalsection}. We will start with a brief introduction to the formalism of weak gravitational lensing and the aperture statistics here employed.  Unless otherwise stated we use a $\\Lambda$CDM model with \\mbox{$\\Omega_{\\rm m}=0.3$}, \\mbox{$\\Omega_{\\Lambda}=0.7$} and \\mbox{$H_{0}=h\\,100\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$} with \\mbox{$h=0.7$}.  A scale-invariant (\\mbox{$n=1$}, Harrison-Zel'dovich) spectrum of primordial fluctuations is assumed.  As transfer function, encoding the physical properties of the cold dark matter fluid, we use Bardeen et al. (1986). ",
        "conclusions": "\\label{finalsection}  Observationally, the galaxy-dark matter bias can be probed by means of various methods (see introduction). Gravitational lensing provides a promising new method in this respect.  It is special because it allows for the first time to map the total matter content (mainly dark matter) with a minimum of assumptions and independent of the galaxy distribution.  Such a map can be compared to the distribution of galaxies, or particular types of galaxies, in order to investigate the galaxy bias.  In particular, correlations between galaxy and dark matter density become directly visible.  For working out the galaxy-dark matter bias, older methods rely on assumptions regarding the growth of dark matter density perturbations, the peculiar velocities of galaxies and their correlation to the dark matter density.  Moreover, they often only allow one to measure the bias on large (linear) scales, \\mbox{$\\gtrsim 8\\,h^{-1}\\rm Mpc$}, whereas the non-linear regime is also accessible with lensing.  However, gravitational lensing has the disadvantage that it is not equally sensitive at all redshifts.  The cosmic shear signal is most sensitive to matter fluctuations roughly half-way between \\mbox{$z=0$} and the mean redshift of the background. This defines a natural best-suited regime for the method at a redshift of about \\mbox{$z\\approx0.5$}, often even slightly lower, considering the depth of current galaxy surveys. It is expected that the most sensitive regime will be shifted towards higher redshifts by future space-based lensing surveys. Furthermore, lensing observables are quite noisy so that large survey areas are required for a good signal-to-noise. Impressively large surveys with instruments such as the CFHT (CFHT-Legacy-Survey, CFHTLS), the VST (Kilo-Square-Degree-Survey, KIDS), Pan-STARRS, or SNAP are either ongoing or about to start within the next years, providing us with plenty of high signal-to-noise information on dark matter and galaxy clustering.  In this paper, we employed aperture statistics to quantify the relation between the dark matter and galaxy density. We tested the evaluation software against Monte Carlo simulated WFI fields, assuming an unbiased galaxy population, and found that the software is working to at least a few percent accuracy (Simon 2005).  The data used is the GaBoDS with restriction to galaxies brighter than $24\\,\\rm mag$ in the $R$-band; this allowed us to estimate the redshift distribution of the galaxies on the basis of three \\mbox{COMBO-17} fields (A901, AXAF/CDFS and S11) for which photometric redshifts in \\mbox{$0\\le z\\lesssim 1.4$} are available. For all the other fields, only $R$-band magnitudes can be used to select galaxies.  For this selection, we defined foreground galaxy samples by choosing galaxies from three $R$-band magnitude bins that have increasingly fainter median magnitudes.  The sample \\mbox{FORE-I} is comparable to the foreground selection in Hoekstra et al. (2002) who applied the same technique as we are using here. By means of the photometric redshifts of the \\mbox{COMBO-17} fields we can translate a GaBoDS R-band magnitude interval into a redshift distribution.  The fainter the bin, the broader the redshift distribution, while the mean redshift moves to larger values.  Therefore, only \\mbox{FORE-I} has a rather sharp peak in redshift, while \\mbox{FORE-III} stretches between redshifts of about \\mbox{$z=0.1$} and \\mbox{$z\\approx 0.9$}.  Hence, \\mbox{FORE-II} and \\mbox{FORE-III} are averages over a relatively wide range of redshifts. In order to get narrower distributions in redshifts with the aim to reconstruct the redshift evolution of biasing, multi-colour lensing surveys are required. \\stressa{Cosmic variance is the main uncertainty in the estimated redshift distribution. Based on the field-to-field variance of the photometric redshift distributions we estimate that this uncertainty translates into a $1\\sigma$-uncertainty of the bias parameters of $\\sim7\\%$, except for the bias factor, $b$, in \\mbox{FORE-I} which has $\\sim16\\%$.}  The B-mode of the aperture statistics $\\ave{NM_{\\rm ap}}$ and $\\ave{M^2_{\\rm ap}}$ are used as an indicator for systematics in the PSF-corrected shapes of the background galaxies (Fig. \\ref{nmapmapsqd}); they cannot be produced by gravitational lensing and should therefore be pure noise.  We find that the B-modes are consistent with zero, maybe leaving a small question mark at $\\theta_{\\rm ap}\\approx3^\\prime$. Note that, at least in principle, physical effects like intrinsic alignments of the source galaxies can be a source of B-modes in $\\ave{M^2_{\\rm ap}}$ (on small scales), so that vanishing B-modes are not the ultimate indicators of PSF systematics. For $\\ave{NM_{\\rm ap}}$, however, the only possible source of B-modes is a violation of a statistical parity-invariance (Schneider 2003). Therefore, $\\ave{NM_{\\rm ap}}$ should always be B-mode free which is clearly the case in our data.  The fit of a theoretical $\\ave{M^2_{\\rm ap}}$ constructed from our fiducial cosmology and redshift distribution of source galaxies to the measured $\\ave{M_{\\rm ap}^2}$ is an important test for the calibration of the bias parameter; $\\ave{M^2_{\\rm ap}}$ is independent of the galaxy bias. Our data points are consistent with the fiducial cosmological model and, therefore, we accept the fiducial cosmological model and the estimated redshift distributions as sufficiently accurate for our purposes. For fiducial cosmological models different from ours (but still flat, \\mbox{$\\Omega_{\\rm m}+\\Omega_\\Lambda=1$}, with negligible baryon density, \\mbox{$\\Gamma=\\Omega_{\\rm m}h$}, and \\mbox{$h=0.7$}, and \\mbox{$\\sigma_8\\propto\\Omega_{\\rm m}^{-0.56}$}) the calibrated bias parameters in Fig. \\ref{biasresult} may be scaled up or down using Fig. \\ref{calibrationfig}.  In a related paper (Hetterscheidt et al. 2006), we discuss in much more detail issues concerning the creation of source galaxy catalogues and their data quality, and we determine constraints on cosmological parameters based on the GaBoDS data. There it can be seen that this cosmic shear analysis supports our adopted fiducial cosmology. \\stressa{An uncertainty in the fiducial cosmology adds an additional uncertainty to the galaxy bias calibration and therefore the inferred bias parameters. For a realistic relative error of $10\\%$ in $\\Omega_{\\rm m}$, we estimate this error to $\\sim7\\%$ for the bias factor, $b$, and to $\\sim2\\%$ for the correlation factor, $r$. Errors given in the following do not include calibration uncertainties.}  The result of the galaxy bias measurement is plotted in Fig. \\ref{biasresult}.  Overall, the galaxy bias factor and the correlation are close to an unbiased population of galaxies, i.e.  \\mbox{$b=1$} and \\mbox{$r=1$}.  A possible scale-dependence is indicated for the bias factor which rises to $b>1$ on scales below \\mbox{$\\theta_{\\rm ap}\\approx 4^\\prime$}, falls below $b=1$ on scales of \\mbox{$\\theta_{\\rm ap}\\approx5^\\prime$} and possibly rises again on larger scales. An aperture radius of $\\theta_{\\rm ap}=4^\\prime$ corresponds to an effective comoving scale of $R=1.4,2.1,2.8\\,h^{-1}\\rm Mpc$ (\\mbox{FORE-I} to \\mbox{FORE-III}) with a relative uncertainty ($1\\sigma$) of about $40\\%$. The origin of this uncertainty is due to the fact that we are actually observing averages of galaxy bias over some redshift (cosmological time) and scale as illustrated by Fig.  \\ref{weight}; the median redshifts for the bias are $\\bar{z}=0.28,0.44,0.59$, respectively.  The median redshift for the correlation parameters are slightly different from those of the bias factor, here $\\bar{z}=0.30,0.38,0.43$, again with relative widths ($1\\sigma$) of about $40\\%$. Thus, the correlation parameters reflect values typical for a slightly different, more recent cosmological time. This mismatch arises if the peak redshift of the lensing efficency, $\\overline{W}(w)$, is displaced with respect to the peak redshift of the foreground sample, as can be seen by Eq.  \\Ref{h3}. An alignment could be achieved by choosing an appropiate background sample for every foreground sample which was not possible in our case, because we did not allow background galaxies fainter than $24\\,\\rm mag$.  Going back to the observed scale-dependence of the bias factor, galaxies become anti-biased on intermediate scales; they are less strongly clustered than the matter. In our data, the minimum value of the bias factor is determined to be \\mbox{$b_{\\rm min}\\sim0.76$}. This kind of scale-dependence has also been detected by Pen et al. (2003) (VIRMOS-DESCART survey) and Hoekstra et al.  (2002) (VIRMOS-DESCART and RCS) which both rely on weak gravitational lensing to probe galaxy bias. While Pen et al.  use $I$-band luminosities to select galaxies, which results in a larger value for the minimum bias factor but at a similar scale of about $R\\approx3\\,h^{-1}\\rm Mpc$ (\\mbox{$k=2\\pi/R\\approx2h{\\rm Mpc}^{-1}$}), the data and sample selection of Hoekstra et al. is relatively similar to our sample \\mbox{FORE-I}; their value of \\mbox{$b_{\\rm min}=0.71_{-0.04}^{+0.06}$} is in agreement ($1\\sigma$) with our measurement, but the quoted scale of $R\\approx1h^{-1}\\rm Mpc$ is different. However, as emphasised before, the position of the minimum is not well defined in our data. Considering the statistical errors one has to admit that the position of the bias minimum is not well determined also in the Pen et al.  analysis (their Fig. 19).  Hence, there is no contradiction between our data and that of the other authors.  An anti-bias on the scales considered here and a characteristic ``dip'' in the functional form of the bias factor is in concordance with recent numerical simulations of dark matter structure formation (Springel et al. 2005; Weinberg et al. 2004; Guzik~\\&~Seljak 2001; Pearce et al.  2001; Yoshikawa et al.  2001; Somerville et al. 2001; Jenkins et al.  1998).  The scale-dependence is due to the fact that the galaxy clustering is a power-law over a wide range of scales, reflected by $\\ave{N^2}$ in Fig.  \\ref{nsqd}, while the dark matter clustering has different shape in CDM simulations and in the observations suggested by, for instance, $\\ave{M_{\\rm ap}^2}$ in Fig. \\ref{nmapmapsqd}.  For the linear correlation parameter, we observe as Hoekstra et al. (2002) and Pen et al. (2003) a high correlation between galaxy and matter distribution. Averaging the measurement of Hoekstra et al. over the range \\mbox{$2^\\prime\\le\\theta_{\\rm ap}\\le19^\\prime$} yields roughly \\mbox{$r\\approx0.8$} which is consistent with our average ($1\\sigma$). \\stressa{Our observed correlations between fluctuations in the galaxy number and mass density appear to be a bit lower, though (Hoekstra, private communication). This could hint to an hitherto undiscovered systematic effect in our data.  However, it should be kept in mind that the statistical errors in $r$ are highly correlated and quite large so that this slightly lower value of $r$ may be just a statistical fluke.}  The clear scale-dependence of the correlation parameter observed by Hoekstra et al. is not visible in our data, because this feature probably gets lost within the statistical uncertainties.  The figures for the correlation parameter -- $r$ is smaller than unity with $68\\%$ confidence -- show that the galaxies are either stochastically or non-linearly biased, or a mixture of both.  To understand what is meant here, imagine that $\\delta_{\\rm g}$ -- the galaxy density contrast -- and $\\delta_{\\rm m}$ -- the dark matter density contrast (both smoothed) -- are quite generally related by \\begin{equation} \\delta_{\\rm g}=f(\\delta_{\\rm m})+\\epsilon(\\delta_{\\rm m})\\;, \\end{equation} where $f$ is some function and $\\epsilon$ a random variable (noise), both satisfying \\begin{equation} \\ave{f(\\delta_{\\rm m})}=\\ave{\\epsilon(\\delta_{\\rm m})}=\\ave{\\epsilon(\\delta_{\\rm m})f(\\delta_{\\rm m})}=0 \\end{equation} owing to the definition of the density contrast and noise ($\\epsilon$ is statistically independent of $f$). In the case of $f$ being linear and $\\epsilon=0$ we have a linear and deterministic relation between $\\delta_{\\rm g}$ and $\\delta_{\\rm m}$, whereas $\\epsilon\\ne 0$ introduces a stochasticity between the density fields. The latter case is called stochastic bias. If $\\delta_{\\rm g}$ and $\\delta_{\\rm m}$ were Gaussian random variables, $f$ would be a linear function. A non-linear function $f$ yields what is commonly called a non-linear bias. The degeneracy between non-linearity and stochasticity arises because a decorrelation -- indicated in the linear stochastic bias scheme by $r\\!\\!<\\!\\!1$ -- can be generated by both a non-linear $f$ and a stochastic component $\\epsilon$: \\begin{equation} r = \\frac{\\ave{\\delta_{\\rm g}\\delta_{\\rm m}}}{\\sqrt{\\ave{\\delta_{\\rm g}^2}\\ave{\\delta_{\\rm m}^2}}}= \\frac{\\ave{\\delta_{\\rm m}f(\\delta_{\\rm m})}}{\\sqrt{\\ave{\\delta_{\\rm m}^2}(\\ave{[f(\\delta_{\\rm m})]^2}+\\ave{[\\epsilon(\\delta_{\\rm m})]^2})}}\\;. \\end{equation}  Discriminating between these two cases requires the additional measurement of the non-linear stochastic bias parameter (Yoshikawa et al. 2001; Dekel~\\&~Lahav 1999) or equivalent quantities. So far, these parameters have only been measured for the relative bias between populations of galaxies (Wild et al. 2005). Therefore, it is unknown from observations whether this decorrelation between galaxies and dark matter is mainly due to a random scatter or a non-linear relation. To resolve this, using approaches similar to ours where statistical moments of the joint PDF of matter and galaxies are measured, one needs to invoke higher-order statistics. As the currently ongoing research is working on the three-point statistics of the aperture mass and the aperture number count (e.g.  Schneider~\\&~Watts 2005; Schneider, Kilbinger~\\&~Lombardi 2005; Jarvis et al.  2004; Schneider~\\&~Lombardi 2003) we can expect to be capable of such a task quite soon.  Within the uncertainties of our measurement we do not see a difference in the biasing parameters between the three foreground bins.  As the three different foreground bins represent different median redshifts of the galaxies, we conclude that on the scales considered the redshift dependence of the (averaged) linear bias for \\mbox{$0.3\\lesssim z\\lesssim0.7$} has to be smaller than about $\\Delta b\\lesssim0.2$ and $\\Delta r\\lesssim 0.4$ ($1\\sigma$) as crudely estimated from the $1\\sigma$-errors of the average bias and correlation parameter; any larger bias evolution should have been detectable despite the relatively large error bars.  These figures are no serious constraints to cosmological models for the bias evolution because all different numerical and analytic models predict evolution rates well below these limits in the redshift range covered here (cf. Magliocchetti et al. 2000).  In a recent paper, Marinoni et al. (2005) measured the non-linear biasing function (Dekel~\\&~Lahav 1999) in the VIMOS-VLT Deep Survey between \\mbox{$0.7\\lesssim z\\lesssim1.5$} and found that the bias evolution is marginal below $z\\sim0.8$ and becomes more pronounced beyond that redshift.  Empirically, they found the redshift dependence of the bias factor on a scale of $8\\,h^{-1}\\rm Mpc$ being described by \\mbox{$b(z)=1+(0.03\\pm0.01)(1+z)^{3.3\\pm0.6}$} which means a change of $b$ of roughly $5\\%$ between $0.3\\le z\\le 0.7$. This figure is in qualitative agreement with our observation.  The bias parameters, no matter whether linear stochastic or non-linear stochastic bias, are just conveniently defined quantities for a comparison of random fields. They bear no obvious relation to the physics of galaxies. In the end, these measurements will need to be interpreted in terms of physical quantities like the halo occupation distribution (Berlind, Weinberg et al.  2003; Berlind~\\&~Weinberg 2002; Peacock~\\&~Smith 2000) in order to learn more about the evolution and formation of galaxies in the environment of their parent dark matter haloes."
    },
    "0606/astro-ph0606308_arXiv.txt": {
        "abstract": "We discuss  three topics related to the neutron star (NS) mass spectrum. At first we discuss the possibility to form low-mass ($ M \\stackrel{<}{\\sim}   1\\,M_{\\odot}$) and suggest this is possible only due to fragmentation of rapidly rotating proto-NSs. Such low-mass NSs should have very high spatial velocities which could allow identification. % A critical assessment of this scenario is given. Secondly, we discuss mass growth due to accretion for NSs in close binary systems. With the help of numerical population synthesis calculations we derive the mass spectrum of massive ($M > 1.8\\, M_{\\odot}$) NSs. Finally, we discuss the role of the mass spectrum in population studies of young cooling NSs. We formulate a kind of {\\it mass constraint} which can be helpful, in our opinion, in discussing different competive models of the thermal evolution of NSs. ",
        "introduction": "\\label{intro}  Mass is one of the key parameters for neutron star (NS) physics and astrophysics. It can be measured with high precision in binary radio pulsar systems. Until very recently all estimates were obtained in the very narrow region of 1.35-1.45~$M_{\\odot}$ (\\citealt{tc1999}). These values lie very close to the Chandrasekhar limit for white dwarfs. Thus, $M=1.4\\, M_{\\odot}$ was considered to be the standard value of the NS mass. Recently, the range widened towards lower masses after the discovery of the double pulsar J0737-3039 (\\citealt{b2003}). One of the NSs in this system has $M=1.25\\, M_{\\odot}$ (\\citealt{l2004}). Several other examples are known (see Sec.~4).  The mass range is extended also towards higher masses, although these results are less certain. There is  one NS in a binary radio pulsar system with  mass significantly higher than the canonical value $1.4\\, M_{\\odot}$. It is the pulsar J0751+1807  with the mass $2.1^{+0.4}_{-0.5}$ (95\\% confidence level) (\\citealt{ns2004}). A recent reanalysis of RXTE data suggests that in low-mass X-ray binaries there are several candidates for neutron stars with masses between 1.9  and 2.1 $M_\\odot$ (\\citealt{bom2005}) as, e.g., 4U 1636-536. The small number of massive radio pulsars, however, can be a result of selection effects.  Cooling curves of NSs are strongly dependent on the star mass, the mass spectrum is one of the most important ingredients of the population synthesis of these sources and, unfortunately, one of the less known. Still, theoretical considerations \\citep{whw2002} can shed light on the general properties of the mass spectrum of NSs. This properties have to be taken into account when confronting theoretical models of the thermal evolution of NSs with observational data.  In this paper we discuss three issues: the formation of low-mass NSs, the mass growth of NSs in binaries, and the role of NS masses in the population synthesis of young cooling objects. ",
        "conclusions": ""
    },
    "0606/astro-ph0606552_arXiv.txt": {
        "abstract": "Details of how the primordial plasma recombined and how the universe later reionized are currently somewhat uncertain. This uncertainty can restrict the accuracy of cosmological parameter measurements from the Cosmic Microwave Background (CMB). More positively, future CMB data can be used to constrain the ionization history using observations. We first discuss how current uncertainties in the recombination history impact parameter constraints, and show how suitable parameterizations can be used to obtain unbiased parameter estimates from future data. Some parameters can be constrained robustly, however there is clear motivation to model recombination more accurately with quantified errors.  We then discuss constraints on the ionization fraction binned in redshift during reionization. Perfect CMB polarization data could in principle distinguish different histories that have the same optical depth. We discuss how well the \\Planck\\ satellite may be able to constrain the ionization history, and show the currently very weak constraints from WMAP three-year data. ",
        "introduction": "We are entering the era of precision cosmology, with future high precision cosmic microwave background (CMB) data in the offing. If the physics governing the evolution of the photon distribution can be modelled reliably, this data offers an almost unique opportunity to measure a large number of cosmological parameters accurately and distinguish models of the early universe. It is widely recognized that small second (and higher) order effects can effect the CMB power spectra at the several-percent level: these include the kinetic and thermal Sunyaev-Zel'dovich (SZ) effects~\\citep{Hu:2001bc} and CMB lensing~\\citep{Lewis:2006fu}. As well as the cosmological parameters of most interest, there are however also other uncertain parameters governing the background evolution of the universe that can be of comparable importance.  Here we focus on the ionization history, parameterized as the ionization fraction as a function of redshift. The broad details of recombination are well understood~\\citep{Peebles:68,Peebles:93,Hu:1995fq}. Direct recombination to the ground state is ineffective as it releases a high energy photon that immediately ionizes another atom: it can only be important if the photon is sufficiently cosmologically redshifted before encountering another atom. The dominant mechanism is in fact capture to an excited state followed by a two-photon transition to the ground state. However there are many excited states in both hydrogen and helium, giving many possible recombination channels. Furthermore level populations may be out of equilibrium, requiring a full multi-level atom evolution to calculate the recombination rate accurately~\\citep{Seager:1999km}. This can only be done to the extent that the different transition rates are known (or can be calculated), in particular excited state two-photon transition rates are not known very well and are potentially important~\\citep{Dubrovich:2005fc}. The two-photon rates can also differ significantly from their empty-space value due to induced decay by CMB photons~\\citep{Chluba:2005uz}.  Most current CMB power spectrum calculations use the effective model used in the code \\RECFAST~\\citep{Seager:1999bc}. However the ionization fraction from this code differs by several percent from a more recent calculation in \\citet{Dubrovich:2005fc} that includes additional transitions. Level splitting and other neglected effects can also have percent-level effects on the CMB power spectra~\\citep{Chluba:2005uz,Leung:2003je,Rubino-Martin:2006ug}. Furthermore there is no full calculation with any quantification of the error. When high resolution CMB data are available we will either need to be able to calculate the recombination history sufficiently accurately that errors can be neglected, or we will need to include uncertainties in any analysis. Not surprisingly assuming an incorrect history can give biased constraints on the cosmological parameters inferred from CMB data~\\citep{Hu:1995fq}. Our aim in this paper is to see whether a crude parameterization of uncertainties in recombination can be used with future data to give reliable parameter constraints when considerable uncertainties remain in the details of recombination. The detailed task of improving recombination models and parameterizing residual uncertainties is an important challenge for the future.  Once recombination has proceeded to reduce the residual ionized density to a low level the details are no longer important for the CMB: the electron density is so low that few CMB photons are scattered. However at some point first collapsed objects will form, and at some later time high-energy photons emitted by (for example) quasars or stars of various populations can cause the universe to reionize \\citep{Loeb:2000fc,Gnedin:2006uz}. Current observations of quasar absorption spectra only put a lower bound on reionization redshift, giving evidence for the first neutral hydrogen on our light cone at $z\\approx 6$ \\citep{Becker:2001ee,Fan:2006dp}. Exactly how the ionization fraction evolved between the low level remaining after recombination and $z\\approx 6$ is unknown in any detail.  The main CMB constraint on reionization comes from the large scale polarization signal generated by scattering of the CMB quadrupole during reionization. This gives a characteristic bump in the large scale polarization power spectra on scales larger than the horizon size at reionization, and the exact details of the shape of the bump can in principle be used to constrain the ionization history \\citep{Kaplinghat:2002vt,Hu:2003gh}. Scales smaller than the horizon size are uniformly damped, leading to a suppression of the acoustic peaks of $e^{-2\\tau}$ where $\\tau$ is the optical depth to reionization; the small scales cannot therefore be used to constrain the details of the history, only the total optical depth. Beyond linear theory there are also tiny secondary anisotropies on  small scales that we do not discuss further here \\citep{Weller:1999ch,Hu:1999vq}.  The current constraint on the optical depth from the WMAP three-year polarization observations is $\\tau = 0.09\\pm 0.03$~\\citep{Page:2006hz,Spergel:2006hy}, significantly smaller than $\\tau\\sim 0.17$ favoured by the one-year data (which had unsubtracted foreground contamination). The new value may be more consistent with simple reionization scenarios, though the lower fluctuation amplitude measured by the three-year data makes this interpretation unclear; for discussion see \\citet{Alvarez:2006ma,Popa:2006rh}.  Since there is currently no convincing model of the reionization history we attempt to constrain it in a relatively model independent way. \\citet{Hu:2003gh} suggest a binned fit and performed a principal component analysis with Fisher matrix techniques. They showed that the first two to three eigenmodes should be fairly well constrained with future observations, but other details are effectively unconstrained by the CMB alone. In the second part of this paper we investigate  how the reionization history binned in redshift can be constrained with current, perfect, and simulated \\Planck~\\citep{unknown:2006uk} data. The general constraints we obtain are rather weak, however distinct models of reionization can be distinguished. Unless there is a convincing physical model for how reionization proceeds, flexible modelling of the reionization history is required in order not to obtain biased constraints on the optical depth (and hence the amplitude of primordial fluctuations and $\\sigma_8$ inferred from the CMB).  Since details of reionization are only important on large scales, and details of recombination are only important on small scales, the two effects are virtually independent. We start by describing our parameter estimation and forecasting methodology, then move on to consider recombination and reionization separately. ",
        "conclusions": "Modelling the ionization history is crucial to interpret high precision observations of the CMB correctly. Uncertainties in the details of recombination affect the small scale anisotropies, and an incorrect model can lead to strongly biased parameter constraints. We have made a first step towards modelling recombination uncertainties using an ad hoc parameterization governing the recombination rate at four different ionization fractions. This was good enough to recover parameter estimates that are consistent, but at the cost of larger error bars. There is clear motivation for future work to pin down the recombination history in more detail theoretically, with quantification of errors on any poorly measured important parameters. In principle future work could do a full multi-atom calculation for each cosmological model, including free parameters with error bars for all rates that are not known very accurately. In practice it is likely to be possible to devise an accurate small set of effective equations that can be evolved quickly, along the lines of the current \\RECFAST\\ code. Using a free parameterization has the advantage of being able to look for any surprises, for example energy injection from annihilating particles, however if the standard model predictions are not well nailed down it may be hard to distinguish subtle unexpected effects from uncertainties in the expected model.  The reionization impacts the large scale CMB polarization anisotropies. Our conclusions effectively agree with previous work, though our direct binning approach is somewhat different. Ideal data can do quite well at distinguishing distinct models, and even \\Planck\\ may be able to separate the most clear-cut cases.  We concentrate on Planck, also there are ground based probes which could provide useful information about reionization. In particular the Background Imaging of Cosmic Extragalactic Polarization (BICEP) has in principle the ability to measure features in the first trough in the E-mode anisotropies, which at least to some degree could constrain the details of the reionization history \\citep{Keating:03}. However the error bars remain quite large even with ideal data, and the CMB is ultimately not going to be the best way to learn about the reionization history observationally: 21cm emission is likely to be a much better tracer. Allowing for different models in a CMB analysis is however crucial to obtain unbiased constraints on the optical depth (and hence the underlying perturbation variance) if the optical depth is quite large. Of course if the optical depth is observed to be low, so we know reionization must all have happened around $z\\sim 6$, the modelling for the history becomes less important as there is then little room more complicated reionization scenarios. Alternatively a consensus may arise that the reionization history is expected to be monotonic (e.g. see \\citet{Furlanetto:2004nt}), in which case modelling is less important and significantly simpler."
    },
    "0606/astro-ph0606078_arXiv.txt": {
        "abstract": "Cosmological consequences of a string-motivated dark energy scenario featuring a scalar field coupled to the Gauss-Bonnet invariant are investigated. We study the evolution of the universe in such a model, identifying its key properties. The evolution of the homogeneous background and cosmological perturbations, both at large and small scales, are calculated. The impact of the coupling on galaxy distributions and the cosmic microwave background is examined. We find the coupling provides a mechanism to viably onset the late acceleration, to alleviate the coincidence problem, and furthermore to effectively cross the phantom divide at the present while avoiding a Big Rip in the future. We show the model could explain the present cosmological observations, and discuss how various astrophysical and cosmological data, from the Solar system, supernovae Ia, cosmic microwave background radiation and large scale structure constrain it. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606287_arXiv.txt": {
        "abstract": "We consider the viability of dark energy (DE) models in the framework of the scalar-tensor theory of gravity, including the possibility to have a phantom DE at small redshifts $z$ as admitted by supernova luminosity-distance data. For small $z$, the generic solution for these models is constructed in the form of a power series in $z$ without any approximation. Necessary constraints for DE to be phantom today and to cross the phantom divide line $p=-\\rho$ at small $z$ are presented. Considering the Solar System constraints, we find for the post-Newtonian parameters that $\\gamma_{PN}<1$ and $\\gamma_{PN,0}\\approx 1$ for the model to be viable, and $\\beta_{PN,0}>1$ (but very close to $1$) if the model has a significantly phantom DE today. However, prospects to establish the phantom behaviour of DE are much better with cosmological data than with Solar System experiments. Earlier obtained results for a $\\Lambda$-dominated universe with the vanishing scalar field potential are extended to a more general DE equation of state confirming that the cosmological evolution of these models rule them out. Models of currently fantom DE which are viable for small $z$ can be easily constructed with a constant potential; however, they generically become singular at some higher $z$. With a growing potential, viable models exist up to an arbitrary high redshift. ",
        "introduction": "% A major turning point in cosmology has been reached with the  {\\it observational} discovery that our Universe is accelerating now (and has been accelerating for several billion years in the past) \\cite{R98,P99}. If interpreted in terms of the Einstein equations for the evolution of a Friedmann-Lema\\^itre-Robertson-Walker (FLRW) cosmological models with the (practically) zero spatial curvature (the latter follows from other arguments), this means that approximately two thirds of the total present energy density of matter in our Universe is due to some gravitationally unclustered component called Dark Energy (DE).  Observations of high redshift supernova, fluctuations of the cosmic microwave background (CMB) temperature and other effects tell us that the effective energy density $\\rho_{DE}$ of DE is very close to minus its effective pressure $p_{DE}$ (see Eqs. (\\ref{E1a},\\ref{E2a}) below for the exact, though conventional, definition of these quantities) and that they are both very weakly changing (if at all) with time and with the expansion of the Universe. The physical nature of DE is unknown at present, and three main logical possibilities exist (see \\cite{SS00,PR03,Pad03,Sah05,CST06} for reviews).\\\\ 1. DE is a cosmological constant, as originally suggested by Einstein, with $\\rho_{DE}= -p_{DE}=\\Lambda/8\\pi G=const$ exactly.\\footnote{$\\hbar=c=1$ is used throughout the paper.}\\\\ 2. Physical DE: DE is the energy density of some new, very weakly interacting physical field (e.g., a quintessence -- a scalar field $\\phi$ with some potential $V(\\phi)$ minimally coupled to gravity). \\\\ 3. Geometric DE: the Einstein general relativity (GR) equations are {\\em not} the correct ones for gravity, but we write them in the Einsteinian form by convention, putting all arising additional terms into the r.h.s. of the equations and calling them the effective energy-momentum tensor of DE.  Of course, this classification is not absolute. In some cases, the difference between physical and geometric DE can, in turn, become conventional. E.g., for a non-minimally coupled scalar field (which constitutes a specific case of scalar-tensor gravity considered in this paper), equations of the model have the same form irrespective of the origin (physical or geometric) of this field. Another way of classifying DE models which is more invariant and, therefore, more important from the point of view of the observational determination of the type of DE is to divide these models into dynamical, if the DE description requires a new field degree of freedom ($\\equiv$ a new particle from the quantum point of view), and non-dynamical in the opposite case. Then all physical DE models and many of the geometric ones belong to the first class, while the cosmological constant itself and some geometric models fall into the second category, i.e., the $F(R)$ model with the Palatini variation of its action where $R$ is the scalar curvature.  At present, all existing observational data are in agreement with the simplest, first possibility (inside $\\sim 2\\sigma$ error bars in the worst case). The case of a cosmological constant is internally self-consistent and non-contradictive. The extreme smallness of the cosmological constant expressed in the either Planck, or even atomic units means only that its origin is {\\em not} related to strong, electromagnetic and weak interactions (in particular, to the problem of the energy density of their vacuum fluctuations) at all. However, in this case we remain with one dimensionless constant only and can not say anything more (at least, at present). That is why the two other possibilities admitting a (slightly) variable dark energy have been also actively studied and compared with observational data recently. Moreover, properties of the present DE are remarkably {\\em qualitatively} similar to those of an \"early DE\" that supported an inflationary stage in the early Universe. But in the latter case, we are sure that this early DE was unstable. So, it is natural to conjecture by analogy that the present DE is not stable, too.  On purely phenomenological grounds, one can consider DE models with a constant equation of state parameter $w_{DE}\\equiv p_{DE}/ \\rho_{DE}$ different from $-1$. However, the latest observational data have already severely restricted this simplest alternative possibility to $|1+ w_{DE}|\\stackrel {<}{\\sim} 0.1$ ($1\\sigma$ error bars) \\cite{Sel05,As06,SPD06}. Therefore, a viable alternative to the cosmological constant has to be looked for among more complicated models with $w_{DE}\\not= const$. Such models include quintessence ones with different potentials $V(\\phi)$ (see \\cite{SS00,PR03,Pad03,Sah05} for numerous references), models with several minimally coupled scalar fields (see, e.g., \\cite{BP04}), those with direct non-gravitational coupling between DE and dark matter (DM) (\\cite{A00} and following papers), unified models of DE and DM (the Chaplygin gas model \\cite{Kam01} and others), etc. Then, however, it becomes very important to investigate if there exist models of DE for which a variable $w_{DE}$ may cross the \"phantom divide\" line $w_{DE}=-1$ (we call DE \"fantom\" if it has $w_{DE}(z) <-1$ for a given $z$ and \"normal\" in the opposite case). Note that the weak energy condition (WEC) is violated for phantom DE. This is not possible at all for quintessence models with the standard kinetic term and hardly possible, i.e., it requires non-generic initial conditions, for scalar field models having a non-standard kinetic term (\"$k$-essence\") \\cite{Vik05} (see also \\cite{ACK05}).  Indeed, analysis of the recent SNe data with redshifts up to $z=1.7$ (the \"Gold dataset\" \\cite{Rie04}) using fits containing at least 2 free parameters, e.g., the linear fit (\\ref{wCPL}) for $w_{DE}$ in terms of the scale factor $a$ \\cite{CP01,L03} or the quadratic polynomial fit (\\ref{wp}) for $\\rho_{DE}$ in terms of the redshift $z$ \\cite{SSSA03}, or with model independent methods, results in best fits to these data having a variable $w_{DE}$ that steadily increases for redshifts $0<z<1$ and crosses the \"phantom divide\" somewhere between $0$ and $0.5$ \\cite{ASS04,HC05,Gong05, DR04,HM04,RAW05,UIS05,LNP05,ER05,XZF06} (see, however, \\cite{WT, Sel05,JBP06,DKC06} for a more conservative view -- in the sense of returning back to a cosmological constant). This statement does not mean that an exact cosmological constant is excluded -- it is still inside $1$ or $2\\sigma$ error bars, as was stated above. Moreover, one should be cautious with this result: it may be a consequence of trying to obtain a too fine-grained graph of $w_{DE}(z)$ from the luminosity distance $d_L(z)$ as has been already emphasized in \\cite{MBMS02}. In this respect, results for $w_{DE}(z)$ averaged over a range of redshifts $>0.3$ (dubbed the \"$w$-probe\" in \\cite{SASS05}) are more statistically reliable, e.g., the result that $\\bar w_{DE}\\approx -1.05_{-0.07}^{+0.09}$ obtained for the same Gold SNe dataset in \\cite{ASS04} for $\\Omega_{m,0}=0.3$ and the averaging redshift interval $(0,0.414)$. However, a remarkable possibility of crossing the phantom divide at recent redshifts, when DE has become the dominant component of the Universe by its effective energy density, still remains viable. The latest Supernovae data with $z\\sim 0.5$ \\cite{CSF05} also admit $w_{DE}(z)<-1$ for $z<0.5$, but not for larger redshifts, as a possible interpretation. Finally, the recent data on acoustic baryon oscillations in the present matter power spectrum (the Sakharov oscillations) \\cite{Eis05,Col05}, while strongly restricting one direction in the plane of parameters for the two-parametric fits (\\ref{wCPL},\\ref{wp}), leave the orthogonal direction practically unconstrained, therefore, permitting the recent phantom divide crossing (see Sec.~6).  So, if this striking behaviour of DE will be confirmed by future data, what is the best way to describe it? One possibility, a {\\em ghost} phantom DE, i.e. a scalar field with the negative sign of its kinetic term, was first proposed in \\cite{Cald02} and triggered a very large wave of publications (see, e.g., the recent papers \\cite{NOT05,STTT05} for a list of references on this topic). However, it has been long known that theories of this type are plagued by quantum instabilities, the most dangerous of those being the process of creation of two particle+antiparticle pairs: one of the ghost field and another of any usual (non-ghost) field (see \\cite{CJM04} for a recent investigation). Moreover, a ghost model of DE is unsatisfactory even at the classical level: it does not explain the observed large-scale isotropy and homogeneity of the Universe! Just the opposite, e.g., for the given Hubble constant $H_0$ averaged over angular directions in the sky, we would expect a universe to be very strongly anisotropic with the anisotropy energy density (i.e., the positive energy density of long-wavelength gravitational waves) being compensated by a large negative energy density of the ghost DE, or of a ghost component of DE in more complicated multi-component models of this type (e.g., in the two-field realization of the ``quintom'' model introduced in \\cite{FWZ05}). It is just a classical analogue of the quantum instability mentioned above, with the \"usual\" field being the gravitational one and with the dynamical quantum instability transformed into the problem of classical initial conditions. For this reason, we are sceptical regarding this approach as a whole.  Fortunately, there is no necessity to introduce a ghost field to explain possible phantom behaviour of DE including its phantom divide crossing. As was first emphasized in \\cite{BEPS00}, scalar-tensor theories of gravity allow for this phenomenon. Scalar-tensor models of DE belong to the third class (geometric DE) according to the classification given above. This class of models is very rich and interesting. It contains the Einstein gravity plus a non-minimally coupled scalar field, as well as the higher-derivative $f(R)$ gravity, where $R$ is the scalar curvature, as particular cases, besides allowing for DE phantom behaviour and transition to a normal DE. In this paper, we will concentrate on scalar-tensor DE models with the latter properties because of tentative observational evidence described above.\\footnote{The third possibility to get a phantom DE including the phantom divide crossing which is based on braneworld models \\cite{SSh03} (see also the recent paper \\cite{CLMQ06}) will not be discussed here.} Note that a possible phantom behaviour of DE in scalar-tensor gravity has a conventional character. The reason is that the effective gravitational constants $G_N$ and $G_{\\rm eff}$ (see Sec.~2 below) are generically time-dependent while the definitions (\\ref{E1a},\\ref{E2a}) of the DE effective energy density and pressure assume some constant $G$ when writing the left-hand side of equations and splitting their right-hand side into energy-momentum tensors of non-relativistic matter (mainly non-baryonic cold dark matter) and dark energy. As a result, a part of the Einstein tensor $G_{\\mu\\nu}$ and the energy-momentum tensor of dark matter multiplied by a change in $G_N$ is conventionally attributed to DE. In other words, phantom behaviour of scalar-tensor DE is always 'curvature-induced', in contrast to the ghost DE models or other models like those considered in \\cite{AKV05,Rub06} where it may occur in the flat space-time already.  The possibility to get both a phantom DE and the phantom divide crossing in scalar-tensor gravity is related to the fact that this theory has two arbitrary and independent functions $F(\\Phi)$ and $U(\\Phi)$ (see the Lagrangian (\\ref{L}) below). Throughout this work we assume spatial flatness though this prior while well motivated theoretically can be challenged \\cite{PR05}. As has been shown in \\cite{BEPS00}, two different types of observations, e.g., determination of the luminosity distance and the inhomogeneity growth factor in the non-relativistic component as functions of redshift, are necessary and sufficient for the total reconstruction of the microscopic Lagrangian (\\ref{L}) of scalar-tensor gravity. However, as shown in \\cite{EP01}, a partial reconstruction using the luminosity distance data only could yield interesting information, too. In that case some assumption about one of the functions $F(\\Phi)$ and $U(\\Phi)$ has to be made, so that only one unknown function in the microscopic Lagrangian (\\ref{L}) remains to be found. Clearly, many different partial reconstruction strategies are possible and we explore some of them here with the aim to investigate which DE models in scalar-tensor gravity are viable.  It has been found in \\cite{EP01} that $\\Lambda$-dominated universes with a vanishing potential $U$ are ruled out as they lead to singular universes already at very low redshifts $z\\sim 0.7$ (see also the recent paper \\cite{Per05}). Of course, completely regular but non-accelerating solutions do exist in this case. Among models with a non-zero potential $U$, one concrete example of a model with the phantom divide crossing was constructed in \\cite{Per05} where the functions $F(\\Phi)$ and $U(\\Phi)$ were given in the parametric form as functions of $z$ up to $z\\approx 3$. More examples for a non-minimally coupled scalar field ($F(\\Phi)=F_0-\\xi\\Phi^2$) with a non-zero potential $U$ were investigated in \\cite{LS05}, also for $z<2$. In the present paper, we make a next step and extend these results in two directions: first, by constructing a generic solution for scalar-tensor DE models for $z<1$ in the form of a power series in $z$; second, by investigating and numerically integrating some of these solutions up to large $z\\gg 1$. The latter task appears to be necessary since, though DE is subdominant for $z\\gg 1$ as compared to non-relativistic non-baryonic dark matter and baryons, its model itself may become intrinsically contradictory for large $z$, namely, $F$ or $\\dot\\Phi^2$ may become negative for an unfortunate choice of $U(\\Phi)$. It is also crucial to check that any DE model has the correct power-law behaviour for large $z$ \\cite{APT06}.  Finally, for the scalar-tensor theory of gravity, it is well known that the present value of ${dF\\over d\\Phi}(z=0)$ is severely restricted from Solar System tests of post-Newtonian gravity (i.e. by the measured values of the post-Newtonian parameters). For this reason, there have been stated that the present phantom behaviour in scalar-tensor models of DE requires large amount of fine tuning and, thus, is unnatural \\cite{T02,SFT05}. Therefore, it is important to investigate this problem in more detail to quantify what amount of fine tuning (and of what kind) is necessary for a significantly phantom behaviour of DE, and to determine the relation between this behaviour and results of Solar System tests of gravity.  The paper is organized as follows. In Sec.~2 we define all quantities related to our scalar-tensor DE model and present the background evolution equations.  In Sec.~3 we derive the general integral solution for the quantity $F(z)$. In Sec. 4 we consider solutions in which DE scales as some power of the FLRW scale factor $a(t)$ and show their existence for DE of the phantom type (the latter requires a non-zero $U(\\Phi)$).  In Sec.~5 we consider the small $z$ behaviour of our model and find conditions for the violation of the WEC by DE today and for the phantom boundary crossing at small $z$. In Sec.~6 the general reconstruction of a microscopic model is considered and the observational constraint from the acoustic oscillations in the matter power spectrum is discussed. In Sec.~7 we consider the reconstruction for a constant potential more specifically and show that the model becomes singular at some redshift that cannot be arbitrarily high once we have chosen a specific equation of state $w_{DE}$. In Sec.~8 non-constant potentials are considered and a model which is asymptotically stable for large $z$ is presented. Sec.~9 contains conclusions and discussion. ",
        "conclusions": "In this work we have considered the viability of scalar-tensor models of Dark Energy. We have used different types of observations: Solar System constraints which constrain the model {\\it today}, and other data like the supernova data which constrain its cosmological evolution, in particular the time evolution of the DE equation of state parameter $w_{DE}$. We were interested specifically in models which violate the weak energy condition on small redshifts $z\\lesssim 0.3$ and in any case today. We have found the formal general integral solution for $F(z)$ when we reconstruct it for given $H(z)$, which can be obtained from the $d_L(z)$ data, and for given $U(z)$. This general solution allows immediately for an integral representation of scaling solutions. We have constructed scaling solutions and shown that they exist in models with $F=\\alpha \\Phi^2$ with $\\alpha=$constant, in these models the Brans-Dicke parameter $\\omega_{BD}$ is constant. Only for nonzero potentials can these models have scaling solutions with constant $w_{DE}<-1$. However, it is shown that for these models $|w_{DE}+1|$ is very small with $|1 + w_{DE}|\\sim \\omega_{BD}^{-1}$. We have further performed systematically the small $z$ expansion of the theory and used it to extract various observational constraints. We recover that a large positive $\\omega_{BD,0}$ requires $|F_1|\\ll 1$ where $F_1$ is the first derivative today of $\\frac{F}{F_0}$. We find that a significantly phantom DE today ($w_0<-1.01$) implies that $F_2$, the second derivative today of $\\frac{F}{F_0}$, must be negative and not small, thus, significantly larger than $F_1$ by modulus. However, while necessary this is not a sufficient condition for phantom DE today, for example for vanishing potential $F_2<0$ whenever $w_{DE,0}<1$.  The condition $|F_2|\\sim 1$ while the Solar System data require $|F_1|<10^{-2}$ (i.e. anomalously small) is the only 'fine tuning' required to get a significantly phantom DE at the present time in scalar-tensor gravity. Note that, since the derivatives $F_i(z)$ are not parameters of the effective microscopic Lagrangian (1) but depend on initial conditions in the early Universe too, it could be even better to call this a \"cosmic coincidence\". Our point of view is that the condition $|F_2|\\sim 1$ cannot be excluded by pure thought, only observations will possibly do it: in the absence of this condition when all $F_i$ are of the same order as $F_1$, the general prediction is that the amount of possible phantomness in scalar-tensor DE models is very small, less than $1\\%$.  As for the solar system constraints, we have shown that the Post-Newtonian parameter $\\gamma_{PN,0}$ must satisfy $\\gamma_{PN,0}-1\\simeq - \\omega_{BD,0}^{-1}< 0$, this is a general requirement for the viability of our model. On the other hand, a significantly phantom DE today implies $\\beta_{PN,0} > 1$. Combining those results, we find that the negative quantity $\\frac{\\gamma_{PN,0}-1}{\\beta_{PN,0}-1}$ does not depend on $F_1$ and can be expressed in function of $F_2$ and $1+w_{DE,0}$. This would enable us to find $F_2$ from the Solar System constraints provided we know $1+w_{DE,0}$ from other cosmological data. Still for phantom DE today, we find that ${\\dot G}_{{\\rm eff},0}$ has the same sign as $F_1$. So, a measurement of the sign of ${\\dot G}_{{\\rm eff},0}$ would give us the sign of $F_1$ if we know we have $1+w_{DE,0}<0$. On the other hand, a measurement of both quantities with opposite signs would rule out phantom DE today though measuring the sign of $F_1$ can be hard to determine observationally. However, due to the already confirmed smallness of $F_1$, connection between cosmological and Solar System tests of dark energy is rather one-way since it requires much greater accuracy from the latter ones for the determination whether DE is phantom at present or not. While a positive detection of $\\beta_{PN,0}>1$ is a strong argument for the phantom DE at present, a negative result (no measurable deviation of $\\beta_{PN,0}$ from unity) tells us nothing regarding DE properties. Also, the Solar System tests are clearly unable to provide $w_1$ in any reasonable future, so no information about the possibility of the phantom divide crossing may be expected from them.  We have also considered numerically the reconstruction of $F$ for various $w_{DE}$, including constant as well as varying equations of state of the phantom type. Generalizing results obtained in \\cite{EP01} for a pure cosmological constant, we find that models with a vanishing potential (see Figure 3) are ruled out and lead to a singular behaviour for $z<0.66$. While it is clearly possible to have phantom DE today with $U=0$ without conflicting with the data, the cosmological evolution of these models rules them out. For models with constant nonvanishing potentials, which can be considered as a good approximation on small redshifts for more general models with varying potentials, we find that it is easy to have models in agreement with observations on small redshifts $z\\lesssim 2$. However, it is interesting that these models generically have a maximal redshift where they become singular. So, to construct a scalar-tensor DE model having a sufficiently long matter-dominated stage, a non-constant potential $U(\\Phi)$ is required, and we have presented an example of such a model where DE tracks matter at large redshifts.  Therefore, the final conclusion is that the generic scalar-tensor gravity (\\ref{L}) with the two functions $F(\\Phi)$ and $U(\\Phi)$, derived from some underlying theory (e.g., from brane cosmology) or taken from observational data, has enough power to provide internally consistent cosmological models with a temporarily phantom DE (at present, in particular) and with a regular phantom divide crossing in the course of the evolution of the Universe."
    },
    "0606/astro-ph0606293_arXiv.txt": {
        "abstract": "A considerable number of astrometric binaries whose positions on the sky do not obey the standard model of mean position, parallax and linear proper motion, were observed by the Hipparcos satellite. Some of them remain non-discovered, and their observational data have not been properly processed with the more adequate astrometric model that includes nonlinear orbital motion. We develop an automated algorithm based on ``genetic optimization'', to solve the orbital fitting problem in the most difficult setup, when no prior information about the orbital elements is available (from, e.g., spectroscopic data or radial velocity monitoring). We also offer a technique to accurately compute the probability that an orbital fit is bogus, that is, that an orbital solution is obtained for a single star, and to estimate the probability distributions for the fitting orbital parameters. We test this method on Hipparcos stars with known orbital solutions in the catalog, and further apply it to 1561  stars with stochastic solutions, which may be unresolved binaries.  At a confidence level of 99\\%, orbital fits are obtained for 65 stars, most of which have not been known as binary. It is found that reliable astrometric fits can be obtained even if the period is somewhat longer than the time span of the Hipparcos mission, i.e., if the orbit is not closed. A few of the new probable binaries with A-type primaries with periods 444--2015 d are chemically peculiar stars, including Ap and $\\lambda$ Boo type. The anomalous spectra of these stars are explained as admixture of the light from the unresolved, sufficiently bright and massive companions. We estimate the apparent orbits of four stars which have been identified as members of the $\\approx 300$ Myr-old UMa kinematic group. Another four new nearby binaries may include low-mass M-type or brow dwarf companions. Follow-up spectroscopic observations in conjunction with more accurate inclination estimates will lead to better estimates of the secondary mass. Similar astrometric models and algorithms can be used for binary stars and planet hosts observed by SIM PlanetQuest and Gaia. ",
        "introduction": "One of the most promising and presently available ways to discover invisible companions (including brown dwarfs and giant planets) to nearby Galactic stars is to analyze high-accuracy astrometric positions of the latter for the presence of the reflex Keplerian motion caused by an orbiting companion. In the relatively near future, with the advance of such space-borne astrometric instruments as SIM and GAIA, it will become one the main instruments in the search for habitable smaller planets, too. For the time being, the capabilities of the method are limited by the moderate precision of the available astrometric data. The Hipparcos Intermediate Astrometry Data (HIAD), at a single point precision of roughly 10--15 mas, is just good enough to detect brown dwarfs around nearby solar-type stars, reliably so in conjuction with spectroscopic measurements \\citep{ha,pou,tor}. Perhaps, giant planets of intermediate orbital periods could also be detected around a few nearest stars.  Obtaining a robust orbital fit for an astrometric binary becomes increasingly difficult when the orbital period $P$ exceeds the time span of Hipparcos measurements, which is about 3.2 years. If the period is 6 years, still more than half of the orbit is represented in the data, and astrometric analysis may provide an independent solution. When the period is longer than 9 years, the astrometric data alone can provide only ambiguous results, since the almost linear segment of the orbit is hard to distinguish from the regular proper motion. Generally, the astrometric orbital fit is a non-linear 12-parameter adjustment problem that includes five astrometric parameters and seven orbital parameters (Appendix A). In some cases, e.g., long-period orbits, or highly inclined orbits aligned with the line of sight, the parameters become so entangled in the non-linear condition equations, that the fit becomes ill-conditioned. Therefore, any orbital solution should be supported by analysis of confidence intervals for all fitting parameters. If the astrometric solution is sufficiently well-constrained, approximate standard deviations of fitting parameters can be derived in the near vicinity of the solution point by covariance analysis, a recipe of which is given in Appendix B. Ill-conditioned solutions will have nearly singular normal matrices that can not be inverted, or large variances for some of the parameters. In more complicated cases, some of which we investigate in this paper, the confidence intervals and probability distributions can be calculated by extensive Monte-Carlo simulations. When the optimization problem is almost singular with respect to particular nonlinear parameters (most often the eccentricity), direct mapping of the objective function on a sufficiently fine grid of parameters is carried out.  Another crucial problem that arises in fitting Keplerian motion for stars previously not known as binaries, is estimation of the probability of a null hypothesis, that is, that the star is single, and the detected perturbation in the astrometric residuals is caused by a chance occurrence of random noise in the data, or other effects, unrelated to binarity. We offer a robust and straightforward (if somewhat computationally heavy) method of confidence level estimation of a binary detection.  These techniques are tested on a set of 1561 stars in the Hipparcos catalog with the so-called stochastic solutions. These data often represent gross errors or failures in the data reduction. Due to the type of detector used in the main instrument of Hipparcos (a grid of long vertical slits and a photomultiplier behind it with a non-uniform response across the pointing field of regard), many of these failed solutions originate in the wrong assumption about the target positions, or in the lack of knowledge of the relative position and brightness of visual binary components. Successful attempts have been made to use more accurate information about stochastic stars from ground based and Tycho-2 data sets to reprocess the published Hipparcos transit data, resulting in accurate solutions for a few hundred stars \\citep{fal,fama}. The remaining stars with stochastic solutions are prime suspects for yet unknown binaries. Orbiting pairs with separations smaller than $0.1$ arcsec were not resolved by Hipparcos; on the other hand, a nonlinear apparent motion of the photocenter as small as several milliarcseconds could break the regular 5-parameter solution. Long-period orbits could be detected as additional acceleration of proper motions or as discrepant proper motions from the short-term Hipparcos data and the long-term Tycho-2 definitions \\citep{maka}. The overlap between such accelerating astrometric binaries and the list of stochastic objects is substantial, but presumably a lot of binaries on shorter orbits are left out of this analysis, as they describe a strongly nonlinear track during the 3.2-3.5 time span of Hipparcos observations. ",
        "conclusions": "We present a method for unconstrained optimization of double stars orbital parameters, based on an adaptation of the genetic evolution algorithm optimization and Monte-Carlo simulations for error estimation. This method can become a useful tool for estimating orbital parameters of binary stars and planet hosts for experiments like SIM, which are expected to observe a limited number of pre-selected targets. Experiments like Gaia are expected to observe much higher numbers of objects and requirements for computing time are likely to become more stringent. However, as the total processing time for the computations in this paper was about two weeks of CPU time on a medium range laptop, we expect the entire Gaia astrometric orbital estimation to become feasible on large scale scientific clusters in the near future. This ``embarrassingly parallel'' algorithm is ideally suited for running in a grid environment and it will take advantage of the ever increasing computer performance. We believe that simplicity and ease of implementation of our algorithm outweigh the relatively high CPU time requirements.  The 65 stars from the Hipparcos stochastic solution annex are good candidates for follow-up spectroscopic studies, which are expected to yield more accurate orbital parameters constraining the astrometric fits.  {\\it Facility:} \\facility{HIPPARCOS}   \\appendix"
    },
    "0606/astro-ph0606546_arXiv.txt": {
        "abstract": "Recent spectroscopic observations of planetary nebulae (PNe) in several elliptical galaxies have revealed structural and kinematical properties of the outer stellar halo regions. In order to elucidate the origin of the properties of these planetary nebula systems (PNSs), we consider the merger scenario  in which an elliptical galaxy is formed by merging of  spiral galaxies. Using numerical simulations, we particularly investigate radial profiles of projected PNe number densities, rotational velocities, and velocity dispersions of PNSs extending to the outer halo regions of elliptical galaxies formed from major and unequal-mass merging. We find that the radial profiles of the project number densities can be fitted to the power-law and the mean number density in the outer halos of the ellipticals can be  more than an order of magnitude higher than that of the original spiral's halo. The PNSs are found to show a significant amount of rotation ($V/\\sigma$ $>$ $0.5$) in the outer halo regions ($R$ $>$ $5R_{\\rm e}$) of the ellipticals. Two-dimensional velocity fields of PNSs are derived from the simulations and their dependences on model parameters of galaxy merging are discussed in detail. We compare the simulated kinematics of PNSs with that of the PNS observed in NGC 5128 and thereby discuss advantages and disadvantages of the merger model in explaining the observed kinematics  of the PNS. We also find that the kinematics of PNSs in elliptical galaxies are quite diverse depending on the orbital configurations of galaxy merging, the mass ratio of merger progenitor spirals, and the viewing angle of the galaxies.  This variation translates directly into possible biases by a factor of two in observational mass estimation.  However, the biases in the total mass estimates can be even larger.  The best case systems viewed edge-on can appear to have masses lower than their true mass by a factor of 5, which suggests that current observational studies  on PNe kinematics  of elliptical galaxies can significantly underestimate their real masses. ",
        "introduction": "Structural, kinematical, and chemical properties of stars in the outer regions of galaxies, their halos, provide vital clues to their global formation and evolutionary histories. Although the physical properties of stellar halos have been extensively investigated mostly for spiral (e.g., the Galaxy and the M31) and dwarf galaxies in the Local Group (e.g., Eggen, Lynden-Bell, \\& Sandage 1962; Mould \\& Kristian 1986; Norris 1986; Freeman 1987; Durrell, Harris, \\& Pritchet 1994, 2001; Pritchet \\& van den Bergh 1998; Reitzel, Guhathakurta, \\& Gould 1998; Grillmair et al. 1996; Chiba \\& Yoshii 1998; Davidge 2002), recent observations have started to reveal physical properties of stellar halos in giant elliptical galaxies, all of which are beyond the Local Group (e.g., Soria et al. 1996; Harris, Harris, \\& Poole 1999 hereafter referred to as HHP; Harris \\& Harris 2000, HH00; Harris \\& Harris 2002, HH02; Marleau et al. 2000; Rejkuba et al. 2002; Gregg et al. 2004). These data prompt us to investigate how the properties of stellar halos can be used to understand the formation of giant elliptical galaxies.  Beyond projected radii of  $\\sim$ $2R_{\\rm e}$, it becomes very difficult to study the integrated properties of stars in ellipticals because of the low stellar surface brightness.  Hence, studies of stellar halos tend to rely on resolving individual stars or star clusters. While there are significant observational difficulties in revealing the three dimensional structure and kinematics of halo red giants and AGB stars, it has been possible to use color-magnitude diagrams to derive their metallicity distribution function (hereafter MDF) and constrain the past star formation and chemical evolution histories of their host elliptical galaxies (e.g., HHP; HH00, HH02). Kinematic properties of these stars, however, which can contain vital clues to both structure and formation of ellipticals, are still extremely difficult to obtain. Furthermore, these previous observations have difficulties in deriving {\\it  global} stellar halo  distributions, which are also considered to have valuable information on the past merging histories of galaxies (e.g., fossil tidal streams), mainly because the halo fields investigated cover only a minor portion of the entire halo.  Planetary nebulae (PNe) are a valuable complement to integrated light and red giant studies.  They are powerful tools for studying the dynamical states of elliptical galaxies, because we can more readily identify PNe through their bright [O III] emissions lines and typically measure their radial velocities to an accuracy of $\\sim$ 15 km s$^{-1}$ (e.g., Peng et al. 2004a, PFF04a). Spectroscopic observations of the radial velocities of PNe in several nearby elliptical galaxies therefore have succeeded in deriving global mass distributions and velocity fields of these galaxies (Ciardullo et al. 1993; Hui et al. 1995; Arnaboldi et al. 1998; Mendez et al. 2001; Romanowsky et al. 2003; Napolitano et al. 2004; PFF04a). For example, kinematic information of PNe in the stellar halo of NGC 5128 at radii up to $\\sim$ 80 kpc allows the authors to derive the zero velocity curve (ZVC) in the two-dimensional velocity field of the PNS and thereby to discuss the triaxial shape of the mass distribution of this galaxy (PFF04a). This observed diversity in kinematics of PNSs is providing fresh clues to the origin of early-type galaxies. (Romanowsky et al. 2003; PFF04a).  In spite of this importance of PNe studies in elliptical galaxies, only a few theoretical and numerical attempts have been made to discuss dynamical properties of PNSs in elliptical galaxies. For example, using kinematical data up to $\\sim$ $6R_{\\rm e}$ for an heterogeneous sample of elliptical galaxies, Napolitano et al. (2004) derived the dependences of the radial gradients of the mass-to-light-ratios on the $B-$band magnitude of the galaxies and thereby discussed whether the derived dependences can be consistent with elliptical galaxy formation models based on a $\\Lambda$CDM model. However, PNSs provide a very rich data set from which to study ellipticals.  In particular, we are interested in exploring {\\it the 2D distributions of structural and kinematical properties of PNSs}, which can provide some vital clues both to the triaxial shapes of the global mass distributions of galaxies and to their formation processes. The origin of 2D structural and kinematical properties of PNSs extending to the outer halo regions of elliptical galaxies remains unclear. Furthermore providing theoretical  predictions on dynamical properties of PNSs will help to interpret the properties of PNSs that will be obtained in future systematic observations for elliptical galaxies.  The purpose of this paper is thus to investigate structural and kinematical properties of PNSs of elliptical galaxies based on numerical simulations. In order to elucidate the origin of the observed properties of the PNSs, we adopt the merger scenario (Toomre 1977) in which elliptical galaxies are proposed to be formed by major merging of two spiral galaxies. The present numerical investigation is two-fold as follows. We first describes radial profiles of structural and kinematical properties of PNSs and two-dimensional (2D) velocity fields of PNSs in elliptical galaxies and their dependences on model parameters. We then compare the simulated kinematics of the PNSs  with the corresponding observations for NGC 5128 and thereby try to provide the best model for the PNS in NGC 5128 and discuss advantages and disadvantages of the model in explaining the kinematics of the PNS self-consistently. The essential reason for our choice of the NGC 5128 PNS is that PFF04a have recently investigated radial velocities of the NGC 5128's 780 PNe (among 1141 PNe), which represents the largest kinematical study of an elliptical galaxy to date and thus can be compared with our simulations in the most self-consistent manner. Furthermore, the present study is complimentary to those by  Bekki et al. (2003, BHH03) and Beasley et al. (2003) which numerically and semi-analytically investigated the MDF of the stellar halo of NGC 5128 but did not investigate the dynamical properties.  In the present paper, we mainly investigate structural and kinematic properties of PNSs for a large radial extent in elliptical galaxies (0 $\\le$ $R$ $\\le$ $10R_{\\rm e}$), which include the outer faint stellar halos and the main bodies of the galaxies. Although previous numerical simulations investigated spatial distributions and the line-of-sight velocity distributions of stars in elliptical and S0 galaxies formed from merging  for $R$ $<$  $2.5R_{\\rm e}$ (e.g., Bendo \\& Barnes 2000; Cretten et al. 2001; Naab \\& Trujillo 2005), they did not investigate the dynamical properties for the outer halo regions of elliptical galaxies. Therefore, the present numerical results on the dynamical properties of the outer stellar halos including PNe ($5R_{\\rm e}$  $\\le$ $R$ $\\le$ $10R_{\\rm e}$) may provide new clues to elliptical galaxy formation. Physical properties of these halo stars may well have fossil information on angular momentum redistribution of stars (or conversion from galactic orbital angular momentum into intrinsic angular momentum) which is an  essential physical process of galaxy merging. We therefore can expect that dynamical properties of stellar halos, which can be provided by studies of PNSs, may depend strongly on physics of galaxy merging.  Through the course of this paper, we thus interpret our simulations in the context of the capabilities of present-day PN surveys in order to facilitate the comparison between theory and observations.  The plan of the paper is as follows: In the next section, we describe our numerical models  for the formation of PNSs extending to the outer halo regions in galaxy mergers. In \\S 3, we present the numerical results mainly on the final 2D distributions  of structural and kinematical properties in  merger remnants (i.e., elliptical galaxies) for variously different merger models. In \\S 4, we compare the simulated results of PNSs with observations of the PNS in NGC 5128 (Cen A). In \\S 5, we discuss whether physical properties of stellar halos can give any constraints on galaxy formation, based on the present numerical results. In this section, we also discuss possible different properties of PNSs in galaxies with different Hubble types. We summarize our  conclusions in \\S 6.       \\begin{table*} \\centering \\caption{Model parameters} \\begin{tabular}{ccccccccccl} Model no. & Orbital type & $m_{2}$   & Comments \\\\ M1 & -- & -- & isolated disk  \\\\ M2 & HI & 1.0 & fiducial  model\\\\ M3 & PP & 1.0 &   \\\\ M4 & RR & 1.0 &   \\\\ M5 & LA & 1.0 &   \\\\ M6 & BO & 1.0 &  \\\\ M7 & HI & 0.1 &   \\\\ M8 & HI & 0.3 &   \\\\ M9 & PO & 0.5 & NGC 5128 model \\\\ M10 & -- & 1.0 & Multiple merger  \\\\ \\end{tabular} \\end{table*} ",
        "conclusions": "We have numerically investigated structural and kinematical properties of PNSs in elliptical galaxies formed by galaxy merging in order to elucidate the origin of the observed physical properties of PNSs. We have mainly investigated the radial profiles of projected number densities, rotational velocities, and velocity dispersion of PNSs and the two-dimensional velocity fields of PNSs in elliptical galaxies. We also compared the simulated dynamical properties of PNSs with the corresponding observations of the PNS in NGC 5128 and thereby tried to provide the best merger model for the PNS in NGC 5128. We summarize our principal results as follows.  (1) The radial densities profiles of PNSs  can be fitted to the power-law in the entire halo regions of elliptical galaxies formed from major merging ($m_{2}$ = 1.0). The projected PNe number densities (${\\rho}_{\\rm PNe}$) at $R$ $\\sim$ 5$R_{\\rm e}$ in elliptical galaxies are more than an order of magnitude higher than those of the halos of the merger progenitor spirals. These results do not depend strongly on the model parameters of major merging. The main reason for  rather high ${\\rho}_{\\rm PNe}$ of PNSs is that a significant fraction ($\\sim$ 10\\%) of disk stars are stripped from the disks and redistributed in the halo regions during major merging. These results suggest that elliptical galaxies have higher PNe number densities (${\\rho}_{\\rm PNe}$) in their halos than spiral ones. \\\\  (2) PNSs of elliptical galaxies formed from major merging can show a significant amount of rotation ($V_{\\rm rot}/\\sigma$ $>$ 0.5) in their outer halo regions ($R$ $>$ $5R_{\\rm e}$). The derived rotation in the outer halos results from the angular momentum redistribution of disk stars during galaxy merging (i.e., conversion of orbital angular momentum of merging two spirals into the intrinsic one of the merger remnant). $V_{\\rm rot}/\\sigma$ of PNSs in ellipticals are more likely to be larger in the outer parts than in the inner ones,  though radial profiles of $V_{\\rm rot}/\\sigma$ are diverse between different models. $V_{\\rm rot}/\\sigma$ at $5R_{\\rm e}$ can weakly correlate with $V_{\\rm rot}/\\sigma$ at $2R_{\\rm e}$ for the PNSs in such a way that PNSs with larger $V_{\\rm rot}/\\sigma$ at $2R_{\\rm e}$ are likely to show larger $V_{\\rm rot}/\\sigma$ at $5R_{\\rm e}$. This result implies that the kinematics of outer stellar halos in elliptical galaxies can correlate with that of the main bodies.\\\\  (3) The two-dimensional (2D) fields of velocity and velocity dispersion of PNSs in elliptical galaxies formed by major merging are  quite diverse depending on the orbital configurations of galaxy merging, the mass ratios of the progenitor spirals, and the viewing angles of the galaxies. For example, PNSs of ellipticals formed from prograde-prograde merging can show more rapid rotation along the major axes of the PNe density profiles in their 2D velocity fields compared with the retrograde-retrograde ones. For most of the models, the 2D velocity fields show minor axis rotation in the outer halo regions of elliptical galaxies. It is doubtlessly worthwhile for future observational studies to systematically investigate dynamical properties of PNS for a larger number of elliptical galaxies and thereby to confirm whether the predicted uniformity and diversity in  the 2D velocity fields of PNSs can be seen in elliptical galaxies. \\\\    (4) ${\\rho}_{\\rm PN}$ of PNSs are  more likely to be lower in E/S0s formed from galaxy merging with smaller mass ratios (e.g., $m_2$ = 0.3, i.e.,  unequal-mass merging). Furthermore, PNSs in the remnants of mergers with smaller $m_2$ have more flattened shapes and a larger amount of rotation ($V_{\\rm rot}$) both in their main bodies and in their outer halos. Therefore these  results suggest that more flattened E/S0s are more likely to show both lower ${\\rho}_{\\rm PNe}$ and higher $V_{\\rm rot}$ (for a given luminosity range of E/S0s). This predicted correlation can  be readily confirmed in future observations of PNSs.\\\\     (5) The observed kinematical properties of the PNS in NGC 5128 (e.g.,  $V_{\\rm rot}/\\sigma$ between 1 and 1.5  and minor axis rotation) can be broadly consistent with the present best model with $m_2$ = 0.5, a small impact parameter, and highly inclined initial disks (similar to a  polar-orbit). The observed ``kink'' in the zero velocity curve (ZVC) of the 2D velocity field of the PNS in NGC 5128 can be reproduced reasonably well, though the location where ZVC begins to twist is different between the best simulation model and the observation. Some disadvantages of the present best model in explaining self-consistently the observed kinematics of the PNS suggest that gas dynamics and star formation may well play an important role in the formation of the PNS observed in NGC 5128.\\\\  (6) The mass estimates of merger remnants viewed face-on are likely to be a factor of two lower than those viewed edge-on.  However, even the mass estimates of systems viewed edge-on can be low by a factor of 5, and in the worst case (face-on E/S0 model) can be low by a factor of 50.  If this conclusion is applicable to real early-type galaxies, it has interesting consequences for current observational work."
    },
    "0606/astro-ph0606636_arXiv.txt": {
        "abstract": "Using analytic calculations and N-body simulations we show that in constant density (harmonic) cores, sinking satellites undergo an initial phase of very rapid (super-Chandrasekhar) dynamical friction, after which they experience no dynamical friction at all. For density profiles with a central power law profile, $\\rho \\propto r^{-\\alpha}$, the infalling satellite heats the background and causes $\\alpha$ to decrease. For $\\alpha < 0.5$ initially, the satellite generates a small central constant density core and stalls as in the $\\alpha = 0$ case.  We discuss some astrophysical applications of our results to decaying satellite orbits, galactic bars and mergers of supermassive black hole binaries. In a companion paper we show that a central constant density core can provide a natural solution to the timing problem for Fornax's globular clusters. ",
        "introduction": "\\label{sec:introduction} In a seminal paper, \\citet{1943ApJ....97..255C} showed that a massive particle moving through an infinite, homogeneous and isotropic background of lighter particles  experiences a force of dynamical friction given by:  \\begin{equation} M_c \\frac{dv}{dt} = -4\\pi G^2 M_c^2 \\frac{v}{|v|^3}\\ln\\left(\\frac{b_\\mathrm{max}}{b_\\mathrm{min}}\\right)\\int_0^{|v|}M(v')dv' \\label{eqn:chandrasekhar} \\end{equation} where $M_c$ and $v$ are the mass and velocity of the in-falling particle, $M(v')dv'$ is the mass density of background objects with speeds $v' \\rightarrow v'+dv'$, and $b_\\mathrm{max}$ and $b_\\mathrm{min}$ are the maximum and minimum impact parameters for the encounters\\footnote{Note that $b_\\mathrm{min} \\rightarrow 0$ can be achieved (see e.g. \\bcite{1976MNRAS.174..467W}; \\bcite{1987gady.book.....B} p. 423), while $b_\\mathrm{max} \\rightarrow \\infty$ cannot. This is because the derivation of equation \\ref{eqn:chandrasekhar} assumes an {\\it infinite} background; $b_\\mathrm{max}$ defines a scale on which the infinite background should be truncated. This is often, reasonably, take to be the radius at which the mean density falls by a factor of two or so.}. From here on, we refer to the massive in-falling object as a `Globular Cluster' (GC) and the background of lighter particles as simply `particles'. However, we could equally refer to, for example, a bar moving in a background of stars and dark matter.  While equation \\ref{eqn:chandrasekhar} is only strictly valid for an infinite, homogeneous and isotropic background, it has been shown to work remarkably well for satellites orbiting in spherical galaxies with more general background distributions\\footnote{When corrected for velocity anisotropies, it has also been shown to work well in aspherical systems (see e.g. \\bcite{1977MNRAS.181..735B}, \\bcite{1991ApJ...375..544S} and \\bcite{2004MNRAS.349..747P}).} (see e.g. \\bcite{1983ApJ...274...53W}, \\bcite{1988MNRAS.235..289Z}, \\bcite{1997MNRAS.289..253C} and \\bcite{1987MNRAS.224..349B}). Such successes make equation \\ref{eqn:chandrasekhar} of great practical value. But they beg the question: why has it been so successful, even when it is used so far beyond its expected regime of validity? Are we missing important physical insight into the dynamical friction process in spherical systems? Does Chandrasekhar fail to work well in some situations?  In order to address some of these issues, \\citet{1984MNRAS.209..729T} (hereafter TW84) and \\citet{1986ApJ...300...93W} (hereafter W86) formulated a perturbative theory of dynamical friction which could be applied to spherical systems. Notice from equation \\ref{eqn:chandrasekhar}, that most of the dynamical friction originates from particles with large impact parameters: it is the accumulation of many long range {\\it small} interactions which leads to most of the dynamical friction; not the large angle scattering of close encounters. This is why perturbative methods can be used. TW84 and W86 consider a general, small, perturbation to a single background particle; and then sum over all particles in the system to obtain to total torque induced on the perturber.  In section \\ref{sec:analytic}, we will briefly summarise the essence of this perturbation method. For now, it is important to note the key assumptions in the method; and the key results. The two main assumptions are: (i) that the perturbation is small and (ii) that the frequency of the perturber changes with time, $\\Omega_s = \\Omega_s(t)$, {\\it faster than the perturbation can grow non-linear}. For most potentials of interest this second assumption is satistifed. The perturber (the GC) will lose angular momentum to the background particles as a result of the dynamical friction, and $\\Omega_s$ will then increase as the GC falls inwards.  Under the above assumptions, the perturbation method gives us new physical insight into the dynamical friction problem. To the order of the perturbation approximation, {\\it all} of the torque comes from background particles which are close to resonance with the perturber. Non-resonant particles do not contribute to the friction at all. This is a key difference between the perturbation solution and that of equation \\ref{eqn:chandrasekhar}. It suggests that if equation \\ref{eqn:chandrasekhar} is ever going to fail, it would do so for background particle distributions which are especially resonant.  In this paper, we describe such a super-resonant potential: that of the constant density (harmonic) core. For this special potential, all particles and the perturber always move with constant angular frequency, $\\Omega$. In this case, perturbation methods can no longer be used. This is because $\\Omega_s = \\mathrm{const}$; assumption (ii), above, is violated; and the perturbations, which are always driven at the same resonant frequency, can grow indefinitely\\footnote{This is true for {\\it any} perturbative scheme (e.g. \\bcite{1999ApJ...525..720C}).}.  To cope with this special case, we develop a non-perturbative analytic model using a 3D driven harmonic oscillator. This essentially generalises an earlier result derived by \\citet{1972gnbp.coll.....L}. Using our analytic model and N-body simulations, we show that in constant density cores, equation \\ref{eqn:chandrasekhar} fails. Sinking satellites undergo an initial phase of very rapid (super-Chandrasekhar) dynamical friction, after which they experience little or no dynamical friction at all.  \\citet{2005astro.ph..8166W} and \\citet{2006astro.ph..1138W}, find similar stalling results for galactic bars (which may be thought of as two diametrically opposed satellites) inside constant density cores.  Constant density cores have recently become interesting in astrophysics. Observations of galaxies on all scales from dwarf spheroidals in the Local Group, up to giant spirals suggest that their central dark matter density has such a constant density core on the scale of $\\sim 1$\\,kpc (see e.g. \\bcite{2003ApJ...588L..21K}, \\bcite{2001ApJ...552L..23D}, \\bcite{2001MNRAS.323..285B} and \\bcite{2001MNRAS.327L..27B}); but see also \\citet{2004MNRAS.355..794H} and \\citet{2004ApJ...617.1059R},  for a discussion of the potential systematic errors in such observations. If cores are present at the centre of galaxies, their resonant properties can significantly affect the dynamics. Bars can be much longer lived\\footnote{\\citet{1998ApJ...493L...5D} and \\citet{2000ApJ...543..704D} show that low central dark matter densities lead to bars which remain fast. Here we discuss the extreme case of constant density cores, in which we show that bars would not slow down at all. This agrees well with earlier findings by \\citet{2005astro.ph..8166W} and \\citet{2006astro.ph..1138W}.}, while in-falling satellites and GCs will stall at the core radius. In a companion paper, \\citet{Goerdt:2006rw}, we investigate this last idea further (see also \\bcite{1998MNRAS.297..517H} and \\bcite{2006astro.ph..1490S}). The Fornax dwarf spheroidal galaxy in the Local Group has 5 GCs at a range of projected radii. Application of Chandrasekhar dynamical friction suggests the clusters should rapidly fall to the centre of Fornax from their current positions; fine tuning is required to have them arrive at their present positions at the current epoch. In \\citet{Goerdt:2006rw} we show that a small core of radius greater than 0.24\\,kpc can solve this problem by causing some, or all, of Fornax's GC to stall.  This paper is organised as follows: in section \\ref{sec:analytic} we briefly review the perturbative method for calculating dynamical friction and demonstrate that it fails for the special case of a constant density core. We show that insight can be gained from a non-perturbative approach by modelling the system as a driven harmonic oscillator. In section \\ref{sec:simulations} we describe our semi-analytic and full N-body simulations. In section \\ref{sec:results}, we test our analytic model against these high resolution ($\\sim 10^7$ particles) simulations of satellites sinking in harmonic cores. We demonstrate that such high resolution is required in order to reduce numerical precession of the GC orbit plane, but that near-converged results for the GC orbit can be obtained at lower resolution with O($10^6$ particles). We discuss the importance of the initial GC orbit, mass, the underlying gravitational potential and the particle-particle interactions. Finally, in section \\ref{sec:conclusions} we briefly discuss the implications of these results and present our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Using analytic calculations and N-body simulations we have shown that in constant density harmonic cores, sinking satellites undergo an initial phase of very rapid (super-Chandrasekhar) dynamical friction, after which they experience no dynamical friction at all. This occurs because, for the special case of harmonic potentials, there are stable solutions where the background particles move on epicycles about the in-falling satellite. The system moves rapidly into this stable configuration. In doing so, the satellite experiences a brief moment of enhanced friction. Once in equilibrium, there is no net momentum transfer between the background particles and the satellite and friction ceases. For density profiles with a central power law profile, $\\rho \\propto r^{-\\alpha}$, the infalling satellite heats the background and causes $\\alpha$ to decrease. For $\\alpha < 0.5$ initially, the satellite generates a small central constant density core and stalls as in the $\\alpha = 0$ case.  Our results concerning dynamical friction stalling in constant density cores are of broad astrophysical interest. Recent observational work suggests that galaxies may have central dark matter density cores, rather than the $r^{-1}$ density cusps predicted by numerical simulations. Galactic bars orbiting in such potentials will experience very weak dynamical friction and can be very long-lived (in fact central density distributions do not need to be pure harmonic to see this effect, low-density will also lead to very little friction -- see e.g. \\bcite{1998ApJ...493L...5D} and \\bcite{2000ApJ...543..704D}). Satellites falling into such galaxies will stall at the core radius and never make it to the centre. This point was investigated in a companion paper \\citep{Goerdt:2006rw}, where we suggested that a constant density core could solve the `timing problem' for the GCs in the Fornax dwarf galaxy. Finally, recent work on merging black holes suggests that they can form a central constant density core in the background distribution (see e.g. \\bcite{2002MNRAS.331L..51M} and \\bcite{2002ApJ...566..801R}) prior to forming a hard binary. If true, our results suggest that this could further exacerbate the well-known problem of getting the binaries to coalesce. Their rate of hardening will stall, even before the majority of stars and dark matter have been ejected from the core, if the background distribution is close to constant density. This may point towards gas playing a more important role in bringing supermassive black holes together at the centre of galaxies (see e.g. \\bcite{2000ApJ...532L..29G})."
    },
    "0606/astro-ph0606366_arXiv.txt": {
        "abstract": "We have developed a completely lithographic antenna-coupled bolometer for CMB polarimetry. The necessary components of a millimeter wave radiometer --- a beam forming element, a band defining filter, and the TES detectors --- are fabricated on a silicon chip with photolithography. The densely populated antennas allow a very efficient use of the focal plane area. We have fabricated and characterized a series of prototype devices. We find that their properties, including the frequency and angular responses, are in good agreement with the theoretical expectations. The devices are undergoing optimization for upcoming CMB experiments. ",
        "introduction": "\\label{sect:intro}  %   The primary science goals of the next generation Cosmic Microwave Background (CMB) polarization experiments are to produce a high fidelity E-mode power spectrum and to search for the B-mode polarization. A detailed E-mode spectrum will help constrain the cosmological parameters, and test crucial hypothesis in standard cosmology. The measurements of B-mode polarization at high multipole value ($\\ell$) can be used to study the matter power spectrum through gravitational lensing. At low $\\ell$, B-mode polarization can only be produced by gravitational waves. The detection of this signature would provide invaluable information on the fundamental physics in the very early Universe, such as the energy scale of Inflation. To achieve these goals in the presence of the astronomical foregrounds, the instruments will require wide frequency coverage, a large number of sensitive microwave detectors, and exquisite control of systematics\\cite{weiss05}.  Bolometers can provide photon noise-limited sensitivity over a wide frequency range. However, the existing feedhorn-coupled micromesh bolometers \\cite{bock95,turner01,jones03} face difficulties in extending the frequency coverage and scaling to kilo-pixel instruments. The micromesh bolometers with low thermal conductance (both the ``spiderweb''\\cite{bock95} and the ``PSB''\\cite{jones03} type) have currently been demonstrated with architectures suitable for frequencies $\\nu>$60 GHz. The metallic corrugated feedhorns used for beam collimation are heavy and expensive to implement in systems with thousands of pixels at cryogenic temperatures. Semiconductor bolometers with high impedances ($> M\\Omega$ ) pose a challenge for kilo-pixel readout. Each of the bolometers requires one impedance-matching JFET, which dissipates significant amount of electrical power. The noise contribution from standard JFET is prohibitively large for use in multiplexers. So far, the number of elements in a semiconductor bolometric receiver is limited to a few hundreds \\cite{turner01}.  \\begin{figure}[tbp] \\begin{center} \\includegraphics[scale=0.6]{figure1.eps} \\caption{{\\em Left}: The focal plane of the ESA CMB satellite mission {\\em Planck}, to be launched in 2008. {\\em Right}: A 150 GHz prototype antenna-coupled bolometer array with $8\\times 8$ spatial pixels and $256$ bolometers. The new architecture is advantageous in weight, cost, and detector density.} \\label{focalplane} \\end{center} \\end{figure}  We describe a CMB polarimeter based on antenna-coupled superconducting transition edge sensors (TES) that can avoid all these difficulties. In this design, the beam forming slot-antenna array, microstrip band defining filter, and TES detectors are fabricated on a silicon chip. These completely photolithographic detectors provide intrinsic polarization sensitivity and purity, collimated beams, simultaneous measurement of Stokes I and Q parameters in a single spatial pixel, high focal-plane packing density, and massively parallel fabrication. The migration from semiconductor bolometers to superconducting TES bolometers enables the readout of thousands of pixels with moderate electronics complexity. This is achieved by superconducting quantum interference device (SQUID) multiplexers \\cite{dekorte03,lanting05}. In a microstrip-coupled bolometer, only the mechanically robust slot antenna scale with wavelength, therefore the entire frequency range ($\\sim$ 30 GHz to 500 GHz) of interest in CMB science can be covered by the same technology.  The densely populated antennas allow a very efficient use of the focal plane area. The excellent uniformity of the photolithography provides intrinsically good alignment and orthogonality of the two polarization modes, uniformity of the in-line filter band-passes, and gain matching of all elements for two orthogonal polarizations. Besides TES, the microwave antenna-coupled architecture is completely compatible with other detector technology, such as indium bump-bonded NTD chips and microwave kinetic inductance detectors.  The paper is organized as follows. We describe the components of a polarization sensitive antenna-coupled bolometer in \\S\\ref{sec:components}. The design, fabrication, and measured optical properties of prototype detectors are presented in \\S\\ref{sec:prototype}. Issues related to the implementation of arrays of antenna-coupled bolometers in a polarization receiver are discussed in \\S\\ref{sec:system}, including a comparison of on-chip polarization modulators and half-wave plates (\\S\\ref{subsec:mod}). We describe our detector development plan and several proposed CMB polarization experiments that will be using this technology in \\S\\ref{sec:plans}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606150_arXiv.txt": {
        "abstract": "{In order to understand the evolution of molecular clouds it is important to identify the departures from self-similarity associated with the scales of self-gravity and the driving of turbulence.}{ A method is described based on structure functions for determining whether a region of gas, such as a molecular cloud, is fractal or contains structure with characteristic scale sizes. } {Using artificial data containing structure it is shown that derivatives of higher order  structure functions provide a powerful way to detect the presence of characteristic scales should any be present and to estimate the size of such structures. The method is  applied to observations of hot H$_2$  in the Kleinman-Low nebula, north of the Trapezium stars in the Orion Molecular Cloud, including both brightness and velocity data. The method is compared with other techniques such as Fourier transform and histogram techniques.} { It is found that the density structure, represented by H$_2$ emission brightness in the K-band (2-2.5$\\mu$m), exhibits mean characteristic sizes of 110, 550, 1700 and 2700AU. The velocity data show the presence of structure at 140, 1500 and 3500AU. Compared with other techniques such as Fourier transform or histogram, the method appears both more sensitive to characteristic scales and easier to interpret.} {} ",
        "introduction": "The velocity and density fields of molecular clouds have been observed to show very complex structure. This property suggests that structure is dominated by turbulence and this has been verified by both observations and numerical simulations \\citep[e.g.][]{falgarone1990,boldyrev2002apj}; see also the extensive reviews by \\cite{elmegreen2004} and \\cite{maclow2004}. Clouds dominated by turbulence lack any characteristic scale and are inherently self-similar with a variety of fractal dimensions. This has been observed over a broad range of scales, 0.02pc-100pc \\citep[e.g.][]{larson1981,chappell2001b,ossenkopf02}.  Self-similarity and lack of characteristic scale in molecular clouds must however break down at some dimension or range of dimensions associated with star formation, as demonstrated in \\cite{vannier} and \\cite{gustafsson2006}. Turbulence creates local clumps of gas and this process, known as turbulent fragmentation, is considered the root cause of observed density structure and may give rise to large density contrasts. Star formation begins to take place at smaller scales in gravitationally unstable clumps.  The distribution of sizes of these clumps is directly related to the initial mass function (IMF) for star formation \\citep{padoan2002}. The predominantly turbulent era, characterized by lack of preferred scale, is expected to end when gravitation becomes dominant in clumps. At this stage, the density and velocity fields are not fractal, but will contain characteristic scales. Hence certain scales may be over-populated or under-populated relative to the model of a simple Kolmogorov-type cascade \\citep{kolmogorov1941,frisch1995}. The range of scales, at which clumps become gravitationally bound and turbulence looses its kinematically dominating effect, is important for both theories of star formation and of turbulence.  Characteristic scales may also be imposed on the medium by other mechanisms than gravitational collapse. In particular energy may be injected into the medium at the largest scales by supernova explosions or at the tens of thousands of AU scale by outflows from massive stars or from low mass stars at scales of a few thousand AU. In order to understand the evolution of molecular clouds it is essential to identify the presence of these energy injection processes and the scales at which they occur.   The purpose of the present paper is to introduce a technique, based on structure functions (SFs), to identify characteristic scales in the interstellar medium - or indeed in any medium. The technique appears more sensitive to characteristic scales, for example in spatial brightness or velocity, than methods presently used. The method is relatively easy to interpret in terms of preferred scale sizes and is as easy to implement as existing methods.  Experience has proven that it is non-trivial to establish the scale(s) at which self-similarity begins to break down and an array of techniques as broad as possible should be made available. For example, four different methods were used in \\cite{vannier} in order to demonstrate that self-similarity breaks down at star-forming scales of around 500 - 1000 AU in a small portion of the Orion Molecular Cloud (OMC1). The present contribution should be seen in this light, not as a panacea for structure identification but rather as another tool for this purpose.   The methods currently available are those of (i)~\\cite{blitz}, who introduced a method of degrading the spatial resolution and establishing the presence of preferred scales from histograms of the degraded images, a method used for example in~\\cite{vannier} and~\\cite{lacombe2004}, (ii) Fourier transform power spectra, used to pick out frequencies and hence scales in images~\\citep[e.g.][]{vannier}, (iii) clump decomposition \\citep{stutzki1990}, (iv) $\\Delta$-variance analysis \\citep[][]{stutzki1998, bensch2001}, (v) wavelet transformation~\\citep[e.g.][]{farge1992}.  The non-fractal nature of molecular clouds at the star-forming scale of $\\sim$~500-1000 AU was first detected in \\cite{vannier} for Orion. Essentially the same preferred scales between 500 and 1100 AU are also reported in \\cite{lacombe2004}. Non-fractal nature has also been reported at larger scales in \\cite{blitz} who found a preferred scale of 0.25-0.5 pc ~(5-10$\\times 10^4\\,$AU) in Taurus. Material may however be self-similar down to a few x 10$^{-4}$pc ($<$100 AU) in diffuse material, with clumps of mass~$\\sim$~10$^{-7}$M$_{\\odot}$, as found in the photodissociation region around the B-star HD37903 in NGC2023 \\citep{rouan1997}. Here we show that structure functions (SFs) or local derivatives of SFs provide a powerful way to detect preferred scales in a medium. The method is tested on observations of H$_2$ emission in OMC1. ",
        "conclusions": "In conclusion the technique presented here, whilst suffering from some degree of uncertainty and imprecision to an extent common to all methods, has proved valuable for analysis of images in our own work. With the present and future advent of higher spatial resolution maps in the radio region of the spectrum, probing the star-forming scale, for example obtained with the Plateau de Bure interferometer, the Submillimeter Array, CARMA and ALMA, this tool may prove useful to an increasing community of observers."
    },
    "0606/astro-ph0606016_arXiv.txt": {
        "abstract": "We present spectroscopic redshifts for the first 466 X-ray and radio-selected AGN targets in the 2 deg$^2$ COSMOS field.  Spectra were obtained with the IMACS instrument on the Magellan (Baade) telescope, using the nod-and-shuffle technique.  We identify a variety of Type 1 and Type 2 AGN, as well as red galaxies with no emission lines.  Our redshift yield is 72\\% down to $i_{\\rm AB}=24$, although the yield is $>90\\%$ for $i_{AB}<22$.  We expect the completeness to increase as the survey continues.  When our survey is complete and additional redshifts from the zCOSMOS project are included, we anticipate $\\sim 1100$ AGN with redshifts over the entire COSMOS field.  Our redshift survey is consistent with an obscured AGN population that peaks at $z \\sim 0.7$, although further work is necessary to disentangle the selection effects. ",
        "introduction": "The Cosmic Evolution Survey \\citep[COSMOS, ][]{cosmos} is an HST Treasury project to fully image a 2 deg$^2$ equatorial field.  The 590 orbits of HST ACS $i$-band observations have been supplemented by observations at wavelengths from X-ray to radio and a major galaxy redshift survey \\citep[zCOSMOS, ][]{zcosmos} carried out with VLT/VIMOS.  The details of the COSMOS AGN survey are found in a companion paper \\citep{cosmosagn}.  Here we present the first X-ray and radio-selected active galactic nuclei (AGN) candidates observed with the IMACS instrument on the Magellan (Baade) telescope.  X-ray observations provide the most efficient method to find Type 1, Type 2, and particularly obscured AGN.  The XMM-Newton observations of the COSMOS field are expected to reach an AGN surface density of $\\sim$1000 deg$^2$.  The current COSMOS X-ray catalog is presented by \\citet{cosmosxray} and has a 0.5-2 keV flux limit of $1 \\times 10^{-15}~{\\rm erg~cm^{-2}~s^{-1}}$ and a 2-10 keV flux limit of $3.3 \\times 10^{-15}~{\\rm erg~cm^{-2}~s^{-1}}$.  The identification of optical counterparts, based on the ``likelihood ratio'' technique, is presented by \\citet{brusa}.  The X-ray selected targets for our IMACS survey were the $i_{\\rm AB}<24$ optical counterparts of \\citet{brusa} that were X-ray point sources with detection in either the 0.5-2 keV or 2-10 keV bands, available at the time of our IMACS observations. Multiple X-ray observations over most of the COSMOS field mitigate the effects of vignetting in the outer region of the XMM-Newton field of view.  The edges of the COSMOS field, however, are observed only once by XMM-Newton and so our observations in these regions must sample a lower density of X-ray selected AGN candidates.  Radio-selected AGN candidates were our second-highest priority targets for IMACS observations.  The COSMOS VLA survey is described by \\citet{cosmosradio}; we use a preliminary VLA catalog with a $4\\sigma$ flux limit of 0.1-0.4 mJy at 1.4 GHz and full coverage across the COSMOS field.  Approximately 20\\% of the radio selected AGN candidates overlapped with the X-ray sample.  We observed only radio sources with radio peak flux $S/N\\ge 4$ and unambiguous optical counterparts within 1'' of the radio peak of magnitude $i_{\\rm AB}<24$.  In \\S 2 we present the details of our observing strategy and set-up, as well as the reduction and calibration of the observations.  We present the classifications and redshifts of our targets in \\S 3, along with estimates of our completeness and other properties of the sample.  We summarize our results in \\S 4 and discuss our timeline for completing the survey. ",
        "conclusions": "The COSMOS AGN survey will provide a large sample of AGN with bolometric measurements from radio to X-ray and supplementary observations of their hosts and local environments.  Here we have presented spectra and redshifts for the first 466 X-ray and radio-selected AGN targets: we have discovered 86 new Type 1 AGN and 130 new Type 2 AGN with high-confidence redshifts and reliable classification.  Our overall redshift yield is $72\\%$, although we are $90\\%$ complete to objects of $i_{\\rm AB} < 22$.  We expect this yield to increase as refurbishments to IMACS take place.  While the survey may be affected by redshift-dependent selection effects for $i_{\\rm AB} > 22$, our findings support an obscured AGN population that peaks at $z \\sim 0.7$.  Our observations with IMACS are designed to cover the entire COSMOS field over three seasons, and a high overall yield will be obtained thanks to spectra taken by VLT/VIMOS during the zCOSMOS redshift survey \\citep{zcosmos}.  With 7/16 IMACS pointings successfully observed, we are on schedule to complete our AGN survey in early 2007."
    },
    "0606/astro-ph0606700_arXiv.txt": {
        "abstract": "We present a catalog of 9316 spectroscopically confirmed white dwarfs from the Sloan Digital Sky Survey Data Release 4.  We have selected the stars through photometric cuts and spectroscopic modeling, backed up by a set of visual inspections.  Roughly 6000 of the stars are new discoveries, roughly doubling the number of spectroscopically confirmed white dwarfs.  We analyze the stars by performing temperature and surface gravity fits to grids of pure hydrogen and helium atmospheres. Among the rare outliers are a set of presumed helium-core DA white dwarfs with estimated masses below $0.3\\msun$, including two candidates that may be the lowest masses yet found.  We also present a list of 928 hot subdwarfs. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS) is greatly expanding the number of spectroscopically confirmed white dwarf stars. \\citet{Har03} presented a catalog of 269 white dwarfs from the Early Data Release, and \\citet{Kle04} followed this with a catalog of 2551 white dwarfs from SDSS Data Release 1. More specialized SDSS catalogs have been presented for magnetic white dwarfs \\citep{Sch03,Van05}, white dwarf plus M star binary systems \\citep{Ray03,Smo04,Sil05}, and cataclysmic variables \\citep{Szk02,Szk03,Szk04,Szk05}. The size of the SDSS white dwarf catalog has significantly expanded samples of rare white dwarfs, including ultra-cool white dwarfs \\citep{Har01,Gat04,Kil05}, stars with atomic carbon and oxygen \\citep{Har03,Lie03}, pulsating DAV stars \\citep{Muk04,Mul05} and DBV stars \\citep{Nit05}, hot DO stars \\citep{Krz04}, very low accretion rate magnetic cataclysmic variables \\citep{Szk03b}, low-mass degenerate helium core stars \\citep{Lie04}, and hot DB stars in the ``DB gap'' \\citep{Eis05db}.  Here we present a catalog of white dwarfs and hot subdwarf stars from the SDSS Data Release 4 \\citep{Aba05}, which presents 800,000 spectra from 4783 square degrees. We use automated techniques supplemented by visual classification to select 13,000 candidates. An extensive analysis of these objects yields 9316 white dwarfs (WD), including 8000 DAs, 713 DBs, 41 DOs or PG1159 stars, 289 DCs, 104 DQs, and 133 DZs.  We find 928 hot subdwarf stars (SD), building on the samples from \\citet{Har03} and \\citet{Kle04}. We fit the WDs to grids of model atmospheres and present temperatures and surface gravities of the DA and DB stars. We also present 774 duplicate spectra of WDs and 60 duplicate spectra of SDs.  This catalog is not complete, even within the SDSS spectroscopic sample, as we have intentionally focused on the ``easy'' white dwarfs, namely those that fall in the blue color region of DA and DB stars. Low temperature stars that fall too near the colors of A and F stars have not been searched.  Cataclysmic variables have been excluded, as they have been separately cataloged.  White dwarf plus M star binaries are included where easily identified, but the catalog of \\citet{Sil05} is more complete and better analyzed.  Magnetic WDs are sufficiently heterogeneous to fool our automated techniques, but we include them where possible.  DC stars are difficult to positively identify, although we do find a fair number.  Although completeness was not our goal, we expect that the catalog does include nearly all of the SDSS white dwarf spectra for the common classes of hotter DA, DB, and DO. We have attempted to find DQ and DZ stars, although this effort is often limited by the signal-to-noise ratio of the spectra. ",
        "conclusions": "We have identified 9316 spectroscopically confirmed white dwarfs and 928 hot subdwarfs from the 800,000 spectra of the SDSS Data Release 4. Rolling together all secondary classifications, the sample includes 8000 DA, 713 DB, 289 DC, 41 DO/PG1159, 104 DQ, and 133 DZ stars.  There are of course many different mixed and subclassifications.  More than 6000 of these WDs are new, so this catalog roughly doubles the number of spectroscopically confirmed white dwarfs.  We have fit model atmospheres to the DA and DB stars using our autofit method.  We find the expected strong peak in surface gravity, along with interesting outliers, such as the low-mass candidates discussed in \\S~\\ref{sec:lowmass}.  The fitting results are validated for internal consistency using repeat observations and by correlations with the photometry.  There is a systematic bias toward higher gravities at lower temperatures of which users should beware.  This appears to be an artifact of the model atmospheres.  There is also a bias toward over-estimating the temperature of DA stars above 30,000~K.  This is caused a subtle systematic error in the SDSS spectrophotometry causing the broad Balmer lines in the blue to be slightly filled in.  Among our fitted stars, we continue to find a rare population of low surface-gravity candidates that are very likely helium-core, low-mass DA white dwarfs, building on the sample of \\citet{Lie04}. Two of these candidates may be the lowest mass DA stars yet found, although this awaits confirmation.  These stars are very similar to the two companions of massive pulsars \\citep{Van96,Nic05}, opening the hypothesis of a generic evolutionary role for neutron star companions in the formation of these white dwarfs.  We cannot claim that this catalog is complete, particularly for difficult cases such as DC stars, low-temperature DA stars, binary systems, and DH stars.  For mundane warm to hot isolated DA, DB, and DO stars (above $\\sim\\!8,000~K$ for DA stars), we believe that the catalog should be reasonably complete (95\\% or more) within the SDSS spectroscopic catalog and down to spectral signal-to-noise ratios where one can expect to classify the stars.  In addition to searching for unusual classes of white dwarfs, large catalogs open the possibility of large-number statistical studies of the distribution of white dwarf properties as well as their Galactic structure.  Although the SDSS spectroscopic target selection is not complete for these stars, the incompletenesses can be modeled.  We have discussed the key aspects of this modeling in \\S~\\ref{sec:complete}.  The SDSS DR4 represents roughly half of the final spectroscopic sample for the original SDSS survey.  In addition, the new Sloan Extension for Galactic Understanding and Exploration (SEGUE) is conducting an extensive study at lower Galactic latitudes with both imaging and spectroscopy.  Although white dwarfs are rare within the SDSS catalogs, the available numbers are sufficient to open a key new opportunity to study this diverse category of stars.   \\bigskip  We thank Detlev Koester for providing grids of his recent model atmospheres. We thank Betsy Green and Gary Schmidt for useful discussions.  We thank David Johnston and Ryan Scranton for use of their stacked catalogs from the SDSS equatorial region and Jay Holberg for help matching to the \\citet{McC99} catalog. DJE was supported by NSF AST-0098577 and AST-0407200 and by an Alfred P.\\ Sloan Research Fellowship.  JL was supported by NSF AST-0307321.  Funding for the Sloan Digital Sky Survey (SDSS) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0606/astro-ph0606470_arXiv.txt": {
        "abstract": "{The COMPTEL unidentified source GRO~J1411-64 was observed by INTEGRAL, and its central part, also by XMM-Newton.} {The aim of these observations and subsequent analysis and theoretical investigations pursued is to shed light upon the nature of such source.} {Usual data analysis techniques of INTEGRAL and  XMM-Newton were used in this work.} {The data analysis shows no hint for new detections at hard X-rays. The upper limits in flux herein presented constrain the energy spectrum of whatever was producing GRO~J1411-64, imposing, in the framework of earlier COMPTEL observations, the existence of a peak in power output located somewhere between 300-700 keV for the so-called low state. The Circinus Galaxy is the only source detected within the 4$\\sigma$ location error of GRO~J1411-64, but can be safely excluded as the possible counterpart:  the extrapolation of the energy spectrum is well below the one for GRO~J1411-64 at MeV energies. 22 reliable and statistically significant sources (likelihood $> 10$) were extracted and analyzed from XMM-Newton data. Only one of these sources, XMMU J141255.6-635932, is spectrally compatible with GRO~J1411-64 although the fact the soft X-ray observations do not cover the full extent of the COMPTEL source position uncertainty make an association hard to quantify and thus risky. }{ The unique peak of the power output at high energies (hard X-rays and gamma-rays) resembles that found in the SED seen in blazars or microquasars, and might suggest that a similar scenario is at work. However, an analysis using a microquasar model consisting on a magnetized conical jet filled with relativistic electrons which radiate through synchrotron and inverse Compton scattering with star, disk, corona and synchrotron photons shows that it is hard to comply with all observational constrains. This fact and the non-detection at hard X-rays introduce an a-posteriori question mark upon the physical reality of this source, which is discussed in some detail. } ",
        "introduction": "The COMPTEL experiment onboard the Compton Gamma-Ray Observatory (CGRO) surveyed the sky in the $0.75-30$ MeV energy range and detected about a dozen sources of Galactic origin. Among them, pulsars, stellar black-hole candidates, and supernova remnants were identified by the telescope (Sch\\\"onfelder et al. 2000). In addition, nine unidentified sources were also reported in the first COMPTEL Catalog. Zhang et al. (2002) have reported the discovery of a new source belonging to this group. It is a variable source located near the Galactic plane. The fluxes of this COMPTEL unidentified source along the first five phases of the Compton satellite were derived by  Zhang et al. (2002), and are shown, together with the corresponding confidence level of each detection, in Table 1. Time periods and effective exposures for different observations are also listed there. The best detection significance, at more than 7$\\sigma$, was obtained in the 1-3 MeV band by combining 7 viewing periods (VPs, the periods of observations in the Compton satellite) during 1995 March-July, VP 414-424. The probability of randomly detecting a source with such high level of significance in the COMPTEL 1-3~MeV band was estimated to be 5.4$\\times$10$^{-10}$; obtained by taking into account the trials for searching in all individual VPs in four energy bands from the beginning of the Compton mission in 1991 to its second reboost in 1997 (Zhang et al. 2002). The most likely source location, i.e. the maximum of the likelihood ratio distribution in the 1-3 MeV band for the combination of VPs in 1995,  is  $(l,b) = (311.5$\\deg ,$-2.5$\\deg$)$. Therefore, the source was referred to as GRO J1411-64. The flare flux (where following the report by Zhang et al. the detection significance of 7.2$\\sigma$ was measured, see the discussion below) was about 0.3 Crab, which would make of GRO J1411-64 one of the strongest sources at low galactic latitudes, second only to Crab itself. The {\\it flare duration} was several months, the longest one compared to all other $\\gamma$-ray sources that were monitored almost continually by COMPTEL. The rather steep spectral shape obtained while the source was flaring would predict a bright, hard X-ray source, if there is no break in the spectrum, which is explored here.  \\begin{table*}[tbh] \\begin{center} \\caption{The GRO~J1411-64 observations: effective exposures, fluxes including detection significance ($\\chi_1^2$) or 2$\\sigma$ flux upper limits for the observations reported by Zhang et al.} \\begin{tabular}{ccccccc}\\hline \\multicolumn{1}{c}{TJD}&\\multicolumn{1}{c}{Effective exposure }&\\multicolumn{4}{c}{Flux (10$^{-5}$ ph cm$^{-2}$ s$^{-1}$)}\\\\ \\multicolumn{1}{l}{ }&\\multicolumn{1}{c}{days}&\\multicolumn{1}{c}{0.75-1 MeV}&\\multicolumn{1}{c}{1-3  MeV}&\\multicolumn{1}{c}{3-10 MeV}&\\multicolumn{1}{c}{10-30 MeV}\\\\  \\hline 8392-8943 (Phase 1)     &9.29   &6.0$\\pm$3.1 (2.3$\\sigma$)  &$<$6.0   &1.7$\\pm$1.6 (1.1$\\sigma$)  &0.8$\\pm$0.6 (1.1$\\sigma$)\\\\ 8943-9216 (Phase 2)     &5.87   &4.2$\\pm$3.5 (1.3$\\sigma$)  &$<$10.1   &$<$3.9   &1.8$\\pm$0.9 (2.0$\\sigma$)\\\\ 9216-9629 (Phase 3)     &7.45   & $<$12.5   &6.7$\\pm$3.8 (1.8$\\sigma$)  &2.2$\\pm$1.6 (1.9$\\sigma$)  &$<$1.6 \\\\ 9629-9993 (Phase 4)     &8.42   &13.7$\\pm$3.8 (4.0$\\sigma$)   &21.2$\\pm$3.5 (5.9$\\sigma$)  &4.6$\\pm$1.6 (2.9$\\sigma$)  &$<$1.2 \\\\ 9993-10371(Phase 5)     &4.23   &18.6$\\pm$4.7 (4.4$\\sigma$)  &$<$14.7   &$<$4.8   &$<$2.6 \\\\ 9797-9923 (VPs 414-424) &5.44   &19.9$\\pm$4.6 (4.6$\\sigma$)  &36.6$\\pm$4.7 (7.7$\\sigma$)  &5.4$\\pm$2.3 (2.1$\\sigma$)  &$<$1.8 (2.0$\\sigma$)\\\\ 8392-10371(Phases 1-5) &35.26  &8.4$\\pm$1.7 (5.1$\\sigma$)  &6.1$\\pm$1.6 (4.0$\\sigma$) &1.6$\\pm$0.8 (2.4$\\sigma$)  &0.5$\\pm$0.3 (2.0$\\sigma$)\\\\  \\hline \\end{tabular}\\end{center} \\label{tab:flux} \\end{table*}  That the source was variable within timescales of several months could be shown by two different methods by Zhang et al. (2002). On one hand, a direct statistical approach shows that if the fluxes of the individual viewing periods  of the first five Phases are fitted with a constant flux, a $\\chi^{2}$ value of 94.6 is derived (1-3~MeV band). That value corresponds to a significance of 4.1$\\sigma$ for a variable source; alternatively, to a probability of $\\sim$4.8$\\times$10$^{-5}$ for a constant flux. On the other hand, the variability index $I$ (Torres et al. 2001), used to normalize the measured variability of this source to the measured variability of other COMPTEL sources considered to be non-variable (so as to encompass systematics), gives a value of 5.2. This is above the 3$\\sigma$-level and also classifies the source as likely variable. However, it is unclear whether this second method would apply to COMPTEL sources, due to the larger spread of indices for detected COMPTEL pulsars. In what follows, we present the results of the INTEGRAL observations of this source, as well as of XMM-Newton observation of its best location. We also  comment on what can be inferred out of these observations concerning models of emission based on microquasars. ",
        "conclusions": "\\begin{figure}[hb] \\centering \\includegraphics[clip=true,scale=.4]{hecomptel.eps} \\caption{ A microquasar model in the light of observational constraints. See text for details. } \\label{hecomptel} \\end{figure}   We have applied to the non-simultaneous data from XMM-Newton (treating them as an upper limit\\footnote{This is actually a not well-justified a priori assumption, since the coverage of XMM-Newton observations is much smaller than that of INTEGRAL, and it is imposed by fiat, in the knowledge that the constraints imposed below do not crucially depend on the XMM-Newton point, but rather on the steepness of the INTEGRAL upper limits.}), INTEGRAL and COMPTEL the microquasar model presented by Bosch-Ramon et al. (2006). This model consists of a magnetized conical jet filled with relativistic electrons which radiate through synchrotron and inverse Compton scattering with star, disk, corona and synchrotron photons. These electrons lose energy mainly through adiabatical losses, as a first order approximation. At the top of Table~\\ref{common}, parameter values for a typical microquasar system and jet geometry are given (Bosch-Ramon et al. 2006). We have considered different values within the range open for the free parameters finding that it is not possible to obtain a simple microquasar model that could fit the SED. In particular, in Figure~\\ref{hecomptel} we show a test case with the free parameters fixed to the values presented in Table~\\ref{common}, at the bottom. { We note that, since the computed SED in Fig. 10 is dominated in the gamma-ray band by SSC emission, the model results would also apply for a low mass microquasar.}  \\begin{table*}[] \\caption[]{Parameter values for GRO~1411$-$64} \\label{common} \\begin{center} \\begin{tabular}{cl} \\hline\\noalign{\\smallskip} Parameter &  values \\\\ \\hline\\noalign{\\smallskip} Stellar bolometric luminosity [erg~s$^{-1}$] & $10^{38}$ \\\\ Distance from the apex of the jet to the compact object [cm] & $5\\times10^7$ \\\\ Initial jet radius [cm] & $5\\times10^6$ \\\\ Orbital radius [cm] & $3\\times10^{12}$  \\\\ Viewing angle to the axis of the jet [$^{\\circ}$] & $45$ \\\\ Jet Lorentz factor & 1.2 \\\\ \\end{tabular} \\begin{tabular}{cl} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} Jet leptonic kinetic luminosity [erg~s$^{-1}$] & $3\\times10^{35}$ \\\\ Maximum electron Lorentz factor (jet frame) & 5$\\times10^2$ \\\\ Maximum magnetic field [G] & 8000 \\\\ Electron power-law index & 1.5 \\\\ Total corona luminosity [erg~s$^{-1}$] & $3\\times10^{33}$ \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{center} \\end{table*}   If the COMPTEL source is to be among the group of low luminosity soft X-ray sources and be detectable at energies about 1~MeV, a very hard photon index of about 0.67 (i.e. $n_{\\nu}\\propto \\nu^{-0.67}$) will be required in the  energy range 10--1000~keV assuming a power-law energy distribution for the radiation. Even for a strongly absorbed soft X-ray counterpart (see, e.g. Walter et al. 2003), or a soft X-ray counterpart out of the XMM-Newton field observed, the INTEGRAL upper-limits impose also a very hard photon index of about 1 in the range 100-1000~keV, whereas radiation at higher energies is strongly constrained by EGRET observations. In the context of microquasars, if the particle spectrum follows a power-law, it is difficult to explain such narrow band emission concentrated in the MeV range.  { Concerning other transients spatially coincident with the INTEGRAL observations we mention 2S 1417-624, a Be/X-ray binary pulsar system, with a Be star of $>5.9$ M$_\\odot$ for a neutron star of 1.4 M$_\\odot$. BATSE observed a huge outburst of this source at hard X-rays (pulsations up to 100 keV were detected), which started on August 26, 1994 and lasted 110 days (Finger et al., 1996). COMPTEL observed the source twice (VPs 402, 402.5) during this period but did not detect it. During the next 200 days, 2S 1417-624 showed 5 smaller outbursts (Finger et al., 1996), which coincide in time with the significant COMPTEL detection of the MeV source in 1995. During the huge outburst in hard X-rays, the pulsed luminosity was found to be much less than it would be estimated from the spin-up rate. This implies that most of the power output is either unpulsed or outside of the hard X-ray range (Finger et al., 1996). Assuming 2S 1417-624 as counterpart of the COMPTEL source, its 0.75-10 MeV luminosity would be $2.5 \\times 10^{37}$ erg s$^{-1}$, obtained by using its upper-limit distance of 11.1 kpc. This luminosity is comparable to the pulsed one of $2.2 \\times 10^{37}$ erg s$^{-1}$. 2S 1417$-$624 as counterpart might show an anticorrelation of the X- and $\\gamma$-ray emission as proposed for other sources (Romero et al. 2001). GS 1354-645, also in the field, is a transient black-hole X-ray binary (XRB) system. However, no X-ray observations of GS 1354-624 during the activity period reported for the COMPTEL source are reported. Since INTEGRAL did not detect these sources during our campaign, we cannot open judgment about the possibility for an outbursting system such as these to be the counterpart of the reported detection. }    Failing to associate a plausible counterpart of GRO~J1411-64 at the hard X-rays from our INTEGRAL observations, we have looked again at the initially reported gamma-ray properties of GRO~J1411-64. In Table 1, we have already expanded beyond the initially reported source characteristics from Zhang et al. (2002) by including the individual detection significances. If we apply a rather conservative detection criterion (significance threshold) of at least 4 $\\sigma$ in an individual measurement, a firm detection can be claimed only in the so-called flare state, which was attributed to occur within CGRO viewing periods 414-424. During this flare period as well as any other investigated time span, a detection threshold of at least 4$\\sigma$ in the individual measurement allows the determination of the photon spectrum from significantly measured data points, i.e., in two energy bands only, from 0.75-1 MeV and 1-3 MeV, respectively. An accordingly determined spectrum yields an index of 2.18 $\\pm$ 0.2, substantially harder than the previous reported index of 2.53 $\\pm$ 0.2 taking into account lower confidence data points. If we fit the range from 0.75 to 10 MeV with a statistically weighted power law fit, an index of 2.48 $\\pm$ 0.2 will be determined. Therefore, the dependence from the individual measurements becomes crucial for determining a spectral slope.  If we compare the observations of GRO~J1411-64, as made in Figure \\ref{XXX}, indicating the measured flux/upper limit in the four energy bands as a function from the respective exposure, other characteristics can be seen. Here, we convert the 2$\\sigma$ upper limits from Zhang et al. (2002) into 1$\\sigma$ upper limits, in order to prevent quoting upper limits based on higher statistical confidence than certain flux values from Table 1. One clearly sees that the flaring phenomenon related to the CGRO observation periods 414-424, is not based on any extraordinary long or short exposure when compared with other observations. It can also be seen that the best signals originated from the two lower energetic band in COMPTEL, as already indicated from the individual detection significances. Finally, whereas there is a clear indication of the spectral behavior being a steep spectrum source in the phase 4 observations, there is a remarkable difference in the phase 5 observation: Here, the largest flux has been quoted from the 0.75 to 1 MeV band. However, in the immediate neighboring energy band (1 -3 MeV), only an upper limit has been obtained. Comparing for the so-called flare phase, where the flux in the 0.75-1 MeV band is clearly exceeded by the 1-3 MeV band data, we should perhaps reconsider the flare phenomenon: On one hand, the actual low energy band data point in phase 5 is measured at higher flux and better statistical significance, on the other hand the 1 -3 MeV flux measurement is missing, indicating that the spectrum was even more extreme in the phase 5, and clearly different from phase 4, and any other considered time span quoted in Zhang et al. (2002). This is also seen in the total superposition, where the four energy bands line up in an more usual behavior, i.e., any subsequent higher energetic band exhibit a lower flux. In the light of such reassessment we note that a flare phenomenon is not consistently to be described as either spectral hardening or softening compared to the sparse overall data from the region. We note, however, that the indication for this phenomenon is primarily based on data with the lowest effective exposure. Thus, the distinction of different state becomes blurry.  \\begin{figure}[hb] \\centering \\includegraphics[clip=true,scale=1.]{exp_vs_flux.eps} \\caption{The original observations of GRO~J1411-64 as reported in Zhang et al. (2002), see Table 1. The dependence from the individual measurements from the underlaying exposure is easily seen, as is the different character in the phase 4 vs phase 5 measurements, which exhibit the largest flux values, but missing a 1-3 MeV flux in case of the phase 5 measurement.} \\label{XXX} \\label{exposurevsflux} \\end{figure}     GLAST observations would help improving the location of the MeV source if radiation at higher energies is not completely suppressed, and would open the door for more efficient multiwavelength searches of the counterpart. However, it is true that the nature of this COMPTEL source might not be constrained further if this detection was a one-time only transient phenomena. GLAST will only be able to help if a candidate counterpart is caught in the act (flaring/quiescent state of an AGN or a more rare galactic object). Having at hand GLAST observations, in any case, will make our currently reported investigation to naturally fit into the testing of any hypothesis on the nature of GRO J1614-64."
    },
    "0606/astro-ph0606193_arXiv.txt": {
        "abstract": "The U.S.~Naval Observatory Robotic Astrometric Telescope (URAT) project aims at a highly accurate (5 mas), ground-based, all-sky survey. Requirements are presented for the optics and telescope for this 0.85 m aperture, 4.5 degree diameter field-of-view, specialized instrument, which are close to the capability of the industry. The history of the design process is presented as well as astrometric performance evaluations of the toleranced, optical design, with expected wavefront errors included. ",
        "introduction": "From 1998 to 2004 the U.S.~Naval Observatory (USNO) conducted an astrometric, all-sky survey to 16th magnitude with its 20 cm aperture Twin Astrograph, resulting in the USNO CCD Astrograph Catalog (UCAC), with its second release (UCAC2) made public in 2003 \\cite{ucac2}. Around the year 2000 the USNO began preparations for a follow-up project, the USNO Robotic Astrometric Telescope (URAT), a dedicated, astrometric, 1-meter size telescope with a wide field of view for an all-sky survey going much deeper than UCAC and on a positional accuracy level aiming at 5 milliarcsecond (mas) standard error per coordinate.  The idea of such a telescope goes back to the late 1980s \\cite{cdv89}, then envisioned for photographic plates. The purpose of such a telescope is threefold. First URAT aims at the densification of the celestial reference frame, providing a large number of accurate reference stars \\cite{a-surveys} to support, for example, general, deep sky surveys like PanSTARRs and LSST, and to support astrometry in our solar system, as outlined in the recent decadal survey \\cite{decadal}. Second, URAT provides absolute proper motions on an inertial system, enabling significant galactic dynamics studies in the pre-Gaia era. Third, URAT will be able to observe trigonometric parallaxes of many thousands of stars unbiased by selection effects like high proper motion targets.  More details about the URAT project in general can be found elsewhere \\cite{potsdam,lowell1}. ",
        "conclusions": ""
    },
    "0606/astro-ph0606646_arXiv.txt": {
        "abstract": "{} {We study the synchrotron flaring behaviour of the blazar 3C~279 based on an extensive dataset covering 10 years of monitoring at 19 different frequencies in the radio-to-optical range.} {The properties of a typical outburst are derived from the observations by decomposing the 19 lightcurves into a series of self-similar events. This analysis is achieved by fitting all data simultaneously to a succession of outbursts defined according to the shock-in-jet model of Marscher \\& Gear (1985).} {We compare the derived properties of the synchrotron outbursts in 3C~279 to those obtained with a similar method for the quasar 3C~273. It is argued that differences in the flaring behaviour of these two sources are intrinsic to the sources themselves rather than being due to orientation effects. We also compare the start times and flux densities of our modelled outbursts with those measured from radio components identified in Very Long Baseline Interferometry (VLBI) images. We find VLBI counterparts for most of our model outbursts, although some high-frequency peaking events are not seen in the radio maps. Finally, we study the link between the appearance of a new synchrotron component and the EGRET gamma-ray state of the source at 10 different epochs. We find that an early-stage shock component is always present during high gamma-ray states, while in low gamma-ray states the time since the onset of the last synchrotron outburst is significantly longer. This statistically significant correlation supports the idea that gamma-ray flares are associated with the early stages of shock components propagating in the jet. We note, however, that the shock wave is already beyond the broad line region during the gamma-ray flaring.} {} ",
        "introduction": "It is generally accepted that radio outbursts in active galactic nuclei (AGN) are triggered by growing shocks in a relativistic jet.  There have been many efforts to identify the individual components in radio to submillimeter (submm) lightcurves. Litchfield et al. (1995) and Stevens et al. (1995, 1996, 1998) managed to follow the early evolution of individual outbursts by subtracting the quiescent contribution (assumed to be constant) from the total spectrum. Valtaoja et al. (1999) identified outbursts by decomposing the variations in millimeter (mm) and centimeter lightcurves into exponential flares. T\\\"urler et al. (1999) measured the properties of synchrotron outbursts by decomposing the radio-to-submillimeter wavelength lightcurves into a series of self-similar flaring events.  The method was further developed by T\\\"urler et al.  (2000) who used the shock model of Marscher \\& Gear (1985) to describe both the average evolution of the outbursts and their individual characteristics. This model was found to provide a good description of the lightcurves of 3C~273. In this work, we decompose the lightcurves of a second source, 3C~279. Since our adopted methodology is analogous to that used by T\\\"urler et al. (2000) the results are directly comparable to those found for 3C~273.   While in previous studies the outbursts were identified in the radio-to-submillimeter regime, we here follow the synchrotron spectrum up to infrared and optical frequencies. Searches for correlations in the radio-to-optical emission from AGN have been conducted since the 1970s. One of the most recent works was presented in Hanski et al. (2002). In their sample of 20 AGNs they found a clear radio-to-optical correlation in seven cases, a possible correlation in six cases and no correlation for the remainder. This agrees well with the findings of previous studies: there appears to be some kind of correlation between the radio and optical emission but this correlation is not seen in all sources at all epochs. The general trend seems to be that all radio outbursts are accompanied by optical outbursts but conversely not all optical outbursts have radio counterparts.  The source 3C~279 is one of the brightest and most variable blazars at radio wavelengths. As such, it is monitored regularly in all radio wavebands from the centimeter to the submillimeter domain. In the optical, the historical lightcurve shows variations with a typical amplitude of about 2 magnitudes, but reaching 8 magnitudes during flares (Webb et al. 1990). The weak blue bump of 3C~279 allows us to follow the synchrotron spectrum up to infrared and optical wavelengths. Indeed, 3C~279 is one of the few objects for which a clear correlation is found between the radio and optical wavebands (Tornikoski et al. 1994).  In this paper we present results of a multifrequency lightcurve decomposition of 3C~279. The method used is based on T\\\"urler et al. (1999; 2000) with some modifications for the specific case of 3C~279, including an extension of the analysis to infrared and optical wavelengths. The average evolution of outbursts is compared to that of 3C~273. We also compare our results with (1) VLBI maps from Wehrle et al. (2001) and Jorstad et al. (2004) to study the connection between the model outbursts and VLBI components, and (2) gamma-ray data obtained by Hartman et al. (2001) to study the connection between the synchrotron spectra and the gamma-ray state.  \\begin{figure*} \\centering \\includegraphics[angle=0,width=17cm]{visu10fit.ps} \\caption{Ten out of nineteen radio-to-optical lightcurves of 3C~279. The points indicate the observed flux density and the solid line our best fit which is a sum of the underlying jet (long dashed line) and the fifteen outbursts (dotted lines). The first dotted line represents the global decay of all outbursts peaking before 1989.0} \\label{FigFinal}% \\end{figure*} ",
        "conclusions": "We have studied synchrotron flaring in the jet of 3C~279 by decomposing the multifrequency lightcurves into a series of outbursts. The analysis is performed within the generally accepted shock-in-jet scenario, with the evolution of the model outbursts following the Marscher \\& Gear (1985) prescription. Similar methods have been used previously to study the quasar 3C~273 (T\\\"urler et al. 1999, 2000) and the micro-quasar GRS~1915+105. For 3C~279 we find:  1. The lightcurve decomposition supports the Marscher \\& Gear model. The smooth low-frequency lightcurves are especially well reproduced by the model. In the high-frequency (millimeter to optical) domain, the sharp peaks of the observed lightcurves are less well reproduced by the model lightcurves during the strongest outbursts. Although the model used here makes a lot of assumptions and simplifications which might not hold in detail, it is still clear that there is some mechanism at work which is not well described by the Marscher \\& Gear (1985) model, and which produces the sharp peaks of outbursts at high frequencies. This is in accordance with previous findings (e.g. Valtaoja et al. 1999): both the theoretical and numerical jet models always produce smooth peaks in model lightcurves while the observed outbursts have exponential rises, sharp peaks and exponential decays.  2. The outburst evolution in 3C~279 follows the characteristic three stage pattern but instead of the synchrotron plateau stage seen in 3C~273, the flux density is already decreasing during the second stage. This might be explained by a highly non-adiabatic jet flow, where radiative losses play a significant role. We also find that the lightcurve of an average outburst is different for 3C~273 and 3C~279. The differences found between these two sources appear to be mainly intrinsic and not due to different orientations. The borderline blazar 3C~273 is thus not simply 3C~279 seen from a larger viewing angle.  3. We find a rather good agreement between the zero-separation epoch of VLBI components and the start times of model outbursts. However, some high-frequency peaking outbursts do not have a corresponding VLBI component which can be explained by the long delay before their spectral turnover reaches VLBI frequencies.  4. When comparing the gamma-ray state at different epochs from Hartman et al. (2001) to the time elapsed since the start of the last synchrotron outburst we find that the shortest time intervals (corresponding to the shortest distances from the apex of the jet to the shock) are seen during the highest gamma-ray states, and the longest ones during the lowest gamma-ray states. This supports previous findings (e.g. Valtaoja \\& Ter\\\"asranta 1996; Jorstad et al. 2001; L\\\"ahteenm\\\"aki \\& Valtaoja 2003) that the gamma-ray flares are connected to shocks propagating in the jet. Although during the highest states the shock components are rather new, they are still well outside the outer limits of the broad-line region. This suggests that external Compton models based on the soft-photon field originating from the BLR are unlikely to be able to explain the gamma-ray flares in this source."
    },
    "0606/astro-ph0606120_arXiv.txt": {
        "abstract": "{}% {We study the stability/establishment of dipolar magnetostatic equilibrium configurations in new--born neutron stars (NSs) in dependence on the rotational velocity $\\Omega$ and on the initial angle $\\alpha$ between rotation and magnetic axis.}% {The NS is modeled as a sphere of a highly magnetized ($B \\sim 10^{15}$G) incompressible fluid of uniform density which rotates rigidly.  For the initial dipolar background magnetic field, which defines the magnetic axis, two different configurations are assumed. We solve the 3D non--linear MHD equations by use of a spectral code. The problem in dimensionless form is completely defined by the initial field strength (for a fixed field geometry), the magnetic Prandtl number $\\Pm$, and the normalized rotation rate. The evolution of the magnetic and velocity fields is  considered for initial magnetic field strengths characterized by the ratio of ohmic diffusion and initial \\Alf{} travel times  $\\ttOhm/\\ttAO \\approx 1000$, for $\\Pm = 0.1, 1, 10$, and  the ratio of rotation period and initial \\Alf{} travel time, $P/\\ttAO = 0.012, 0.12, 1.2, 12$.}% {We find hints for the existence of a unique stable dipolar magnetostatic configuration for any specific $\\alpha$, independent of the initial field geometry. Comparing NSs possessing the same field structure at the end of their proto--NS phase, it turns out that sufficiently fast rotating NSs ($P\\la6\\,$ms) with $\\alpha \\la 45^0$ retain their magnetar field, while the others lose almost all of their initial magnetic energy by transferring it into magnetic and kinetic energy of relatively small--scaled fields and continue their life as radio pulsars with a dipolar surface field of $10^{12\\dots13}$G.}{} ",
        "introduction": "The strength of a neutron star's (NS's) dipolar  magnetic field, its inclination angle with respect to the NS's rotation axis, $\\alpha$, and the rotational velocity $\\Omega$ are the quantities determining the observability of pulsars as well as the power of their magneto--dipole radiation and/or the pulsar wind, which eventually set the spin--down behavior of an isolated pulsar \\citep{ST83}. Therefore, the determination of these (more or less exactly observable) quantities in their temporal evolution is subject of scientific debates since the discovery of pulsars.  From the observational point of view there exists only little information about the magnetic inclination $\\alpha$ immediately after birth of a NS. This is owed to the poor observability of new--born NSs, which are surrounded by opaque material. Recently, \\citet{HAGC03} and \\citet{H06} investigated the emission geometry of very young and energetic pulsars and found that five out of six have $\\alpha \\lesssim 30^\\circ$. The 1600 years old pulsar \\object{J1119-6127}, e.g., has an inclination angle $\\alpha = 19^\\circ$ and a dipolar magnetic field strength of $4.2\\cdot 10^{13}$G. Age and dipolar magnetic field strength correspond to the observed spin period of $0.42$ s, provided the field has been constant during the lifetime of the NS. Given that constancy an initial spin period of less than a millisecond can be estimated \\citep[see, e.g.,][]{CGP00}.  Unfortunately, there is even less observational evidence about $\\alpha$ in magnetars because their slow rotation impedes the detection of open field line regions. The typically wide X--ray pulses come likely from a very wide beam. If the X--ray emission region is not too far above the surface and the pulse peak corresponds to the magnetic pole, then that wide beam may indicate that the inclination angle $\\alpha$ is small \\citep{Z06}.  In order to understand those, yet vague, observational evidences of relatively small values of $\\alpha$ for magnetars and young energetic pulsars, and to get at all an idea about the ``initial'' magnetic field structure of NSs, the study of the establishment/stability of magnetostatic equilibrium configurations immediately after a NS's birth is necessary. In this very early phase, when the core is not yet superfluid, consists of npe--matter, and when the crust is completely liquid, an equilibrium configuration may indeed be reached. It defines the initial magnetic field structure and strength, especially that of the dipolar component, which plays a decisive role for spin--down behavior and observability of the NS.  In new--born NSs, the rotation velocity exceeds the hydromagnetic (or \\Alf{}) one, $v_\\tAlfven$, in large regions, even for magnetic field strengths in the order of magnetar fields ($B \\lesssim 10^{15}$G). Therefore, the effect which rotation may have on the establishment of a magnetostatic equilibrium in young NSs has to be studied in detail.  It was \\citet{W73}, who postulated the stability of a magnetic field configuration the poloidal and toroidal constituents of which are of comparable strength. This conclusion was supported by subsequent studies, e.g., \\citet{MT73} or \\citet{BN05}. Tayler and co--authors performed a series of analytical stability analyses of magnetic fields in stars: \\citet{T73} considered the stability of a purely toroidal axisymmetric field without rotation and found that it is always unstable with respect to azimuthally large--scale ($m=1$) perturbations. The growth rate of this instability  is on the order of the \\Alf{} frequency $\\Omega_A= v_A/r$, where $v_A = B/\\sqrt{4\\pi\\rho}$, $B$ the magnetic field strength,  $\\rho$ the density and $r$ the radial coordinate. This ``Tayler--instability'', meanwhile basic ingredient of the ``Spruit--dynamo'' \\citep[]{S99, S02}, acts locally in meridional ($r,\\theta$) planes but globally in the azimuthal ($\\varphi$) direction. \\citet{MT73} and \\citet{W73} performed similar adiabatic stability analyses for an axisymmetric poloidal field and concluded its general instability if at least some of its field lines are closed within the star. The instability generates small scale field structures on the same time scale as for a toroidal field which dissipate rapidly. \\cite{MT73}  discuss the possibility that sufficiently rapid rotation and/or the simultaneous existence of equipollent  toroidal and poloidal components may stabilize the axisymmetric background fields.  \\citet{PT85} explored in ideal MHD and cylindrical geometry the effect of rigid rotation on the stability of a field consisting of a toroidal component which increases with the distance from the magnetic axis, and/or a uniform $z$ (i.e., poloidal) component for different analytically feasible special cases: For $\\alpha=0$ both the purely toroidal field  and the combined toroidal--poloidal one show qualitatively the same stabilizing effect of the rotation. For $\\alpha=90^\\circ$, however, a purely toroidal field reveals a much less stabilizing effect of the rotation. They studied further the stabilizing effects of a stable stratification of the star's density profile in the case $\\alpha = 0^\\circ$ and for a purely toroidal field which may increase non--linearly with axial distance. As already shown by \\citet{T57}, such a field is stable even in the absence of rotation for sufficient small wavelengths of a perturbation along the magnetic axis. Larger scaled perturbations, however, cause instabilities which can be counteracted by sufficiently rapid rotation. On the other hand, \\citet{PT85} found that at the same rotation rate, which stabilizes the original instabilities, rotation introduces a new instability. Therefore, they concluded, that ``rotation is unlikely to lead to complete stability of general magnetic field configurations''.  In 2005 the interest in the stability of magnetostatic equilibria in stars became revitalized by a series of papers: \\citet{BS05a,B05,BN05}. The latter authors  studied the stability of magnetic fields in the radiative interiors of Ap stars and found, that  any random field is generally unstable but evolves in a stably stratified star towards a ``twisted torus'' configuration with approximately equipollent toroidal and poloidal components. \\citet{BS05a} considered the stability of MHD configurations in ultra--magnetized NSs, the magnetars. Although they found that, quite similar to the Ap stars, stable magnetic fields exist when being concentrated to a relatively small region around the center of the star and after having evolved into a poloidal--toroidal torus shape, these fields cannot have magnetar strength unless their initial strength exceeds 10$^{18}$ G. This estimate relies on two field--diminishing effects in the simulations of \\citet{BN05}: Firstly, a decrease of typically $\\lesssim 10$ appears during the relaxation into the torus--like structure; secondly, these fields are strongly concentrated in the central region and decay by a factor of up to 1000 towards the surface. However, NSs born with magnetic fields of $\\sim10^{18}$ G are unlikely. On the other hand an internal field of (only) magnetar strength would need a much too long time to diffuse to the surface.  Until now the effect of rotation on the stability of purely poloidal field configurations has not been considered in general. The results of \\citet{PT85}, \\citet{L88} and \\citet{CP95} do not apply directly since these works consider combined toroidal and poloidal fields, the latter being homogeneous. In the sole case in which \\citet{CP95} consider an inhomogeneous poloidal field, the Tayler instability was not found. The work of Braithwaite and Spruit doesn't include rotation except in \\citet{B05} who considers toroidal fields only.  Since rotation in new--born NSs can for sure not be neglected and the dipolar magnetic field determines both observability and rotational evolution, we will study the consequences of rotation for the stability of such fields. In the next section we describe and justify our model assumptions, present the set of equations to be solved together with the initial conditions for the models under consideration and sketch our numerical approach. After presenting the results for the different MHD configurations in \\sectref{results}, we finally discuss the results. ",
        "conclusions": " Never again in an isolated NS's life its magnetic field may suffer  so drastic changes as during the first at most $1000$ seconds after birth, provided the dipolar field present immediately after the proto--NS phase has magnetar strength. Depending on the rotational velocity and the inclination of the field with respect to the rotation axis, the inborn field will either survive that period as a magnetar field or lose almost all of its initial magnetic energy by transferring it into magnetic and kinetic energy of quite small--scaled fields  eventually being consumed by ohmic and viscous dissipation. After the end of the unstable phase, large--scaled magnetic fields dominate again, but have a much lower strength.  We found that purely dipolar fields in highly magnetized new--born NSs are unstable on a time--scale of a  few \\Alf{} times unless fast rotation, i.e., $\\Omega_\\tAlfven/\\Omega \\la 0.1$, acts stabilizing. The latter condition, however, is only a necessary one, fulfilled for NSs which have at the end of their proto--NS phase a poloidal field strength of $\\sim 10^{15}$G  and a rotation period of $1\\dots 6$ ms.  As a sufficient condition, the above presented results yield that rotation is only efficiently stabilizing as long as the inclination angle between rotation and magnetic axis $\\alpha <  45^\\circ$. Therefore, we may speculate that those NSs whose inborn field is not too much bent away from the rotation axis and which rotate with a initial period $\\la 6$ms will continue their life as magnetars, while for quite similar NSs, which, however, are born as orthogonal rotators, even rotation periods  $\\approx 6$ms  will not stabilize the poloidal field against its substantial reduction down to standard pulsar strengths. As another conclusion we suggest that for each inclination angle $\\alpha$ there exists a ``final'', apart from ohmic decay, stable field configuration, whatever the inborn field shape has been. Indeed, they seem to be identical up to a rotation by $\\alpha$.  \\citet{FR77} consider qualitatively the relaxation of a uniform field into its lowest energetic state. Assuming  a relatively large \\Alf{} time scale of $\\sim 100$ s (in 1977 magnetar field strengths were  inconceivable!) and a rotation period of the new--born NS $\\approx\\!6$  ms,  they conclude that the time for the readjustment of that field is about 10 days, clearly beyond the onset of crystallization which impedes the relaxation already after a few 1000 seconds. Their relaxation time scale is (except for a factor of 2) that of the Tayler instability, $\\ttA$,  enhanced by a factor $\\ttA/P$ due to the strong effect of the Coriolis force \\citep{S99}. Making use of this argumentation with respect to magnetar field strengths, however, this estimate for the relaxation time scale reduces to $\\sim 1$ s, i.e.,  the initially uniform magnetic field would have ample time to establish its energetically lowest state which is supposed to be no longer dominated by a large scale field. In contrast,  our study shows that this will not happen for a rotation period of $\\lesssim 6$ ms unless $\\alpha\\ga 45 ^\\circ$. For an understanding of this discrepancy the behavior of the so--called Levitron\\textsuperscript{TM} can be instructive \\cite[cf. e.g.][] {JWG97}: Similar to the setup of \\citet{FR77} it consists of two permanent, mutually repelling magnets where a stable (levitation)  situation can be achieved only via rotation. However, it seems questionable at all whether the analogy with permanent magnets  is appropriate for NSs because the latter can be described in terms of ideal MHD characterized by a frozen--in field whereas the former are characterized by frozen--in currents.  In this paper we consider the stability of purely poloidal fields. However, it is very unlikely, that in nature such a configuration exists. More likely is that the poloidal field is accompanied by a toroidal one, which, if it is of comparable strength, may act stabilizing as well \\citep{T80,BS04,BN05}. The dependence of the stability of such a configuration  on rotation and inclination will be studied in a forthcoming paper.  It is well--known, that the effect of a NS's stably stratified density profile exerts another stabilizing effect \\citep[see, e.g.,] []{PT85}. This is due to the strong suppression of motions in the direction of gravity being an additional constraint for the development of an instability. Hence, the neglect of this effect in the present study justifies the more our conclusions about the stabilizing action of rapid rotation. This could, however, imply that the marginal $P$, which discriminates between stability and instability, shifts to somewhat larger values. In principle, compressibility, which is the cause for stratification, affects any MHD process and may give rise to specific instabilities, e.g. the buoyancy (Parker) one, not existent in an incompressible setup. Thus, deviation from the assumption of uniform density may add both a stabilizing and a destabilizing tendency. Although we are, guided by our estimate of the ratio of sound speed to fluid velocity, quite sure that the stabilizing effect prevails, a quantitative certainty can only be gained on the basis of a numerical study with an (at least) anelastic model.  Apart from the magnetic field the vigorous flows occurring in the unstable situations may result in observable consequences: When extrapolating the maximum velocities found in the simulations to lower, more realistic viscosities, it seems likely that in some regions of the NS velocities of $\\ga 10^8$ cms$^{-1}$ can be reached. Supposed that the initial perturbation for the instability considered in this paper  is given by the (decaying) PNS convection, it is conceivable that the geometry of the latter is somehow ``transferred'' to the one of the Lorentz--force induced flow. Moreover, as the rotation rates of the PNS and the newly born NS are almost identical,  it is allowed to speculate that similar to the PNS convection the demonstrated unstable decay of a strong magnetic field may be a source of gravitational waves.  Applying likewise the estimates of \\cite{MJ97}, we conclude that for a duration from seconds to a few minutes a detectable gravitational wave signal can be expected rising at about 20\\ldots 30 s after bounce.   Our study is based on the assumption, that newly born NSs inherit from their proto--NS phase surface magnetic fields  $\\gtrsim10^{15}$G. This state can be reached either by flux conservation if the progenitor was already highly magnetized, by a PNS dynamo, or by other, still unknown, processes.  If the NSs are born with a strong dipolar field then \\begin{enumerate} \\item those whose rotation period is less than $\\sim 6$ms and whose magnetic inclination angle $\\alpha$ is smaller than $\\sim45^\\circ$ will retain their extremely large surface field strength and appear, after a rapid spin down, as magnetars; \\item those which rotate less rapidly, say with $P \\ga 6$ms and/or for which is $\\alpha \\ga 45^\\circ$ will lose almost all of their inborn magnetic energy and appear as radio pulsars. \\end{enumerate} It turns out that rotation is likely to be the only stabilizing agent which allows of the existence of magnetars whereas the stable configurations found by \\citet{BS05a} are less suited to support such strong surface fields.  In the light  of the present study, the much smaller number of observed magnetars in comparison with that of all the other observed realizations of NSs seems to be a consequence of the fact that only a small fraction of all new--born NSs are rotating as fast as or faster than $P\\sim6$ms. Of course, this conclusion is still premature, lacks broad observational support, does not consider other selection effects, and has to be confirmed by the study of both more realistic NS models and a denser mesh of initial $\\alpha$ and $P$ values as well as by studies of the evolutionary stages `supernova progenitor' -- `supernova' -- `proto--NS'. These considerations should definitely include the effects of magnetic fields and return a better idea about the distribution of magnetic field strengths, rotation periods and inclination angles at the beginning of a fully evolved NS's life.  Let us now relax the assumption that neutron stars have at the end of the PNS phase a field of $\\ga 10^{15}$ G and adopt instead fields at that moment which are two to three orders of magnitude smaller. When additionally conceding that the initial rotation period could be as large as, say, 60 ms, the ratio $q_P$ would be still at most $\\sim 10^{-2}$.  As our results demonstrate, such fields would be always stable irrespective of the magnetic inclination $\\alpha$."
    },
    "0606/astro-ph0606316_arXiv.txt": {
        "abstract": "Recently, a new non-Standard Model neutrino interaction mediated by a light scalar field was proposed, which renders the big-bang relic neutrinos of the cosmic neutrino background a natural dark energy candidate, the so-called Neutrino Dark Energy. As a further consequence of this interaction, the neutrino masses become functions of the neutrino energy densities and are thus promoted to dynamical, time/redshift dependent quantities. Such a possible neutrino mass variation introduces a redshift dependence into the resonance energies associated with the annihilation of extremely high-energy cosmic neutrinos on relic anti-neutrinos and vice versa into Z-bosons. In general, this annihilation process is expected to lead to sizeable absorption dips in the spectra to be observed on earth by neutrino observatories operating in the relevant energy region above $10^{13}$~GeV. In our analysis, we contrast the characteristic absorption features produced by constant and varying neutrino masses, including all thermal background effects caused by the relic neutrino motion. We firstly consider neutrinos from astrophysical sources and secondly neutrinos originating from the decomposition of topological defects using the appropriate fragmentation functions. On the one hand, independent of the nature of neutrino masses, our results illustrate the discovery potential for the cosmic neutrino background by means of relic neutrino absorption spectroscopy. On the other hand, they allow to estimate the prospects for testing its possible interpretation as source of Neutrino Dark Energy within the next decade by the neutrino observatories ANITA and LOFAR. ",
        "introduction": "According to Big Bang Cosmology, in an expanding universe the freeze-out of a particle species occurs, when its interaction strength is too small to keep it in thermal equilibrium. Neutrinos, being the particles with the weakest known interactions, therefore are assumed to have already decoupled when the universe was just $\\approx 1$~ s old, thereby guaranteeing a substantial relic neutrino abundance today with average number density $n_{\\nu_{0,i}}=n_{\\nubar_{0,i}}=56\\,{\\rm{cm}}^{-3}$ per neutrino species $i=1,2,3$. Yet, in turn, the weakness of the neutrino interactions so far has spoilt all attempts to probe this $1.95$ K cosmic neutrino background (\\CnuB), which is the analog of the $2.73$ K cosmic microwave background (CMB) of photons, in a laboratory setting~\\cite{Hagmann:1999kf,Ringwald:2004te,Gelmini:2004hg,Ringwald:2005zf}. However, other cosmological measurements, such as the light element abundance, large scale structure (LSS) and the CMB anisotropies are sensitive to the presence of the \\CnuB and therefore have provided us at least with indirect evidence for its existence (see e.g. Ref.~\\cite{Hannestad:2004nb} for a review).  Independently, Type Ia Supernova (SNIa) results (e.g.~\\cite{Riess:2004nr}), supported by CMB~\\cite{Spergel:2006hy} and LSS data (e.g. Refs.~\\cite{Dodelson:2001ux,Szalay:2001rw}), strongly suggest the existence of an exotic, smooth energy component with negative pressure, known as dark energy, which drives the apparent accelerated expansion of our universe. Recently, Fardon, Nelson and Weiner~\\cite{Fardon:2003eh} have shown that the relic neutrinos, which constitute the \\CnuB, are promoted to a natural dark energy candidate if they interact through a new non-Standard-Model scalar force -- an idea which has great appeal. Neutrinos are the only Standard Model (SM) fermions without right-handed partners. Provided lepton number is violated, the active (left-handed) neutrinos are generally assumed to mix with a dark right-handed neutrino via the well-known seesaw mechanism~\\cite{Gell-Mann,Yanagida,Minkowski,Mohapatra:1979ia}, thus opening a window to the dark sector. Therefore, it would not seem to be surprising if neutrinos, whose interactions and properties we know comparably little about, were sensitive to further forces mediated by dark particles. Moreover, the scale relevant for neutrino mass squared differences as determined from neutrino oscillation experiments, $\\delta {\\mnu}^2\\sim (10^{-2}\\,\\,\\eV)^2$, is of the order of the tiny scale associated with the dark energy, $(2\\times 10^{-3}\\,\\,\\eV)^4$.  As a consequence of the new interaction in such a scenario, an intricate interplay links the dynamics of the relic neutrinos and the mediator of the dark force, a light scalar field called the acceleron. On the one hand, the neutrino masses $\\mnui$ are generated by the vacuum expectation value $\\A$ of the acceleron, $\\mnui(\\A)$. Correspondingly, the $\\A$ dependence of the masses $\\mnui(\\A)$ is transmitted to the neutrino energy densities $\\rho_{\\nu_i}(\\mnui(\\A))$ since these are functions of $\\mnui(\\A)$. On the other hand, as a direct consequence, the neutrino energy densities $\\rho_{\\nu_i}(\\mnui,\\A)$ can stabilize the acceleron by contributing to its effective potential $\\Veff(\\A,\\rhoi)$, which represents the total energy density of the coupled system. Moreover, cosmic expansion manifests itself in the dilution of the neutrino energy densities $\\rhoi(z)\\sim (1+z)^3$. Therefore, it crucially affects the effective acceleron potential $\\Veff(\\A,\\rhoi(\\mnui,z))$ by introducing a dependence on cosmic time, here parameterized in terms of the cosmic redshift $z$. For a homogeneous configuration, the equilibrium value of the acceleron instantaneously minimizes its effective potential $\\Veff(\\A,\\rhoi(\\mnui,z))$ and therefore also $\\A(z)$ varies on cosmological time scales. Finally, in turn, since the neutrino masses $\\mnui(\\A)$ are sensitive to changes in $\\A$, they are promoted to dynamical quantities depending on $z$, $\\mnui(z)$, that is depending on cosmic time. To summarize, the variation of the neutrino masses represents a clear signature of the so-called Neutrino Dark Energy scenario.  In a subsequent work Fardon, Nelson and Weiner~\\cite{Fardon:2005wc} presented a supersymmetric Mass Varying Neutrino (MaVaN) model in which the origin of dark energy was attributed to the lightest neutrino $\\nu_1$ and the size of the dark energy could be expressed in terms of neutrino mass parameters. By naturalness arguments the authors concluded that the lightest neutrino still has to be relativistic today, thereby preventing potential instabilities~\\cite{Afshordi:2005ym,Takahashi:2006jt,Spitzer:2006hm} which could occur in highly non-relativistic theories of Neutrino Dark Energy.  The rich phenomenology of the MaVaN scenario has been explored by many authors. The cosmological effects of varying neutrino masses have been studied in Refs.~\\cite{Brookfield:2005td,Brookfield:2005bz} and were elaborated in the context of gamma ray bursts~\\cite{Li:2004tq}. Apart from the time variation, the conjectured new scalar forces between neutrinos as well as the additional possibility of small acceleron couplings to matter lead to an environment dependence of the neutrino masses governed by the local neutrino and matter density~\\cite{Fardon:2003eh,Kaplan:2004dq,Weiner:2005ac}. The consequences for neutrino oscillations in general were exploited in Refs.~\\cite{Kaplan:2004dq,Zurek:2004vd} and in particular in the sun~\\cite{Cirelli:2005sg,Gonzalez-Garcia:2005xu,Barger:2005mn}, in reactor experiments~\\cite{Schwetz:2005fy,Gonzalez-Garcia:2005xu} as well as in long-baseline experiments~\\cite{Gu:2005pq}.  In light of the possible realization of Neutrino Dark Energy in nature, a (more) direct detection of the \\CnuB should be thoroughly explored with special attention turned to possible new physics beyond the SM. By this means, a time evolution of neutrino masses could be revealed which would serve as a test of Neutrino Dark Energy. In addition, the general importance of a (more) direct evidence for the existence of the \\CnuB lies in a confirmation of standard cosmology back to the freeze-out of the weak interactions and therefore thirteen orders of magnitudes before the time when photons where imprinted on the last scattering surface.  An appealing opportunity to catch a glimpse of the \\CnuB as it is today emerges from the possible existence of extremely high-energy cosmic neutrinos (\\UHEnu's). Such \\UHEnu's can annihilate with relic anti-neutrinos (and vice versa) into $Z$ bosons, if their energies coincide with the respective resonance energies $E_{0,i}^{\\res}$ of the corresponding process $\\nu\\nubar\\rightarrow Z$~\\cite{Weiler:1982qy,Weiler:1983xx,Gondolo:1991rn,Roulet:1992pz,Yoshida:1996ie,Eberle:2004ua,Barenboim:2004di,D'Olivo:2005uh}. These energies, \\be \\label{Eres0} E_{0,i}^{\\res}=\\frac{M^2_Z}{2m_{\\nu_{0,i}}}=4.2\\times 10^{12}\\,\\,\\left(\\frac{\\eV}{\\mnui}\\right) {\\rm GeV} \\ee in the rest system of the relic neutrinos, are entirely determined by the $Z$ boson mass $M_Z$ as well as the respective neutrino masses $\\mnui$. An exceptional loss of transparency of the \\CnuB for cosmic neutrinos results from the fact that the corresponding annihilation cross-section on resonance is enhanced by several orders of magnitude with respect to non-resonant scattering. As a consequence, the diffuse \\UHEnu flux arriving at earth is expected to exhibit absorption dips whose locations in the spectrum are determined by the respective resonance energies of the annihilation processes. Provided that the dips can be resolved on earth, they could produce the most direct evidence for the existence of the \\CnuB so far. Furthermore, as indicated by the resonance energies in Eq.~(\\ref{Eres0}), the absorption features depend on the magnitude of the neutrino masses and could therefore reflect their possible variation with time. Moreover, they are sensitive to the flavor composition of the neutrino mass eigenstates as well as to various cosmological parameters. Accordingly, the possibility opens up to perform relic neutrino absorption spectroscopy as an independent means to probe neutrino physics and cosmology.  The existence of \\UHEnu's is theoretically well motivated and is substantiated by numerous works on possible \\UHEnu sources of astrophysical nature (bottom-up) (see e.g.~\\cite{Torres:2004hk} for a review) or so-called top-down sources (see e.g. Ref.~\\cite{Bhattacharjee:1998qc} for a review). In the latter case, \\UHEnu's with energies well above $10^{11}$ GeV are assumed to be produced in the decomposition of topological defects (TD's) which originate from symmetry breaking phase transitions in the very early universe.  \\begin{figure} \\vspace* {0.0in} \\begin{center} \\includegraphics*[bbllx=12pt,bblly=208pt,bburx=577pt,bbury=611pt,height=9.5cm,width=14.5cm]{e2flux_all_0805GeV.eps} \\caption[]{Current status and next decade prospects for EHEC$\\nu$ physics, expressed in terms of diffuse neutrino fluxes per flavor, $F=F_{\\nu_\\alpha}+F_{\\bar\\nu_\\alpha}$, $\\alpha =e,\\mu,\\tau$. The upper limits from  AMANDA~\\cite{Ackermann:2005sb}, see also Ref.~\\cite{Halzen:2006ic}, ANITA-lite~\\cite{Barwick:2005hn}, FORTE~\\cite{Lehtinen:2003xv}, GLUE~\\cite{Gorham:2003da}, and RICE~\\cite{Kravchenko:2003tc} are plotted. Also shown are projected sensitivities of ANITA\\cite{Barwick:2005hn}, EUSO~\\cite{Bottai:2003ep}, IceCube~\\cite{Ahrens:2002dv}, LOFAR~\\cite{Scholten:2005pp}, OWL~\\cite{Stecker:2004wt}, the Pierre Auger Observatory in $\\nu_e$, $\\nu_\\mu$ modes and in $\\nu_\\tau$ mode (bottom swath)\\cite{Bertou:2001vm}, SalSA~\\cite{Gorham:2001wr}, and WSRT\\cite{Scholten:2005pp}, corresponding to one event per energy decade and indicated duration. Also shown are predictions from astrophysical Cosmic Ray (CR) sources\\cite{Ahlers:2005sn}, from inelastic interactions of CR's with the cosmic microwave background (CMB) photons (cosmogenic neutrinos)\\cite{Ahlers:2005sn,Fodor:2003ph}, and from topological defects\\cite{fodor:tbp}.} \\end{center} \\label{fig1} \\end{figure}  Depending on the underlying \\UHEnu sources the \\UHEnu fluxes could be close to the current observational bounds set by existing \\UHEnu observatories such as AMANDA~\\cite{Ackermann:2005sb} (see also Ref.~\\cite{Halzen:2006ic}), ANITA-lite\\cite{Barwick:2005hn}, BAIKAL~\\cite{Wischnewski:2005rr}, FORTE~\\cite{Lehtinen:2003xv}, GLUE~\\cite{Gorham:2003da} and RICE~\\cite{Kravchenko:2003tc} which cover an energy range of $10^7\\,\\,{\\rm GeV}<E_0<10^{17}\\,\\,{\\rm GeV}$ (cf. Fig.~\\ref{fig1}). Promisingly, the sensitivity in this energy range will be improved by orders of magnitude (cf. Fig.~\\ref{fig1}) by larger \\UHEnu detectors such as ANITA, EUSO~\\cite{Bottai:2003ep}, IceCube~\\cite{Ahrens:2002dv}, LOFAR~\\cite{Scholten:2005pp}, OWL~\\cite{Stecker:2004wt}, SalSA~\\cite{Gorham:2001wr} and WRST~\\cite{Scholten:2005pp} which are planned to start operating within the next decade (cf. Fig.~\\ref{fig1}). Accordingly, the prospects of confirming the existence of the \\CnuB by tracking its interaction with \\UHEnu's have substantially improved since the original proposal in 1982~\\cite{Weiler:1982qy}. Moreover, in the likely case of appreciable event samples the valuable information encoded in the absorption features of the \\UHEnu spectra could be revealed within the next decade (cf. Fig.~\\ref{fig1}), rendering the theoretical exploration of relic neutrino absorption spectroscopy a timely enterprise.  Note that the scenario introduced above has also attracted attention for another reason than the possible detection of the \\CnuB -- namely for the controversial possibility of solving the so-called GZK-puzzle to be discussed briefly in the following. Beyond the Greisen-Zatsepin-Kuzmin (GZK) energy, $E_{\\rm GZK}=4\\times 10^{10}$ GeV, ultra-high energy nucleons rapidly lose energy due to the effective interaction with CMB photons (predominantly through resonant photo-pion production)~\\cite{Greisen:1966jv,Zatsepin:1966jv}. In the so-called $Z$-burst scenario, the observed mysterious cosmic rays above $E_{\\rm GZK}$ were associated with the secondary cosmic ray particles produced in the decays of $Z$ bosons. The latter were assumed to originate from the neutrino annihilation process outlined above~\\cite{Fargion:1997ft,Weiler:1997sh,Yoshida:1998it,Fodor:2001qy,Fodor:2002hy,Gelmini:2004zb}.  However, recently, ANITA-lite~\\cite{Barwick:2005hn} appears to have entirely excluded the $Z$-burst explanation for the GZK-puzzle at a level required to account for the observed fluxes of the highest energy cosmic rays. We would like to stress, that this only means that the GZK-puzzle stays unsolved (if there is any) . Moreover, this neither restricts the possible success of producing evidence for the \\CnuB by means of detecting absorption dips in the \\UHEnu spectra nor does it spoil the possibility of gaining valuable information from performing relic neutrino absorption spectroscopy.  The goal of this paper is to carefully work out the characteristic differences in the \\UHEnu absorption features which result from treating the neutrino masses as time varying dynamical quantities in comparison to constants. In our analysis, we incorporate the full thermal background effects on the absorption process whose impact grows for smaller neutrino masses~\\cite{Barenboim:2004di,D'Olivo:2005uh}. This means, that in general relic neutrinos cannot be assumed to be at rest. Instead, they have to be treated as moving targets with a momentum distribution, if their mean momenta turn out to be of the order of the relic neutrino masses.  We illustrate our results for the diffuse neutrino fluxes to be observed at earth firstly by considering astrophysical \\UHEnu sources. Secondly, we calculate the neutrino spectrum (both for varying and constant neutrino masses) expected from the decomposition of a topological defect using the appropriate fragmentation functions and including the full thermal background effects. By this means, for a given \\UHEnu injection spectrum, we present all technical tools to interpret \\UHEnu absorption dips as soon as they are observed at earth. Thereby the possibility opens up to test the \\CnuB and its interpretation as source of Neutrino Dark Energy.  The paper is organized as follows. In Sec.~1 we discuss the MaVaN scenario and focus on a determination of the time dependence of the neutrino masses. Sec.~2 reviews the treatment of the absorption of an \\UHEnu by the \\CnuB in terms of the damping factor comprising the thermal background effects. Furthermore, in order to make contact to treatments in the literature~\\cite{Weiler:1982qy,Weiler:1983xx,Roulet:1992pz,Eberle:2004ua,Barenboim:2004di,D'Olivo:2005uh}, we include in the discussion common approximations~\\cite{Roulet:1992pz,Eberle:2004ua,Barenboim:2004di} which neglect part or all of the dependence of the damping on the relic neutrino momenta. Moreover, we extend the complete analysis to incorporate a possible variation of the neutrino masses with time. In Sec.~3 we present and compare our results for the survival probabilities of \\UHEnu's with varying and constant masses which encode the physical information on all possible annihilation processes on their way from their source to us, again taking into account the thermal motion of the relic neutrinos. In order to gain more physical insight, in addition, we disentangle the characteristic features of the absorption dips caused by the mass variation by switching off all thermal background effects. Sec.~4 illustrates the discovery potential of neutrino observatories for the \\CnuB and gives an outlook for the testability of the MaVaN scenario. Therefore, both for astrophysical sources and for a topological defect scenario, we calculate the expected observable \\UHEnu flux arriving at earth which results from folding the survival probabilities with the corresponding \\UHEnu source emissivity distribution. In the latter case, we perform the full state-of-the-art calculation with the help of fragmentation functions and by the inclusion of all thermal background effects. In Sec.~5 we summarize our results and conclude. ",
        "conclusions": "} %  In light of the number of extremely high-energy neutrino (\\UHEnu) observatories in operation and under construction with a combined sensitivity ranging up to $10^{17}$ GeV, the prospects for establishing the existence of \\UHEnu fluxes appear to be very promising. As a next step, the exciting possibility opens up to trace the annihilation of \\UHEnu's and relic anti-neutrinos (and vice versa) into $Z$ bosons by localizing absorption dips in the \\UHEnu spectra at energies set by the neutrino masses. On the one hand, their detection could furnish the most direct evidence for the \\CnuB so far and thereby confirm standard cosmology back to the time of light neutrino decoupling. On the other hand, the shape of the absorption lines could reveal a variation of neutrino masses with time and thus verify the interpretation of the \\CnuB as source of Neutrino Dark Energy.  We therefore considered a viable Mass Varying Neutrino (MaVaN) model with the following features entering our analysis on relic neutrino absorption. By the requirement that the lightest neutrino still has to be moderately relativistic today the neutrino mass scale is set to be low, which leads to very conservative predictions. Furthermore, the evolving neutrino masses $\\mnui(z)$, which we determined numerically as functions of redshift assuming $m_{\\nu_{0,1}}=10^{-5}$ eV, turned out to be well approximated by simple power laws $(1+z)^{-1}$ and $(1+z)^{-1/2}$ in the low and in the high redshift regime, respectively. Accordingly, as a generically important feature, they are decreasing functions of redshift as in all standard MaVaN scenarios.  In order to provide all technical tools to interpret \\UHEnu absorption dips for a given injection spectrum and to extract valuable information on neutrino physics, cosmology and possibly physics beyond the Standard Model, we proceeded in the following way. We considered in parallel the neutrino masses to be functions of cosmic time as well as to be constants. In our analysis we took into account the full thermal background effects which result from the relic neutrino motion according to their phase space distribution. In order to compare our results to the literature, we included in our discussion common approximations~\\cite{Roulet:1992pz,Eberle:2004ua,Barenboim:2004di} which neglect part or all of the dependence of the damping on the relic neutrino momenta.  On the level of the survival probabilities which govern the \\UHEnu fluxes, we found the following results: For low emission redshifts ($z\\sim {\\cal{O}}(5)$), the absorption dips produced by the varying neutrino masses $\\mnui(z)$ for $i=2,3$ exhibit narrow absorption minima, which do not suffer a distortion to lower energies as the corresponding dips of constant mass neutrinos. As a consequence, for MaVaNs, the absorption dips of the flavor components $\\nu_{\\mu}$ and $\\nu_{\\tau}$ (which are mostly composed of the heavier two mass eigenstates) are clearly deeper and shifted to higher energies by almost an order of magnitude with respect to the corresponding constant mass minima. For an increased emission redshift $z\\gg 5$, these features become somewhat less pronounced but essentially prevail. A better understanding of the characteristic signatures caused by the mass evolution was obtained by switching off the superposing thermal wash-out caused by the relic neutrino motion. After neglecting the relic neutrino momenta for this purpose, we found that the crucial deviations result from the dependence of the corresponding resonance energies on the neutrino masses $E^{\\res}_{i}\\sim 1/ \\mnui$ for $i=1,2,3$. In the case of MaVaNs, the mass variation $\\mnui(z)$ induces a dependence on the annihilation redshift $z$ according to $E^{\\res}_{i}(z)\\sim E^{\\res}_{0,i}(1+z)$ for all neutrino species $i=1,2,3$ in the low redshift regime. Accordingly, this $z$ dependence of the resonance energies compensates for the energy loss of the \\UHEnu due to cosmic redshift proportional to $(1+z)^{-1}$, resulting in narrow absorption spikes at constant energies $E^{\\res}_{i}(z)/(1+z)=E^{\\res}_{0,i}$ (like one would expect for constant neutrino masses in a non-expanding universe). In contrast, for constant neutrino masses $\\mnui=m_{\\nu_{0,i}}$ the absorption dips are broadened, since the redshifted resonance energies to be measured on earth are given by $E^{\\res}_{i}/(1+z)=E^{\\res}_{0,i}/(1+z)$, for $z$ taking values between $0$ and the \\UHEnu emission redshift. Since in the standard MaVaN scenario the neutrino masses are decreasing functions of redshift, they generically reduce the effect of cosmic redshift on the \\UHEnu survival probabilities. As a result, they always produce deeper absorption minima, which, in addition, are shifted to higher energies in comparison to the dips caused by constant neutrino masses.  In order to illustrate the discovery potential for absorption dips in the \\UHEnu spectra to be observed at earth and to estimate the prospects of testing scenarios of Neutrino Dark Energy, we considered plausible \\UHEnu fluxes originating from astrophysical acceleration sites or from topological defect sources. We presented our results both for the energy squared times the flavor summed \\UHEnu flux $E^2_0 F$  with $F=\\sum F_{\\nu_{\\alpha}}+\\sum F_{\\bar{\\nu}_{\\alpha}}$ and for $E^2_0 F_{\\tau}$ with $F_{\\tau}=F_{\\nu_{\\tau}}+F_{\\bar{\\nu}_{\\tau}}$, where the latter can at best be identified by LOFAR~\\cite{Private}. Despite the adopted low neutrino mass scale, we found both for varying and constant neutrino masses that for topological defect and for astrophysical \\UHEnu sources at $z_{\\max}>5$, LOFAR and ANITA promise a statistically significant evidence for absorption dips (even if the underlying fluxes are well below the cascade limit). Accordingly, the most direct detection of the \\CnuB so far seems to be in reach within the next decade.  Furthermore, the flux dips of varying and constant mass \\UHEnu's expected from astrophysical sources retain the characteristic differences induced by the survival probabilities. Besides being clearly shifted to higher energies, the MaVaN dips are deeper and therefore even facilitate a resolution of absorption features in the \\UHEnu spectra in comparison to constant mass neutrinos both in the case of $E^2_0 F$ and of $E^2_0 F_{\\tau}$. As a main result of our analysis, these deviations of the MaVaN and constant mass absorption curves for astrophysical sources turned out to be statistically significant, yet a decent energy resolution seems necessary for their detection. Given an energy resolution of $\\Delta E/E\\sim 30\\%$ as at best achievable for LOFAR~\\cite{Private}, relic neutrino absorption spectroscopy could reveal a variation of neutrino masses and therefore possibly the nature of Dark Energy.  As concerns topological defect sources, the absorption lines in the \\UHEnu fluxes for time dependent and constant neutrino masses altogether are more similar in shape, however, somewhat deeper for constant neutrino masses. Furthermore, they extend to much lower energies than for astrophysical \\UHEnu sources and their minima are considerably deeper. All of these features are a result of the much higher annihilation redshifts $z_s\\gg 1$ possible for \\UHEnu's originating from the decomposition of topological defects in comparison to \\UHEnu's from astrophysical acceleration sites. At high redshifts, the \\UHEnu's are absorbed by a hotter bath of relic neutrinos. Consequently, in the energy region spanned by the absorption dips where $\\mnui/T_{\\nu} \\ll 1$, thermal background effects wash out any features produced by the neutrino mass or its possible variation. Since the MaVaN masses are decreasing functions of redshift, they reach this limit for much smaller redshifts than the corresponding constant masses. Only the mass of the heaviest constant mass eigenstate is sufficiently large, $m_{\\nu_3}/T_{\\nu}(z)\\gg 1$, in the relevant energy region, leading to a deeper absorption curve than the one produced by all of the other MaVaN and constant mass eigenstates. Since ${\\nu_{\\tau}}$ is mostly composed of the heaviest mass eigenstate, $F_{\\tau}$ exhibits deeper constant mass dips than $F$. Accordingly, for $F_{\\tau}$ the signatures of varying neutrino masses can more easily be distinguished from those of constant masses than in the case of $F$. However, in order to reveal a neutrino mass variation, it seems necessary both to identify the tau neutrino flavor and to have a good energy resolution.  Recently, the authors of Ref.~\\cite{Takahashi:2006jt} claimed that certain constraints on the fundamental scalar acceleron potential $V_0(\\A)$ (cf. Sec.~\\ref{sec:MaVaNs}) can also lead to MaVaN models stable against the growth of inhomogeneities~\\cite{Afshordi:2005ym} even in the highly non-relativistic regime. Accordingly, one could construct other MaVaN models than the one under consideration in this paper, whose viability would not rely on a low neutrino mass scale (cf. Sec.~\\ref{sec:MaVaNs}). As we pointed out, the characteristic absorption signatures of any standard MaVaN scenario (cf. Sec.~\\ref{sec:SurvivalProb} and Sec.~\\ref{sec:AbsorptionDips}) are essentially generic apart from details. Yet, a higher neutrino mass scale would even increase the overall dip depth in comparison to our rather conservative predictions and also reduce the importance of the thermal background effects on the absorption features. Accordingly, we would expect the features produced by a possible mass evolution not to be washed out by the temperature effects for a wider energy range of the dips. Thus the deviations with respect to the corresponding constant mass curves would even be more prominent."
    },
    "0606/astro-ph0606585_arXiv.txt": {
        "abstract": "The luminous infrared galaxy Arp~299 (IC~694+NGC~3690) is studied using optical integral field spectroscopy obtained with the INTEGRAL system, together with archival \\textit{Hubble Space Telescope} WFPC2 and NICMOS images. The stellar and ionized gas morphology shows $\\lambda$-dependent variations due to the combined effects of the dust internal extinction, and the nature and spatial distribution of the different ionizing sources. The two-dimensional ionization maps have revealed an off-nuclear conical structure of about 4 kpc in length, characterized by high excitation conditions and a radial gradient in the gas electron density. The apex of this structure coincides with B1 region of NGC~3690 which, in turn, presents Seyfert-like ionization, high extinction and a high velocity dispersion. These results strongly support the hypothesis that B1 is the true nucleus of NGC~3690, where an AGN is located. In the circumnuclear regions H\\,{\\sc ii}-like ionization dominates, while LINER-like ionization is found elsewhere. The H$\\alpha$ emitting sources with ages from 3.3 to 7.2$\\times$10$^{6}$ years, have masses of between 6 and 680$\\times$10$^{6}$ M$\\odot$ and contribute (extinction corrected) about 45$\\%$ to the bolometric luminosity. The ionized (H$\\alpha$) and neutral (NaD) gas velocity fields show similar structure on scales of several hundred to about 1 kpc, indicating that these gas components are kinematically coupled. The kinematic structure is complex and on scales of about 0.2 kpc does not appear to be dominated by the presence of ordered, rotational motions. The large velocity dispersion measured in NGC~3690 indicates that this galaxy is the most massive of the system. The low velocity amplitude and dispersion of the interface suggest that the ionized gas is slowly rotating or in a close to quiescent phase. ",
        "introduction": "Arp~299 is a nearby luminous infrared galaxy  (d$\\sim$43 Mpc for H$_{0}$=70 km~s$^{-1}$Mpc$^{-1}$, $\\Omega_{M}$= 0.7, $\\Omega_{\\Lambda}$= 0.3) with  an infrared (8-1000 $\\mu$m) luminosity of L$_{IR}$=5.7$\\times$10$^{11}$L$_{\\odot}$\\footnote{IRAS fluxes obtained from Moshir et al. (1993). L$_{IR}$ derived from the equation detailed on Sanders \\& Mirabel (1996).}.  It  is an interacting system (Fig. \\ref{galaxia}) composed of two individual galaxies (IC~694+NGC~3690\\footnote{Through this paper, we use the notation introduced by Gehrz et al (1983) for the two components of Arp~299: the nucleus of IC~694 (eastern component) is region A, while the main sources in NGC~3690 (western component) are B1, B2, C(C1+C2+C3) and C$'$ (see also Alonso-Herrero et al 2000).}) in an early dynamical stage. Hibbard \\& Yun (1999) have discovered an H\\,{\\sc i} tidal tail of 180 kpc in length, with no evidence for tidal dwarf galaxies.  Powerful starburst regions with star formation rates of up to about 100 M$_{\\odot}$ yr$^{-1}$ have been identified (Alonso-Herrero et al. 2000, hereafter AAH00). Moreover, the system presents several regions where supernova remnants have been detected at radio wavelengths (Neff et al. 2004). More recently, Mattila et al. (2005) have discovered a new supernova using near-IR images, reinforcing the prediction that this galaxy has a high rate of supernova explosions (Mannucci et al. 2003).  Optical spectroscopic studies have classified IC~694 as starburst, and NGC~3690 as starburst/LINER (e.g., Coziol et al. 1998), and mid-IR observations classified the system as starburst (e.g., Laurent et al. 2000). Using \\textit{Hubble Space Telescope} (HST) imaging and near-infrared spectroscopy, AAH00 modeled the star formation properties of the system without the need for an AGN. The first hint for the presence of an AGN in Arp~299 came from X-ray observations with the \\textit{BeppoSAX} satellite by Della Ceca et al. (2002). These authors concluded that a highly absorbed (N$_{H}$~$\\simeq$~2.5$\\times$10$^{24}$ cm$^{-2}$) AGN, with an intrinsic luminosity of L$_{0.5-100 keV}$ $\\simeq$ 1.9$\\times$10$^{43}$ erg s$^{-1}$, is present in this galaxy, although due to the low spatial resolution its exact location could not be determined. Using \\textit{Chandra} and \\textit{XMM-Newton} observations, recent studies (Ballo et al. 2004) have inferred the presence of two AGNs, one located in NGC~3690 (source B1; 6.4 keV line), and another in IC~694 (source A; 6.7 keV Fe-H$\\alpha$ line). In addition, Gallais et al. (2004) suggest the presence of an AGN in NGC~3690 based on the shape of the mid-IR continuum, while in IC~694 they only detect spectral features characteristic of massive, dust enshrouded star formation.  This paper presents a detailed study of the two-dimensional star formation, kinematics, and ionized gas properties of Arp~299, covering the two individual galaxies of the system and the interface area. At a distance of about 43 Mpc, Arp~299 is one of the nearest interacting, luminous infrared galaxies, and therefore these observations provide high spatial resolution (about 0.205 kpc arcsec$^{-1}$) and allow us to investigate in detail the structure and physical properties of this galaxy. The paper is organized as follows: \\S 2 describes the observations and data reduction, while the selection of regions under study is described in \\S 3. The next sections present a discussion of the results, with special emphasis on the stellar and ionized gas morphology (\\S \\ref{sec4}), the internal dust distribution and extinction effects (\\S \\ref{sec5}), the extinction estimates using optical and near-IR emission lines (\\S \\ref{sec54}), the two dimensional ionization structure (\\S \\ref{sec6}), the electron density (\\S \\ref{electron}), the origin of the LINER-like ionization (\\S \\ref{sec8}), the star-forming properties (\\S \\ref{sec9}), and the kinematics and dynamical mass derivations (\\S \\ref{sec-kin}). Finally \\S \\ref{sec11} presents a short summary of the main new results. ",
        "conclusions": "\\label{sec11} A detailed IFS study of the LIRG Arp~299 (IC~694+NGC~3690) system covering a region of 7.0$\\times$6.2 kpc has been presented. This work illustrates the flexibility and power of the IFS in studying  in two dimensions and simultaneously the different physical properties, such as stellar populations, internal dust extinction, ionization structure and gas kinematics, of complex systems like Arp~299. In what follows only the main results of this study are summarized:  \\begin{enumerate} \\item The observed overall two dimensional stellar and ionized gas structure differs due to internal extinction effects caused by the dust, and to the different spatial distribution of the diverse (AGN, stars, shocks) ionizing sources. The stellar continuum and emission line distribution peaks are not spatially coincident but separated up to 1.4 kpc. The peaks of the ionized gas traced by different emission lines (e.g. hydrogen recombination, high excitation [OIII], and shock-tracer [OI]) are also not coincident indicating that the ionizing mechanisms change on scales of about 200 pc or less.   \\item The region B1 is identified as the true nucleus of NGC~3690. The presence of an AGN in B1 is inferred from the detection of a high-excitation Seyfert-like conical structure of about 1.5 kpc in size, an opening angle  of about 54\\,$^{\\circ}$, and its apex being located in B1. This region also presents the additional characteristics of a dust -enshrouded nucleus with high extinction (A$_{V}$=3.7$\\pm$0.6), red colors (F160W-F222M=2.06$\\pm$0.1) indicating hot dust emission, and high velocity dispersion ($\\sigma$=172$\\pm$20 km s$^{-1}$).  \\item The observed near-IR colors reveal the presence of hot dust over the entire structure. According to he F160W-F222M color, the upper limit for the mean hot dust contribution is about 30$\\%$ at 600 K for the majority of the star-forming regions. The nuclei A and B1, and region C1 are the most affected by the hot dust component, with an upper limit of about 70$\\%$, 80$\\%$ and 90$\\%$ at 600 K respectively.  \\item The high surface brightness regions are mainly dominated by star formation, while the low surface brightness regions are dominated by LINER-like ionization process. The H$\\alpha$-emitting star-forming regions present an extinction of between A$_{V}$=1$-$6 mag, with ages between 3.3 and 7.2$\\times$10$^{6}$ years, and masses that range between about 6 and 680$\\times$10$^{6}$ M$_{\\odot}$. The total contribution (extinction corrected) of these sources to the bolometric luminosity is about 45\\%.   \\item The structure of the ionized (H$\\alpha$) and neutral (NaD) gas two dimensional velocity fields show similar features, indicating that both components are kinematically coupled on scales of hundreds of pc and larger. The velocity fields are very complex, and in general the galaxy does not appear to be dominated by ordered virialized simple (e.g. rotation) motions.  The morphology of the velocity field does not correlate with the distribution of mass as traced by the continuum images, that is, no spatial coincidence between the kinematic and the photometric centers exists.  On average, the interface region presents the lowest value of velocity dispersion ($\\sigma$$\\sim$45 km s$^{-1}$), IC~694 presents the intermediate values ($\\sigma$$\\sim$65 km s$^{-1}$), and NGC~3690 the highest ($\\sigma$$\\sim$120 km s$^{-1}$). This indicate that NGC~3690 is the warmest body of the entire system, and probably the most massive. On the contrary, the low velocity amplitude and velocity dispersion of the interface suggests that the ionized gas is slowly rotating or in a close quiescent phase.  The velocity dispersion measurements over the nuclei and the brightest emitting regions,  imply a mean value for the dynamical mass of about 1.2$\\times$10$^{9}$ M$_{\\odot}$ and 1.5 $\\times$10$^{9}$ for nuclei A and B1 respectively, within a mean radius of 58 pc and 28 pc. For regions B2, C1 and C$'$ the mass ranges between $\\sim$0.2$\\times$10$^{9}$ M$_{\\odot}$ and $\\sim$0.6$\\times$10$^{9}$ M$_{\\odot}$, within a radius of about 30 pc.  \\end{enumerate}  \\smallskip \\textbf"
    },
    "0606/astro-ph0606250_arXiv.txt": {
        "abstract": "In the last few years, new Adaptive Optics [AO] techniques have emerged to answer new astronomical challenges: Ground-Layer AO [GLAO] and Multi-Conjugate AO [MCAO] to access a wider Field of View [FoV], Multi-Object AO [MOAO] for the simultaneous observation of several faint galaxies, eXtreme AO [XAO] for the detection of faint companions. In this paper, we focus our study to one of these applications : high red-shift galaxy observations using MOAO techniques in the framework of Extremely Large Telescopes [ELTs]. We present the high-level specifications of a dedicated instrument. We choose to describe the scientific requirements with the following criteria : $40$\\% of Ensquared Energy [EE] in $H$ band ($1.65\\mu$m) and in an aperture size from $25$ to $150$ mas. Considering these specifications we investigate different AO solutions thanks to Fourier based simulations. Sky Coverage [SC] is computed for Natural and Laser Guide Stars [NGS, LGS] systems. We show that specifications are met for NGS-based systems at the cost of an extremely low SC. For the LGS approach, the option of low order correction with a faint NGS is discussed. We demonstrate that, this last solution allows the scientific requirements to be met together with a quasi full SC. ",
        "introduction": "\\label{sect:intro}  % A new era of astronomical telescopes is about to start with diameters reaching $30$ to $60$ meters. This concept of ELTs will provide a dramatic advance in our understanding of the primordial universe. By accommodating 3D spectroscopy devices and relatively large FoV ($10$ arc minutes), ELTs will be a privilaged tool for the study of formation and evolution of galaxies.\\\\ The spatial resolution of ground based telescopes is limited by the presence of the atmospheric turbulence which leaves us with image resolution of $0.5$ arc second at best. Current Very Large Telescopes [VLTs] ($8$-$10$ meter-class) are now equipped with AO which provides a real time correction of turbulence and enables large telescopes to give almost diffraction limited images. But because of anisoplanatism effects, the AO corrected FoV is small, and a good correction requires bright (Natural or Laser\\cite{Foy85}) Guide Stars near the scientific object. Typically, the compensated FoV around the Guide Star is of the order of a few times the isoplanatic patch $\\theta$$_0$ (a few tens of arcsec in the Near InfraRed). In extragalactic studies, to avoid contamination of light by the interstellar medium, it is necessary to observe in a direction far of our galactic plane. In that case, the surface density of stars becomes very small, and classical AO working with NGS can not be used. Even close to the galactic plane, because of  the required star brightness, SC would be lower than $5$\\%.\\\\ To overcome this problem, new AO techniques (GLAO\\cite{Hubin05}, MOAO\\cite{Hammer01,Gendron05}) have emerged in the last few years to increase the corrected field. These new methods are based on a full measurement of the 3D turbulent volume using several Guide Star. The correction is then applied using one or more Deformable Mirrors [DMs] and the corrected FoV can reach few arcmins. To improve the sky coverage, LGS have been proposed\\cite{LeLouarn98}. This solution, that create an artificial star in the sodium layer of the atmosphere, could be used where the density of stars becomes too small for a good turbulence analysis even with GLAO or MOAO. Should some critical issues be solved (cone effect, spot elongation, tilt indetermination), the gain achievable by such a system could be very promising. The ultimate solution would be to associate LGS with multi-analysis concepts\\cite{Lelouarn02} . \\\\ These instruments are now under development for VLTs. The AO design rules and choices to meet a given performance are relatively well understood for this kind of telescope diameter. However, for ELTs the whole instrumentation has to be reconsidered and redefined. \\\\ \\\\ In section $2$ we present the instrumental specifications imposed by astrophysical goals. Section $3$ is dedicated to a preliminary estimation of an AO design and the SC issue is discussed. In Section $4$ we analyze the performance of an AO system using laser probe, and in Section $5$ we extend this study to Wide Field AO concepts. ",
        "conclusions": "We have presented here some results concerning AO concepts for futures ELTs and particularly for 3D spectroscopy. We have shown that Sky Coverage issue will be a critical point when working with NGS-based system and the LGS solution seems inevitable for observations at high galactic latitudes. To solve the Tilt determination problem and to obtain a full Sky Coverage, we propose a new concept of rather simple Laser Guide Stars systems without Tip-tilt correction. For median atmospheric conditions this concept fulfills the scientific requirements. For better performance, the use of a Tilt reference NGS even weak or far from optical axis reduces significantly the Tilt variance, and still allow a good Sky Coverage. Application to MOAO systems shows that thanks to partial Tilt correction, it could be possible to fulfill the scientific specifications in a large range of observational conditions while preserving a significant Sky Coverage."
    },
    "0606/astro-ph0606299_arXiv.txt": {
        "abstract": "{Mass loss plays a dominant role in the evolution of low mass stars while they are on the Asymptotic Giant Branch (AGB). The gas and dust ejected during this phase are a major source in the mass budget of the interstellar medium. Recent studies have pointed towards the importance of variations in the mass-loss history of such objects.} {By modelling the full line profile of low excitation CO lines emitted in the circumstellar envelope, we can study the mass-loss \\emph{history} of AGB stars.} {We have developed a non-LTE radiative transfer code, which calculates the velocity structure and gas kinetic temperature of the envelope in a self-consistent way. The resulting structure of the envelope provides the input for the molecular line radiative calculations which are evaluated in the comoving frame. The code allows for the implementation of modulations in the mass-loss rate. This code has been benchmarked against other radiative transfer codes and is shown to perform well and efficiently.} {We illustrate the effects of varying mass-loss rates in case of a superwind phase. The model is applied to the well-studied case of \\object{VY CMa}. We show that both the observed integrated line strengths as the spectral structure present in the observed line profiles, unambiguously demonstrate that this source underwent a phase of high mass loss ($\\sim 3.2 \\times 10^{-4}$\\,\\Msun\\,yr$^{-1}$) some 1000\\,yr ago. This phase took place for some 100\\,yr, and was preceded by a low mass-loss phase ($\\sim 1 \\times 10^{-6}$\\,\\Msun\\,yr$^{-1}$) taking some $800$\\,yr. The current mass-loss rate is estimated to be in the order of $8 \\times 10^{-5}$\\,\\Msun\\,yr$^{-1}$.} {In this paper, we demonstrate that both the relative strength of the CO rotational line profiles and the (non)-occurrence of spectral structure in the profile offer strong diagnostics to pinpoint the mass-loss history.} ",
        "introduction": "\\label{introduction} When low and intermediate mass stars approach the end of their lives and ascend the Asymptotic Giant Branch (AGB), mass loss dominates the subsequent evolution ultimately leading to the removal of the entire envelope. In this process, the star develops a dust-driven wind with a velocity of 10 -- 15\\,km\\,s$^{-1}$ \\citep[][and references therein]{Habing2003agbs.conf.....H}. AGB stars lose $\\sim 35 - 85$ per cent of their mass through the stellar wind before ending up as a white dwarf \\citep{Marshall2004MNRAS.355.1348M}. Likewise, more massive stars (M$_{\\rm{ZAMS}} \\ga 8$\\,\\Msun) may pass through a red supergiant phase and lose mass in a similar manner.  The mass-loss properties change with time as these stars ascend the AGB and become more luminous. It is often postulated and in rare cases observed \\citep[e.g.]{Justtanont1996ApJ...456..337J, vanLoon2003MNRAS.341.1205V, Marshall2004MNRAS.355.1348M} that the AGB evolution ends in a very high mass-loss phase, the so-called superwind phase \\citep{Iben1983ARA&A..21..271I}. The nature and onset of such a superwind are far from understood.  This superwind phase is not the only type of variable mass loss. Recent observations of (post) AGB objects and Planetary Nebulae (PNe) show that winds from late type stars are far from being smooth. The shell structures found around e.g. \\object{CRL 2688} \\citep[the Egg Nebula,][]{ Ney1975ApJ...198L.129N, Sahai1998ApJ...493..301S}, \\object{NGC 6543} \\citep[the Cat's Eye Nebula,][]{Harrington1994AAS...185.9003H} and the AGB star \\object{IRC +10 216} \\citep{Mauron1999A&A...349..203M, Mauron2000A&A...359..707M, Fong2003ApJ...582L..39F}, indicate that the outflow has quasi-periodic oscillations, with density contrasts corresponding to mass-loss variations of up to a factor $\\sim$100$-$1000 \\citep{Schoier2005A&A...436..633S}. The time scale for these oscillations is typically a few hundred years, i.e.\\ too long to be a result of stellar pulsation, which has a period of a few hundred days, and too short to be due to nuclear thermal pulses, which occur once in ten thousand to hundred thousand years. Presently, the physical mechanism responsible for these variations, their rate of occurrence and their importance in terms of the amount of ejected matter involved are unknown. Possibly they are linked to the hydrodynamical properties of a dust-driven wind. For example, \\citet{Simis2001A&A...371..205S} find that a feedback mechanism between the wind-acceleration and dust growth at the base of the wind may cause modulations in the mass-loss rate.  Already in the late seventies, low excitation rotational carbon-monoxide (CO) emission lines emerging from the circumstellar envelopes (CSEs) around AGBs were observed and analysed to estimate the mass-loss rates during the AGB phase \\citep[e.g.][]{Zuckermann1977ApJ...211L..97Z, Morris1980ApJ...236..823M, Knapp1980ApJ...242L..25K, Knapp1982ApJ...252..616K, Knapp1985ApJ...292..640K}. CO has proven to be instrumental for this since (almost) all the elemental carbon is locked up in CO throughout most of the envelope. While previous studies have assumed a constant mass-loss rate in the analysis of rotational CO lines, the goal of our study is \\emph{to exploit the rotational CO line profiles of different excitation levels as a complementary diagnostic to probe variations in mass-loss rate}. With the arrival of HIFI and PACS the powerful spectrometers on the Herschel Space Observatory --- ESA's far infrared mission to be launched in 2008 --- and ALMA --- the submillimeter interferometer which is being built by ESO and NRAO in Chile --- new high-quality data of CO emission lines will put new light on this ongoing research of understanding the geometry, physical conditions and mass-loss history of AGB stars.  In this paper, we describe our non-LTE (non- Local Thermodynamic Equilibrium) radiative transfer code developed to predict the CO molecular line strengths. In Sect.~\\ref{temperature_velocity}, the equations governing the temperature and velocity structure in the CSE are presented. Since the CO rotational line profiles are very sensitive to the temperature structure in the CSE, the kinetic temperature is calculated self-consistently taking the main heating and cooling processes into account (Sect.~\\ref{temperature}). The code can treat any density profile, allowing us to deduce the amplitude of the mass-loss variations (Sect.~\\ref{var_mass1}). The non-LTE radiative transfer code is described in Sect.~\\ref{radtrans}, and the effect of variations in mass loss on the predicted CO profiles is studied. In Sect.~\\ref{vycma}, we apply our theoretical code to model the observed CO rotational emission lines in the supergiant \\object{VY CMa}. We summarise the results in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{summary} We have developed a code to derive the gas mass-loss \\emph{history} from a non-LTE analysis of the full line profiles of CO rotational transitions with different excitation energies, including a self-consistent computation of the velocity structure and gas kinetic temperature, taking several heating and cooling mechanisms in the CSE into account. The code has been benchmarked to other existing codes, demonstrating good performance.  The code is applied to the analysis of the wind structure of \\object{VY CMa}. For the first time, we could pinpoint the mass-loss \\emph{history} of this supergiant from the analysis of the CO(1--0), (2--1), (3--2), (6--5), and (7--6) line profiles and could demonstrate that the mass-loss rate underwent substantial changes during the last 1400\\,yr. We especially note a very strong peak of mass loss of $\\sim 3.2 \\times 10^{-4}$\\,\\Msun\\,yr$^{-1}$ some 1000\\,yr ago having lasted $\\sim 100$\\,yr. This period of high mass loss was preceded by a long period of low mass loss, in the order of $1 \\times 10^{-6}$\\,\\Msun\\,yr$^{-1}$ taking some 800\\,yr. The current mass loss is estimated to be in the order of $5 \\times 10^{-5}$ to $1 \\times 10^{-4}$\\,\\Msun\\,yr$^{-1}$. The mass-loss history that we derive is significantly different from what one would estimate from a constant mass-loss rate model, and therefore paints a much more complete picture on the evolution of the mass loss of AGB and red supergiant stars."
    },
    "0606/astro-ph0606244_arXiv.txt": {
        "abstract": "We present far-ultraviolet (FUV) imaging of the Hubble Deep Field North (HDF-N) taken with the Solar Blind Channel of the Advanced Camera for Surveys (ACS/SBC) and the FUV MAMA detector of the Space Telescope Imaging Spectrograph (STIS) onboard the {\\it Hubble Space Telescope}.  The full WFPC2 deep field has been observed at 1600 Angstroms.  We detect 134 galaxies and one star down to a limit of $FUV_{AB}\\sim 29$.  All sources have counterparts in the WFPC2 image.  Redshifts (spectroscopic or photometric) for the detected sources are in the range $0<z<1$.  We find that the FUV galaxy number counts are higher than those reported by GALEX, which we attribute at least in part to cosmic variance in the small HDF-N field of view.  Six of the 13 {\\it Chandra}\\ sources at $z<0.85$\\ in the HDF-N are detected in the FUV, and those are consistent with starbursts rather than AGN.  Cross-correlating with {\\it Spitzer}\\ sources in the field, we find that the FUV detections show general agreement with the expected $L_{\\rm{IR}}/L_{\\rm{UV}}$\\ vs.\\ $\\beta$\\ relationship.  We infer star formation rates (SFRs), corrected for extinction using the UV slope, and find a median value of 0.3 \\Myr\\ for FUV-detected galaxies, with 75\\%\\ of detected sources have SFR$<1$\\ \\Myr.  Examining the morphological distribution of sources, we find that about half of all FUV-detected sources are identied as spiral galaxies.  Half of morphologically-selected spheroids at $z< 0.85$\\ are detected in the FUV, suggesting that such sources have significant ongoing star-formation in the epoch since $z\\sim 1$. ",
        "introduction": "The star formation rate density of the Universe, integrated over all galaxy populations, shows a sharp decline since redshifts near unity \\citep[e.g., ][]{Madau 1996, Madau 1998}.  While the precise shape of the decline with redshift is still uncertain \\citep{Lilly 1996, Hogg 1998, Flores 1999, Wilson 2002}, its existence points to a ``downsizing'' in galaxies that host most of the star formation at $z<1$\\ \\citep{Cowie 1996}.  The characteristics (morphology, mass, luminosity) of these low redshift starbursts may explain the global decline in star formation.  \\cite{Wolf 2005}\\ suggest that the decline is dominated by decreasing star formation in normal spiral galaxies rather than, for example, the decreasing rate of major mergers.  Ultraviolet (UV) emission is an indication of recent star formation in a galaxy.  Despite absorption by dust, the rest-frame UV is strong enough in the majority of star-forming galaxies to be detected in current surveys \\citep{Adelberger and Steidel 2000}.  UV detection can distinguish star-forming from quiescent systems, and indicates the amount of recent star formation, subject to the effects of extinction by dust \\citep[e.g.,][]{Calzetti 1994, Fitzpatrick 1986, Meurer 1999, Buat 2005}.  Far-ultraviolet (FUV) surveys, in particular, can provide direct evidence of recently formed, massive stars in the dominant populations at $z<1$\\ \\citep{Schiminovich 2005}.  Of particular interest is the the star-formation activity present in early type galaxies.  Recent studies have shown that some galaxies which appear morphologically to be spheroids have, nonetheless, substantial ongoing star formation.  \\cite{Yi 2005} find that at least 15\\%\\ of bright, local ellipticals show evidence of recent star formation, ruling out pure monolithic collapse histories for at least that fraction of such sources.  Similarly, studies of the internal color variations in elliptical galaxies have shown almost one third of them show gradients inconsistent with passive evolution \\citep{Abraham 1999, Menanteau 2001, Papovich 2003}.  The formation of spheroid galaxies, then, is a crucial component of successful hierarchical models.  \\cite{Conselice 2005}\\ find that the massive galaxies at $z<1$, both spirals and ellipticals, likely have major-merger progenitors at higher redshifts.  Nonetheless, a significant fraction of stellar mass must still have formed since $z\\sim 1$\\ \\citep{Bell 2004, Dickinson 2003}.  The assembly of that additional stellar mass should be detectable in UV surveys.   We present a far-ultraviolet (1600 \\AA; FUV) imaging survey of the Hubble Deep Field North \\citep[HDF-N;][]{Williams 1996}.  The data were taken in two surveys.  The first utilized the FUV camera of the Space Telescope Imaging Spectrograph \\citep[STIS;][]{Kimble 1998, Woodgate 1998}\\ to survey a small part of the field.  That portion of the data set has been been used to measure galaxy number-magnitude counts \\citep{Gardner 2000 GBF}\\ and the diffuse FUV background emission \\citep{Brown 2000}.  We surveyed the remaining area with the Solar Blind Channel (SBC) of the Advanced Camera for Surveys \\citep{Ford 1998}.  We outline the survey and data reduction in Section \\ref{sec: obs}\\ and present the catalog in Section \\ref{sec: results}.  In Section \\ref{sec: discussion}, we discuss the properies of FUV-detected sources at other wavelengths, the inferred star formation rates, and the implications of the detection of elliptical galaxies.  Throughout, we assume a $\\Lambda$-dominated flat universe, with $H_0=70$\\ km s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7, \\Omega_{m}=0.3$.  Photometry is presented with magnitudes on the AB system which is defined by $AB = -2.5 log F_{\\nu} - 48.6$, where $F_{\\nu}$\\ is given in units of ergs cm$^{-2}$\\ s$^{-1}$\\ Hz$^{-1}$\\ \\citep{Oke 1971}. ",
        "conclusions": "\\label{sec: discussion} \\subsection{Number Counts}   We measure galaxy number-magnitudes for the ACS sources and compare them to the published STIS counts (Figure \\ref{fig: nc}). Dark current variation complicates the measurement of the counts.  We use the procedure outlined in \\cite{Gardner 2000 GBF}.  First, we use the variance map to determine the area over which each galaxy would have been detected.  The STIS and ACS counts are generally consistent. Only one object was measured in common.  In the figure, we also compare the HST number counts to the FUV number counts measured by GALEX \\citep{Xu 2005}, XMM \\citep{Sasseen 2002}, and FOCA \\citep{Milliard 1992}.  These other counts reach measurement reached $FUV_{AB}\\le 24$.  We correct the XMM and FOCA counts from their central wavelength of 2000 \\AA\\ to 1500 \\AA\\ following Xu et al.\\ by assuming a UV slope, $\\beta=-0.8$.  We do not include a color correction for the difference between the ACS and GALEX filters, but we estimate that such a correction could be substantial for distant sources (see Figure \\ref{fig: galex_color_corr}).  The difference for $z>0.5$\\ is the result of the bluer wavelength coverage of the GALEX filter, which is more strongly affected by the 912 \\AA\\ limit (see Figure \\ref{fig: filter_curve}). The redshifting of the Lyman limit combined with the redder transmission of the ACS filter causes it to be sensitive to a larger volume than the GALEX filter, by a factor of $\\sim 30$\\%.  For $z\\sim 0.15$\\ sources, the color correction is reversed for sources with strong \\lya\\ emission lines falling in the GALEX filter but below the blue end of the F150LP filter.  About half of the ACS sources lie at $z>0.5$.  The color correction would be less extreme for the STIS filter, which had a red-end cutoff between that of GALEX and ACS.  The HST counts are higher (by a factor of $\\sim 2$) than both the GALEX counts and the model which fits them.  At $FUV_{AB}\\le 24$, the descrepancy is only marginally significant as the ACS counts lie within 1-2 $\\sigma$\\ of the the GALEX counts.  The XMM and FOCA counts are also higher than GALEX at these magnitudes.  The difference with the model is more significant, as it is repeated over a larger number of bins.  The GALEX counts are fit by a model which assumes essentially pure luminosity evolution, $L_*\\sim (1+z)^{2.5}$, and a starburst SED that is flat between 1000 and 1200 \\AA\\ \\citep{Xu 2005}.  The HST counts are significantly higher than the model.  Some of this difference is the result of the filter difference discussed above, which causes a $\\sim 30$\\% difference in the volume surveyed and potentially a half magnitude of color-correction.  Thus, there cannot be much more than a factor of $\\sim 2$\\ between the counts and the model. The most likely explanation for this difference is cosmic variance.  The HST counts are dominated by the very small field of view of the HDF-N, as the STIS counts include only a single pointing in the HDF-S and seven in the HDF-N.  The northern sightline is known to have source overdensities at $z\\sim 0.45$\\ and $z\\sim 0.8$\\ \\citep{Cohen 2000}. So, it may not be surprising that the HST counts are higher. \\cite{Somerville 2004}\\ estimate that in a typical area the size of the HDF, the cosmic variance of highly clustered sources is a factor of $\\sim 2$.  We also note that the higher HST counts could indicate that the pure luminosity evolution model is not sufficient at the faintest UV fluxes, and perhaps number density evolution is required as well.  \\cite{Gardner 2000 GBF}\\ report that the FUV number counts measured from the STIS subset of the data are surprisingly flat compared to the predicted counts \\citep{Granato 2000}.  We see no significant change in the slope with the addition of the ACS data.  Finally, we also examined the HST data to determine whether the difference in spatial resolution between GALEX and HST could result in source confusion.  There is no evidence of confusion in the HDF at the depth of the \\cite{Xu 2005}\\ counts, $FUV_{AB}<24$.  However, at the depth of the GALEX ultradeep surveys, $FUV_{AB} \\sim 26$, some individual sources would be confused.  The 34 sources brighter than 26 in the ACS area would correspond to $\\sim 20$\\ beams per source at the GALEX resolution.      \\subsection{Star Formation Rates and Comparison to Other Wavelengths} \\label{sec: SFR}  The detection of HDF-N sources in the FUV provides a sample of starbursts, and other star-forming galaxies, out to redshifts near unity.  We can compare their FUV properties to the extensive data available in the field outside of the HST wavelength range.  Strong starbursts should also be mid-infrared (MIR) bright galaxies. UV light absorbed by dust is re-radiated in the far-IR; and heated dust grains themselves, both small grains and polycyclic aromatic hydrocarbons (PAHs), emit in the mid-IR.  The ``Great Observatories Origins Deep Survey'' (GOODS; Dickinson et al.\\ 2005, in prep.) Spitzer Legacy Program has obtained ultra-deep observations of the field with the Infrared Array Camera \\citep[IRAC;][]{Fazio 2004} and the 24 $\\mu$m array of the far-IR photometer \\citep[MIPS][]{Rieke 2004}.  Only 56 of the FUV sources are detected by IRAC and 18 by MIPS at $\\ge 2\\sigma$.  For comparison, we examine the IRAC catalog for the {\\it Chandra}\\ Deep Field South (Dickinson et al.\\ 2006, in preparation) with the publically available GALEX catalog\\footnote{http://galex.stsci.edu/GR1}.  We find that the average FUV-IRAC1 color is 2.1 magnitudes.  The HDF FUV image reaches $AB \\ge 29$ while the IRAC channel 1 image reaches $AB \\sim 25$\\ (completeness limit due to confusion), so UV-luminous objects with typical colors will be more easily detected in the FUV.  The MIR luminosity of local galaxies in the IRAS bright galaxy sample \\citep{Soifer 1987}\\ has been found to correlate strongly with their far-infrared luminosity which is dominated by large, cool dust grains \\citep{Chary and Elbaz 2001}.  This correlation has been applied to develop a library of model templates of the mid- and far-infrared SED of galaxies.  The library consists of template SEDs across a range of luminosities, which can be redshifted to predict the MIR flux of a source with given luminosity at a redshift of interest. For each source in the FUV sample, we select the template for which the library predicts the closest 24 \\mic\\ flux density at the appropriate redshift to apply a bolometric correction; we do not use the shorter wavelength IRAC measurements.  The corrections based on the \\cite{Chary and Elbaz 2001}\\ and \\cite{Dale 2002}\\ template are used to derive an infrared luminosity (\\lir$=8-1000$\\ \\mic).  This technique of deriving the bolometric luminosity from the rest-frame MIR luminosity is shown to be accurate to 40\\%\\ in the local Universe \\citep{Chary and Elbaz 2001}.  The difference between the derived \\lir\\ from the two templates is assumed to be representative of the systematic uncertainty in the bolometric correction.  Statistical uncertainties are assumed to correspond to the signal to noise ratio of the source at 24 \\mic.  The validity of the mid- to far-infrared correlation and the one-to-one correlation between the bolometric correction and the MIR luminosity has been tested for field galaxy samples out to $z\\sim 1$\\ \\citep{Appleton 2004, Marcillac 2005}.  The MIPS detected sources have a median luminosity of $\\sim 10^{10}~L_{\\odot}$ with four sources falling in the class of luminous infrared galaxies (LIRGs; \\lir$\\simgt 10^{11}~ L_{\\odot}$).  Figure \\ref{fig: lirz}\\ shows the inferred luminosity of FUV sources detected by MIPS.  \\cite{Meurer 1999}\\ find a correlation between the the UV slope $\\beta$ and ratio of infrared luminosity to UV luminosity for starburst galaxies, where $f_{\\lambda} \\propto \\lambda^{\\beta}$. We measure this slope by combining the FUV data with WFPC2 photometry in order to measure the slope of the UV continuum from the available data.  We estimate the value of $\\beta$\\ following the technique used by \\cite{Meurer 1999}. We begin with 17 spectra from the catalog of \\cite{Kinney 1993}, spanning a range in $\\beta$\\ as measured by Meurer et al.  After ``redshifting'' each spectrum from $z=0$\\ to $z=0.85$\\ in steps of $\\delta z=0.1$, we integrate under the filter transmission curves for the ACS F150LP, and the WFPC2 F300W and F450W at each redshift.  For each redshift, this allows us to define a linear relation between the color of the object and its intrinsic FUV slope, yielding a function for $\\beta(z, \\mbox{color})$.  At low redshift we use the F150LP-F300W color.  At $z>0.4$, the F150LP filter contains little information redward of restframe 1000 \\AA, so we use the F300W--F450W color.  We estimate the error in $\\beta$\\ to be $\\sim 30$\\%\\ from the photometric uncertainty combined with the formal error in the linear fit to the template values.  In the redshift range $0.2<z<0.4$\\ the F150LP filter includes the \\lya\\ line, which can strongly affect the estimate of $\\beta$; furthermore, at these redshifts the FUV filter is stronly affected by the flattening of the $<1200$\\ \\AA\\ continuum \\cite{Buat 2002}.  Thus, we will consider values for objects at that redshift to be 50\\%\\ more uncertain; there are only 6 objects with IRAC counterparts in that range.  We obtain $\\beta$\\ values with a median of $-1.2\\pm 0.6$, with no clear trend in redshift or FUV magnitude.  \\cite{Schiminovich 2005}\\ measure the UV slope for nearly one thousand galaxies out to $z\\sim 1$\\ with $FUV_{\\mbox{AB}}<24$, using GALEX photometry and VVDS redshifts.  They report a median slope $\\beta_{GLX}=-1.44\\pm 1$\\ for sources where confusion is not an issue, in good agreement with other estimates \\citep{Treyer 2005,Adelberger and Steidel 2000}.  For each source, we also estimate the FUV luminosity ($\\nu L_{\\nu}$) at rest-frame 1600 \\AA, using the derived value of $\\beta$\\ and the flux density in either F150LP or F300W, whichever is closer in the rest frame.  In Figure \\ref{fig: LIR}, we plot the ratio of inferred \\lir\\ to $L_{\\mbox{FUV}}$ (hereafter $IRX$) against the slope of the UV continuum and indicate the inferred \\lir\\ of the sources.  We find that the relationship of \\cite{Meurer 1999}\\ is generally reproduced for the more luminous objects.  Less luminous sources tend to fall below the line, a trend already noted in GALEX results by \\cite{Seibert 2005}.  In the figure, we also plot the $\\beta$\\ relationship measured by \\cite{Cortese 2005}\\ for normal star-forming galaxies which are less luminous than the Kinney et al.\\ starbursts. The less luminous objects in our sample approximately follow the normal galaxy relation.  Alternately, these sources may have a contribution from older stars to the filters used in the fit, causing their $\\beta$\\ values to appear redder.  This latter explanation may be less likely, as our sample is FUV-selected, while the Coresse et al.\\ sample is not.  \\cite{Burgarella 2005}\\ find that the influence of older stars to the UV colors of galaxies in a UV-selected sample is small.  At the other extreme, infrared luminous sources have been observed to generally have higher $IRX$\\ for a given $beta$\\ than the starburst galaxies in the Meurer et al.\\ sample \\citep{Goldader 2002}. The brightest LIRG in our sample does lie slightly above the line.    \\subsubsection{Star Formation Rates}   FUV imaging provides a powerful tool for measuring the star formation in normal galaxies, but is strongly affected by extinction. We can derive the star formation rate (SFR) for each galaxy from the detected FUV flux.  \\cite{Kennicutt 1998}\\  gives the calibration from the 1500 \\AA\\ continuum to the SFR:  \\begin{displaymath} \\rm{SFR} (M_{\\odot}~\\rm{yr}^{-1}) = 1.4\\times 10^{-28}L_{\\rm{FUV}}~(\\rm{ergs}~\\rm{s}^{-1}\\rm{Hz}^{-1}), \\end{displaymath}  \\noindent assuming a Salpeter IMF \\citep{Salpeter 1955} with mass limits of 0.1-100 \\Msun and continuous star formation.  The \\cite{Meurer 1999}\\ IRX relation suggests a calibration for the dereddening factor of $A_{\\rm{FUV}} = 4.43 + 1.99(\\beta)$, expressed in magnitudes.  For the FUV-detected sources, we obtain a median extinction factor, as a multiple rather than in magnitudes, of $\\sim 6$.  The relationship assumes, however, that the UV flux is entirely the product of young stars.  Our sample likely includes sources with only moderate star formation, ordinary spirals and even elliptical galaxies, so the flux within the filter wavelength range (particularly in the WFPC2 filters) may include some contribution from an older (or aging) stellar population.  The star formation rates (SFRs) that we infer will thus be upper limits.  In Figure \\ref{fig: SFR}, we show the inferred star formation rates for FUV sources as a function of redshift and morphology.    \\subsubsection{X-ray Properties}   The Chandra 2 Ms catalog \\citep{Alexander 2003}\\ contains 13 sources within the FUV survey area with spectroscopic redshifts of $z<0.85$\\ from the catalog of \\cite{Barger 2003}.  Six optical counterparts to these X-ray sources are detected in the FUV.  These objects lie at redshifts 0.089, 0.139, 0.475, 0.556, and 0.752.  The object at 0.321 is described by \\cite{Barger 2003} as a possible multiple structure contaminated by a foreground object; the optical/FUV counterpart is more than an arcsecond from the X-ray position. Three of the FUV counterparts are spatially extended, and the other three are extremely faint in the FUV but extended in the F450W filter. None of them are detected in the hard band, none have broad optical emission lines \\citep{Barger 2003}\\ and all are near the detection limit of the softband (0.5-2 keV), with fluxes $SB \\le 0.08$\\ ergs cm$^{-2}$ s$^{-1}$.  These properties are consistent with the interpretation that the source of the X-rays is star formation rather than active galactic nuclei.  Similarly, the \\cite{Ranalli 2003}\\ calibration of X-ray luminosity as a star formation indicator yields rates generally in agreement (within a factor of a few) with the UV-inferred SFR.     \\subsection{Morphological Distributions} \\label{sec: morph}  FUV imaging picks out the location of the most recent star formation. Photometry tells us the total SFR; morphology tells us where it occurs within the galaxies.  As a result, the appearance of a galaxy can vary dramatically in different passbands even in the absence of dust, \\citep[a ``morphological k-correction'', see ][and the references therein]{Papovich 2003}.  In Figure \\ref{fig: morph kcorr}, we compare the morphologies of selected UV-bright galaxies in the FUV.  As expected, some galaxies appear similar across wavelengths while others show substantial differences.  Truly irregular or morphologically disturbed galaxies tend to appear similar across wavelength, as do some elliptical galaxies \\citep[see counter examples in][]{Windhorst 2002}.  The morphological K-correction is most pronounced for early to mid-type spirals, in which a substantial population of old stars defines the optical shape, but regions of recent star formation are ``lit up'' across the galaxy \\citep{Windhorst 2002}.   We avoid the K-correction by matching the FUV detections to a morphological catalog of galaxy types \\citep{Conselice 2005}\\ and CAS parameters \\citep[see ][]{Conselice 2003 CAS}. The catalog includes galaxy morphologies in the rest-frame B-band for the 200 HDF galaxies out to $z=0.85$\\ that are bright enough for visual classification (out of 240 possible).  We restrict our analysis to $z < 0.85$ because the redshifting Lyman break leaves little flux within the F150LP filter at higher redshift.  We exclude 10 sources that are not in the \\cite{Conselice 2005} catalog, mostly due to the lack of NICMOS counterparts.   This gives us some indication of the types of galaxies that are emitting in the FUV at redshifts $z < 1$. We find that the galaxies detected in the FUV span all the major morphological types, as also seen by \\cite{de Mello 2004}.  Figure \\ref{fig: zhist morph}\\ shows the morphological break down for systems based on their apparent morphological types as classified by \\cite{Conselice 2005}.  As can be seen the spiral galaxies dominate the number counts for the FUV sources, although spheroids make up a significant fraction of the detections at $z > 0.6$ and irregulars are also represented. We find, in fact, that a significant fraction of all spheroids (30/56) and a similar fraction of spirals (50/87) at $z < 0.85$\\ are detected in the FUV in the HDF-N.  The relative distributions of FUV emitting types with redshifts can be seen in Figure \\ref{fig: fuv mb}, which plots the absolute M$_{\\rm B}$ magnitude as a function of redshift.  From this diagram, there is a broad range of absolute magnitudes for the FUV sources at all redshifts.  The median luminosity for UV-detected spheroids, spirals, and irregulars is $-17.3$, $-18.2$, and $-16.2$, respectively, compared to median values of $-17.7$, $-17.8$, and $-16.4$\\ for all galaxies of the three types at $z<0.85$\\ in the HDF.  Interestingly, while the median values agree, the most luminous ellipticals in the HDF (M$_{\\rm B}<-19$) are not generally detected in the FUV (3 out of 11).  The figure shows that less than half of peculiar/irregular at $z<0.85$\\ are detected in the FUV; only 38\\% (18/47) have FUV detections. This is due in large part to their intrinsic faintness, rather than unusually red color.  Half of the irregulars have V-band magnitudes fainter than 27, which makes them undetectable in some or all areas of the FUV image. The FUV-detected irregulars are somewhat bluer than the rest of the UV sample, with a median value of $FUV_{AB}-V_{AB}\\sim 1.3$\\ compared to a median color of 1.5 for the entire FUV catalog.  These same objects are optically blue, with a median $V-I\\sim 0.5$\\ compared to 0.6 for the all $z<0.85$\\ sources in the HDF.   \\subsection{Star formation in Spheroids}  We find evidence for star formation in $\\sim 50$\\%\\ of spheroids at $z<0.85$.  These objects are typically less massive than $10^{10}$\\ \\Msun\\ (see Figure \\ref{fig: massdisp}) and less luminous than $M_{B}=-19$.  Their sizes (half light radii) are similar to other spheroids in the HDF.  We find a median SFR of 0.25 $M_{\\odot}$\\ yr$^{-1}$, after extinction correction (see Section \\ref{sec: SFR}).  So far, we have only considered morphologically-selected spheroids. However, few of these objects have the SED of purely old stellar populations, even without including the FUV.  \\cite{Stanford 2004}\\ show that morphologically- and spectroscopically-identified spheroids in the HDF are not necessarily the same population.  A morphological selection identifies sources which have SEDs similar to local spheroids and identifies additional sources that are bluer, less massive, and less luminous than those.  Only one of the objects in the \\cite{Stanford 2004}\\ sample of spectroscopic ellipticals is detected in the FUV.  Their sample of morphologically selected ellipticals is not identical to that of \\cite{Conselice 2005}, but the fraction of FUV detections is similar.  The FUV detection, then, supports the \\cite{Stanford 2004}\\ conclusion that some morphological spheroids have recent or ongoing star formation.  This effect has also been seen in the ``blue-core'' ellipticals \\citep{Menanteau 2001}.  These objects were initially identified by strong color gradients in the WFPC2 images, which show significant bluing towards the center.  The presence of the blue cores suggested a population of young ($<1$Gyr) ellipticals which may have undergone recent merger activity or some type of residual star formation. Ten of the 21 sources at $z<0.85$\\ in the Menanteau et al.\\ sample are detected in the FUV.  In Figure \\ref{fig: blue-core}, we show that the sources with the bluest cores are the ones most likely to be detected in the FUV.  The detection of FUV flux near the core of the sources confirms that these objects have small amounts of ongoing star formation.  It is highly unlikely that the FUV flux detected in spheroids is the result of the ``UV upturn'' that arises from a minority population of hot horizontal branch stars \\cite[e.g.][]{Brown 1997}.  For example, cluster ellipticals at $z= 0.33$\\ and $z=0.55$\\ have been observed to show small amounts of UV emission \\citep{Brown 2000b, Brown 2003}.  Unlike the HDF spheroids, these objects have optical SEDs broadly consistent with old stellar populations.  The UV emission in UV-upturn galaxies is a small fraction of the total luminosity. \\cite{Brown 2003}\\ find $m_{1500}-V\\sim 4$\\ for UV-upturn galaxies, while the median color of the HDF spheroids is $\\sim 2$\\ mag.\\ and only one of them has a color greater than 3.  The differences in the STIS and ACS filters could account for ACS objects being $\\sim 0.4$\\ magnitudes brighter than their STIS counterparts at $z>0.2$, but the HDF spheroids are still significantly brighter in the UV.  Similarly, \\cite{Brown 2003}\\ estimate that the flux associated with the UV upturn would correspond to $SFR \\sim 0.005-.02$\\ \\Myr\\ if it arose instead from star formation.  The inferred SFR in HDF spheroids is a factor of several higher even without extinction correction, and significantly higher with the correction.  However, we note that an old stellar contamination of the redder filters could result in an overestimate of the SFR.  Thus it is likely, though not certain, that most of the UV flux in HDF spheroids is the result of star formation and not the UV upturn.  The more massive and luminous ellipticals in the HDF appear not to be forming stars at rates similar to the smaller and fainter ones that we detect in the FUV.  This could be a direct indication that lower luminosity ellipticals in the field form later than the giant ellipticals.  This is consistent with the widely varying ages measured for local ellipticals \\citep{Trager 2000}.  Such a distinction may be evidence of downsizing in the galaxy formation process, which may be directly related to the rate of merging which is seen to be high for lower luminosity and lower mass galaxies at $z < 1$ \\citep{Conselice 2003}.  Another way to investigate the star-formation nature of early type galaxies is by examining their location in the concentration-asymmetry diagram for galaxies at $z < 1$ (Figure \\ref{fig: CAS}).  The most evolved spheroids which have had no star formation in the recent past, should contain a high concentration and a low asymmetry.  These objects typically do not have FUV emission. On the other hand, morphologically-identifiable spheroids with high asymmetries, that indicate a recent evolution, are more likely than not to have FUV emission.  This result supports the conclusion that the FUV emission is originating from star formation, which produces the structural asymmetries. Furthermore, many of the FUV-detected spheroids have relatively low concentrations, consistent with their morphology tracing the regions of young stars as well as the underlying older population \\citep[e.g.][]{Windhorst 2002}.  The inferred SFR in HDF spheroids is not high enough for them to be the progenitors of local giant ellipticals, but it may suggest that a significant fraction of the stars in lower luminosity and lower mass ellipticals form at $z < 1$.  There is no evidence that the FUV detected spheroids are undergoing the last years of a final episode of star formation.  Instead, we might estimate a duty cycle of star formation episodes.  The detection fraction ($\\sim 50$\\%) suggests that these objects could spend as much as half the time producing small amounts of stars.  With a median SFR of 0.3 \\Myr, this duty cycle would allow as much as $\\sim 10^9$\\ \\Msun\\ to form between $z=0.85$\\ and $z=0$.  The median stellar mass estimated by \\cite{Conselice 2005}\\ for the spheroids is $\\sim 10^{8.5}$\\ \\Msun, so they could double or triple in size by present day.  They would still remain much less massive than the spectroscopically-identified ellipticals, which typically have stellar masses greater than $10^{10}$\\ \\Msun\\ \\citep{Stanford 2004}."
    },
    "0606/astro-ph0606591_arXiv.txt": {
        "abstract": "We compare the properties of galaxy groups extracted from the Updated Zwicky Catalogue (UZC) with those of groups extracted from $N$-body simulations of the local Universe, in a $\\Lambda$CDM and a $\\tau$CDM cosmology. In the simulations, the initial conditions of the dark matter density field are set to reproduce the present time distribution of the galaxies within $80 h^{-1}$ Mpc from the Milky Way. These initial conditions minimize the uncertainty originated by cosmic variance, which has affected previous analyses of this small volume of the Universe. The simulations also model the evolution of the photometric properties of the galaxy population with semi-analytic prescriptions. The models yield a galaxy luminosity function sensibly different from that of the UZC and are unable to reproduce the distribution of groups and their luminosity content. The discrepancy between the model and the UZC reduces substantially, if we redistribute the luminosity among the galaxies in the simulation according to the UZC luminosity function while preserving the galaxy luminosity rank. The modified $\\Lambda$CDM model provides the best match to the UZC: the abundances of groups by harmonic radius, velocity dispersion, mass and luminosity are consistent with observations. We find that this model also reproduces the halo occupation number of groups and clusters. However, the large-scale distribution of groups is marginally consistent with the UZC and the redshift-space correlation function of galaxies on scales larger than $6 h^{-1}$ Mpc is still more than 3-$\\sigma$ smaller than observed. We conclude that reproducing the properties of the observed groups certainly requires a more sophisticated treatment of galaxy formation, and possibly an improvement of the dark matter model. ",
        "introduction": "The formation and evolution of galaxies is one of the major challenges of cosmology. Within the framework of the hierarchical structure formation by gravitational instability, galaxies properties are strongly affected by their environment. Most galaxies reside in groups and much work on compact (\\citealt{hickson97} and references therein; \\citealt{barton03}; \\citealt{mendes03}; \\citealt{lee04}; \\citealt{kelm04}; \\citealt*{tovmassian06}) and loose groups (\\citealt{postman84}; \\citealt{allington93}; \\citealt{tran01}; \\citealt{martinez02a},b; \\citealt{tanvuia03}; \\citealt{girardi03}; \\citealt{balogh04}; \\citealt{weinmann05}; \\citealt*{zandivarez06}; \\citealt{martinez06}; \\citealt{braugh06}) shows that these systems can be more effective than clusters at shaping the properties of galaxies, namely morphology, luminosity, color, star formation rate and cold gas content (see also \\citealt{kodama01}; \\citealt{ellingson04}; \\citealt{verheijen04}; \\citealt{tanaka05}).  Galaxy groups trace the large-scale distribution of galaxies (e.g. \\citealt{berlind06}), and thus describe the distribution of the optical light from sub-megaparsec scales to the largest scale probed. Models of galaxy formation should be able to reproduce the properties of galaxy groups. For an appropriate comparison of models with observations, however, the surveyed volume of the real groups must be sufficiently large to be representative of the Universe. Over the last years, the compilation of deep galaxy redshift surveys has provided galaxy group catalogues with an increasing number of systems (\\citealt{ramella99}; \\citealt{eke04}; \\citealt{berlind06} and references therein). These catalogues are deep enough to average out the large-scale structure fluctuations and obtain a robust estimate of the light distribution.  In current models, large-scale structure forms by the gravitational collapse of dark matter, which dominates the dynamics of galaxy systems. Comparing models with data therefore requires a valuable treatment of the dynamics of the baryonic matter within the dark matter halos. \\citet{diaferio1999} were the first to attempt to compare real redshift surveys with mock surveys where galaxies were formed and evolved with semi-analytic prescriptions applied to merger trees of dark matter halos extracted from $N$-body simulations.  They compared the model with the northern sector of the Updated Zwicky Catalog (UZC) redshift survey \\citep{falco1999}. This comparison resulted in a general agreement between the now standard Cold Dark Matter model with a non-zero cosmological constant ($\\Lambda$CDM) and observations, although the mean luminosity of the simulated groups was smaller than observed. A further  topological analysis of the mock redshift surveys and the UZC quantified how the large-scale structure was much fuzzier in the mock catalogues than in the UZC \\citep{schmalzing2000}. Both disagreements were tentatively attributed to cosmic variance, responsible for the lack, in the mock surveys, of a coherent structure as thin and wide as the Great Wall. The UZC is only $150 h^{-1}$ Mpc deep and indeed its volume is not large enough to suppress the fluctuations of the large-scale structure. Therefore a fair comparison of the model with the UZC requires a not-random choice of the volume in the simulations.  \\citet{mathis2002} perform $N$-body simulations where the initial conditions are constrained to reproduce the large-scale distribution of galaxies observed in the IRAS 1.2 Jy survey \\citep{fisher1995} within $80 h^{-1}$ Mpc of the Milky Way. In their simulations, \\citet{mathis2002} also include semi-analytic recipes to form and evolve galaxies. Thus, these simulations seem appropriate to test whether the current galaxy formation models yield the observed clustering and luminosity properties of groups once the large scale distribution of galaxies faithfully mirror the real one. This paper is devoted to this test.  Sect. 2 summarizes the properties of the constrained simulations; sects. 3 and 4 describe our real and mock galaxy redshift surveys and our group catalogues in redshift space. In sect. 5 we investigate the large-scale distribution of groups and in sect. 6 we discuss the distribution of light among and within groups. We conclude in sect. 7. ",
        "conclusions": "We have compared the properties of groups and clusters of galaxies extracted from the UZC \\citep{falco1999} with those of systems extracted from mock redshift surveys. We have compiled these mock surveys from $N$-body simulations constrained to reproduce the large-scale distribution of galaxies in the nearby Universe \\citep{mathis2002}. In the simulations, galaxies are formed and evolved according to a semi-analytic procedure and we are thus able to constrain both the group clustering and the group luminosity content.  By using simulations with constrained initial conditions, we minimize the possible role played by cosmic variance, and differences between real and mock catalogues should mostly originate from the galaxy formation recipes. Our approach thus differs from previous attempts of comparing simulations with real group catalogues, where either cosmic variance was an issue \\citep{diaferio1999} or the real catalog was large enough that cosmic variance was naturally suppressed (\\citealt{eke04}; \\citealt{berlind06}).  The gross large-scale distribution of galaxies, including the location of the major nearby clusters, is very similar in the mock and real surveys, as mostly imposed by the initial conditions. However, the simulated large-scale structures are not as sharply defined as in the UZC, confirming earlier results of mock surveys extracted from unconstrained simulations \\citep{schmalzing2000}. This disagreement amplifies substantially when we consider the properties and the large-scale distribution of groups.  The group-finder algorithm strongly depends on the galaxy luminosity function. The numerical recipes adopted to form and evolve galaxies in the simulations provide a galaxy luminosity function which is considerably different from that of the UZC. We test that this difference is the major responsible for the model inability at reproducing the group properties. We assign new luminosities to the simulated galaxies according to the luminosity function of the UZC, while preserving the luminosity rank predicted by the model. We thus have new mock catalogues where galaxy luminosities distribute according to the UZC luminosity function (UZCLF), besides the catalogues where galaxy luminosities distribute according to the semi-analytic luminosity function (SALF).  Unlike the groups extracted from the SALF and the $\\tau$CDM-UZCLF mock catalogues, the groups extracted from the $\\Lambda$CDM-UZCLF catalogue have statistical properties in satisfactory agreement with observations: group abundances by luminosity, harmonic radius, velocity dispersion and mass are generally within 3-$\\sigma$ errors, or less, from the UZC group abundances. These groups also reproduce the UZC relations between the group mass-to-light ratio and mass, and the galaxy number and mass, namely the halo occupation number. Finally, the $\\Lambda$CDM-UZCLF groups, similarly to the UZC groups and unlike the SALF groups, show a weak correlation between the luminosity difference between the two brightest galaxies in a group and the total group luminosity. This result indicates that the two semi-analytic prescriptions of allowing merging of satellite galaxies only with the central galaxy of the dark matter halo and gas cooling only onto this same central galaxy produces a too large luminosity difference between the two brightest galaxies in a group.  The success in the statistical properties of groups, obtained by adopting the observed luminosity function, is not shared by  the large-scale distribution of groups. Specifically, the simulated groups in the North Galactic Cap of the $\\Lambda$CDM-UZCLF catalogue trace the large-scale distribution of galaxies at significance levels as much as seven times smaller than those of the UZC. This is a consequence of the looser large-scale structures in the models: in fact, the redshift-space correlation function of galaxies is more than 3-$\\sigma$ below the observations on scales larger than $6 h^{-1}$ Mpc (Figure \\ref{xigg}; \\citealt{mathis2002}) and the number of groups is 25\\% smaller than in the UZC (Table \\ref{groups}).  We therefore conclude that (1) the semi-analytic recipes used in \\citet{mathis2002} have seriuos difficulties at distributing the luminosity among galaxies correctly; and (2) even the most successful $\\Lambda$CDM model does not yield cosmic structures as coherent as observed. Whereas the first problem might be solved, in principle, with more sophisticated semi-analytic modellings (e.g. \\citealt{croton06}; \\citealt{weinmann06b}) the second problem appears to be more fundamental, although the disagreement with observation is less dramatic: in this case, both different dark matter models and an appropriate treatment of the interplay between gas dynamics and dark matter might be necessary to reconcile the model with group properties."
    },
    "0606/astro-ph0606072_arXiv.txt": {
        "abstract": " ",
        "introduction": "Strong collisionless shocks are present in various astrophysical objects, and under a wide range of conditions. These shocks effectively heat the gas and are also believed to accelerate a fraction of particles up to very high energies. Analyses based on multi-frequency observations allow one to determine properties of both thermal and nonthermal component. In the case of supernova remnants (hereafter SNRs), optical and UV emission are typically thermal, radio is nonthermal, while thermal and nonthermal emission may coexist in X-rays. The momentum distribution of the accelerated particles is locally well approximated by a power law; this can be inferred from the power-law synchrotron spectra, in the case of electrons; while in the case of ions it can be measured directly in the energy distribution of cosmic rays (if one accepts the ``SNR Paradigm'' for the origin of galactic cosmic rays).  Diffusive acceleration (Fermi acceleration) is believed to be the dominant process that allows particles to gain energies in excess of typically thermal values. The standard theory of diffusive acceleration, in test-particle approximation (see e.g.\\ the review of Jones \\& Ellison \\cite{Jones-Ell-rev-91}), shows that a power-law distribution develops at high energies. The spectral index of this high-energy population depends on the shock compression ratio, while it does not depend at all on the original energy distribution of the injected particles. In fact, a simplified way to describe the overall particle evolution, from nearly thermal to very high velocities, is to treat particle {\\it injection} and {\\it acceleration} as two separate problems. The {\\it injection} problem consists in finding out the initial momentum distribution of that fraction of (originally thermal) particles that can enter the acceleration process, i.e. to make at least one acceleration cycle. The {\\it acceleration} problem, instead, consists in following the evolution of the distribution of these particles along all next acceleration cycles.  In order to keep the number of free parameters low, while modelling more effectively the emission in all observed spectral ranges, one needs to introduce a physically self-consistent scenario for the thermal and nonthermal populations. For instance, the standard model of particle acceleration constrains the slope of the electron distribution at high velocities, but does not predict its normalization: in other terms, the injection efficiency (i.e.\\ the fraction of particles that enter the acceleration process) is poorly known, because this process is sensitive to physical details not included in the standard model of Fermi acceleration. The level of electron-ion equilibration or, alternatively, the electron temperature is another key quantity hard to determine ``a priori'' in collisionless shocks. While in models of SNR shocks injection and equilibration efficiencies are taken as independent free parameters, in the reality both depend on the physical conditions within the shock transition, and therefore they are not independent. Goal of this paper is in fact to investigate, in nonrelativistic SNR shocks, a possible connection between electron injection and thermal equilibration.  A self-consistent treatment of injection and acceleration must include a microphysical model of particle-wave interactions in the plasma. A few physical processes have been proposed to account for the electron injection (see Malkov \\& Drury \\cite{Malkov-Drury-2001} for a review). The scattering of electrons is suggested to be due to some ion-generated instabilities (Bykov \\& Uvarov \\cite{Byk-Uv-99}; hereafter BU99), whistler waves \\cite{Levinson-92} and lower-hybrid waves from ions \\cite{Leroy-et-al-82}. These models mainly address the plasma microphysics; while only BU99, to our knowledge, are able to model the formation of the post-shock electron distribution.  In this paper, we approach the problem in a simplified way. We assume the presence of scattering centres, without concentrating on their nature, but simply assuming that they match the following requirements: i) the interaction with these scattering centres generates a nearly isotropic, Maxwellian velocity distribution of particles on timescales not longer than one collision time; ii) the timescales for (wave-mediated) isotropization and energy exchange between electrons are both smaller than the (wave-mediated) electron-ion equilibration time; iii) the scattering centres play at the same time the role of thermalising, within the shock, the incoming particle population to the post-shock temperature and that of driving the process of diffusive acceleration. When we will need to use more specific properties of the wave-particle interaction (in Sect.~3.2), we will refer to the results of BU99 on the electron kinetics.  There are in general three ways to calculate the post-shock momentum distribution of particles: either by solving the kinetic equations, or by making a hybrid simulation (which is however unable to model the momentum distribution of electrons because electrons are treated as a fluid), or finally by extending the individual particle approach of Bell \\cite{Bell1978a}. He has estimated the probability for a particle to {\\em return} to the shock from downstream and has shown that in this way one obtains a power-law distribution for the accelerated particles with velocities $v\\gg\\Vs$ (where $\\Vs$ is the shock velocity).  In the present paper we propose to extend the Bell approach to the problem of injection by introducing the probability to {\\em recross} the shock from downstream to upstream. This probability is connected to the process of thermalisation of the incoming flow within the shock. This fact has been shown by Malkov \\cite{malkov-98} in the case of protons. Namely, the idea that ions are prevented from backstreaming by the self-generated waves (which also participate in thermalisation of ions) has allowed Malkov to obtain an analytic solution of the injection problem for protons. The main point in his thermal leakage theory is that only those protons that can ``leak'' upstream are injected into the Fermi process.  We use the same idea in our approach to electrons, even though we are not tight to any specific kind of interaction. We consider only the case of parallel shocks, namely when the ambient magnetic field is parallel to the shock normal.  The plan of the present paper is as follows. In Sect.~\\ref{sect2} we deal with the injection problem by developing an analytic treatment for calculating the injection efficiency as well as for determining of the distribution of particles which are able to be accelerated. Sect.~\\ref{sect5} deals with the process of thermalisation and its influence on injection and, using results from BU99 model, give quantitative estimations on the injection efficiency. Sect.~\\ref{sect4} concludes. ",
        "conclusions": "\\label{sect4}  The electron injection and thermalisation are not independent processes. This is clearly outlined by Fig.~\\ref{fig-a} where the probabilities $\\Pret$ and $\\Pcro$ as a function of reduced momentum are shown for two values of $\\chis$ together with the initial distribution (\\ref{fini}) of injected particles. The hybrid electron distribution $\\nes\\fH(y)dy$, Maxwellian up to $\\yb$ and power-law above \\cite{Porq-01}, is also shown on the figure to see the differences. The break momentum $\\yb$ is given by the assumption that all injected particles obtain momenta higher than $\\yb$ after acceleration i.e. is defined by\\ $\\vsg=\\int_{\\yb}^\\infty\\fH(y)dy$.  It is a common believe that only particles from the energetic tail of Maxwellian distribution are capable to be accelerated. On the contrary, the distribution $\\finj$ shows that thermal particles with velocities $v>\\vmin$ have the possibility to participate in acceleration process, although with different probability. The minimum velocity $\\vmin=0.07\\left(\\chis/{0.03}\\right)^{-1/2}\\vth$ may be considerably less than the thermal velocity (see Eq.~(\\ref{ymin})). The most probable velocity $v\\rs{*}$ at which the maximum of the distribution $f\\rs{inj}$ occurs is $v\\rs{*}\\approx (2\\div3)v\\rs{th}$ for a wide range of injection fractions $\\vsg=10^{-3}\\div10^{-6}$ (Fig.~\\ref{fig-4} and compare with Fig.~\\ref{fig-c}). In other words, most of the electrons are injected with velocities $v\\rs{*}\\simeq 50\\chis^{1/2}\\Vs$.  \\begin{figure} \\centering \\includegraphics{fig4.eps} \\caption{Velocity corresponding to the maximum of the distribution function $f\\rs{inj}$ versus $\\chis$ for a few choices of the shock Mach number (see curve labels). Model of electron kinetics is the same as in the previous figure. } \\label{fig-4} \\end{figure}  The injection efficiency $\\vsg$ of electrons in a collisionless shock is associated to the process of electron heating within the shock, through the competition of two effects. On one side, the higher the post-shock electron temperature, the higher the energy of thermal electrons and the higher the fraction of those which are ready to cross the shock from downstream to upstream (this is given by the probability $\\Pret$, lines 2 on Fig~\\ref{fig-a}a,b). On the other side, however, the higher the temperature, the higher the number of scattering centers. Electrons traversing the shock from downstream to upstream also interact with these sites and the more such interactions the less electrons are able to cross the shock and to enter into the Fermi acceleration loop (see probability $\\Pcro$, lines 3 on Fig~\\ref{fig-a}a,b). In this paper we show that, for a given ${\\cal M}$, the combined effect of these two processes is that the quantity $\\vsg$ decreases with increasing of $\\chis$ (Fig.~\\ref{fig-c}).  Both injection and thermalisation are sensitive to the Mach number. It is shown (see Sect.~\\ref{BUmodel}) that, for a standard range of parameters, $\\vsg(\\chis,{\\cal M})$ is a decreasing function of a single argument $\\vsg=\\vsg({\\cal M}^2\\chis)$. Theoretical models show that in high-velocity shocks the energy of the shock is transferred to the thermal electrons less efficiently, so that $\\chis\\propto{\\cal M}^{-m}$ with $0\\leq m\\leq 2$ (BU99). Observations favour a dependence $\\chis\\propto{\\cal M}^{-1}$, suggesting a Bohm-like type of diffusion. Our calculations show that the level of electron-ion equilibration is expected to depend on the injection fraction as well, so that the approximate relation between these three parameters is $\\chis\\propto {\\cal M}^{-2}\\vsg^{-1/q}$ (Sect.~\\ref{BUmodel}). The smaller the Mach number, the higher the level of electron-ion equilibration for a given injection efficiency. On the other hand, for a given thermalisation level, the stronger the shock the less particles can be injected.  To conclude, we would like to review the assumptions used in the present paper.  Our approach is in test-particle approximation. Actually, it is known that, in young SNRs, shocks could be strongly modified; and the inclusion of nonlinear effects could change our results significantly. The usage of a test-particle approach in the present paper is however consistent with what done by BU99, and the results of their analysis are valid at least for shocks with Alfv\\'en Mach number less than $\\simeq 43$. Thus our results should be applicable at least to SNRs either in the late adiabatic phase or beyond. Our opinion is that, together with using nonlinear treatments, it is valuable investigating what happens in the linear (test-particle) approximation, also in consideration that a nonlinear theory has anyway to give, as limit case, the linear results.  Eqs.~(\\ref{iso-Maxwel}) and (\\ref{P1-a}) assume nearly isotropic distribution of particles. In general, this is not fully true for the thermal population right after the shock. In order to overcome this difficulty, in numerical calculations these formulae are assumed to apply a few mean free paths downstream, in order to insure that the distribution is isotropic in the local frame (e.g. \\cite{Ell-et-al-1990}). Within our approach, this implies some restrictions on the underlying physics. Our assumptions about properties of the scattering centers in our model (see Introduction) require that the timescale for isotropisation is not larger than the timescale for one interaction. This means that, in order to assume isotropy of particles velocities, we would in principle need to increase the number of interactions $N\\rs{c}$ at least by one (see Eq.~(\\ref{Pcro})). Since already $N\\rs{c}\\gg 1$ for, say, $\\chi\\rs{s}>10^{-3}$ (see estimations after Eq.~(\\ref{Pcro})), this increment would not change much our results, in the case of shocks producing an electron population thermalised up to the level $\\chi\\rs{s}$ higher than $10^{-3}$. Since $\\chis\\propto{\\cal M}^{-1}$, our models are limited again to the shocks with moderate Mach numbers.  {\\small {\\bf Acknowledgements} We are grateful to A.~Bykov for helpful discussion on implementation of their results as well as S.~Reynolds for valuable discussions. This work has been supported by MIUR under grant Cofin2001--02--10, by MIUR under grant Cofin2002 and by UFRF 02.07/00430 grant. OP acknowledges grant MIUR Cofin2001--02--10 for his fellowship.}"
    },
    "0606/astro-ph0606558_arXiv.txt": {
        "abstract": "{ We compare the number of detected 22 GHz H$_2$O masers in the Local Group galaxies M31, M33, NGC\\,6822, IC\\,10, IC\\,1613, DDO\\,187, GR8, NGC\\,185, and the Magellanic Clouds with the water maser population of the Milky Way. To accomplish this we searched for water maser emission in the two Local Group galaxies M33 and NGC\\,6822 using the Very Large Array (VLA) and incorporated results from previous studies. We observed 62 H{\\sc ii} regions in M33 and 36 regions with H$\\alpha$ emission in NGC\\,6822. Detection limits are 0.0015 and 0.0008 L$_\\odot$ for M33 and NGC\\,6822, respectively (corresponding to 47 and 50 mJy in three channels with 0.7 km~s$^{-1}$ width). M33 hosts three water masers above our detection limit, while in NGC\\,6822 no maser source was detected. We find that the water maser detection rates in the Local Group galaxies M31, M33, NGC\\,6822, IC\\,1613, DDO\\,187, GR8, NGC\\,185, and the Magellanic Clouds are consistent with expectations from the Galactic water masers if one considers the different star formation rates of the galaxies. However, the galaxy IC\\,10 exhibits an overabundance of masers, which may result from a compact central starburst. ",
        "introduction": "Galactic water maser emission, discovered by \\citeN{CheungRankTownes1969}, is mainly associated with star forming regions and late type stars. In star forming regions the masers are found in the vicinity of compact H{\\sc ii} regions (\\citeNP{GenzelDownes1977}). Observations of these maser sources trace extremely dense ($>10^7$cm$^{-3}$) and warm ($>400$ K) molecular gas and yield important information about physical and dynamical properties of their environments. They can also be used to measure accurate annual parallaxes and proper motions of star forming regions using Very Long Baseline Interferometry (VLBI) (e.g. \\citeNP{HachisukaBrunthalerReid2006}). While more than one thousand water masers are known in the Milky Way \\cite{ValdettaroPallaBrand2001}, only a small number are known in other Local Group galaxies.  The first extragalactic water maser was discovered by \\citeN{ChurchwellWitzelHuchtmeier1977} in IC\\,133, a star forming region in the Local Group galaxy M33. Further water masers in  Local Group galaxies were found towards the Magellanic Clouds (e.g.~\\citeNP{ScaliseBraz1981}) and IC\\,10 (e.g. \\citeNP{BeckerHenkelWilson1993}). Searches toward M31 \\cite{GreenhillHenkelBecker1995,ImaiIshiharaKameya2001} and other Local Group galaxies \\cite{HuchtmeierRichterWitzel1980,GreenhillMoranReid1990} have not yet yielded a detection. Beyond the Local Group, water masers are found mainly toward active galactic nuclei (see \\citeNP{ZhangHenkelKadler2006}, and \\citeNP{Kondratko2006} for a list, and http://www.cfa.harvard.edu/wmcp/surveys/survey.html for an updated list).  Proper motion measurements and accurate distances of Local Group galaxies are important for our understanding of the dynamics and evolution of the Local Group. With the NRAO\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.} Very Long Baseline Array (VLBA) \\citeN{BrunthalerReidFalcke2005} measured the proper motions of two regions of water maser activity on opposite sides of M33. The comparison of the relative proper motion between the two regions and the expected motion from the known rotation curve and inclination of M33 led to a determination of a rotational parallax (730 $\\pm$ 168 kiloparsecs) of this galaxy. This distance is consistent with recent Cepheid and tip of the red giant branch estimates (\\citeNP{LeeKimSarajedini2002}; \\citeNP{McConnachieIrwinFerguson2005}) and earlier distance estimates using the internal motions of water masers in IC\\,133 \\cite{GreenhillMoranReid1993,ArgonGreenhillMoran2004}. Since the proper motion measurements were made relative to a distant extragalactic background source, the proper motion of M33 itself could also be determined. This measured proper motion of M33 is a first important step towards a kinematical model of the Local Group and was used to constrain the proper motion of the Andromeda Galaxy M31 \\cite{LoebReidBrunthaler2005}.  The interpretation of the proper motions of water masers in Local Group galaxies may be limited by the unknown three dimensional peculiar motions of the star forming regions. In the Milky Way, peculiar motions of star forming regions can be 20 km~s$^{-1}$ as seen in W3(OH) (\\citeNP{XuReidZheng2006}). Hence it is important to use multiple regions of maser activity for proper motion studies to constrain the systematic effect of peculiar motions. This lead us to search for water vapor masers towards 62 H{\\sc ii} regions in M33 with the NRAO Very Large Array (VLA).  A detection of maser sources in other Local Group galaxies would allow us to extend such studies considerably. Since the strongest maser emission from high-mass star froming regions in our Galaxy is typically 100 times stronger than emission from late-type stars, we targeted NGC\\,6822, a northern Local Group member with prominent star formation activity. ",
        "conclusions": "We have searched for water maser emission with the VLA toward 62 and 36 H{\\sc ii} regions in M33 and NGC\\,6822, respectively. M33 hosts three water masers above our detection limit (all previously identified), while in NGC\\,6822 no maser source was detected. IC\\,133 and M33/V077 show velocity components that are displaced by several 10 km~s$^{-1}$ from previously observed maser features. We also provide an accurate position for M33/V077, while several previously reported marginal detections in M33 are not confirmed.  The number of known maser sources in the Local Group galaxies M31, M33, NGC\\,6822, SMC, LMC, IC\\,1613, DDO\\,187, NGC\\,185, and GR8 is consistent with the expected number of maser sources based on the Galactic water maser luminosity function and the current star formation rates. Only the galaxy IC\\,10 shows a much higher detection rate than expected.  Based on this analysis one can expect to find 17 and 5 water masers with flux densities above 1 mJy in M31 and NGC\\,6822, respectively, that could be used to measure the proper motions and geometric distances of these galaxies with future radio telescopes like the Square Kilometre Array (SKA).\\footnotemark[2]   \\footnotetext[2]{See http://www.skatelescope.org.}"
    },
    "0606/astro-ph0606302_arXiv.txt": {
        "abstract": "% {}% {To characterize the molecular environment of classical T Tauri stars in Lupus observed with the Spitzer Space Telescope and to search for gas-rich disks toward these sources.} {Submillimeter observations of $^{12}$CO, $^{13}$CO and C$^{18}$O in the $J$=3-2 and 2-1 lines were obtained with the James Clerk Maxwell Telescope toward a sample of 21 T Tauri stars with disks in the Lupus molecular clouds. Pointings at the sources and at selected off-positions are presented in order to disentangle material associated with disks from ambient cloud material.} {One source, IM Lup (Sz 82), has been found with double-peaked $^{12}$CO and $^{13}$CO profiles characteristic of a large rotating gas disk. The inclination of the disk is found to be $\\sim 20^\\circ$, with an outer radius of 400-700 AU. For most other sources, including GQ Lup with its substellar companion, the single-dish $^{12}$CO lines are dominated by extended cloud emission with a complex velocity structure.  No evidence for molecular outflows is found. Compact C$^{18}$O emission due to dense circumstellar material has been detected toward only two sources. Future searches for gas-rich disks in Lupus should either use interferometers or perform very deep single-dish integrations in dense gas tracers to separate the cloud and disk emission.} {} ",
        "introduction": "A general characteristic of pre-main sequence (PMS) stars is the presence of circumstellar disks, which have been observed for a wide range of evolutionary stages associated with low-mass star formation \\citep[e.g.,][]{Lada87,Andre93,Greaves05}. These disks contain all the necessary ingredients for forming complex planetary systems similar to our own Solar system, now also seen around other stars \\citep[e.g.,][]{Ruden99,Marcy95}.  The gas-rich phase of disks is crucial in the evolution, since gas is needed to build Jovian-type planets. Also, the presence of gas affects the dynamics and evolution of the dust in the system and the migration of any planets embedded in the disk \\citep[e.g.,][]{Gorti04}. With time, the disk will loose its gas through processes such as photoevaporation, stellar winds, formation of gas-rich planets and accretion onto the star \\citep[e.g.,][]{Hollenbach00,Alexander05}. However, the timescale for gas dissipation in disks is still poorly constrained observationally and it is even unclear whether gas and dust are lost simultaneously from disks when they evolve from the massive optically thick to the tenuous optically thin `debris' phase.  To constrain these timescales, surveys of gas in disks in a variety of environments are needed with sensitivities sufficient to detect a Jupiter mass of gas ($\\sim$10$^{-3}$ M$_\\odot$). Surveys to date have focussed on CO millimeter observations, mostly of disks in the Taurus molecular cloud \\citep[e.g.][]{Thi01}.  Although CO is known to be a poor tracer of the gas mass due to the combined effects of photodissociation in the upper layers and freeze-out in the midplane \\citep[e.g.,][]{Zadelhoff03}, it is much simpler to observe with current instrumentation than other tracers including H$_2$ \\citep{Thi01} or atomic fine-structure lines \\citep{Kamp03}. Interferometer surveys have detected gas-rich disks around a large fraction ($>$60\\%) of classical T Tauri stars in Taurus-Auriga with ages up to a few Myr and gas masses of at least a few times $10^{-4}$ M$_\\odot$, depending on assumptions about the CO abundance \\citep{Koerner95,Dutrey96,Dutrey03}. In contrast, the detection rate is only 10\\% for the more evolved weak-line T Tauri stars \\citep{Duvert2000}. Similar statistics are found in a single-dish submillimeter survey of CO in a spatially more distributed sample of Herbig Ae and Vega excess stars \\citep{Dent05} and in CO infrared surveys \\citep{Najita03,Blake04}.  Surprisingly little is known about the presence of gas-rich disks around T Tauri stars in other nearby star-forming regions such as Corona Australis, Lupus and Chamaeleon. This is largely due to the absence of interferometers in the Southern sky, and even Ophiuchus has been poorly sampled with current facilities. We present here an initial search for gas-rich disks in the Lupus molecular clouds, using a single dish telescope.  The Lupus clouds, located at around galactic coordinates 335$^\\circ \\le l \\le 341^\\circ$ and $7^\\circ \\le b \\le 17^\\circ$ \\citep{Hughes94}, are among the closest star-forming regions to the Sun at $\\sim$150 pc \\citep[see reviews by][]{Krautter91,Comeron06}.  Mid M-type stars dominate the stellar population of Lupus, which contains no confirmed O or B-type stars. In total 69 Pre-Main Sequence (PMS) objects were found by \\citet{Schwartz77}, distributed over four clouds now known as Lupus 1--4 \\citep[see also][]{Hughes94}. In addition, some 130 new weak-line T Tauri stars were detected in ROSAT images \\citep{Krautter94,Krautter97}. The most massive stars in the complex are two A-type stars, HR 5999 and HR 6000, present in the richest sub-grouping in stars and complexity, Lupus 3. A good overview of the clouds is given by the extinction map made by \\citet{Cambresy99}, based on star counts. Recently a higher resolution ($\\sim 30''$) map has been published by \\citet{Teixeira05}.  At millimeter wavelengths the Lupus clouds have been studied both in molecular emission lines and in the continuum. $^{12}$CO \\citep{Tachihara01}, $^{13}$CO \\citep{Tachihara96} and C$^{18}$O \\citep{Hara99} $J$=1--0 large-scale maps were made at 2\\farcm6 resolution with the NANTEN telescope. Higher resolution observations were done with the Swedish-ESO Submillimeter Telescope (SEST) over more limited regions with $\\sim$45$''$ resolution by \\citet{Gahm93} (mostly $^{12}$CO 1--0 in Lupus 2), \\citet{Rizzo98} ($^{12}$CO 1--0 in Lupus 1 and 4 filaments) and \\citet{Vilas-Boas00} ($^{13}$CO and C$^{18}$O 1--0 in localized dark cores in Lupus 1-4). \\citet{Nuernberge97} performed a continuum survey of 32 T Tauri stars in Lupus at 1.3 mm with the SEST, probing the cold circumstellar dust.   The Lupus star-forming region appears to differ in several aspects from other low-mass star-forming clouds. In contrast with Ophiuchus and Taurus, which contain a significant fraction of deeply embedded Class 0 and Class I objects, Lupus has at most a few embedded objects \\citep{Krautter91, Tachihara96, Comeron06}. This could be an indication that Lupus is more evolved than other star-forming regions. The peak in the age distribution is $3\\times10^6$ yr for stars in Lupus 3 and 4, calculated from the PMS evolutionary tracks of \\citet{dantona94} for a distance of 150 pc \\citep{Hughes94}. Under the same assumptions, Lupus 1 and 2 PMS stars are younger with a peak of $1\\times10^6$~yr. These age estimates are comparable to those for other star-forming regions such as Chamaeleon \\citep{Hartigan93}, but are about double that of the Taurus cloud \\citep{Simon93}.  Thus, a study of disks in Lupus is interesting because this is precisely the age range in which significant evolution of the gas disk is expected.  A second characteristic of the Lupus clouds is the distribution of spectral types. Compared to the Taurus cloud, Lupus is dominated by lower mass stars. The peak of the distribution in spectral type is M0, with very few stars having spectral types higher than K7. In contrast, the peak of the spectral type distribution in Taurus is at K7.     In this work, single-dish submillimeter observations of CO toward a sample of 21 classical T Tauri stars with disks are presented obtained with a beam of 14$''$, much smaller than that of previous data.  The sources are a subset of the PMS stars found by \\citet{Schwartz77}, \\citet{Krautter91} and \\citet{Hughes94}, and they have been observed with the {\\it Spitzer Space Telescope} in the context of the `From Molecular Cores to Planet Forming Disks' (c2d) Legacy survey \\citep{Eva03}.  The aim of the CO observations is to characterize the molecular environment of the T Tauri stars at high angular resolution and to search for gas-rich disks suitable for future follow-up interferometer observations. ",
        "conclusions": "The main conclusions of this work are as follows: \\begin{description} \\item{--} A $^{12}$CO $J$=3--2 survey of 21 classical T Tauri stars with dust disks in Lupus taken in frequency-switched mode has revealed mostly extended diffuse cloud material with a few magnitudes of extinction and complex velocity structure.  \\item{--} A C$^{18}$O $J$=3--2 and 2--1 survey of the same 21 sources taken in beam-switched mode shows dense compact circumstellar material associated with only 2 sources, HT Lup and Sz 73.  \\item{--} Only a single source, IM Lup (Sz 82), has been found with clear evidence for a gas-rich disk. Modeling of the line intensity and line profile indicates a large disk with radius $\\sim$600 AU and inclination 20$^\\circ$. This source will be a prime target for future interferometer observations in the southern sky, especially with the {\\it Atacama Large Millimeter Array}. \\end{description} The lack of detection of gas disks in Lupus in the current survey does not imply that the gas has been dissipated from these disks; rather, it shows the limitations of single-dish CO line observations to reveal them.  Future searches need to either focus on high density single-dish tracers such as HCO$^+$ $J$=4--3 or HCN $J$=4--3 to distinguish the warm, dense gas in the disk from the cold, diffuse cloud material or use newly commisioned interferometers such as the SubMillimeter Array (SMA).  From comparison with other disks in Taurus \\citep{Thi04,Greaves04}, these lines are expected to have peak intensities of at most 0.1 K in a single-dish 15$''$ beam and will be very difficult to detect.  Increased sensitivity to compact emission is only possible with interferometer studies.  Interferometric surveys will therefore be essential to determine the fraction of gas-rich disks in Lupus and thus the timescale for gas dissipation in this star forming region compared with other regions such as Taurus."
    },
    "0606/astro-ph0606628_arXiv.txt": {
        "abstract": "The RApid Temporal Survey (RATS) is a survey to detect objects whose optical intensity varies on timescales of less than $\\sim$70 min.  In our pilot dataset taken with the INT and the Wide Field Camera in Nov 2003 we discovered nearly 50 new variable objects. Many of these varied on timescales much longer than 1 hr. However, only 4 objects showed a modulation on a timescale of 1 hour or less. This paper presents followup optical photometry and spectroscopy of these 4 objects. We find that RAT J0455+1305 is a pulsating (on a period of 374 sec) subdwarf B (sdB) star of the EC\\,14026 type. We have modelled its spectrum and determine $\\teff = 29,200\\pm 1900$\\,K and $\\log g = 5.2 \\pm 0.3$ which locates it on the cool edge of the EC\\,14026 instability strip. It has a modulation amplitude which is one of the highest of any known EC 14026 star. Based on their spectra, photometric variability and their infra-red colours, we find that RAT J0449+1756, RAT J0455+1254 and RAT J0807+1510 are likely to be SX Phe stars - dwarf $\\delta$ Sct stars. Our results show that our observing strategy is a good method for finding rare pulsating stars. ",
        "introduction": "The aim of the RApid Temporal Survey (RATS) is to discover objects whose optical intensity varies on timescales of a few mins to several hours (Ramsay \\& Hakala 2005). The prime aim is to discover interacting ultra-compact binary systems - systems consisting of two degenerate (or semi-degenerate) stars orbiting around a common center of gravity - with binary orbital periods less than $\\sim$70 mins. Our pilot set of data was obtained from La Palma in Nov 2003 using the Isaac Newton Telescope and the Wide Field Camera: it covered 3 square degrees and reached a depth of $V\\sim$22.5. Nearly 50 sources were found to show significant intensity variations: none were previously known variable objects (Ramsay \\& Hakala 2005). However, only 4 objects showed modulations which varied on periods of approximately 1 hour or less. Table \\ref{sources} shows the positions and modulation periods, while in Figure \\ref{finding} we show the finding chart for each object. This paper reports followup observations of these 4 systems, the aim being to determine their nature. ",
        "conclusions": "We have obtained followup optical photometry and spectroscopy of 4 sources which were found to show intensity modulations on timescales of less than $\\sim$1 hr in our pilot RATS dataset. We find that the source which shows the shortest period (374 sec), RAT J0455+1305, is an EC 14026 star.  We have modelled its Balmer absorption lines and find a temperature of 29200K$\\pm$1900 K and a gravity of $\\log g=5.2\\pm0.3$. This places RAT J0455+1305 at the cool side of the EC 14026 instability strip. The other sources have spectra which indicate temperatures between 7200-7500K. This together with their optical light curves and infra-red colours strongly suggest they are SX Phe stars.  There are currently around 20 EC 14026 stars, while there are less than 20 currently known SX Phe stars which are not associated with globular clusters, open clusters or external galaxies. The usual method of discovering EC 14026 stars is a three-step process. Candidate stars are discovered using colour or prism surveys. Spectroscopy is then required to confirm their classification and determine their temperature and gravity. Lastly, high speed photometry is needed to search for variability. We have found that our `reverse' strategy is a good method of discovering variable EC 14026 stars.  Further, our strategy appears to be a good method of discovering SX Phe stars, which should help determine their space density and period distribution."
    },
    "0606/astro-ph0606134_arXiv.txt": {
        "abstract": "{ We present the first results of our deep survey of lensing clusters aimed at constraining the abundance of star-forming galaxies at $z\\sim6-10$, using lensing magnification to improve the search efficiency and subsequent spectroscopic studies. Deep near-IR photometry of two lensing clusters (A1835 and AC114) was obtained with ISAAC/VLT. These images, combined with existing data in the optical bands including HST images, were used to select very high redshift candidates at $z\\gtrsim 6$ among the optical-dropouts. Photometric selection criteria have been defined based on the well-proven dropout technique, specifically tuned to target star-forming galaxies in this redshift domain.  We have identified 18(8) first and second-category optical dropouts in A1835 (AC114), detected in more than one filter up to $H$ (Vega) $\\sim 23.8$ (AB $\\sim 25.2$, uncorrected for lensing). Among them, 8(5) exhibit homogeneous SEDs compatible with star-forming galaxies at $z\\gtrsim 6$, and 5(1) are more likely intermediate-redshift EROs based on luminosity considerations. We have also identified a number of fainter sources in these fields fulfilling our photometric selection and located around the critical lines. We use all these data to make a first attempt at constraining the density of star-forming galaxies present at $6\\lesssim z \\lesssim10$ using lensing clusters. Magnification effects and sample incompleteness are addressed through a careful modeling of the lensing clusters. A correction was also introduced to account for the expected fraction of false-positive detections among this photometric sample.  It appears that the number of candidates found in these lensing fields, corrected for magnification, incompleteness and false-positive detections, is higher than the one achieved in blank fields with similar photometric depth in the near-IR. The luminosity function derived for $z\\gtrsim 6$ candidates appears compatible with that of LBGs at $z\\simeq 3$, without any renormalization. The turnover observed by Bouwens et al. (2005) towards the bright end relative to the $z\\sim 3$ LF is not observed in this sample. Also the upper limit for the UV SFR density at $z\\sim6-10$, integrated down to $L_{1500}=0.3\\ L^{*}_{z=3}$, of $\\rho_\\star=7.4\\ 10^{-2}$ M$_{\\odot}$ yr$^{-1}$ Mpc$^{-3}$ is compatible with the usual values derived at $z \\simeq 5-6$, but higher than the estimates obtained in the NICMOS Ultra Deep Field (UDF). The same holds for the upper limit of the SFR density in the $z \\simeq 8-10$ interval ($\\rho_\\star=1.1\\ 10^{-1}$). This systematic trend towards the bright end of the LF with respect to blank fields could be due to field-to-field variance, a positive magnification bias from intermediate-redshift EROs, and/or residual contamination. Given the low S/N ratio of the high-$z$ candidates, and the large correction factors applied to this sample, increasing the number of blank and lensing fields with ultra-deep near-IR photometry is essential to obtain more accurate constraints on the abundance of $z \\gtrsim 6$ galaxies. ",
        "introduction": " ",
        "conclusions": "\\label{conclusions}  We have obtained deep $JHK$ near-IR photometry of two well-known lensing clusters, A1835 and AC114, plus $z$ and $SZ$ imaging of A1835 with FORS and ISAAC at the VLT. Our photometric depth reached $SZ\\sim 25.6$, $J\\sim 24.4$, $H\\sim 23.5$ and $K_s\\sim 23.3$ (Vega system), in addition to a minimum magnification factor of 1 magnitude over half of the ISAAC field of view. These images, combined with existing data in various optical bands including HST images, have been used to select galaxy candidates at very high redshift ($z \\sim 6-10$). The candidates have been selected with the dropout technique and two-color selection criteria appropriate for high-$z$ galaxies.  From our $H$ band selected sample we have identified 18 (8) ``first and second-category'' optical dropouts in A1835 (AC114) up to $H_{\\rm Vega} \\sim 23.9$ (AB $\\sim 25.3$, uncorrected for lensing). Second category is defined here as objects detected in $\\ge 2$ near-IR bands, the best-detected sources being defined as first priority. Among them, 8(5) exhibit homogeneous SEDs compatible with star-forming galaxies at $z\\gtrsim 7$, and 5(1) are more likely intermediate-redshift EROs. In both fields we have also identified a few additional dropouts detected only in the $H$ band (``third category'' objects), which satisfy our photometric selection criteria. We have estimated the fraction of spurious sources expected in the different filter combinations, and corrected all the relevant derived quantities for this effect.  Typically our candidates are magnified by a factor of 1.5 ($\\sim$ 0.44 mags) to 10 (2.5 mags), with average (median) values of the order of $6.5-7.9$ ($2.3-3.5$) for the two clusters. All high-$z$ candidates turn out to be fainter than $H_{\\rm Vega} \\sim$ 23, with typical effective (i.e.\\ lensing corrected) magnitudes $H_{\\rm Vega} \\sim  24-25$ ($H_{\\rm AB} \\sim$ $25.4-26.4$) and fainter in some cases. Assuming standard SFR(UV) conversion factors, which may however be questionable for galaxies of such presumably young age, the SFR is found to be between few units and $\\sim$ 20 \\msun\\ yr$^{-1}$. Their UV restframe spectrum, measured by the $H-K$ color, seems to be very blue -- a trend also reported for other high-$z$ galaxy samples (e.g.\\ Papovich et al.\\ 2004, Bouwens et al.\\ 2004b).  Taking into account the gravitational lensing effects, sample incompleteness and expected spurious detections, a first attempt was made to constrain the density of star-forming galaxies present at $z \\sim 6-10$ based on lensing data. Integrated quantities, such as LFs and SFR densities, are affected by non-negligible uncertainties due to the large correction factors applied to this sample. The LF measured for LBGs at $z\\simeq 3$ seems to be slightly higher (by $\\sim 0.5$ dex) but still consistent with the LF derived for our sample. The turnover observed by Bouwens et al. (2005) towards the bright end relative to the $z\\sim 3$ LF is not observed in this sample.  We have also estimated an upper limit for the cosmic SFR density from these data. Our values in the $z \\sim 6-10$ domain are higher than the estimates obtained in the NICMOS UDF, even when the most conservative assumptions and corrections are applied to the data. This difference is related to the bright end of the UV LF for our candidates, i.e. $L_{1500}\\gtrsim 0.3 L^*_{z=3}$. This systematic trend with respect to blank fields could be due to to field-to-field variance, a positive magnification bias from intermediate-redshift EROs, and/or residual contamination by spurious sources. Given the error bars, residual magnification bias should be negligible for high-$z$ sources, but a population of faint intermediate-redshift interlopers affected by positive magnification bias cannot be excluded. At least one of such intermediate sources was spectroscopically identified in the field of A1835 (Richard et al. 2003).  According to our simulations, in agreement with the first photometric results presented in this paper, the use of lensing clusters as gravitational telescopes seems to significantly improve the survey efficiency of $z\\gtrsim 6$ galaxies compared to blank fields.   Given the uncertainties involved in the candidate-selection process, and the faint fluxes observed for our photometric candidates, the present results are to be considered as a first attempt to constraint the population of $6 \\lesssim z \\lesssim 10$ star-forming galaxies using lensing clusters. The present results and conclusions have to be confirmed and improved. Spectroscopic follow-ups are underway to determine the efficiency of our selection technique, and the contamination level by intermediate-redshift interlopers. Additional deep photometry in various bands are being secured with HST, IRAC/Spitzer, and from the ground to improve the SEDs characterization of the high-$z$ candidates. Increasing the number of lensing fields with ultra-deep near-IR photometry is essential to obtain more accurate constraints on the abundance and physical properties of $z\\gtrsim 7$ starburst galaxies."
    },
    "0606/astro-ph0606378_arXiv.txt": {
        "abstract": "Recent site testing (see: http://www-luan.unice.fr/Concordiastro/indexantartic.html) has shown that Dome C in Antarctica might have a high potential for stellar interferometry if some solutions related to the surface atmospheric layer are found. A demonstrator interferometer could be envisioned in order to fully qualify the site and prepare the future development of a large array.  We analyse the performances of a prototype interferometer for Dome C made with 3 telescopes of 40 cm diameter. It assumes classical Michelson recombination. The most recent atmospheric and environmental conditions measured at Dome C are considered (see K. Agabi  \"First whole atmosphere night-time seeing measurements at Dome C, Antarctica\"[\\cite{Agabi}]). We also study the possible science reachable with such a demonstrator. Especially we evaluate that even such small aperture interferometer could allow the detection and low resolution spectroscopy of the most favourable pegaside planets. ",
        "introduction": "\\label{sect:intro}  % Astronomical observations in infrared are mainly limited by the atmospheric background. This flux depends strongly on the atmospheric temperature and emissivity. That is why the Antarctic should be an interesting place for astronomical observations (Chamberlain, 2000 [\\cite{astoo}]). Since the forthcoming of permanent stations in Antarctica, more and more research groups are testing atmospheric characteristics there and expect to build science instruments on Antarctic plateau. The size of the projects ranges begin from a small telescope to an ELT or to an optical interferometric array like Kiloparsec Explorer for Optical Planet Search (KEOPS)(Vakili, 2005 [\\cite{KEOPS}]). Since a few years the Laboratoire Universitaire Astrophysique de Nice (LUAN) has access to the station Concordia. The site characteristics were investigated. A strong potential of this site for interferometry was pointed out. ",
        "conclusions": ""
    },
    "0606/astro-ph0606652_arXiv.txt": {
        "abstract": "{} {As part of the Red MSX Source Survey of Massive Young Stellar Objects (MYSOs) we have conducted multi-wavelength follow up observations of the well-known object.  We present our data on this object, whose near-infrared spectrum is exceptional and place these in context with previous observations} {Our observations of V645 Cyg included near/mid infrared imaging observations, $^{13}$CO 2-1 line observations and high signal-to-noise velocity resolved near-infrared spectroscopy.}  {The spectrum shows P-Cygni hydrogen Brackett emission, consistent with a high velocity stellar wind.  A red-shifted emission component to a number of near-IR emission lines was also uncovered. This is associated with a similar component in the H$\\alpha$ line. V645 Cyg is also found to have variable CO first overtone bandhead emission}{The data clearly indicate that the outflow of V645 Cyg is variable. The unidentified feature in a previously published optical spectrum is identified with a receding outflow at 2000 kms$^{-1}$. The nature of this feature, which is found in hydrogen and helium atomic lines and CO molecular lines remains a puzzle.  }  \\titlerunning{RMS observations of V645 Cyg} \\authorrunning{Clarke et al.} ",
        "introduction": "The variable star V645 Cyg has been subject to detailed study since its identification with the strong infrared source AFGL 2789 by \\cite{walker} and \\cite{lebofsky}. The optical follow up observations of \\citet{cohen} revealed the star to be a young O7e type star within a polarised optical nebula. The presence of P-Cygni profile hydrogen recombination lines in the optical spectrum was interpreted as evidence for a stellar wind with velocities as high as 1000~\\rm{kms}${^{-1}}$. Further analysis by \\cite{goodrich} and \\cite{hamann} showed the star to have similar P-Cygni profile He {\\sc i}, O {\\sc i}, Si {\\sc ii} and Ca {\\sc ii} transition lines and concluded that the optical spectrum was more typical of a Herbig Ae/Be type star. In addition to the lines associated with the stellar wind, the optical spectrum of V645 Cyg also possesses emission lines of C {\\sc i}, [S{ \\sc ii}], K {\\sc i}, Fe{\\sc i}, Fe {\\sc ii} and [Fe {\\sc ii}]. The presence of blue-shifted forbidden [S {\\sc ii}] lines had been suggested by \\citet{hamann} to be evidence for a dense circumstellar disc which obscures the receding, red-shifted part of a bipolar flow around V645 Cyg.  Near-IR spectroscopic observations have also been performed on V645 Cyg by \\cite{harvey}, \\cite{geballe}, \\cite{carr} and \\cite{biscaya} who found H {\\sc i} emission and  that the CO first-overtone bandheads may be variable respectively.  V645 Cyg is located within a high density molecular cloud and has been associated with a low velocity moderately collimated bipolar molecular outflow. The CO millimetre line observations of \\cite{rodriguez} and \\cite{montenegro} showed the outflow to be oriented in the north-south direction at small scales ($\\approx$15 arcsec), coinciding with the optical outflow observed by \\cite{goodrich}. Whilst at larger scales ($\\approx$1$\\arcmin$) the CO outflow is reoriented into the southeast-northwest direction.  Studies of the radio continuum emission found weak ($\\approx$0.8 mJy) extended ($\\approx$10 arcsec) emission in the north-south direction (\\citealt{curiel}; \\citealt{skinner}). This emission was also determined to be variable by \\cite{girart}, who speculated that the variability could be due to the episodic ejection of material or recombination of the ionised gas on a very short timescale.  The object is also associated with OH (1665 MHz) \\citep{morris}, H$_2$O \\citep{lada} and methanol masers \\citep{slysh}.  Therefore, the picture of V645 Cyg that has emerged is that of a relatively unembedded young massive star, with a high velocity stellar wind and an associated optical and molecular outflow.  We are currently undertaking a survey aimed at finding massive young stellar objects (MYSOs) in the Galaxy (\\citealt{lumsden}; \\citealt{hoare}), known as the Red MSX Survey (RMS). We have selected the MSX source G094.6028-01.7966 as having infrared colours indicative of a MYSO. This source is identified as V645 Cygni. In order to have complete datasets to study the properties of MYSOs, we conducted a series of follow-up observations.  Here, we present millimetre, mid-IR, near-IR and optical observations of V645 Cyg. In Section \\ref{observations} we describe the observations and the procedure taken to produce the final spectrum, in Section \\ref{results} we present the data and discuss the features present and in Section \\ref{discuss} we discuss our observations in the context of previous work. ",
        "conclusions": "\\label{discuss}  We have presented new multi-wavelength observations of V645 Cyg.  Our mid/near infrared imaging data imply that the stellar N0 component of V645 Cyg is the location of the mid infrared emission from this source.  N0 is also the exciting star of the diffuse reflection nebulae V645 Cyg N1, as previously shown by \\cite{minchin}.  The P-Cygni like absorption of the Brackett hydrogen lines is consistent with a powerful stellar wind.  The velocity of this wind is typical of that expected from a young OB type star but stands out for a massive young stellar object, where much lower outflow velocities are the norm (\\cite{drew}: many massive young stellar objects show no P-Cygni features at all). Indeed, when taking the crude spectral classification of an O8.5 supergiant at face value, the spectral class itself would be consistent with low outflow velocities, as typically, supergiants have lower escape velocities than dwarf stars. The fact that the classification is based on absorption lines extending to 2000 kms$^{-1}$ strongly suggests, however, we are dealing with a rapidly expanding pseudo-photosphere.  This notion is corroborated by the luminosity estimate for V645 Cyg, which indicates it to be a high mass object but not a supergiant. For the accepted distances (roughly 3-6 kpc), the luminosity falls within the range 2$-$7$\\times 10^4$L$_{\\odot}$. This is consistent with early B/late O star zero-age-main-sequence bolometric luminosities, but not with a supergiant luminosity.  The fact that the high excitation He {\\sc i} lines do not show visible emission indicates that the acceleration region is very small, and that the terminal velocities are reached quickly.  These high velocities are not often observed in MYSOs. Low outflow velocities in MYSOs can be explained by radiatively driven disk-winds (cf. \\citealt{drewproga}). Were V645 Cyg to belong to the MYSO class, it would would appear that the circumstellar disk has disappeared and the wind is a purely radiatively driven stellar wind. It may thus well be a transitionary object between the MYSO phase and the ensuing H {\\sc ii} region phase.  In this case the wind must still be thick enough to absorb the ionising radiation from the central star since there is as yet no evidence for an H {\\sc ii} region being present, and a star of this luminosity should be able to generate such a region (unlike lower mass Herbig type stars).  The relatively low extinction towards the source is consistent with this picture of an `evolved' MYSO.  Intriguingly, the spectrum of V645 Cyg shows red-shifted emission components to the He {\\sc i}, hydrogen recombination and CO lines. The velocity of this redshifted component is remarkably similar to the blueshifted velocities measured in the P Cygni absorption troughs. We therefore conclude that the redshifted features correspond to the hitherto unidentified receding part of the flow.  The relative weakness of the emission could be explained by obscuration due to a dense disk or simply dust extinction within the nebula.  The observation of variable CO bandhead emission is not particularly unusual, as \\cite{biscaya} found many YSO's show CO bandhead variability on a range of timescales. The variability of CO bandhead emission is related to the special conditions that are required to produce it. The warm, dense and neutral material that is required to produce CO bandhead emission is generally only found in dense rotating disks or mildly shocked regions.  The observation of a high velocity red-shifted component to the CO emission would appear to imply that its emission is related to the stellar wind of the star rather than a rotating disk. This introduces another puzzle, as an 1800\\rm{kms}$^{-1}$ flow shocking the interstellar medium would most likely destroy any molecular CO.  It may be possible that we are witnessing the complex transition zone between the stellar wind and the larger scale molecular outflow.  In this case the CO emission should be tied strongly to the strength of the stellar wind, and the existing data do suggest this may be true.  V645 Cyg therefore presents a variety of problems as regards its exact nature.  Further observations of the clearly variable wind would be helpful in illuminating some of these.  Overall it may represent a relatively rare class of transition objects between a genuinely massive young stellar object and a normal young Oe type star in a weak H{\\sc ii} region."
    },
    "0606/astro-ph0606187_arXiv.txt": {
        "abstract": "We present the results of numerical experiments, in which we have investigated the influence of the inelastic neutrino-helium interactions on the standing accretion shock instability supposed to occur in the post-bounce supernova core. The axisymmetric hydrodynamical simulations of accretion flows through the standing accretion shock wave onto the protoneutron star show that the interactions are relatively minor and the linear growth of the shock instability is hardly affected. The extra heating given by the inelastic reactions becomes important for the shock revival after the instability enters the non-linear regime, but only when the neutrino luminosity is very close to the critical value, at which the shock would be revived without the interactions. We have also studied the dependence of the results on the initial amplitudes of perturbation and the temperatures of mu and tau neutrinos. ",
        "introduction": "Most of the supernova modelers are currently concerned with the multi-dimensional aspects of the dynamics, pushed by the accumulating observational evidences that the core-collapse supernovae are generally aspherical \\citep{wang96,wang01,wang02}. Various mechanisms to produce the asymmetry have been considered so far. Among them are the convection (e.g., \\citet{herant_94,burrows_95,jankamueller96}), growth of asymmetry seed generated prior to core-collapse \\citep{burohey,fryerkick}, rotation and magnetic fields (see \\citet{yamasawa,kotakemhd,takiwaki,sawa,ard} and \\cite{kotake_rev} for collective references).  Recently the standing accretion shock instability (SASI) is attracting interest of researchers. The instability was originally studied for transonic accretion flows to black holes (e.g., \\citet{foglizzo01,foglizzo02}) and was re-discovered by \\citet{blondin_03} in the context of core-collapse supernovae. \\citet{blondin_03} found in their 2D numerical simulations of the spherically symmetric, isentropic, steady accretion flows, that a standing shock wave are unstable to non-spherical perturbations and that the perturbations grow up to the nonlinear regime with the clear dominance of the $\\ell=1$ mode at first and $\\ell=2$ mode later, leading to the global deformation of the shock wave as has been observed. Here $\\ell$ stands for the azimuthal index of the Legendre polynomials.  The mechanism of the instability is still controversial. \\citet{blondin_05} took coolings into account in a simple analytic way following \\citet{HC} and claimed that the repeated propagations of pressure fluctuations are responsible for the instability. On the other hand, \\citet{ohnishi_1} did 2D numerical experiments, implementing more realistic heatings as well as coolings by neutrino and found that the original idea by \\citet{foglizzo01,foglizzo02} that the non-radial instability is driven by the cycle of the advection of entropy and velocity fluctuations and propagation of pressure perturbations seems to be more appropriate. Obviously detailed linear analyses of the instability are needed~\\citep{ga05,yy06b}.  It is important that such a low-$\\ell$-mode deformation of the shock wave has been also observed in more realistic simulations \\citep{scheck_04} and that the asymmetric explosion following the instability might reproduce various non-spherical features of SN1987A~\\citep{kifo}. It should be also emphasized that the instability is helpful for the shock revival just as the convection is~\\citep{ohnishi_1}. The problem is whether the instability enhances the neutrino heating sufficiently to revive the stalled shock. Recent numerical investigations by~\\citet{ja05} seem to say ``no'' to this question. It is true that we have to wait for detailed 3D simulations before drawing a conclusion, but we had better continue to seek for some other processes that would further facilitate the shock revival. As such a potential boost, we will focus in this paper on the interplay between the inelastic neutrino-nucleus interactions and the SASI.  The potential importance of the inelastic neutrino-nucleus interactions was first pointed out by \\citet{haxton}, who paid attention to the heating and dissociations of nuclei in the matter ahead of the shock wave via these reactions, the so-called preheating. Taking the reactions into account in their 1D spherically symmetric numerical simulations, \\citet{bruenn_haxton} found that the effect of preheating is quite minor mainly because $\\nu_{\\rm e}$ energies obtained in the simulation was lower than the one assumed in \\citet{haxton}. They also discussed the possibility of shock revival by enhanced heating of the postshock material by inelastic $\\nu_{\\tau, \\mu}-{}^4{\\rm He}$ scatterings, the idea similar to ours pursued in this paper. They found that the reactions are not very important. It should be emphasized, however, that the conclusion will be sensitive to the background model and they considered a single snapshot after the bounce. Furthermore, in the spherically symmetric models, most of the nuclei are photodisintegrated after passing through the shock and the reactions will scarcely occur anyway. The situation may be different in non-spherical cases. Since the shock wave hovers at larger radius in general, not all the nuclei are dissociated and the heating region will be wider. In this paper, we pay particular attention to the interplay between the SASI and the inelastic reactions.  The plan of this paper is as follows. We describe the numerical methods, input physics and models in section \\ref{s2}. The main numerical results are shown in section \\ref{s3}. We conclude this paper with section \\ref{s4}. ",
        "conclusions": "}  We have investigated the possible effects of the inelastic interactions of neutrino with helium on the shock revival in the post-bounce supernova core. In particular, we have paid attention to their influence on the SASI, one of the major causes for the asymmetry of dynamics and a possible trigger of explosion. For the spherically symmetric models, \\citet{bruenn_haxton} found that both the preheating of matter ahead of the shock and heating of matter behind the shock via these reactions are quite minor. In fact, most of the nuclei are photodisintegrated after passing through the shock and the reactions will scarcely occur anyway in the spherically symmetric models. The situation may be different in non-spherical cases, where the shock wave hovers at larger radius in general and not all the nuclei are dissociated and the heating region is wider. In fact, these reactions have never been explored in the multi-dimensional context so far. We have done 2D numerical experiments on the post-bounce accretion flows through the stalled shock wave onto the protoneutron star, systematically changing the luminosity and temperature of neutrino and the initial amplitude of perturbation as well as the reaction rates.  We have found that the incorporation of the inelastic interactions has essentially no influence on the growth of the SASI, since very little helium is existing in the post-shock matter initially. However, these reactions become appreciable later when the SASI enters the non-linear regime and the shock oscillates with large amplitudes. It has been shown that the extra heating by these interactions is helpful for the shock revival in principle. This is, however, true in practice only when the shock revival is not obtained with a slight margin without the interactions. In fact, we have observed that a small ($\\sim$10\\%) reduction of the neutrino luminosity makes the interactions entirely negligible. Hence it is understandable that larger initial amplitudes of perturbation make the interactions more important for shock revival. It is, however, mentioned that even if the luminosity is very close to the critical value, the cross sections estimated by \\citet{haxton} seem to be too small. Although it is not easy to evaluate the uncertainties of the theoretical prediction, recent new calculations by \\citet{gazit04} may be used as a guide. They predicted a bit larger $\\beta$ than \\citet{haxton}. However, the enhancement of the heating rate for our reference model is only 15\\%, much too small for the interactions to have some influence on the shock revival. The fact that the dynamics is rather insensitive to the neutrino temperature is not encouraging, either. Hence we conclude that the inelastic interactions of neutrino with helium will be important only in determining the shock-revival-point precisely. It is, however, noted finally that the inelastic neutrino-nuclei interactions should be incorporated in the realistic simulations, since, as we have seen, the preheating may suppress the non-spherical oscillations of the shock wave."
    },
    "0606/astro-ph0606714_arXiv.txt": {
        "abstract": "In a previous  paper, we outlined a new  Bayesian method for inferring the  properties of extended  gravitational lenses,  given data  in the form  of resolved  images.  This  method  holds the  most promise  for optimally  extracting  information  from  the observed  image,  whilst providing reliable uncertainties in all parameters. Here, we apply the method  to the  well  studied optical  Einstein  ring 0047-2808.   Our results are in broad agreement with previous studies, showing that the density profile of the lensing  galaxy is aligned within a few degrees of  the light  profile,  and  suggesting that  the  source galaxy  (at redshift 3.6) is a  binary system, although  its size is  only of order 1-2 kpc. We also find  that the mass of  the elliptical lensing galaxy  enclosed  by  the  image  is  (2.91$\\pm$0.01)$\\times$10$^{11}$ M$_{\\sun}$. Our method is able to achieve improved resolution for the source reconstructions, although we also find that some of the uncertainties  are greater than has been found in previous analyses, due to the inclusion of extra pixels and a more general lens model. ",
        "introduction": "Gravitational  lenses have  long  held the  promise  of being  natural telescopes, probing  structural scales smaller than  the resolution of current technologies  while simultaneously revealing  the distribution of   matter  within   the   lensing  galaxy   \\citep[For  reviews   of gravitational   lensing,  see][]{1992grle.book.....S,kochreview}.   In practice, this promise can only  be truly realised in systems in which the lensed images are extended,  as these can offer substantially more constraints than multiply imaged  point-like quasars. For this reason, significant  attention has  been paid  to extended  sources  lensed by intervening  galaxies  to produce  roughly  circular  images known  as Einstein rings \\citep[][]{2001ApJ...547...50K}.  A long-standing  issue has  focused upon the  question of what  is the {\\it  best} approach  to invert  such extended  gravitationally lensed images, providing the unlensed source brightness distribution and mass in the lensing  galaxy. Addressing this question has  spawned a number of   seemingly   distinct   methods,    such   as   the   ring   Cycle \\citep{1989MNRAS.238...43K},           Semi-Linear           Inversion \\citep{2003ApJ...590..673W}        and        Genetic       Algorithms \\citep{2005PASA...22..128B}. Given the  extended nature of the source, reconstructions have been based upon  a pixellated source plane, so as to make minimal assumptions about the source brightness distributions. However, due  to the  ill-posed nature of  such an  inversion, authors have tended to use low  numbers of pixels, and/or regularization. In a previous contribution,  we demonstrated  that other approaches  can be unified in terms of a Bayesian interpretation, with each corresponding to  the use of  particular, often  unjustified, assumptions  about the nature of the  source \\citep[][hereafter \\BL]{ours}. Furthermore, \\BL\\ presented a general Bayesian approach to the question of gravitational lens inversion,  using realistic prior distributions  and Markov Chain Monte Carlo (MCMC) methods to recover the properties of the lensing system.  In this present contribution, the approach detailed in \\BL\\ is applied to   the   well-studied   optical   Einstein   ring   0047-2808.    In Section~\\ref{image},  the  details  of   this  system,  and  the  data employed,  are  presented,  while Section~\\ref{approach}  details  the approach to  the problem.  Section~\\ref{results}  outlines the results of  this study,  with  a  comparison to  other  techniques, while  the conclusions are presented in Section~\\ref{conclusions}. ",
        "conclusions": "In this paper,  we have presented an analysis of the  HST image of the Einstein  ring  0047-2808 using  a  Bayesian  procedure.  Most of  the results are qualitatively consistent with previous work, leading us to conclude that the   various   different   methods   available   for   reconstructing gravitational  lens systems  all  give satisfactory  results for  this system,  despite  the  disagreement  about the  uncertainties  in  the parameter estimates.   However, this may  not always be the  case with other data sets \\citep[see][for an example]{2005ApJ...631.1198G}. However, our use of smaller pixels and a more reasonable prior for the source has allowed us to reconstruct the source with a greater resolution than has been possible with other methods. As a result, we have found that the second component of the source is not just a single patch of light, but has a narrow structure protruding from it; given structure formation in $\\Lambda$CDM cosmologies, we should expect such objects in the early universe to show significant substructures indicative of major mergers and accretion.  Discovered in continuum radio sources \\citep{1988Natur.333..537H}, the number,  diversity  and observational  detail  of  Einstein rings  has continued          to         grow~\\citep[e.g.][]{2003ApJ...583...67C, 2003MNRAS.343L..29T,    2003Sci...300..773C}.   Furthermore,   optical Einstein   rings   should  be   uncovered   in   a  upcoming   surveys \\citep{1992MNRAS.259P..31M}, particularly the SLACS survey \\citep{2006ApJ...638..703B},  and  the  number  of  cases  is  steadily increasing \\citep{2003A&A...406L..43S,  2005A&A...436L..21C}. Hence it is vital that progress is made in developing and understanding optimal and  reliable  methods  for  analyzing  them.   Model  selection  will probably  play  an   important  role  in  this  area   in  the  future \\citep{suyu}, however we  doubt that it will be  useful to compare the evidence for  one ad-hoc regularization formula  against another, when it is  known in advance that  neither of them  accurately describe the prior information that is  available.  Despite these concerns, the level of qualitative agreement between the  results presented in this  paper and in  others show that these debates may only be of minor importance in practice."
    },
    "0606/astro-ph0606522_arXiv.txt": {
        "abstract": "In this paper we present the characterization of all the principal meteorological parameters (wind speed and direction, pressure, absolute and potential temperature) extended up to $25$ km from the ground and over two years (2003 and 2004) above the Antarctic site of Dome C. The data set is composed by {\\it 'analyses'} provided by the General Circulation Model (GCM) of the European Center for Medium Weather Forecasts (ECMWF) and they are part of the catalog MARS. A monthly and seasonal (summer and winter time) statistical analysis of the results is presented. The Richardson number is calculated for each month of the year over $25$ km to study the stability/instability of the atmosphere. This permits us to trace a map indicating where and when the optical turbulence has the highest probability to be triggered on the whole troposphere, tropopause and stratosphere. We finally try to predict the best expected isoplanatic angle and wavefront coherence time ($\\theta_{0,max}$ and a $\\tau_{0,max}$) employing the Richardson number maps, the wind speed profiles and simple analytical models of $\\CN2$ vertical profiles. ",
        "introduction": "The Antarctic Plateau has revealed to be particularly attractive for astronomy since already several years Fossat (2005), Storey et al. (2003). It is extremely cold and dry and this does of this site an interesting candidate for astronomy in the long wavelength ranges (infrared, sub-millimeter and millimeter) thanks to the low sky brightness and high atmospheric transmission caused by a low temperature and concentration of the water vapour in the atmosphere (Valenziano \\& Dall'Oglio 1999, Lawrence 2004, Walden et al. 2005). The Antarctic Plateau is placed at high altitudes (the whole continent has an average height of $\\sim$ $2500$ m), it is characterized by a quite peculiar atmospheric circulation and a quite stable atmosphere so that the level of the optical turbulence ($\\CN2$ profiles) in the free atmosphere is, for most of the time, lower than above whatever other mid-latitude sites (Marks et al. 1996, Marks et al. 1999, Aristidi et al. 2003, Lawrence et al. 2004). Gillingham (1991), suggested for the first time, such a low level of the optical turbulence above the Antarctic Plateau. Atmospheric conditions, in general, degrade in proximity of the coasts. A low level of optical turbulence in the free atmosphere is, in general, associated to large isoplanatic angles ($\\theta_{0}$). The coherence wavefront time ($\\tau_{0}$) is claimed to be particularly large above the Antarctic Plateau due to the combination of a weak $\\CN2$ and a low wind speed all along the whole troposphere. Under these conditions, an adaptive optics system can reach better levels of correction (minor residual wavefront perturbations) than those obtained by an equivalent AO system above mid-latitude sites. Wavefront correction at high Zernike orders can be more easily reached over a large field of view, the wavefront-corrector can run at reasonably low frequencies and observations with long exposure time can be done in closed loop. This could reveal particularly advantageous for some scientific programs such as searches for extra-solar planets. Of course, also the interferometry would benefit from a weak $\\CN2$ and $\\tau_{0}$. \\newline  In the last decade several site testing campaigns took place, first above South Pole (Marks et al., 1996, Loewenstein et al. 1998, Marks et al. 1999, Travouillon et al. 2003a, Travouillon 2003b) and, more recently, above Dome C (Aristidi et al. 2003, Aristidi et al. 2005a, Lawrence et al. 2004). Dome C seems to have some advantages with respect to the South Pole: {\\bf (a)} The sky emission and atmospheric transparency is some order of magnitude better than above South Pole (Lawrence 2004) at some wavelengths. The sensitivity (depending on the decreasing of sky emission and increasing of transparency) above Dome C is around 2 times better than above South Pole in near to mid-infrared regions and around 10 times better in mid to far-infrared regions. {\\bf (b)} the surface turbulent layer, principally originated by the katabatic winds, is much more thinner above Dome C (tens of meters - Aristidi et al. 2005a, Lawrence et al. 2004) than above South Pole (hundreds of meters - Marks et al. 1999). The thickness and strength of the surface turbulent layer is indeed tightly correlated to the katabatic winds, a particular wind developed near the ground characterizing the boundary layer circulation above the whole Antarctic continent. Katabatic winds are produced by the radiative cooling of the iced surface that, by conduction, cools the air in its proximity. The cooled air, in proximity of the surface, becomes heavier than the air in the up layers and, for a simple gravity effect, it moves down following the ground slope with a speed increasing with the slope. Dome C is located on the top of an Altiplano in the interior region of Antarctica and, for this reason, the katabatic winds are much weaker above Dome C than above other sites in this continent such as South Pole placed on a more accentuated sloping region. \\newline  At present not much is known about the typical values of meteorological parameters above Dome C during the winter (April-September) time i.e. the most interesting period for astronomers.  \\noindent The goals of our study are the following.\\newline {\\bf (i)} We intend to provide a complete analysis of the vertical distribution of the main meteorological parameters (wind speed and direction, absolute temperature, pressure) in different months of the year using European Center For Medium Weather Forecasts (ECMWF) data. A particular attention is addressed to the wind speed, key element for the estimate of the wavefront coherence time $\\tau_{0}$. The ECMWF data-set is produced by the ECMWF General Circulation Model (GCM) and is therefore reliable at synoptic scale i.e. at large spatial scale. This means that our analysis can be extended to the whole troposphere and even stratosphere up to $20$-$25$ km. The accuracy of such a kind of data is not particularly high in the first meters above the ground due to the fact that the orographic effects produced by the friction of the atmospheric flow above the ground are not necessarily well reconstructed by the GCMs\\footnote{It was recently shown (Masciadri 2003) that meso-scale models can estimate the near ground wind above astronomical sites better than the GCMs.}. We remind to the reader that a detailed analysis of the wind speed near the ground above Dome C extended over a time scale of $20$ years was recently presented by Aristidi et al. (2005a). In that paper, measurements of wind speed taken with an automatic weather station (AWS) ($http://uwamrc.ssec.wisc.edu$) are used to characterize the typical climatological trend of this parameter. In the same paper it is underlined that estimates of the temperature near the ground are provided by Schwerdtfeger (1984). The interested reader can find information on the value of this meteorologic parameter above Dome C and near the surface in these references. Our analysis can therefore complete the picture providing typical values (seasonal trend and median values) of the meteorological parameters in the high part of the surface layer, the boundary layer and the free atmosphere. Thanks to the large and homogeneous temporal coverage of ECMWF data we will be able to put in evidence typical features of the meteorological parameters in the summer and winter time and the variability of the meteorological parameters in different years. The winter time is particularly attractive for astronomical applications due to the persistence of the {\\it 'night time'} for several months. This period is also the one in which it is more difficult to carry out measurements of meteorological parameters due to logistic problems. For this reason ECMWF data offer a useful alternative to measurements for monitoring the atmosphere above Dome C over long time scales in the future. \\newline\\newline {\\bf (ii)} We intend to study the conditions of stability/instability of the atmosphere that can be measured by the Richardson number that depends on both the gradient of the potential temperature and the wind speed: $R_{i}$$=$$R_{i}$(${\\partial \\theta }/{\\partial h}$,${\\partial V}/{\\partial h}$). The trigger of optical turbulence in the atmosphere depends on both the gradient of the potential temperature (${\\partial \\theta }/{\\partial h}$) and the wind speed (${\\partial V}/{\\partial h}$) i.e. from the $R_{i}$. This parameter can therefore provide useful information on the probability to find turbulence at different altitudes in the troposphere and stratosphere in different period of the year. Why this is interesting ? At present we have indications that, above Dome C, the optical turbulence is concentrated in a thin surface layer. Above this layer the $r_{0}$ is exceptionally large indicating an extremely low level of turbulence. The astronomic community collected so far several elements certifying the excellent quality of the Dome C site and different solutions might be envisaged to overcome the strong surface layer such as rising up a telescope above $30$ m or compensating for the surface layer with AO techniques. The challenging question is now to establish more precisely how much the Dome C is better than a mid-latitude site. In other words, which are the {\\it typical} $\\varepsilon$, $\\tau_{0}$ and $\\theta_{0}$ that we can expect from this site ?  We mean here as {\\it typical}, values that repeat with a statistical relevance such as a mean or a median value. For example, the gain in terms of impact on instrumentation performances and astrophysical feedback can strongly change depending on how weak the $\\CN2$ is above the first $30$ m. In spite of the fact that $\\CN2$$=$$10^{-18}$, $\\CN2$$=$$10^{-19}$ or $\\CN2$$=$$0$ are all small quantities, they can have a different impact on the final value of $\\varepsilon$, $\\tau_{0}$ and $\\theta_{0}$. Only a precise estimate of this parameter will provide to the astronomic community useful elements to better plan future facilities (telescopes or interferometers) above the Antarctic Plateau and to correctly evaluate the real advantage in terms of turbulence obtained choosing the Antarctic Plateau as astronomical site. With the support of the Richardson number, the wind speed profile and a simple analytical $\\CN2$ model we will try to predict a $\\tau_{0,max}$ and a $\\theta_{0,max}$ without the contribution of the first $30$ m of atmosphere. \\newline\\newline {\\bf (iii)} Data provided by ECMWF can be used as inputs for atmospheric meso-scale models usually employed to simulate the optical turbulence ($\\CN2$) and the integrated astroclimatic parameters (Masciadri et al. 2004, Masciadri \\& Egner 2004, Masciadri \\& Egner 2005). Measurements of wind speed done during the summer time have been recently published (Fig. 1 - Aristidi et al. 2005a). We intend to estimate the quality and reliability of the ECMWF data comparing these values with measurements from Aristidi et al. so to have an indication of the quality of the initialization data for meso-scale models. We planned applications of a meso-scale model (Meso-Nh) to the Dome C in the near-future. As a further output this model will be able to reconstruct, in a more accurate way than the ECMWF data-set, the meteorologic parameters near the ground.\\newline  The paper is organized in the following way. In Section \\ref{meteo} we present the median values of the main meteorological parameters and their seasonal trend. We also present a study of the Richardson number tracing a complete map of the instability/stability regions in the whole $25$ km on a monthly statistical base. In Section \\ref{rel} we study the reliability of our estimate comparing ECMWF analysis with measurements. In Section \\ref{disc} we try to retrieve the typical value of  $\\tau_{0,max}$ and $\\theta_{0,max}$ above Dome C. Finally, in Section \\ref{conc} we present our conclusions. ",
        "conclusions": "\\label{conc}  In this paper we present a complete study of the vertical distribution of all the main meteorological parameters (wind speed and direction, pressure, absolute and potential temperature) characterizing the atmosphere above Dome C from a few meters from the ground up to $25$ km. This study employs the ECMWF {\\it 'analyses'} obtained by General Circulation Models (GCM); it is extended over two years 2003 and 2004 and it provides a statistical analysis of all the meteorological parameters and the Richardson number in each month of a year. This parameter provides us useful insights on the probability that optical turbulence can be triggered in different regions of the atmosphere and in different periods of the year. The Richardson number monitors, indeed, the conditions of stability/instability of the atmosphere from a dynamic as well as thermal point of view. The main results obtained in our study are:\\newline \\begin{itemize}  \\item The wind speed vertical distribution shows two different trends in summer and winter time due to the {\\it 'polar vortex'} circulation. In the first $8$ km above the ground the wind speed is extremely weak during the whole year. The median value at $5$ km, correspondent to the peak of the profile placed at the interface troposphere/tropopause, is $12$ m/sec. At this height the 3rd quartile of the wind speed is never larger than $20$ m/sec. Above $5$ km the wind speed remains extremely weak (the median value is smaller than $10$ m/sec) during the summer time. During the winter time the wind speed increases monotonically with the height and with an important rate reaching, at $25$ km, median values of the order of $30$ m/sec. A fluctuation of the order of $20$ m/sec is estimated at $20$ km between the summer and winter time.  \\item The atmosphere above Dome C shows a quite different regime of stability/instability in summer and winter time. During the summer time the Richardson number indicates a general regime of stability in the whole atmosphere. The turbulence can be triggered preferably at [$2$-$5$] km from the ground. During the winter time the atmosphere shows a more important turbulent activity. In spite of the fact that the analysis of the Richardson number in different months of the year is qualitative\\footnote{It does not provide a measure of the $\\CN2$ profiles but the relative probability to trigger turbulence in the atmosphere.} our predictions are consistent with preliminary measurements obtained above the site in particular period of the year. Considering the good reliability of the meteorological parameters retrieved from the ECMWF analyses the Richardson maps shown here should be considered as a reference to check the consistency of further measurements of the optical turbulence in the future.  \\item With the support of a simple model for the $\\CN2$ distribution, the Richardson number maps and the wind speed vertical profile we calculated a best $\\theta_{0,max}$$\\sim$$10$\\arcsec$-11$$\\arcsec$ and $\\tau_{0,max}$$\\sim$$16$ msec above Dome C during the summer time.  \\item The vertical distribution of all the meteorological parameters show a good agreement with measurements. This result is quite promising for the employing of the ECMWF analyses as initialization data for meso-scale models. Besides, it opens perspectives to employ ECMWF data for a characterization of meteorologic parameters extended over long timescale.  \\end{itemize}"
    },
    "0606/gr-qc0606014_arXiv.txt": {
        "abstract": "This Letter considers the generalized second law of gravitational thermodynamics in two scenarios featuring a phantom dominated expansion plus a black hole. The law is violated in both scenarios. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606053_arXiv.txt": {
        "abstract": "{ We studied the relation between the distribution of sunspot groups and the Gleissberg cycle. As the magnetic field is related to the area of the sunspot groups, we used area-weighted sunspot group data. On the one hand, we confirm the previously reported long-term cyclic behaviour of the sum of the northern and southern sunspot group mean latitudes, although we found a somewhat longer period ($P\\sim 104$ years). We introduced the difference between the ensemble average area of sunspot groups for the two hemispheres, which turns out to show similar behaviour. We also investigated a further aspect of the Gleissberg cycle where while in the $19$th century the consecutive Schwabe cycles are sharply separated from each other, one century later the cycles overlap each other more and more. ",
        "introduction": "The activity of the Sun has been studied in several ways for a long time. The oldest observed and recorded features on the solar disk are the sunspots. Regular sunspot observations were started by Galileo in 1610 and since then the data sets have been being continuously extended. Using these databases one has the opportunity to study the statistical properties of solar activity \\citep[see e.g.][]{vitinsky_etal86}. The most prominent cycle in the sunspot series and in all solar activity is the 11-year Schwabe cycle. The long-term trend in the amplitude of the Schwabe cycle is known as the secular Gleissberg cycle \\citep{gleissberg45}, which has a varying time scale of 80-120 years.  This paper aims to study the relation between the distribution of sunspot groups and the Gleissberg cycle. Although the origin of the Gleissberg cycle is unclear, it is generally accepted that the interaction of the solar differential rotation and the magnetic field plays a basic role in the generation of all solar activity. From the observational data sets a correlation can be seen between secular variations in magnetic activity and differential rotation variations \\citep{vitinsky_etal86,Yoshimura_Kambry93,Javaraiah03}. Although the role of the differential rotation in the long-term variation is not yet clear, several authors have investigated the interplay of differential rotation and the large-scale magnetic field, which drives the solar grand minima. Different modulational mechanisms have been proposed in the recent literature, namely (1) fluctuations in the $\\alpha$-effect which could cause the dynamo to stochastically flip and lead to modulation of magnetic energy; (2) a dynamic $\\alpha$-effect forced modulation; (3) two-layer dynamos with meridional circulation; (4) nonlinear dynamics via macroscopic (e.g.~Lorentz force) or microscopic (e.g.~$\\Lambda$-quenching) feedback \\citep[see a detailed review in][]{tobias02}.  The following study may provide some additional observational background for the investigation of solar dynamo models dealing with the description of the long-term evolution of solar activity. ",
        "conclusions": "In this paper we studied the long-term evolution of the distribution of sunspot groups from different aspects, e.g.~latitude, area and extension of activity. We applied two new methods, which in our opinion give a more realistic picture of the long-term variation of the toroidal magnetic field. Our method is based on the results of \\citet{Petrovay_Szakaly99}. They found that the latitudinal distribution of the magnetic field at the surface reflects the conditions at the bottom of the convective zone, i.e.~in this regard the convective zone behaves as a ``steamy window''. Although their original statement is restricted to the poloidal magnetic field, it can be assumed that this is also valid for the toroidal component (Petrovay, 2004, private communication). This is why we applied an area-weighted method, since the sunspot area probably reflects the toroidal magnetic flux.  First, we confirmed the results of \\citet{pulkkinenetal99} using the area-weighted mean latitude data. However, we found a longer period applying the discrete Fourier-transform and phase-dispersion method, respectively. Based on the $\\chi^2$ test, the period of $P_\\mathrm{PHD}=37\\,900$ days with an amplitude of $1.04\\degr$ turned out the best. We also studied the difference between the ensemble average area of sunspot groups for the two hemispheres, i.e.~ $(A_{30}/N_{30})_{\\mathrm{North}}-(A_{30}/N_{30})_{\\mathrm{South}}$, which also indicates that the dominance of the toroidal magnetic flux oscillates between the southern and northern hemispheres in a long time scale. For this dataset we a found somewhat longer period ($P=48\\,500$ days). Although the deviation of the two periods is greater than the formal errors ($\\pm100$ days), the real uncertainty is significantly larger due to the shortness of the observing window with respect to the length of the cycle and the large scatter. Thus, we do not think that the difference of the periods implies real physical effects.  \\citet{pulkkinenetal99} interpreted this latitudinal variation as a mixed parity mode in which a quadrupolar component is oscillating with this period. They examined the toroidal components of both the quadrupolar and the dipolar fields based on the paper of \\citet{Brooke_etal98}. They claimed that multiplying the absolute value of the total toroidal field by the signed mean latitude of the activity belt, this product should oscillate about zero. As the difference between the ensemble average area of sunspot groups for the two hemispheres, i.e.~$(A_{30}/N_{30})_{\\mathrm{North}}-(A_{30}/N_{30})_{\\mathrm{South}}$ is related to the strength and pattern of the toroidal magnetic field, our results may indicate observational evidence of such an oscillation. We note that \\citet{pipin99} also found an asymmetry of magnetic activity.  Beside studying the secular variation of the magnetic activity, we also investigated a further aspect of Gleissberg cycle. In Fig.~\\ref{fig:cshfull_intime}-\\ref{fig:avszfull_intime} it is visible that the sequential Schwabe cycles overlap each other at certain times. To characterize the rate of the overlapping we defined the daily latitudinal area-weighted extension of activity. We examined its mean and its standard deviation around the minima of the 11-year solar cycle. These quantities are increasing as the Gleissberg cycle tends to its maximum (see in Fig.~\\ref{fig:avsz_mean_stdev}). While in the northern hemisphere this tendency is obvious, in the southern one it is less so, but the tendency is visible. Nevertheless, as our investigation is limited to one Gleissberg cycle, we cannot declare unambiguously that this phenomenon correlates with the Gleissberg cycle. \\citet{Pelt_etal00} found a similar tendency in the distances in between nearby cycles \\citep[see Fig.~11 in][]{Pelt_etal00}. Hence, the continuous production of detailed, homogenous catalogues (e.g.~Debrecen Photoheliographic Data successor of the GPR) is crucial from this point of view. On the other hand, the interpretation of the variable amount of overlap beetween the Schwabe cycle is a challenge for future solar dynamo models."
    },
    "0606/astro-ph0606265_arXiv.txt": {
        "abstract": "The  mean electromotive force (MEMF) in a rotating stratified magnetohydrodynamical turbulence is studied. Our study rests on the  mean-field magnetohydrodynamics framework and $\\tau$ approximation. We compute the effects of the large-scale magnetic fields (LSMF),  global rotation and large-scale shear flow on the different parts of the MEMF (such as $\\alpha$ - effect, turbulent diffusion, turbulent transport, etc.) in an explicit form. The influence of the helical magnetic fluctuations which stem from the small-scale dynamo is taken into account, as well. In the paper,  we derive the equation governing the current helicity evolution. It is shown that the joint effect of the differential rotation and magnetic fluctuations in the stratified media can be responsible for the generation, maintenance and redistribution of the current helicity. The implication of the obtained results to astrophysical dynamos is considered as well. ",
        "introduction": "The mean-field magnetohydrodynamics presents one of the most powerful tools for exploring the nature of the large-scale magnetic activity in cosmic bodies \\citep{moff:78,park,krarad80}. It is widely believed that magnetic field generation there is governed by interplay between turbulent motions of electrically conductive fluids and global rotation. The growth and evolution of the large-scale magnetic fields (LSMF) in cosmic plasma strongly depends on the mean electromotive force, $\\mathbf{\\mathcal{E}}=\\left\\langle {\\mathbf{u}\\times\\mathbf{b}}\\right\\rangle $, which is given by the correlation between the fluctuating components of the velocity field of plasma, $\\mathbf{u}$, and the fluctuating magnetic fields, $\\mathbf{b}$.  The global rotation, stratification and the strong LSMF can substantially modify the structure and amplitude of the mean electromotive force (hereafter, MEMF) leading to the rich variety of the turbulence effects driving the evolution of the LSMF in cosmic bodies, e.g., the $\\alpha$-effect \\citep{rob-saw,moff:78,krarad80,park,kit-rud:1993b}, the rotationally-induced anisotropy of turbulent diffusion and effective drift of LSMF \\citep{rob-saw,krarad80,kit-pip-rud}, etc.. Generally speaking, the nonlinear effects of the small-scale Lorentz forces on the MEMF and LSMF evolution stem from two sources. One is driven by perturbations of the LSMF due to turbulent motions and another one is due to magnetic fluctuations, which are maintained by the small-scale dynamo action in a turbulent medium. The role of the small-scale dynamo in the LSMF evolution is still insufficiently understood. Numerous contributions to this subject can be found in the modern literature, e.g., \\citep{moff:78,pouquet-al:1975a,pouquet-al:1975b,bra-sub:04}. According to the mentioned studies the most important effect of the growing magnetic fluctuations on the LSMF evolution is caused by the helical part of magnetic fluctuations. The magnetic helicity conservation law, if applied to the mean-field magnetohydrodynamics, requires that the amount of helicity contained in the LSMF (controlled mostly by $\\alpha$-effect) should be roughly the same and opposite in sign to its counterpart in the small scales, see \\citep{kleruz82}. In this way the helical part of magnetic fluctuations, which is excited both due to shredding the LSMF by turbulent motions and due to small-scale dynamo, effectively saturates the generation of the LSMF by $\\alpha$-effect (\\citet{vain-kit:83,bran:01,bf:02a,black-bran:02}). Further discussions on this subject can be found in above cited papers. Their main lesson is that the construction of the realistic mean-field dynamo theory requires the evolution of the small scale magnetic (or current-) helicity to be taken into account.  Currently, there are two basic schemes for computing the MEMF of turbulent fields. One is the quasi-linear approximation (the same approximation is called the FOSA or SOCA in literature). A comprehensive discussion about its applicability and validity in astrophysics can be found at papers by \\citet{moff:78,park,krarad80,bra-sub:04}. This scheme remains one of the main tools of the mean-field magnetohydrodynamics. However, one of unfortunate problem of SOCA is that the contribution of the magnetic fluctuations (and the corresponding magnetic helicity) driven by the small-scale dynamo is hardly possible to include in the theory in self-consistent way. The third order closure scheme based on $\\tau$-approximation (\\citet{ors:70,vain-kit:83,rad-kle-rog,bra-sub:04}) gives a chance to consider, roughly, the effects of the small-scale dynamo on the MEMF. Following \\citep{bra-sub:04} (hereafter BS05), I will call it MTA (minimal tau approximation). Different kinds of this approximation are used in the literature, see \\citep{vain:83,vain-kit:83,rad-kle-rog,bra-sub:04,kle-rog:04a,bf:02,bf:02a}. In the paper we follow procedure described in BS05. Furthermore, the variant of tau approximation with a scale-independent relaxation time, $\\tau$, is applied. For this reason, some results obtained in the paper can be different of those that are given elsewhere: \\citep{kle-rog:04a,kle-rog:04b,kle-rog:04c,rad-kle-rog}.  The main purpose of this paper is to compute the MEMF via MTA taking into account the influence of the global rotation and LSMF on the turbulence. The stratification of the medium and the large-scale shear are taken into account as well. The influence of rotation, LSMF and uniform shear on the different parts of the MEMF (such as $\\alpha$ - effect, turbulent diffusion, turbulent transport and etc.) is explicitly defined via factors describing the efficiency of rotational and LSMF feedback on the turbulent flows. The influence of rotation is measured by the Coriolis number, $\\Omega^{\\ast}=2\\Omega\\tau_{c}$ , where $\\Omega$ is the solid body angular velocity and $\\tau_{c}$ - the typical correlation time of turbulent flows. The influence of LSMF is measured by $\\beta=\\bar{B}/\\left(u_{c}\\sqrt{\\mu\\rho}\\right)$, where $\\bar{B}$ is the strength of the LSMF, $u_{c}$ is a typical rms velocity of turbulent flows and $\\mu$, $\\rho$ are the magnetic permeability and the density of the media, respectively. Following the basic approach developed in above cited papers we derive the equations governing the evolution of the current helicity both in rotating and in magnetized turbulent flows with imposed uniform shear.  The paper is structured as follows. In the next section we shortly outline the basic equations, assumptions and the computational scheme for derivation of the MEMF and the evolutionary equation for current helicity. Section 3 is devoted to the results of calculations of the MEMF for different situations (slow rotation, strong LSMF, vice versa and etc.). In section 4 we derive the evolutionary equation for current helicity. In section 5 we summarize the main results of the paper. ",
        "conclusions": "In the paper the mean electromotive force of turbulent flows and magnetic fields is computed analytically using the framework of mean-field dynamo theory and MTA (minimal $\\tau$ approximation). There is an overlap in  results obtained with SOCA and with MTA approximation. The two approximations give qualitatively the same results about nonlinear dependence of mean electromotive force on the strength of LSMF or on the Coriolis number. Also there is a difference between predictions of  SOCA and  MTA for mean-electromotive force expressions if the shear is taken into account. This difference can be explained, in part, by the crudeness of the given version of tau approximation.  The deficient accuracy of calculations of shear contribution is due to an assumption about the scale-independent $\\tau$. In whole, the accuracy of the theory presented in the paper is comparable with the mixing-length approximation. It has no the firm grounds and should be considered with caution.  Finally, I would like to focus on the new findings of the paper. In this study it is shown that the new interesting component of transport of LSMF appears due to joint contribution of current helicity and shear. The effect does not disappear in the strong LSMF limit, $\\beta\\gg1$. It may be important near the base of solar CZ where the influence of rotation and shear on the turbulence is quite strong. Furthermore, the analysis, which we carried out for the current helicity evolution, suggests that the shear and rotation may redistribute the helicity in solar CZ amplifying it (in amplitude) at the near equatorial regions in agreement with observations. Beside, the effect of rotation and stratification on the $h_{\\mathcal{C}}$ evolution is calculated explicitly. Basically, the equation for current helicity is obtained using the same approach as for the mean electromotive force and on the base of quantities which are explicitly gauge invariants. Therefore, we can expect that the dynamo model based on the above approach could be capable for meeting the requirements of both solar and stellar dynamo simulations.   \\paragraph"
    },
    "0606/physics0606128_arXiv.txt": {
        "abstract": "Differential cross sections for electron collisions with the O$_2$ molecule in its ground ${X}^{3}\\Sigma_g^-$ state, as well as excited ${a}^{1}\\Delta_g$ and ${b}^{1}\\Sigma_g^+$ states are calculated. As previously, the fixed-bond R-matrix method based on state-averaged complete active space SCF orbitals is employed. In additions to elastic scattering of electron with the O$_2$ ${X}^{3}\\Sigma_g^-$, ${a}^{1}\\Delta_g$ and ${b}^{1}\\Sigma_g^+$ states, electron impact excitation from the ${X}^{3}\\Sigma_g^-$ state to the ${a}^{1}\\Delta_g$ and ${b}^{1}\\Sigma_g^+$ states as well as '6 eV states' of ${c}^{1}\\Sigma_u^{-}$, ${A'}^{3}\\Delta_u$ and ${A}^{3}\\Sigma_u^{+}$ states is studied. Differential cross sections for excitation to the '6 eV states' have not been calculated previously. Electron impact excitation to the ${b}^{1}\\Sigma_g^+$ state from the metastable ${a}^{1}\\Delta_g$ state is also studied. For electron impact excitation from the O$_2$ ${X}^{3}\\Sigma_g^-$ state to the ${b}^{1}\\Sigma_g^+$ state, our results agree better with the experimental measurements than previous theoretical calculations. Our cross sections show angular behaviour similar to  the experimental ones for transitions from the ${X}^{3}\\Sigma_g^-$ state to the '6 eV states', although the calculated cross sections are up to a factor two larger at large scattering angles. For the excitation from the ${a}^{1}\\Delta_g$ state to the ${b}^{1}\\Sigma_g^+$ state, our results marginally agree with the experimental data except for the forward scattering direction. ",
        "introduction": "A detailed knowledge of electron collisions with the oxygen molecule is important for the physics and chemistry of both laboratory and astrophysical plasmas. In particular, recent attempts to understand the electrical discharge oxygen-iodine laser have suggested that excited electronic states of O$_2$ molecule play an important role \\cite{Sh96,Gu04}. In a previous paper \\cite{Ta06} (henceforth denoted I), we studied integral cross sections for electron collisions with the O$_2$ molecule in its excited $a{}^{1} \\Delta_g$ and $b{}^{1} \\Sigma_g^{+}$ states, in addition to the much studied electron scattering by the O$_2$ ${X}^{3} \\Sigma_g^{-}$ ground state. We used the fixed-bond R-matrix method with 13 target states represented by valence configuration interaction wave functions. State-averaged complete active space SCF (SA-CASSCF) orbitals with Gaussian type basis functions were employed. The calculated cross sections for electron impact excitation from the ${a}^{1}\\Delta_g$ state to the ${b}^{1}\\Sigma_g^{+}$ state at 4.5 eV agree well with the available experimental data of Hall and Trajmar \\cite{Ha75}. Although elastic scattering of electrons by the ${a}^{1}\\Delta_g$ and ${b}^{1}\\Sigma_g^{+}$ states was also studied, we could not find any experimental data for comparison.  In I, theoretical and experimental integral cross sections were compared. However, differential cross sections (DCSs) provide a more stringent test of theory and are often easier to measure reliably than integral cross sections. For electron impact electronic excitations, calculations which give good integral cross sections often give DCS's which differ significantly from those observed experimentally. In this paper, we present DCSs for the corresponding processes calculated using the same R-matrix model.  Previous experimental and theoretical studies in the field of electron O$_2$ collisions are well summarized by Brunger and Buckman \\cite{Br02}. Here we only review  works relevant to this paper. The DCSs of electron collisions with O$_2$ molecule have been measured by many experimental groups. In particular, electron impact excitations to the low-lying ${a}^{1}\\Delta_g$ and ${b}^{1}\\Sigma_g^{+}$ states have been studied experimentally by Trajmar et al. \\cite{Tr71}, Shyn and Sweeney \\cite{Sh93}, Allan \\cite{Al95}, Middleton et al. \\cite{Mi94}, and Linert et al. \\cite{Li04b}. In contrast to these experimental works, only  Middleton et al. \\cite{Mi94} report calculations of DCSs for these excitation processes. Some of the more recent measurements have focused on electron impact excitations from the ${X}^{3} \\Sigma_g^{-}$ ground state to the `6 eV states', i.e., the ${c}^{1}\\Sigma_u^{-}$, ${A'}^{3}\\Delta_u$ and ${A}^{3}\\Sigma_u^{+}$ states which are also called the Herzberg pseudocontinuum \\cite{Ca00,Sh00,Gr02}. Although these DCSs are not state-resolved in most case, Shyn and Sweeney \\cite{Sh00} obtained cross sections for excitation to the individual electronic state within the `6 eV states'. In this paper, we also calculate the DCSs of this process using the fixed-bond R-matrix method, since no previous theoretical calculation exists. Up to now, there is only one measurement of DCSs for electron collisions with electronically excited O$_2$ molecule. Hall and Trajmar \\cite{Ha75} obtained excitation cross sections from the O$_2$ ${a}^{1}\\Delta_g$ to the ${b}^{1}\\Sigma_g^{+}$ state at electron impact energy of 4.5 eV. Their integral cross section was compared with our R-matrix results in I.  In this paper, details of the calculations are presented in section 2, and we discuss the results in section 3 comparing our results with previous theoretical and available experiments. Then the summary is given in section 4. ",
        "conclusions": "\\subsection{Electron collisions with the O$_2$ ${X}^3\\Sigma^{-}_{g}$ ground state}  Figure \\ref{fig1} shows DCSs for elastic electron scattering from the O$_2$ ${X}^3\\Sigma^{-}_{g}$ state compared with previous theoretical and experimental results. Our results are very similar to the previous R-matrix cross sections of {W\\\"oste} et al. \\cite{Wo95}. The cross sections of Machado et al. \\cite{Ma99} were calculated using the Schwinger variational iterative method combined with the distorted-wave approximation. Their results at 5 eV are much lower than the R-matrix results at low scattering angle below 50 degrees. Our results agree reasonably well with the experimental cross sections at 10 eV, including the recent results of Linert et al. \\cite{Li04} for backward scattering. At 5 eV, our model significantly overestimates the cross sections for forward scattering compared to the experimental values. For example, our result is twice as large as the experimental values at 10$^\\circ$. This situation is the same in the previous R-matrix calculation of {W\\\"oste} et al. \\cite{Wo95}. As discussed by {W\\\"oste} et al., this deviation can be attributed to a lack of long-range polarizability in the scattering model. For example, Gillan et al. \\cite{Gi88} introduced polarized pseudostates to account for the long-range polarizability in electron-N$_2$ scattering and reduced the cross sections by 50\\%\\ in the threshold energy region. The interaction potential of Machado et al. \\cite{Ma99} includes the correlation-polarization term based on free-electron-gas model. Probably the polarization introduced by this term is responsible for their better agreement with experiment at 5 eV. Since we are interested in electron collisions with the excited electronic states of O$_2$ in this work, we chose not to pursue precise accuracy further for the ground state elastic scattering. However, we have to be mindful that similar long-range polarizability problems may exist in the other low-energy electron scattering processes, especially elastic electron scattering of the ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ state O$_2$ molecules, which will be discussed below.  The DCSs for excitation to the ${a}^1\\Delta_{g}$ state at electron impact energy 5 and 10 eV are compared in figure \\ref{fig2} with the previous theoretical calculation and the experimental measurements of Middleton et al. \\cite{Mi94}, Shyn and Sweeney \\cite{Sh93}, Allan \\cite{Al95} and Linert et al. \\cite{Li04b}. The cross sections at 5 eV agrees well with the previous calculation and experimental data below 120$^\\circ$. However, our results are much lower than the previous calculation of Middleton et al. at scattering angle above 130$^\\circ$. At an electron scattering energy of 10 eV, our cross section deviates further from the previous calculation of Middleton et al. \\cite{Mi94} especially at scattering angle below 60$^\\circ$ and above 140$^\\circ$. In contrast to the backward enhanced cross sections of Middleton et al., our DCSs have a slightly forward enhanced character. Our results agree better with the experimental data at low scattering angles than the previous calculation of Middleton et al. \\cite{Mi94}. At large scattering angle, our results deviate from the experimental data of Shyn and Sweeney \\cite{Sh93}, but agrees rather well with the recent measurement of Linert et al. \\cite{Li04b}. The R-matrix model of Middleton et al. \\cite{Mi94} included the lowest 9 O$_2$ target states and $l$ = 0 - 5 scattering electron orbitals with $\\sigma$, $\\pi$ and $\\delta$ symmetry. In this work, we included 13 target states and all components of $l$ = 0 - 5 scattering electron orbitals. In addition to these differences, Middleton et al. used HF/STO orbitals where we employed CASSCF/GTO orbitals. We carried out a test calculation with  $l$ = 0 - 3 scattering electron orbitals and got almost the same cross sections as in $l$ = 0 - 5 case, which suggests that difference in the number of target states may be important for the shape of these excitation cross sections.  Figure \\ref{fig3} compares DCSs for excitation to the ${b}^1\\Sigma^{+}_{g}$ state at electron impact energy of 5 and 10 eV with the previous R-matrix calculation and the experimental measurements of Middleton et al. \\cite{Mi94}, Shyn and Sweeney \\cite{Sh93} and Allan \\cite{Al95}. Transitions between $\\Sigma^{+}$ and $\\Sigma^{-}$ target states are forbidden at scattering angles of 0$^\\circ$ and 180$^\\circ$, because the scattered electron wave function vanishes in the plane defined by incident electron beam and the molecular axis for any orientation of the molecule \\cite{Ca71,Go71}. As a consequence, the DCSs decrease to be zero toward 0 and 180$^\\circ$. As is apparent from figure \\ref{fig3}, our cross sections become zero at 0 and 180$^\\circ$, which is consistent with this selection rule. Compared to the previous R-matrix calculations of Middleton et al. \\cite{Mi94}, our cross sections have similar profile, but with slightly smaller magnitude at all scattering angles. Agreement with experiment is good at 5 eV below 120$^\\circ$, although our results underestimate the experimental cross sections at larger scattering angles. At 10 eV, the magnitude of the experimental cross sections of Middleton et al. \\cite{Mi94} and Shyn and Sweeney \\cite{Sh93} do not agree with each other, however our cross sections are closer to the results of Shyn and Sweeney at low scattering angles below 50$^\\circ$. Between 60$^\\circ$ and 90$^\\circ$, our results are closer to the results of Middleton et al.  Figure \\ref{fig4} shows DCSs for excitations to the '6 eV states' for electron impact energies of 10 and 15 eV. Here the '6 eV states' means the group of the O$_2$ ${c}^1\\Sigma^{-}_{u}$, ${A'}^3\\Delta_{u}$ and ${A}^3\\Sigma^{+}_{u}$ states. The cross sections shown in figure \\ref{fig4} are a sum of individual excitation cross sections of these 3 electronic states, in line with most experimental measurements. The figure includes the recent experimental cross sections of Campbell et al. \\cite{Ca00}, Green et al. \\cite{Gr02} and Shyn and Sweeney \\cite{Sh00}. The individual cross sections are shown in figure \\ref{fig5} and \\ref{fig6} for impact energies of 10 and 15 eV, together with the state-resolved experimental cross sections of Shyn and Sweeney \\cite{Sh00}. Our summed cross sections given in figure \\ref{fig4} are backward-enhanced for both the 10 and 15 eV cases, in accordance with the experimental cross sections Campbell et al. \\cite{Ca00} and Shyn and Sweeney \\cite{Sh00}. However, the forward enhancement of the DCSs at 10 eV observed by Green et al. \\cite{Gr02} is not reproduced by our calculation. The individual cross sections in figure \\ref{fig5} and \\ref{fig6} show similar angular behaviour compared to the experimental results of Shyn and Sweeney. However our DCSs for excitation to the ${A'}^3\\Delta_{u}$ state is more steep toward backward direction. Also the peak in the ${A}^3\\Sigma^{+}_{u}$ state cross sections is more pronounced in our calculation. Note that our results for the ${A}^3\\Sigma^{+}_{u}$ state become zero at 0 and 180$^\\circ$ as dictated by the $\\Sigma^{-}-\\Sigma^{+}$ selection rule.  \\subsection{Electron collisions with the O$_2$ ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ excited states }  The DCSs for elastic electron scattering with the excited O$_2$ ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ states are shown in figure \\ref{fig7}. We cannot compare them with previous theoretical or experimental work, since there is no available data. These DCSs show strong similarity with those of the elastic electron scattering with the ${X}^3\\Sigma^{-}_{g}$ ground state in figure \\ref{fig1}. The magnitude of these cross sections is almost the same for the 10 eV case. All of them have a large forward peak at 0$^\\circ$, a small rise in the cross sections at 180$^\\circ$. The location of the minimum moves inward from 140$^\\circ$ to 90$^\\circ$ as the electron scattering energy increases. This similarity is also reflected in the integral cross sections for elastic electron collisions with the ${X}^3\\Sigma^{-}_{g}$, ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ states. The profiles and magnitudes of the integral cross sections are basically the same for all these 3 electronic states as shown in I. The main configuration of these 3 electronic states has the form $\\left( {\\rm core} \\right) \\pi_g^4 \\pi_u^2 $, and this may be responsible for this similarity. Our R-matrix calculations tend to overestimate the elastic scattering cross sections of the ${X}^3\\Sigma^{-}_{g}$ state at low scattering angles, below 50$^\\circ$, compared to the experimental data. Considering the strong similarity of the cross section profiles for elastic scattering from excited states and the ground state, our calculations may also overestimate the cross section at low scattering angle at low electron impact energy.  In table \\ref{tab1}, we show momentum transfer cross sections for electron elastic scattering by the ${X}^3\\Sigma^{-}_{g}$, ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ states. As a consequence of the similarity in DCSs, the momentum transfer cross sections have a similar magnitude. Compared to the experimental data of Shyn and Sharp \\cite{Sh82} and Sullivan et al. \\cite{Su95}, our calculation overestimates the ${X}^3\\Sigma^{-}_{g}$ state momentum transfer cross section at 2 eV by 20 \\%, but underestimates the cross section at 10 eV by 8 \\% of Sullivan et al.'s value or 29 \\% of Shyn and Sharp's value. Our momentum transfer cross sections for the ${a}^1\\Delta_{g}$ and ${b}^1\\Sigma^{+}_{g}$ states may similarly be overestimates or underestimates depending on the electron impact energy.  Figure \\ref{fig8} shows DCSs for electron impact excitation from the ${a}^1\\Delta_{g}$ state to the ${b}^1\\Sigma^{+}_{g}$ state. The figure also includes the experimental data of Hall and Trajmar \\cite{Ha75} at impact energy of 4.5 eV. Our cross section profiles have characteristic features of minima around 10$^\\circ$ and 90$^\\circ$ and maxima around 50$^\\circ$ and 150$^\\circ$. They agree with the experimental cross sections of Hall and Trajmar \\cite{Ha75} within their error bars except at 20$^\\circ$ and 30$^\\circ$. The cross sections of Hall and Trajmar appear to increase from 50$^\\circ$ to 0$^\\circ$ whereas our cross sections decrease from 60$^\\circ$ toward 10$^\\circ$. In the 50$^\\circ$ -- 140$^\\circ$ angular region, Hall and Trajmar's cross sections vary less than ours. However, a precise comparison is difficult because of large error bars and lack of other experimental data."
    },
    "0606/astro-ph0606115_arXiv.txt": {
        "abstract": "We present results from the {\\it XMM-Newton} observation of the non-cooling flow cluster A~1060. Large effective area of {\\it XMM-Newton} enables us to investigate the nature of this cluster in unprecedented detail. From the observed surface brightness distribution, we have found that the gravitational mass distribution is well described by the NFW profile but with a central density slope of $\\sim$1.5. We have undoubtedly detected a radial temperature decrease of as large as $\\sim$30\\% from the center to the outer region ($r \\sim13'$), which seems much larger than that expected from the temperature profile averaged over nearby clusters. We have established that the temperature of the region $\\sim 7'$ southeast of the center is higher than the azimuthally averaged temperature of the same radius by $\\sim$20\\%. Since the pressure of this region already reaches equilibrium with the environment, the temperature structure can be interpreted as having been produced between $4 \\times 10^{7}$~yr (the sound-crossing time) and $3 \\times 10^{8}$~yr (the thermal conduction time) ago. We have found that the high-metallicity blob located at $\\sim 1.\\!'5$ northeast of NGC~3311 is more extended and its iron mass of $1.9 \\times 10^{7}$~M$_{\\odot}$ is larger by an order of magnitude than estimated from our {\\it Chandra} observation. The amount of iron can still be considered as being injected solely from the elliptical galaxy NGC~3311. ",
        "introduction": "Clusters of galaxies, being the largest virialized systems in the universe, are filled with the intracluster medium(ICM), which consists of X-ray emitting hot plasma with a typical temperature of a few $\\times 10^{7}$~K.  The ICM gives us rich information about the structure and evolution of clusters of galaxies.  X-ray Spectra of ICM immediately reveal its temperature and metal abundance, and the ICM surface brightness shows us ICM density distribution.  Through the assumption of hydrostatic equilibrium, furthermore, it is possible to measure the distribution of the gravitational mass in clusters.  The cluster mass distribution is an important observational clue in determining the nature of the dark matter and in constraining the structure formation scenario. In particular, the CDM model predicts a significant central cusp in the dark matter profile  \\citep{1996ApJ...462..563N}, in contrast to the flat core profile described by the $\\beta$ model. High-sensitivity X-ray study is a powerful method to look into the gravitational mass distribution (e.g. \\cite{2003ApJ...590..225C}). For clusters characterized by cool central cores, presence of the bright central galaxy greatly hampers the mass determination for the pure cluster component. Therefore, detailed study of a non cool-core cluster has a significant importance.  Recent observational studies of the clusters with {\\it Chandra} and {\\it XMM-Newton}, with their powerful imaging capability and large effective area, unveiled a lot of new interesting features. \\citet{2003ApJ...596..190M} found an unusual X-ray morphology, blob and hole like features, in the central dense region of the 2A~0335+096 cluster.  \\citet{2000ApJ...541..542M} detected sharp brightness edges, called cold fronts, with lower temperature inside the subcluster in A~2142 cluster.  \\citet{2003ApJ...596..181F} found two high-metallicity blobs located symmetrically with respect to the center of the poor cluster AWM~7.  \\citet{2005A&A...433..101P} investigated temperature profiles of 13 nearby cooling flow clusters observed with {\\it XMM-Newton}, and found that the temperature decreases by $\\sim 30$\\% between 0.1~$r_{vir}$ and 0.5~$r_{vir}$. This feature is in good agreement with the result of \\citet{1998ApJ...503...77M} who found that the temperature of 30 nearby clusters show a significant decline at large radii.  Recent {\\it Chandra} observation of A~1060 revealed that the central two elliptical galaxies emit X-ray \\citep{2002ApJ...578..833Y}, and the temperature and the abundance distribution are unexpectedly inhomogeneous \\citep{2004PASJ...56..743H}.  A~1060 is a nearby cluster ($z = 0.0114$) characterized by a smooth and symmetric ICM distribution.  The ASCA and ROSAT observations found that the temperature and the abundance showed very uniform distribution \\citep{1996PASJ...48..671T,2001PASJ...53..421F}.  For these reasons, this cluster was thought to be in a highly relaxed state, and has been used to estimate the gravitational mass profile which can be directly compared with the results of hydrodynamic cluster simulation \\citep{2000ApJ...535..602T}.  In this paper, we report the results of {\\it XMM-Newton} observation of the A~1060 cluster.  Based on the surface brightness distribution, we have computed the total gravitational mass of the cluster.  We confirmed the existence of the high-metallicity blob discovered by our {\\it Chandra} observation but with a larger iron mass content.  Using the high spatial resolution and the large effective area, we searched for a sign of a sub-cluster merger from a two-dimensional temperature map.  Throughout this paper, we assume $H_{0}=75$km~s$^{-1}$~Mpc$^{-1}$ with $q_{0}=0.5$.  Accordingly, an angular size of $1'$ corresponds to 13~kpc.  The solar number abundance of iron relative to H is assumed to be $4.68\\times 10^{-5}$ \\citep{1989GeCoA..53..197A}.  We employ Galactic absorption as $N_{\\rm H}=4.9 \\times 10^{20}$cm$^{-2}$. ",
        "conclusions": "We have presented the {\\it XMM-Newton} observation of the A~1060 cluster of galaxies. Large effective area of {\\it XMM-Newton} enables us to investigate spatial distribution of the surface brightness, the temperature and the metal abundance in unprecedented detail. We have derived a gravitational mass profile based on a model-independent estimation of the density profile. The overall shape can be well represented by Moore's profile with a central density slope $\\propto r^{-1.5}$. The temperature is found to decrease undoubtedly by 30~\\% from the center to the edge of the field of view $r = 13'$. This drop is clearly steeper than the Markevitch's composite profile.  Our spatially-resolved analysis revealed slight temperature enhancement at $r \\simeq 7'$ from NGC~3311 in the southeast direction. It is thought that some heating, like a subcluster merger, has occurred in this region in the past between the sound crossing time $4 \\times 10^{7}$~yr and the thermal conduction time $3 \\times 10^{8}$~yr ago. We note that the deviation of temperature profile of A~1060 from the Markevitch's one can also be explained if the central region is heated like the way seen in the southeast region.  The high resolution metallicity map for the central $r< 5'$ region confirmed the high-metallicity blob discovered with the {\\it Chandra} observation locating at around $1.\\!\\!'5$ northeast of NGC~3311. Its spatial extent is, however, found to be as large as $r \\sim 2'$, which is significantly larger than that estimated with {\\it Chandra} ($r\\sim 40''$). This discrepancy is probably because the {\\it Chandra} observation is photon-limited in the iron line band. The iron mass in this region is revised to be $1.9 \\times 10^{7}$~M$_{\\odot}$. Although this is larger than the {\\it Chandra} value roughly by an order of magnitude, it can still be supplied solely by NGC~3311."
    },
    "0606/astro-ph0606609_arXiv.txt": {
        "abstract": "We report the result of our near-infrared observations ($JH\\Ks$) for type II Cepheids (including possible RV Tau stars) in galactic globular clusters. We detected variations of 46 variables in 26 clusters (10 new discoveries in seven clusters) and present their light curves. Their periods range from 1.2 d to over 80 d. They show a well-defined period-luminosity relation at each wavelength. Two type II Cepheids in NGC~6441 also obey the relation if we assume the horizontal branch stars in NGC~6441 are as bright as those in metal-poor globular clusters in spite of the high metallicity of the cluster. This result supports the high luminosity which has been suggested for the RR Lyr variables in this cluster. The period-luminosity relation can be reproduced using the pulsation equation ($P\\sqrt{\\rho}=Q$) assuming that all the stars have the same mass. Cluster RR Lyr variables were found to lie on an extrapolation of the period-luminosity relation. These results provide important constraints on the parameters of the variable stars.  Using Two Micron All-Sky Survey (2MASS) data, we show that the type II Cepheids in the Large Magellanic Cloud (LMC) fit our period-luminosity relation within the expected scatter at the shorter periods. However, at long periods ($P>40$ d, i.e. in the RV Tau star range) the LMC field variables are brighter by about one magnitude than those of similar periods in galactic globular clusters. The long-period cluster stars also differ from both these LMC stars and galactic field RV Tau stars in a colour-colour diagram. The reasons for these differences are discussed. ",
        "introduction": "\\label{sec:Intro}  Type II Cepheids (hereafter T2Cs) are variables in the Cepheid instability strip, but belong to older populations than classical Cepheids. (see \\citealt{Wallerstein-2002}, and references therein, for a review). They reside in globular clusters, the thick disc, the bulge and the halo, but not in the thin disc or spiral arms. Based on their periods, they are often separated into BL Her stars ($P< 7$ d), W Vir stars ($ 7< P< 20$ d) and RV Tau stars ($P> 20$ d). The main feature of RV Tau stars is alternating deep and shallow minima. However, the classification and the nature of RV Tau stars are ambiguous. Several authors suggested so-called RV Tau stars include some heterogeneous types of variables. Whilst six objects in globular clusters have been claimed to be RV Tau stars, some authors doubted this classification from both the photometric \\citep{Zsoldos-1998} and the spectroscopic point of view \\citep{Russell-1998}. In this paper, we will not make a strict distinction between RV Tau stars and other T2Cs in clusters.  From previous studies of T2Cs in globular clusters, it is known that they obey period-luminosity relation (PLR) in the visible ($BVI$). \\citet{Harris-1985} and \\citet{McNamara-1995} claimed the slope of the PLR steepens for periods longer than about $\\log P= 1$. On the other hand, \\citet{Pritzl-2003} did not find such a feature for the variables in the globular clusters NGC~6388 and NGC~6441. As \\citet{Pritzl-2003} noted, many studies of the T2Cs were based on old photographic data, and we need more investigations with modern CCD photometry. In the near-infrared, no studies have so far been reported.  Studies of variable stars in the near-infrared have become more numerous in recent times. For example, many papers have been published on the infrared properties of RR Lyr variables (e.g. \\citealt{Clement-2001}; Castellani, Caputo \\& Castellani 2003). RR Lyr variables also lie in the Cepheid instability strip but are fainter than T2Cs. One of the important motivations for studies of RR Lyr variables is their application as distance indicators. Whilst a larger number of investigations have been devoted to their absolute visible magnitudes, studies in the infrared have some advantages. Longmore, Fernley \\& Jameson (1986) and Longmore et al. (1990) discovered a well-defined PLR in the near-infrared for the first time. It was suggested that the near-infrared relation is less affected by metal abundance than the visible one, making the near-infrared one a promising distance indicator. This led to further works (\\citealt{Butler-2003}; \\citealt{Dallora-2004}; \\citealt{Storm-2004}; \\citealt{DelPrincipe-2005}). Extensive theoretical studies of the RR Lyr PLR have been also carried out (Bono et al. 2001, 2003; Catelan, Pritzl \\& Smith 2004; Di Criscienzo, Marconi \\& Caputo 2004).  In this paper, we report the result of our near-infrared observations for T2Cs in globular clusters and present their PLR in $JH\\Ks$ filters. We also compare the PLR with that of RR Lyr variables and that of candidate T2Cs in the Large Magellanic Cloud (LMC). ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{Metallicity effect on the PLR}  First, we discuss the metallicity effect on the zero-point of the PLR, by comparing deviations from the PLR (\\ref{eq:PLR_J}) -- (\\ref{eq:PLR_K}) with the metallicity for each object. We simply adopted the metallicity of the globular cluster, in which a T2C is found, as the metallicity of the T2C. The relation in the $\\Ks$ filter is shown in Fig. \\ref{fig:deltaPLR}. The distance moduli we used equation (\\ref{eq:MvHB}) have of course a metallicity dependence by themselves. A linear regression for the data in Fig. \\ref{fig:deltaPLR} has a slope of $-0.10(\\pm 0.06)$ which is hardly significant. It would be reduced to $-0.02$ if we adopted the slope of 0.30 for equation (\\ref{eq:MvHB}) derived earlier by \\citet{Sandage-1993}. Note that adopting the latter slope makes negligible difference (less than 1 percent) to our PLR slopes. \\begin{figure} \\begin{center} \\includegraphics[clip,width=0.99\\hsize]{figure4.ps} \\caption{ The relation between the metallicity and the deviation from the PLR in $\\Ks$. \\label{fig:deltaPLR}} \\end{center} \\end{figure}   \\subsection{Extension of the PLR to RR Lyr region\\label{sec:RRcomp}}   We found that RR Lyr variables also obey the PLR (\\ref{eq:PLR_J}) -- (\\ref{eq:PLR_K}). Plus symbols in Fig. \\ref{fig:RRcomp} indicate RR Lyr variables in NGC~6341 (M~92) taken from \\citet{DelPrincipe-2005}. We adopted a distance modulus of 14.65~mag obtained in the same manner as for other clusters (equation \\ref{eq:MvHB}). This is the only cluster with RR Lyr observations at all three wavelengths. A comparison can be made at $\\Ks$ for a number of other clusters. As shown in Fig. \\ref{fig:compPLRs}, both the slope and the zero-point of the PLR agree satisfactorily with that of the RR Lyr variables in all cases. The data for the RR Lyr variables are from \\citep{Longmore-1990}, \\citep{Butler-2003}, \\citep{Storm-2004}, and \\citep{DelPrincipe-2005}. We averaged the results for the eight globular clusters in \\citet{Longmore-1990}. Although their magnitudes are $K$ (not $\\Ks$) in various photometric system, the differences are negligible (less than 0.01 mag; see \\citealt{Carpenter-2001}, for example). Theoretical studies also provided sufficiently close PLRs. For example, the slope in \\citet{Bono-2001} was $-$2.07 and that in \\citet{Catelan-2004} was $-$2.35. These works showed that there is a small metallicity dependence of the zero-point ($\\sim 0.17 \\log Z$ in $K$), but the effect is not clear in the observational results and must be small (\\citet{Longmore-1990} derived the metallicity-dependent term as $0.04$[Fe/H]). These results carry the implication that stars with the same age and probably the same mass within the instability strip obey the same PLR. We will discuss this in next section. \\begin{figure} \\begin{center} \\includegraphics[clip,width=0.99\\hsize]{figure5.ps} \\caption{ Comparison of the slope and the zero-point of the PLR of T2Cs (this work) and those of RR Lyr variables in references. L90 shows the result of \\citet{Longmore-1990}, B03 \\citet{Butler-2003}, S04 \\citet{Storm-2004} and DP05 \\citet{DelPrincipe-2005}. \\label{fig:compPLRs}} \\end{center} \\end{figure}  Some comments should be provided about NGC~6341 V7 \\citep{DelPrincipe-2005}, which has a period of $\\log P = 0.0259$ and is deviant from the PLR (the triangle in Fig. \\ref{fig:RRcomp}). Unfortunately, NGC~6341 lies too north to be observed by us. \\citet{Kopacki-2001} reported that this star is a BL Her star, but it is apparently brighter than expected from our PLR. Although we need to confirm its membership, the location of NGC6341 in the halo indicates that this star belongs to the cluster (see Section \\ref{sec:density}). In that case, this star may be the second anomalous Cepheid in globular clusters after NGC~5466 V19. An anomalous Cepheid is a more massive variable with $0.3 < P({\\rm d}) < 2$ and is brighter than RR Lyr variables by about one magnitude (\\citealt{Zinn-1976}; \\citealt{Cox-1988}). This star needs further investigation, including a confirmation of cluster membership. \\subsection{Reproduction of the PLR}  Here, we discuss the PLR by using the $P\\sqrt{\\rho} = Q$ relation. The relation can be written as \\begin{eqnarray} \\Mbol &=& -3.33 \\log P - 1.67 \\log M - 10 \\log \\Teff  \\nonumber \\\\ & & + M_{\\rm bol,\\odot} + 10 \\log T_{\\rm eff,\\odot} + 3.33 \\log Q, \\label{eq:Pul} \\end{eqnarray} where $M$ is the mass in units of the solar mass, and $\\Teff$ the effective temperature (e.g. \\citealt{McNamara-1995}). Two linear relations, \\begin{eqnarray} \\log \\Teff &=& -0.058 \\log P + 3.81 , \\label{eq:TP} \\\\ \\log Q &=& 0.24 \\log P -1.39 , \\label{eq:QP} \\end{eqnarray} are adopted from \\citet{McNamara-1994} to derive the PLR. We express the bolometric correction as \\begin{equation} \\Mbol - M_\\lambda = \\alpha _\\lambda \\log \\Teff + \\beta _\\lambda . \\label{eq:BCP} \\end{equation} Using these relations, eq (\\ref{eq:Pul}) can be expressed in the form, \\begin{eqnarray} M_\\lambda &=& -(1.95-0.058 \\alpha _\\lambda) (\\log P-1.2) -1.67\\log {\\rm M} \\nonumber \\\\ & & - (2.70 +3.74\\alpha _\\lambda +\\beta _\\lambda) \\label{eq:Pul2} \\end{eqnarray} with the constants of $M_{\\rm bol,\\odot}=4.75$ and $T_{\\rm eff,\\odot}=5780{\\rm K}$. We define $\\mu _\\lambda$ and $\\eta _\\lambda$ as the dependence on the period (${\\rm d}M_\\lambda /{\\rm d}\\log P$) and the zero-point of the relation at $\\log P=1.2$, respectively [i.e. $M_\\lambda = \\mu _\\lambda (\\log P -1.2) + \\eta _\\lambda$]. If the mass term ($\\log M$) has no dependence on the period, $\\mu _\\lambda$ equals to $-(1.95-0.058\\alpha _\\lambda)$. However, Bono, Caputo \\& Santolamazza (1997) predicted that the mass of T2Cs varies from $0.59M _\\odot$ to $0.52M _\\odot$, decreasing with increasing period from 1 to 10 d. This period dependence increases $\\mu _\\lambda$ by 0.08 compared with the case of the constant mass.  We obtained $\\alpha _\\lambda$ and $\\beta _\\lambda$ in Table \\ref{tab:PLRslope} from the model atmospheres listed in table 1 of Bessell, Castelli \\& Plez (1998). % They listed both the models with overshooting (table 1) and those without overshooting (table2), but the difference between the two sets has negligible effect on our results (up to 0.03 mag). Whilst these models are computed with the solar metallicity, Sandage, Bell \\& Tripicco (1999) computed model atmospheres for Cepheids between [Fe/H]$=0.0$ and $-1.7$. Their results show that the slopes of the $\\log \\Teff$-bolometric correction relation (equation \\ref{eq:BCP}) are within the uncertainty of our adopted values, and the zero-points get slightly smaller for the lower metallicity (about 0.1-mag difference between [Fe/H]$=0.0$ and $-1.7$). In $JH\\Ks$ filters we would expect any effect of metallicity on equation (\\ref{eq:BCP}) to be less than in $V$ and $I$.  The third and fourth columns in Table \\ref{tab:PLRslope} show the  predicted slope $\\mu _\\lambda$ and zero-point $\\eta _\\lambda$ in case of constant mass. They are approximately consistent with the counterparts obtained from the observational data (the fifth and sixth columns). The observational values in $V$ and $I$ filters are taken from \\citet{Pritzl-2003} and those in $JH\\Ks$ filters are obtained by us. Fig. \\ref{fig:PLRslope} shows the relation between $\\alpha _\\lambda$ and $\\mu _\\lambda$. The solid line shows the case of constant mass, and the broken one shows the one shifted by 0.08 for the mass-dependent case. The observational values (filled circles) favour the constant mass model at least in our own data ($JH\\Ks$). In $V$, the slope expected from equation (\\ref{eq:Pul2}) is steeper than the observational value. A linear fit to the bolometric correction (equation \\ref{eq:BCP}) is not very  good in $V$, since a quadratic term is no longer small unlike the case of $JH\\Ks$. The value of $\\alpha _V$ in equation (\\ref{eq:BCP}) can range from 0 to 4 between the extreme values of $\\log \\Teff$, $3.85$ and $3.70$. If we adopt a quadratic relation instead of equation (\\ref{eq:BCP}), the PLR also become quadratic and the slope of a linear fit to the entire period range is about $-1.7$ in $V$, which is close to the observational value ($-1.64$). Since our discussion is based on very simple scheme, more detailed work on both the theoretical and observational side is desirable.  It is worth noting that T2Cs and RR Lyr variables may be unique in that they comprise a group of variables with almost a constant mass obeying PLRs. For instance in the cases of classical Cepheids and Mira variables the mass increases with increasing period. It is therefore interesting that our PLRs can be reproduced by the simple scheme with a constant mass and no need to adopt a mass-luminosity relation. \\begin{table*} \\begin{minipage}{120mm} \\caption{The relation between the adopted bolometric correction ($\\alpha _\\lambda, \\beta _\\lambda$) and the PLR (slope $\\mu _\\lambda$ and zero-point $\\eta _\\lambda$). Observational values for $JH\\Ks$ are from the present paper, and those for $V$ and $I$ filters are from \\citet{Pritzl-2003}. \\label{tab:PLRslope}} \\begin{center} \\begin{tabular}{clrllll} \\hline Filter & \\multicolumn{1}{c}{$\\alpha _\\lambda$} & \\multicolumn{1}{c}{$\\beta _\\lambda$} & \\multicolumn{2}{c}{Equation (\\ref{eq:Pul2})} & \\multicolumn{2}{c}{Observation} \\\\ & & & \\multicolumn{1}{c}{$\\mu _\\lambda$} & \\multicolumn{1}{c}{$\\eta _\\lambda$} &  \\multicolumn{1}{c}{$\\mu _\\lambda$} & \\multicolumn{1}{c}{$\\eta _\\lambda$} \\\\ \\hline $V$   & $+2.1\\pm 1.0$ & $-7.06$ & $-1.85$ & $-1.94$ & $-1.64\\pm 0.05$ & $-1.92$ \\\\ $I$   & $-2.0\\pm 0.5$ &   9.11  & $-2.08$ & $-2.77$ & $-2.03\\pm 0.03$ & $-2.80$ \\\\ $J$   & $-4.75\\pm 0.1$ &  19.40  & $-2.23$ & $-3.34$ & $-2.23\\pm 0.07$ & $-3.54$ \\\\ $H$   & $-7.0\\pm 0.1$ &  27.66  & $-2.36$ & $-3.75$ & $-2.34\\pm 0.06$ & $-3.94$ \\\\ $\\Ks$ & $-7.2\\pm 0.1$ &  28.47  & $-2.37$ & $-3.81$ & $-2.41\\pm 0.06$ & $-4.00$ \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*} \\begin{figure} \\begin{center} \\includegraphics[clip,width=0.99\\hsize]{figure6.ps} \\caption{The relation between the adopted bolometric correction ($\\alpha _\\lambda$) and the slope of the PLR ($\\mu _\\lambda$). The solid line indicates the relation according to equation (\\ref{eq:Pul2}) in case of the constant mass, and the broken one is the one shifted by 0.08 considering the mass dependency mentioned in the text. Filled circles show the observational results. \\label{fig:PLRslope}} \\end{center} \\end{figure}   \\subsection{Comparison with T2Cs in the LMC}  \\citet{Alcock-1998} reported 33 candidate T2Cs with $8< P({\\rm d}) <100$ based on the Massive Compact Halo Object (MACHO) data base in the LMC. We searched for their near-infrared magnitudes in the 2MASS point-source catalog\\citep{Curti-2003}, and found 27 matches among 33 objects (Table \\ref{tab:LMC_Cep}). Fig. \\ref{fig:LMCcomp} shows the PLR in $JH\\Ks$ for the T2Cs in the LMC (crosses). Also plotted (filled circles) are the T2Cs in globular clusters. The absolute magnitudes for the LMC objects were obtained with an assumed distance modulus of 18.50 mag. There are some uncertainties of about $\\pm 0.3$~mag in using the 2MASS data because they are based on single-epoch observations. However, it seems rather clear that longer period LMC variables ($P>40$d) are brighter than the counterparts in globular clusters. The LMC variables at shorter periods fit the PLR of globular clusters within the uncertainties. This feature is also seen in the $\\log P$-$V$ diagram (fig. 9) in \\citet{Pritzl-2003}, but their conclusion is somewhat uncertain because they do not have their own data for the cluster variables with $P > 40$ d [compare the panels (a) and (d) of their fig. 9].  A mass difference could be one of the reasons for the LMC RV Tau stars being about one magnitude brighter than the globular cluster PLRs. Massive variables are expected to be brighter according to equation (\\ref{eq:Pul2}). A difference of about 1 mag corresponds to an increase in mass by a factor of  about 4. Considering that the masses of the globular cluster variables are about 0.5 -- 0.6 $M _\\odot$, this would lead to a mass larger than the Chandrasekhar limit for the LMC RV Tau stars. This is too large if these variables are post-AGB stars which have already gone through their major mass-loss phase (e.g. \\citealt{Jura-1986}; \\citealt{Pollard-1999}). In that case some other parameter is necessary to explain the difference in absolute magnitude.  Fig. \\ref{fig:CCD_RV} shows a colour-colour diagram for the variables with $P> 20$ d in globular clusters (filled circles) and the LMC (crosses) and also for the galactic field RV Tau stars (triangles) taken from \\citet{LloydEvans-1985}. The colours of our cluster objects and the LMC objects were corrected for reddenings, while those of the galactic field objects were not because \\citet{LloydEvans-1985} did not give any estimate of the reddenings. More than half of the galactic field objects have the galactic latitudes of $|b|>5\\degr$ so that the reddening effect on the colour is expected to be not large ($E(H-K) < 0.1$). The large excesses of the $K-L$ colours reported by \\citet{LloydEvans-1985}, which should be smaller than $E(H-K)$ in the case of interstellar reddening, also support that the objects are intrinsically red. In Fig. \\ref{fig:CCD_RV}, the globular cluster sample occupies a rather limited region whilst many of the LMC and the local stars spread to redder $H-\\Ks$ colour. As mentioned in Section \\ref{sec:Intro}, the classification of the longer period variables ($P> 20$ d) in globular clusters as RV Tau stars is unclear. One of the characteristics often seen in RV Tau stars is an infrared excess caused by their circumstellar dust shells \\citep{Jura-1986}. The only globular cluster RV Tau star which has been claimed to have an infrared excess is NGC~6626 V17. \\citet{Nook-1989} found an excess at 10 $\\mu$m in this star. However, NGC~6626 V17 which has $(H-K)_0=0.12$ and $(J-H)_0=0.48$ lies with the other cluster stars in Fig. \\ref{fig:CCD_RV}, and it also lies on our PLRs. We conclude that RV Tau stars in the LMC belong to a different family of variables from the T2Cs of the same periods in globular clusters. Whether they define a PLR is not clear. \\begin{table} \\begin{minipage}{86mm} \\begin{center} \\caption{2MASS magnitudes for candidate T2Cs in the LMC. Star IDs are from \\citet{Alcock-1998}. \\label{tab:LMC_Cep}} \\begin{tabular}{lrrrr} \\hline Star ID & $P$ & $J$ & $H$ & $K_s$ \\\\ \\hline 1.3812.61    &  9.387 & 15.453 & 14.970 & 15.222 \\\\ 10.4040.38   &  9.622 & 14.635 & 14.153 & 13.998 \\\\ 80.6469.135  & 10.509 & 15.749 & 15.483 & 15.125 \\\\ 80.6590.137  & 11.442 & 15.347 & 14.824 & 14.771 \\\\ 3.7332.39    & 12.704 & 15.868 & 15.380 & 15.178 \\\\ 80.6475.2289 & 13.925 & 15.003 & 14.611 & 14.373 \\\\ 81.9006.64   & 14.337 & 15.053 & 14.647 & 14.469 \\\\ 47.2611.589  & 14.469 & 15.793 & 15.326 & 15.183 \\\\ 19.4425.231  & 14.752 & 14.963 & 14.899 & 14.473 \\\\ 2.5877.58    & 14.855 & 15.857 & 15.389 & 15.033 \\\\ 1.3808.112   & 14.906 & 15.209 & 14.758 & 14.703 \\\\ 14.8983.1894 & 15.391 & 14.988 & 14.500 & 14.526 \\\\ 2.5025.39    & 16.602 & 14.723 & 14.369 & 14.368 \\\\ 9.5117.58    & 16.747 & 14.697 & 14.394 & 14.199 \\\\ 10.3680.18   & 17.127 & 14.836 & 14.424 & 14.254 \\\\ 78.6338.24   & 17.560 & 14.354 & 14.120 & 14.002 \\\\ 2.5026.30    & 21.486 & 14.571 & 14.136 & 13.984 \\\\ 78.6698.38   & 24.848 & 14.281 & 13.823 & 13.463 \\\\ 77.7069.213  & 24.935 & 15.364 & 14.676 & 14.457 \\\\ 82.8041.17   & 26.594 & 14.584 & 14.088 & 13.877 \\\\ 19.6394.19   & 31.716 & 14.011 & 13.641 & 13.356 \\\\ 78.5856.2363 & 41.118 & 13.668 & 13.250 & 13.179 \\\\ 81.8520.15   & 42.079 & 13.547 & 13.251 & 13.166 \\\\ 82.8405.15   & 46.542 & 13.114 & 12.884 & 12.569 \\\\ 81.9728.14   & 47.019 & 13.201 & 12.635 & 12.098 \\\\ 79.5501.13   & 48.539 & 13.089 & 12.635 & 12.093 \\\\ 47.2496.8    & 56.224 & 13.124 & 12.709 & 12.512 \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table} \\begin{figure} \\begin{center} \\includegraphics[clip,width=0.99\\hsize]{figure7.ps} \\caption{ The PLR for the globular cluster sample (filled circles) and the LMC candidates (crosses). \\label{fig:LMCcomp}} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[clip,width=0.99\\hsize]{figure8.ps} \\caption{ A colour-colour diagram for variables with $P>20$ d: filled circles for globular cluster stars, crosses for the LMC stars from \\citet{Alcock-1998} and triangles for galactic field stars\\citep{LloydEvans-1985}. The colours of the globular cluster stars and the LMC stars were corrected for the reddenings, while those of the galactic field stars were not (see the text). \\label{fig:CCD_RV}} \\end{center} \\end{figure}"
    },
    "0606/astro-ph0606323_arXiv.txt": {
        "abstract": "We present the first results from the Distant Radio Galaxies Optically Non-detected in the SDSS ({\\em DRaGONS}) Survey. Using a novel selection technique for identifying high redshift radio galaxy (HzRG) candidates, a large sample is compiled using bright ($S_{\\rm 1.4GHz} >\\,100$mJy) radio sources from the Faint Images of the Radio Sky at Twenty centimeters ({\\em FIRST}) survey having no optical counterpart in the Sloan Digital Sky Survey ({\\em SDSS}). Near-infrared (NIR) $K$-band imaging with the {\\em FLAMINGOS} instrument on the 4-meter telescope at Kitt Peak for 96 such candidates allows preliminary identification of HzRG candidates through the well-known $K-z$ relation, and these objects will subsequently be observed spectroscopically to confirm their redshifts. Of the initial candidates, we identify 70 with magnitudes brighter than $K\\,\\approx\\,19.5$, and compute limiting magnitudes for the remainder.  Assigning redshifts based on a linear fit to the $K-z$ Hubble diagram gives a mean redshift for our sample of $z=2.5$ and a median redshift of $z=2.0$, showing that this method should be very efficient at identifying a large number of HzRGs. This selection is also sensitive to a previously unseen population of anomalously red radio galaxies ($r-K\\,>\\,6.5-7$), which may indicate significant obscuration at moderate redshifts.  These obscured objects can be used to test the completeness of QSO surveys to the effects of reddenning.  More than ten percent of our sample falls into this category, which may represent a sizable radio loud population missing from current optically selected AGN samples. We additionally identify 479 bright Extremely Red Objects (EROs) in the fields surrounding our HzRG candidates, to a magnitude of $K=17.5$, within a non-contiguous area of 2.38 square degrees.  This constitutes a small overdensity of EROs surrounding the radio galaxy candidates over random fields, and we see possible evidence for a physical association of the EROs with the radio galaxies. However, examining the clustering of {\\em all} $K\\le19.0$ galaxies around the radio targets reveals no evidence of a global galaxy excess, strengthening our conclusion that the EROs trace an overdensity not evident in the overall galaxy population. ",
        "introduction": "\\label{int} \\footnotetext[1]{Based on observations collected at Kitt Peak National Observatory, National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc. (AURA) under cooperative agreement with the National Science Foundation.} Radio galaxies have long been used to probe the epoch of galaxy formation \\citep[]{Lil:84,Gra:96,Blu:98,Ste:99,Jar:01}. Luminous radio galaxies are known to be highly biased relative to the underlying dark matter, residing in the most overdense regions of the Universe.  Under the standard Cold Dark Matter hierarchy for galaxy formation these galaxies should be the first systems to collapse and, therefore, the site of some of the most evolved stellar populations.  As the likely hosts of the earliest star formation, it may also be possible to probe the epoch of reionization \\citep[]{Bark:05}.  Isolating a sample of radio galaxies at high redshift would allow us to probe the physical processes that drive the formation and evolution of structure and the timescales that govern star formation in the early universe, as well as to distinguish between hierarchical and ``down-sizing\" \\citep[]{Cow:96} formation scenarios \\citep[e.g.,][]{Roc:04,Del:05}, and also the role of feedback on the local radio galaxy environment \\citep[]{Var:05, Del:05}.  While radio surveys can address many fundamental questions in cosmology and galaxy formation the numbers of radio galaxies identified at high redshift remain small in comparison to low redshift surveys \\citep[]{Bra:05,Mag:04}.  Without large, statistically complete and homogeneously selected samples we cannot hope to constrain hierarchical galaxy formation models without the concern that sample variance might bias our analyses.  For example, at redshifts $z>3$ extensive optical and near-infrared (NIR) campaigns  have yielded less than one hundred galaxies \\citep[]{VanB:98,Deb:01,Var:05}.  The reason for the paucity of these samples comes from the necessity of surveying large volumes to identify the most massive systems.  Given the broad redshift distribution of radio galaxies \\citep[e.g.,][]{DP:90}, large numbers of radio targets must be observed in order to extract the high redshift component.  In this paper we describe a novel selection technique that utilizes existing optical and radio surveys to overcome these challenges in order to isolate high redshift candidates for follow up in the NIR. We show that this approach gives more than a factor of ten increase in efficiency over blind radio selected surveys.  The resulting combination of NIR data with existing optical measurements from SDSS allows us to study the environment of the high redshift radio galaxies.  Early in the development of infrared observations it was noted that EROs are often found in the vicinity of high redshift objects \\citep[]{Mcc:92,Gra:94,Dey:95}.  More recently, targeted searches have been undertaken to search for clustering of galaxies \\citep[]{Hal:01} and extremely red objects (EROs) around high redshift quasars and radio galaxies \\citep[][and references therein]{Cim:00,Wold:03,Zhe:05}. The focus of these investigations has been whether the increased surface density of EROs is physically associated with the radio galaxy or quasar, or whether it lies in the foreground, possibly as a cluster that may have gravitational lensing effects on the target.  Also of interest in a near-infrared radio survey is the number of obscured quasars.  It is well know that optical selection is sensitive to a wide range of effects that bias samples against the detection of heavily obscured galaxy populations.  Indeed, heavily obscured quasars may account for a significant fraction of the total population \\citep[]{Web:95,Whi:03,Gli:04}.  Our selection criteria result in the detection of a population of bright ($K\\,<\\,17.5$), red ($r-K\\,>\\,6.5$) objects that may fall in to this category.  The structure of this first in a series of papers is as follows:  The target selection procedure is presented in \\S\\,\\ref{target}.  A description of the observations, astrometry and photometry is given in \\S\\,\\ref{obs}.  The results are presented and discussed in \\S\\,\\ref{res}, and a summary and discussion of future work is presented in \\S\\,\\ref{summ}.  SDSS magnitudes are in the AB system, while we will use Vega magnitudes for all $K$-band measurements for easy comparison with the high redshift radio galaxy literature. A flat Lambda cosmology is assumed throughout, as indicated by numerous recent measurements, with $H_0=70\\,$km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ \\citep[e.g.,][]{Spe:03}. ",
        "conclusions": "\\label{summ} Using a novel radio-optical selection technique, we have obtained NIR imaging for 96 HzRG candidates. Based on redshift estimates from the $K-z$ Hubble diagram, our new selection technique appears to be effective at identifying HzRGs, as more than half of our targets have $K$-band photometry consistent with $z>2$. This technique is not sensitive to radio spectral slope, and avoids the frequency dependence of USS techniques for galaxies with non-power law radio spectra.  Of the 75 target galaxies present in the Texas 365 MHz survey 24 have $\\alpha^{\\rm 1.4GHz}_{\\rm 365MHz} <0.8$. These would be excluded by the USS criteria of \\citet[]{VanB:98}, and all but four of these 75 have $\\alpha^{\\rm 1.4GHz}_{\\rm 365MHz} <1.3$, which would be excluded from more recent USS selections \\citep[e.g.][]{Deb:01,Jar:01}. Comparison with the previous surveys of \\citet[]{Bes:99} and \\citet[]{Deb:01} shows that our technique selects nearly all of their high redshift targets, while also eliminating the low redshift sources. The $K$-band number counts are consistent with previous work, and the large non-contiguous area (2.38 square degrees) covered by the {\\em DRaGONS} survey to date makes it an excellent resource for exploring galaxy properties through the combination of NIR and optical data.  We have uncovered a previously unseen class of radio sources with anomalously red colors ($r-K\\,>\\,6.5-7$), which may be evidence of significant obscuration at moderate redshifts.  These galaxies represent more than ten percent of our observed sample, indicating that they may be a substantial percentage of the radio galaxy population.  Being non-detections in SDSS (with $r\\,>\\,24.1$), current optical AGN selection techniques are insensitive to these sources.  Radio loud QSOs comprise only $\\sim\\,5-10\\%$ of overall QSOs; therefore, these objects could represent a significant contribution to radio loud AGNs that are not counted in current samples.  We have identified a sample of 479 EROs to a depth of $K=17.50$, one of the largest bright ERO samples to date, and have identified a modest overdensity of EROs around our HzRG candidates. We do not detect any similar overdensity when all the $K$-band selected galaxies to $K=19.0$ are considered.  Spectroscopic follow up of the radio galaxy candidates and the surrounding EROs is of the utmost importance, both for validation of our HzRG candidates, as well as determination of the relationship between the radio galaxies and the ERO overdensity. Firm redshifts will also allow us to accurately test the $z-\\alpha$ relation for our non-USS selected sample.  Near infrared spectroscopy of the anomalously red radio galaxies should pinpoint the cause of their extreme colors, as well as determine the amount of obscuration present at intermediate and high redshifts.  Deeper optical imaging is required in order to establish colors for the radio galaxies and explore their morphology in more detail. Deeper optical imaging will also allow us to select a deeper and more complete sample of EROs, allowing for a more detailed study of the radio galaxy environment and the nature of the ERO overdensity. A more quantitative study of ERO clustering will also be possible with deeper optical data (including a two-point angular correlation function analysis).  A future paper will also examine the effects of both our radio flux density cut, as well as our optical color cuts.  In future observing runs we plan to target a small sample of lower flux density radio sources that meet our optical color cuts.  We will supplement this sample with galaxies from the FLAMEX survey that meet our selection criteria.  This dataset will allow us to examine any biases introduced by our radio selection.  Several of the targets in our most recent round of observing lie in the SDSS Southern Survey, an area of SDSS that is more than four times more sensitive than the standard SDSS survey.  This deeper imaging data will allow us to study the optical properties of those targets that are detected."
    },
    "0606/astro-ph0606751.txt": {
        "abstract": "%context {Satellite accretion events have been invoked for mimicking the internal secular evolutionary processes of bulge growth. However, N-body simulations of satellite accretions have paid little attention to the evolution of bulge photometric parameters, to the processes driving this evolution, and to the consistency of this evolution with observations. %\\mbremark{no estoy seguro de lo que quiere decir esto: ''No clear trace of a connection between satellite accretion and secular evolution has been reported to date either.''} } %aims {We want to investigate whether satellite accretions indeed drive the growth of bulges, and whether they are consistent with global scaling relations of bulges and discs. %\\mbremark{quitamos esto tambien? '', and if infall-driven bulge growth is truly related to secular evolution at some point.''} } %methods {We perform $N$-body models of the accretion of satellites onto disc galaxies. A Tully-Fisher ($M\\propto V_{\\rm rot}^\\alphaTF$) scaling between primary and satellite ensures that density ratios, critical to the outcome of the accretion, are realistic.  We carry out a full structural, kinematic and dynamical analysis of the evolution of the bulge mass, bulge central concentration, and bulge-to-disc scaling relations.} %results { %The outcome of the accretion is sensitive to the density ratio between primary and secondary; The remnants of the accretion have bulge-disc structure. Both the bulge-to-disc ratio ($B/D$) and the S\\'ersic index (\\emph{n}) of the remnant bulge increase as a result of the accretion, with moderate final bulge S\\'ersic indices: $n = 1.0$ to 1.9.  Bulge growth occurs no matter the fate of the secondary, which fully disrupts for $\\alphaTF=3$ and partially survives to the remnant center for  \\alphaTF = 3.5 or 4. %$\\alphaTF=3$ leads to complete disruption of the secondary before reaching the primary center, while, for \\alphaTF = 3.5 or 4, the cores of the satellites reach the center. %No matter what the fate of the secondary is, however, Global structural parameters evolve following trends similar to observations. We show that the dominant mechanism for bulge growth is the inward flow of material from the disc to the bulge region during the satellite decay.} %conclusions {The models confirm that the growth of the bulge out of disc material, a central ingredient of secular evolution models, may be triggered externally through satellite accretion.}  % ",
        "introduction": "\\label{Sec:Introduction}  Both observationally and theoretically, the central bulges of disc galaxies have received much attention since the last decade \\citep[see reviews by ][]{Wyse97,Kormendy04}.  However, we are far from understanding the processes that led to their formation.  Several scenarios have been proposed.  The similarities between bulges and ellipticals have prompted models of fast, early formation from the collapse of a protogalactic halo \\citep{Eggen62,Gilmore98}, or as the result of violent interactions and major mergers of collapsed pre-galactic clumps prior to the formation of the thin disc \\citep[][]{Kauffmann93,Kauffmann94,Kauffmann96}. But the realization that the \\emph{bulges are old} paradigm is not entirely true \\citep[][]{Kormendy93} has led to the development of models in which bulges grow over long timescales, after the formation of the disc. Processes that take many galactic rotation periods to complete \\citep[of the order of several Gyr, see KK04; ][]{Athanassoula05} have been termed \\emph{secular evolution} processes. They include bulge growth from an internal disc bar instability followed either by a vertical instability in the bar itself \\citep{Hohl71,Combes81,Pfenniger90, Debattista04}, or by bar dissolution due to a central mass concentration \\citep{Hasan90,Hasan93}. KK04 list  satellite accretion as a secular process as well, triggered by the environment of the galaxy rather than internally.  Bulge formation/growth could come from the accreted satellites themselves, from the disc \\citep{Pfenniger93,Mihos95}, or from a combination of both \\citep{Aguerri01}.  %Gerhard81,Barnes88,Moore96,Walker96,Dubinski98,Naab99,Velazquez99,Bendo00,Cretton01,Naab03,Bournaud04,Bournaud05,Gonzalez05,Aceves05a,Aceves05b,Aceves06}.  Minor mergers have been extensively studied using numerical simulations in the past \\citep[][to name just a few]{Quinn93,Velazquez99,Bournaud04,Bournaud05, Aceves05a,Aceves05b,Aceves06}. However, interest in satellite accretion has focussed more on the effects on the disc, such as thickening and warping \\citep{Walker96,Huang97,Helmi99}, with little attention given to the evolution of the bulge photometric parameters and to the processes that drive this evolution.  Simulations can easily tell us whether minor mergers can indeed drive the growth of early-type bulges.  They can also tell us whether infall-driven bulge growth is consistent with global scaling relations of bulges and discs.  ABP01 investigated the previous questions with N-body simulations of the accretion of dense, spheroidal satellites. After the accretion of the satellite, bulge-disc decompositions showed that the bulge grew and the S\\'ersic index $n$ increased proportionally to the satellite mass. However, most ABP01 satellites were dense spheroids, not representative of the majority of satellites of disc galaxies in the local Universe; those satellites reached the center of the remnant with little tidal stripping, yielding bulge growth measures that may have been unrealistic.  Accretion-driven bulge growth is sensitive to the density contrast between the primary and the merging satellite.  While high-density satellites such as those of the ABP01 study deposit most of their mass, orbital energy and angular momentum in the center of the remnant, low-density satellites disrupt during the merger, and their mass gets deposited at intermediate radii.  The fate of satellites with realistic densities cannot be ascertained a priori or with analytical means, but it can with simulations.  We have expanded on ABP01 work by running N-body simulations of the accretion of satellites that have lower densities than ABP01 experiments.  Our satellites have more complex inner structure than ABP01 spheroidal satellites: they comprise bulge, disc and halo.  Density ratios between primary and satellite derive from using the Tully-Fisher (TF) relation in the scaling of the primary galaxy and the satellites.  We use our models to study whether satellites with realistic densities drive bulge growth as compared with the ABP01 experiments, and to study what physical processes are involved in the evolution of the central component of the galaxies. We will show that the fate of the satellites depends critically on the exponent \\alphaTF\\ of the TF relation used for density scaling: varying \\alphaTF\\ over a realistically motivated range the fate varies from survival to full disruption. %with the lowest central densities and mass ratios in our experiments do not reach the galactic center of the primary galaxy, as commented later: they disrupt completely. Our models suggest that TF-scaled accretion events produce an increase of both the $B/D$ and the S\\'ersic index $n$ of the bulge, as in the high-density models of ABP01,  but the mechanisms that drive these changes are different in both cases. We find that TF-scaled satellite accretion processes alter galaxies systematicaly, giving place to remnants whose properties are scaled between them, depending on the mass ratios and the orbit of the encounter. We will show that the growth of the bulge out of disc material, a central ingredient of secular evolution models,  may be triggered externally through satellite accretion.  We will indistinctly refer to our models as \"satellite accretions\" and \"minor mergers\" \\citep{Bertschik03,Bertschik04b}.  This paper is organized as follows. Models are described in \\S\\ref{Sec:models}. We explain the bulge-disc decomposition procedure in \\S\\ref{Sec:profiles}, while results on growth of bulges are presented in \\S\\ref{Sec:growthofbulges} and \\S\\ref{Sec:whynincreases}. Population mixing is analysed in \\S\\ref{Sec:populationmixing}, while section \\S\\ref{Sec:thickening} comments the vertical structure of the final remnants. We analyse phase-space evolution and model kinematics in \\S\\ref{Sec:dynamics} and \\S\\ref{Sec:kinematics}, respectively. Scaling relations of discs and bulges in the remnants and the comparison to observed correlations are given in \\S\\ref{Sec:scaling}. Limitations of the current models are analysed in \\S\\ref{Sec:limitations}. We discuss the connections between remnants of satellite accretion and pseudobulges  in \\S\\ref{Sec:discussion}, and finally, we list the main conclusions in \\S\\ref{Sec:conclusions}. ",
        "conclusions": "\\label{Sec:conclusions}   Our simulations of TF-scaled satellite accretion experiments show that:  \\begin{enumerate}  \\item The TF-scaled satellite accretion drives pre-existing discs to a secular evolution towards earlier types (higher $B/D$ and $n$). More massive satellites produce higher changes in $n$, $B/D$ ratio, and disc thickening.  \\item  For the $B/D$=0.5 primaries modeled here, $B/D$ and $n$ grow such that,  roughly, $n$ increases by a factor of $\\sim 2.3$ when $B/D$ doubles.  Steeper growth of $n$ is expected in primary galaxies with smaller initial bulges.  \\item The evolution of the bulge component is driven foremost by the injection of disc material to the center through transitory non-axisymmetric distortions, and to a lesser extent by expansion of the pre-existing bulge and by deposition of satellite material.  Vertical heating thickens the disc.  \\item The tidal radius of the satellite depends critically on the choice of TF exponent assumed for the scaling of galaxy densities.  For the $B/D$=0.5 primaries modeled here, $\\alphaTF = 3$ leads to complete disruption and to a torus of accreted material, while the satellites reach the center for $\\alphaTF= 4$.  \\item  The satellite disc deposits its mass onto the remaining primary disc. Particles from the satellite bulge settle into a rather cold inner disc or ring, due to orbit circularization.  \\item The radial and vertical redistribution of material by heating produces a population mixture.  Part of the material initially belonging to the primary disc is interpreted as bulge in the decompositions. The photometrically identified final bulge is partially supported by rotation, as it is observed in several galaxies.   \\item Final rotation curves exhibit symmetric shapes, and a gentle outward decline typical of early-type disc galaxies. The satellite material does not dominate the central parts of the final remnants in any case, and hence, central counterrotation is not visible in the total rotation curves of retrograde mergers.  \\item Accreted TF-scaled satellites cause systematic structural transformations in the primary galaxy, which evolves towards higher $B/D$, higher $n$, higher $\\sigma _{\\mathrm{0}}$, higher $h_{\\mathrm{Disc}}$ and lower $\\mu _{\\mathrm{0,Disc}}$, all following trends with the bulge magnitude similar to the  correlations observed by BGP04b. On the other hand, some observational trends that are observed to not correlate with the bulge magnitude do it strongly in our remnants, as the ratio between the central intensities of the bulge and the disc, and the bulge effective surface density. We fail to reproduce the trends involving the bulge effective radius $r_{\\mathrm{e,Bul}}$.  \\item Secular evolution of galaxy bulges, currently being discussed as the outcome of bar instabilities in the disc due to gas cooling, can be due to satellite accretion as well, through the inward transport of disc material to the region of the bulge.  \\end{enumerate}"
    },
    "0606/astro-ph0606445_arXiv.txt": {
        "abstract": "We have mapped the \\Cone\\ emission at 492~GHz toward the supernova remnant Cas~A. We detect \\CI\\ emission from the periphery of the diffuse Photon Dominated Region (PDR) covering the disk of Cas~A, as traced by the carbon recombination lines, as well as from the denser PDRs associated with the molecular clouds towards the south-east.  \\CI\\ emission is detected from both the Perseus and Orion arm molecular clouds, with the $-47$~\\kms\\ Perseus arm feature being strong enough to be detected at all positions. We estimate the C/CO relative abundance to be 0.2 at the position of the identified CO clouds and $>1$ for most of the cloud. Here we show that the distribution of \\CI\\ emitting regions compared to the \\cplus\\ region and molecular cloud is consistent with a scenario involving PDRs. Using physical models for PDRs we constrain the physical properties of the \\CI\\ line-forming regions. We estimate  the densities of the \\CI\\ emitting regions to be between 10$^2$ and 10$^3$~\\cmcub. Based on rather high volume filling factors ($\\sim 50$\\%) we conclude that \\CI\\ emission mainly arises from diffuse neutral gas in the Perseus arm. ",
        "introduction": "Molecular clouds are embedded in the diffuse interstellar medium (ISM) of the Galaxy where most of the gas is in an atomic form. The atomic gas is mostly located in two phases that are in pressure equilibrium, the cold neutral medium (CNM) with temperatures of $T\\sim 100$~K and the warm neutral medium (WNM) with $T\\sim 10^4$~K \\citep{kulkarni1987,dickey1990,wolfire2003,mckee2004}. Some of the cold atomic gas is also associated with the denser molecular clouds, possibly forming an extended halo or envelope \\citep{andersson1993,moriarty1997}.  Far-ultraviolet (FUV) photons ($6.0~<h\\nu<~ 13.6$~eV) from OB stars produce Photon Dominated Regions (PDRs) either at the interface between the \\HII\\ region and the molecular cloud (classical high density PDRs) or in neutral components (atomic or molecular) of the diffuse interstellar medium (diffuse PDRs).  PDRs are defined as neutral regions where the chemistry and heating are regulated by the FUV photons  \\citep{hollenbach1999}. Current understanding of PDRs suggest chemical stratification in which with increasing depth from the surface of the PDR, the dominant carbon-bearing species changes from \\cplus\\ through C$^0$ to CO. Observations of \\CIIl, \\CI\\ and CO lines tracing the different layers of PDRs is thus a useful tool to constrain various conditions of the PDRs.   Ionized carbon in dense star forming regions is mainly traced using the  $^2$P$_{3/2}$--$^2$P$_{1/2}$ fine structure line at 158~\\micron\\ \\citep{howe1991,mookerjea2003}.  The radio observations of carbon recombination lines (RRLs) provide a very good alternative way to trace the spatial distribution of ionized carbon.  Classical \\CII\\ regions which form part of high-excitation PDRs, are usually identified through the observation of narrow (4 -- 10 \\kms) carbon RRLs at frequencies $>$ 1 GHz toward \\HII\\ regions \\citep{pankonin1977,wyrowski2000} and have been well studied.  The diffuse \\CII\\ regions which form part of low-excitation PDRs, are identified through  observations of carbon RRLs in absorption at frequencies below $\\sim$ 150 MHz and in emission above $\\sim$ 200 MHz \\citep{payne1989}.  Diffuse \\CII\\ regions were discovered more than two decades ago through the detection of carbon RRL in absorption at 26.3 MHz (C631$\\alpha$) in the direction of Cas A \\citep{konovalenko1980, blake1980}. Since then several RRLs spanning over 14 to 1400 MHz have been observed in this direction \\citep[][and references therein]{kantharia1998} and it remains the only diffuse \\CII\\ region which has been studied and modelled in such detail using C RRL.  A smooth transition of lines in absorption at frequencies below 115 MHz to lines in emission at frequencies above 200 MHz has been detected \\citep[][see Fig.  1]{payne1989}.  Moreover, this is the only direction in the Galaxy where high resolution ($\\sim$ 1\\arcmin) image of the diffuse \\CII\\ region in carbon RRLs near 332 MHz exist \\citep[][see Fig 2]{kantharia1998}. Modeling the line emission in this direction shows that the carbon RRLs originate in cold regions with $T_e$ = 75 K and $n_e$ = 0.02~\\cmcub\\ \\citep{kantharia1998}.   Several molecular line studies have been performed towards the direction of Cas~A. These observations do not detect much CO emission from the disk of Cas~A, but identify molecular clumps located primarily to the south east of the disk \\citep{troland1985,wilson1993,liszt1999}.  The diffuse \\CII\\ regions could  be coexistent with the cold neutral medium (CNM) that produce \\HI\\ absorption toward Cas A \\citep{payne1994} or with the molecular component of the interstellar medium (ISM) \\citep{ershov1987}.  In the direction of Cas~A, a morphological comparison of C RRL distribution with that of \\HI\\ and \\CO\\ indicate that the C RRL emission is more likely associated with \\HI\\ gas \\citep{kantharia1998}.  The upper state of the \\Cone\\ transition at 492~GHz is only 23~K above ground. For this transition, the critical density for collisions with H$_2$ molecules is only 1000~\\cmcub\\ \\citep{schroeder1991}.  This implies that \\CI\\ is easily excited and also the line is easily detectable even when emitted by moderate density interstellar gas exposed only to a radiation field equal to the mean interstellar radiation field in the solar neighborhood. Thus \\CI\\ can be used as a reliable tracer of the diffuse PDRs as well.  Here we present \\CI\\ mapping observations at 492~GHz in the direction of Cas~A in order to probe whether the \\CI\\ emission like the C RRL emission arise solely from the atomic CNM or the molecular phase also contributes to it. We have compared the \\CI\\ observations with the C270$\\alpha$, \\HI, \\twCO\\ 2--1 and \\thCO\\ 1--0 observations available in literature and have used PDR models to explain the neutral carbon and CO emission wherever the two spatially overlap. ",
        "conclusions": "We have used the \\CI\\ emission at 492~GHz to probe the PDRs seen in the direction of the supernova remnant Cas~A. We detect \\CI\\ emission at velocities around $-47$ and at $-39$~\\kms\\ from the Perseus arm clouds, and at $-1$~\\kms\\ from the local Orion arm cloud.  The \\CI\\ emission does not show any strong morphological correlation with the continuum emission from Cas~A.  Further, we do not detect any obvious broadening or strengthening of spectral lines closer to the continuum peak. Both of these conclusively rule out that the \\CI\\ emitting gas is interacting with the SNR in Cas~A.  We detect \\CI\\ emission from both the diffuse CNM (as traced by the C RRLs) and the denser molecular cloud.  Morphological comparison of the \\CI\\ emission with the carbon RRL (C270$\\alpha$), \\HI\\ optical depth and CO distribution appear to be consistent with a PDR scenario in which the dominant carbon-bearing species changes from \\cplus\\ in the diffuse region through C$^0$ to CO in the more dense molecular clumps. We note with caution that despite the morphological and kinematic evidences the present observations cannot rule out the possibility that the \\CI\\ and C~RRL emitting regions may not completely occupy the same volume of space.  We estimate the C$^0$ column density, N(C), to be between 2~10$^{16}$ and $1.3~10^{17}$~\\cmsq. At the position of the previously identified CO clouds we estimate the  C/CO abundance ratio to be around 0.2, similar to the typical value found in dense cloud cores. Most of the \\CI\\ emitting region however shows C/CO abundance ratios $>1$, typical of diffuse, translucent clouds.  We have derived hydrogen volume densities of the \\CI\\ emitting regions from the \\CI/\\twCO 2--1 intensity ratios using PDR models. We find that the local hydrogen densities vary from $\\sim 10^2$~\\cmcub\\ near the \\CI\\ peaks to $10^{3.5}$~\\cmcub\\ around CO peaks. This also shows a continuity in the density structure in the region between the the diffuse \\CII\\ region and the denser molecular clouds. We estimate the area filling factor of the \\CI\\ emitting regions to be  $\\sim 0.3$, while the volume filling factor is on an average $>30$\\%. This further suggests that most of the \\CI\\ emission stems from the diffuse PDR also traced by the carbon recombination lines."
    },
    "0606/astro-ph0606745.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} %  {The evidence of dusty disks in many Planetary Nebulae (PNs) with complex shape is growing and an equatorial disk is a key ingredient for most of the physical processes proposed for the shaping of PNs. However, the characteristics of such disks are still poorly known by lack of high spatial resolutions observations taken in the near to mid-infrared domain.} % aims heading (mandatory) {We present high spatial resolution observations of the dusty core of the Planetary Nebula with Wolf-Rayet central star CPD-56$^\\circ$8032, for which indications of a compact disk have been found by HST/SITS observations. } % methods heading (mandatory) {These observations were taken with the mid-infrared interferometer VLTI/MIDI in imaging mode providing a typical 300\\,mas resolution and in interferometric mode using UT2-UT3 47m baseline providing a typical spatial resolution of 20\\,mas. We also made use of unpublished HST/ACS images in the F435W and F606W filters.} % results heading (mandatory) {The visible HST images exhibit a complex multilobal geometry dominated by faint lobes. The farthest structures are located at 7$\\arcsec$ from the star. The mid-IR environment of CPD-56$^\\circ$8032 is dominated by a compact source, barely resolved by a single UT telescope in a 8.7$\\mu$m filter ($\\Delta \\lambda$=1.6$\\mu$m, contaminated by PAH emission). %The measured position angle (PA) of the flattened structure at 103.2$^\\circ %\\pm$2.8$^\\circ$ does not correspond to obvious structures in the HST images, though nearby bow-shaped %features %can be seen in that direction. The infrared core is almost fully resolved with the three 40-45m projected baselines ranging from -5$^\\circ$ to 51$^\\circ$ but smooth oscillating fringes at low level have been detected in spectrally dispersed visibilities. This clear signal is interpreted in terms of a ring structure which would define the bright inner rim of the equatorial disk. Geometric models allowed us to derive the main geometrical parameters of the disk. For instance, a reasonably good fit is reached with an achromatic and elliptical truncated Gaussian with a radius of 97$\\pm$11AU, an inclination of 28$\\pm$7$^\\circ$ and a PA for the major axis at 345$^\\circ \\pm$7$^\\circ$. Furthermore, we performed some radiative transfer modeling aimed at further constraining the geometry and mass content of the disk, by taking into account the MIDI dispersed visibilities, spectra, and the large aperture SED of the source. These models show that the disk is mostly optically thin in the N band and highly flared. As a consequence of the complex flux distribution, an edge-on inclination is not excluded by the data.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "CPD-56$^\\circ$8032 (He~3-1333) is the Wolf-Rayet central star (CSPN) of a young planetary nebula (PN) named PN G332.9-09.9. Wolf-Rayet CSPNe ([WR] stars), which represent about 10\\% of all PNe, are H-deficient stars (mostly [WC] type) that have dense stellar winds and share many characteristics with their higher-mass counterparts from population I (WR stars). CPD-56$^\\circ$8032, which is among the coolest CSPNe in this group, is classified as [WC10] (Crowther, De Marco, \\& Barlow 1998). There is only a handful of known [WC10] stars and it it must be noted that another member of this small group, He~2-113 (He~3-1044), exhibits striking similarities in terms of spectra, fluxes and nebula extension (i.e. age and distance, D=1.3-1.5\\,kpc). An extensive study of the CSPNe of CPD-56$^\\circ$8032 and Hen\\,2-113 has been conducted by De Marco and collaborators (De Marco, Barlow, Storey 1997 hereafter DBS97; De Marco \\& Crowther 1998, hereafter DC98, and De Marco, Storey \\& Barlow 1998, hereafter DSB98). This implies that the two stars and their nebulae have been observed generally with the same instruments involving similar spectral and spatial resolving power but the lack of spatial resolution prevented De Marco and collaborators from performing an in-depth study of the nebulae whereas their study of central stars  has been exhaustive.  The ISO spectrum of these objects is rich in emission features (Cohen et al. 1999b, 2002). The Red Rectangle, CPD-56$^\\circ$8032 and Hen\\,2-113 (Lagadec et al. 2006) are archetypal polycyclic aromatic hydrocarbon (PAH)-band sources and have amongst the highest ratios of PAH-band fluxes to total IR flux known. ISO spectroscopy also found a chemical dichotomy with evidence of crystalline silicates at longer wavelengths (Waters et al. 1998). The initial interpretation of the cold crystalline silicates was that they probably were located at large distances from the star, in order to be cool enough. However, the rate of occurrence of dual dust chemistry objects is now too high to be consistent with objects that happened to be caught during the brief transition from O-rich to C-rich chemistry. Binarity and the presence of circumstellar disks or tori in which O-rich grains from a previous evolutionary phase are increasingly believed to be responsible for the dual dust chemistry (Waters et al. 1988, Cohen et al. 1999b, 2002). The binary companion would, in this scenario, form the O-rich disk during the early AGB and play a role in the heavy mass-loss which leads to the formation of the [WC] CSPN.  De Marco \\& Soker (2002) proposed that the correlation between Wolf-Rayet characteristics and dual-dust chemistry points to a common mechanism within a binary evolution framework (see also Zijlstra et al., 1991). In their first scenario a low-mass main-sequence star, brown dwarf, or planet spirals into the asymptotic giant branch star, inducing extra mixing, a chemistry change, and terminating the asymptotic giant branch evolution. In the second scenario a close binary companion is responsible for the formation of a disk around either the binary or the companion. This long-lived disk harbors the O-rich dust. These models are mostly speculative owing to the rarity of spatial information on the dust environment at the center of these systems and the scarcity of of binary companion detection.  Because of the quasi-periodic light variations shown by CPD-56$^\\circ$8032, Cohen et al. (2002) suspected the presence of a precessing disk around it. De Marco, Barlow \\& Cohen (2002), presented the first direct evidence for an edge-on disk/torus around CPD-56$^\\circ$8032, revealed by HST/STIS spectroscopy. Their HST spectra show a spatially resolved continuum split into two bright peaks separated by 0$\\farcs$1 and interpreted to be stellar light reflected above and below an obscuring dust disk.  The discovery of this disk would probably put this object in the frame of an evolutionary sequence after objects like the archetype post-AGB Red Rectangle (Soker 1997, Waters et al. 1998, Men'shchikov et al. 2002) or IRAS 16279-4757 (Matsuura et al. 2004).  Mid-infrared interferometric observations with the MIDI/VLTI instruments are particularly well suited for studying spatially compact dusty environments. We observed Hen\\,2-113 and CPD-56$^\\circ$8032 with MIDI in order to constrain the extension of inner regions of their putative disks or rings. Whereas the environment around Hen2-113 appeared sufficiently extended to be fully resolved with the spatial resolution of a single 8m telescope (about 250mas at 10$\\mu$m) revealing a 0.5$\\arcsec$ wide ring (Lagadec et al. 2006), fringes were obtained from CPD-56$^\\circ$8032 demonstrating that the inner parts of the disk are much closer to the star. In addition, the outer environment in the 8.7$\\mu$m MIDI acquisition images also appeared well resolved by a single dish telescope, providing further constraints on the dust content of the inner nebula.  These observations, complemented by HST images are described in Sect.2. The images are described in Sect.3 and the MIDI spectra in Sect.4. Sect.5 is devoted to the extraction of the flux distribution parameters in the N band by means of simple geometrical models. Based on this study, an attempt to derive the physical parameters of the dusty environment by means of more sophisticated radiative-transfer models is presented in Section 6. Finally we discuss the constraints brought by these observations on the geometry and mass of the disk. In particular, we perform a critical comparison of the dusty environment of CPD-56$^\\circ$8032 and Hen\\,2-113. ",
        "conclusions": "We have presented new high spatial resolution observations of the dusty environment of the PN CPD-56$^\\circ$8032, from HST/ACS and MIDI/VLTI instruments. These observations confirm that a compact disk is present at the center of the PN with a position angle of its projected major axis close to the one inferred by De Marco et al. (2002). The signature of the inner rim of the disk is quite clear, evidenced by the oscillations detected in the MIDI dispersed visibilities and the radius of this rim has been determined to be of the order 75\\,mas that, at a distance of 1.35kpc translates to about 100\\,AU. The recorded MIDI continuum flux is about half of the ISO flux and we have shown that in the 8.7$\\mu$m filter, the compact core is resolved by a single dish 8.2\\,m telescope.  The parameters extracted in this paper have used a very limited sample of interferometric measurements taken in the context of testing the feasibility of observing this kind of object with MIDI, with the help of a complex HST image taken in visible light and a well-constrained large-aperture SED thanks to the ISO spectrum. In particular, the inclination of the disk is not uniquely constrained by these observations. The next natural step is, as for Hen2-113, to observe CPD-56$^\\circ$8032 in the K, L and M bands with the VLT/NACO adaptive optics instrument, which is able to provide images at the scale of 100\\,mas. The bright disk (F$_{12\\mu m}$=143\\,Jy) can also be studied in depth with MIDI thanks to the new 1.8m VLT Auxiliary Telescopes, providing a great wealth of short (8-32m) baselines although we note that the fringes from this extended object will be recorded from a much larger beam. However, the present study has shown that most of the flux originated from regions unresolved by a 8m telescope, and this problem should easily be solved by the large number of baselines expected for these observations.  %__________________________________________________________________"
    },
    "0606/astro-ph0606390_arXiv.txt": {
        "abstract": "{} {We present molecular line observations of the southwestern part of the IC 348 young cluster, and we use them together with NIR and mm continuum data to determine the distribution of dense gas, search for molecular outflows, and analyze the ongoing star formation activity in the region.} {Our molecular line data consists of C$^{18}$O(1--0) and N$_2$H$^+$(1--0) maps obtained with the FCRAO telescope at a resolution of about $50''$ and CO(2--1) data obtained with the IRAM 30m telescope at a resolution of $11''$. } {The dense gas southwest of IC 348 is concentrated in two groups of dense cores, each of them with a few solar masses of material and indications of CO depletion at high density. One of the core groups is actively forming stars, while the other seems starless. There is evidence for at least three bipolar molecular outflows in the region, two of them powered by previously identified Class 0 sources while the other is powered by a still not well characterized low-luminosity object. The ongoing star formation activity is producing a small stellar subgroup in the cluster. Using the observed core characteristics and the star formation rate in the cluster, we propose that that similar episodes of stellar birth may have produced the subclustering seen in the halo of IC 348. } {} ",
        "introduction": "IC 348 is one of the most studied young clusters. Its stellar population has been observed at different wavelengths from the IR to the X rays \\citep{lad95,luh98,her98,naj00,mue03,pre04}, and its surrounding gas, part of the Perseus molecular cloud, has been mapped with different resolutions in mm-wave lines and continuum \\citep{bac86,hat05,eno06,rid06,sun06,kir06}. The cluster lies at a distance of 320 pc (e.g., \\citealt{her98}) and consists of more than 300 stars distributed with a core-halo structure over a region 20 arcminutes in diameter \\citep{mue03}. Superposed to the smooth distribution of stars, there is a population of stellar groups (``subclusters'') each containing 5-20 stars within a radius of 0.1-0.2 pc \\citep{lad95}. Most stars in IC 348 have formed at a close-to-constant rate over the last 2-3 Myr, although some cluster members may be significantly older \\citep{her98,luh98}.  Star formation in IC 348 seems to have ceased toward the center but it continues at some level near its southern border, where the cluster meets the molecular cloud. \\citet{str74} identified an IR object to the SW of IC 348 (IC348-IRS1 hereafter), and further observations in the visible and IR by \\citet{bou95} have resolved it into a bipolar nebula likely due to an embedded star with a disk. \\citet{bac87} identified several dense cores in its vicinity, and \\citet{mcc94} found additional signposts of recent star formation in the form of shock-excited H$_2$ emission. The most prominent of these H$_2$ features is the HH211 flow, which originates from source HH211-MM and is associated with a highly collimated molecular outflow mapped in CO and SiO \\citep{gue99,cha01,pal06,hir06}. Another prominent H$_2$ feature has been recently associated with an outflow by \\citet{eis03}, who found an additional mm source that seems to power it. The NIR and optical observations of \\citet{eis03} and \\citet{wal05,wal06} reveal multiple HH objects in the region and suggest that additional young stellar objects (YSOs) lie embedded in the dense gas.  In this paper we present the results from a survey of the southwest region of IC 348 (IC348-SW hereafter) aimed to study its star formation activity and the relation between the dense gas and the molecular outflows. These data reveal the presence of several aggregates of dense cores, some of them starless, together with at least three molecular outflows (two newly mapped in CO), and help clarify the kinematics of both the outflow gas and the dense material. Combining our observations with published IR data, we show that the subclusters in the halo of IC 348 first identified by \\citet{lad95} could have originated from star formation episodes inside small core aggregates like the one currently active in IC348-SW. ",
        "conclusions": "We have mapped the molecular gas south-west of the IC 348 young cluster in C$^{18}$O(1--0), N$_2$H$^+$(1--0), and CO(2--1). Combining these data with observations of the young stars and YSOs in the NIR and mm continuum, we have reached the following conclusions.  1. The dense gas south-west of IC 348 is concentrated in two regions, referred here as IC348-SW1 and IC348-SW2, each containing several cores of a few solar masses of material. The two regions are embedded is lower density gas and they seem located in the front part of the IC 348 cluster. While IC348-SW1 is actively forming stars, IC348-SW2 seems quiescent.  2. The distribution of N$_2$H$^+$(1--0) emission is similar to the distribution of NH$_3$ and submm continuum mapped by previous authors, but differs significantly from that of C$^{18}$O(1--0). As the rare isotope is not optically thick, the disagreement suggests that the CO molecule is depleted at the core centers (likely due to freeze out onto grains) as previously seen in the more quiescent cores of Taurus.  3. The fast CO emission ($|\\Delta V| > 5$ km s$^{-1}$ with respect to the ambient cloud) is dominated by two highly collimated outflows, each associated with a bright mm/submm source. The first outflow is the molecular counterpart of HH211 and has already been mapped using interferometers by a number of authors. The second outflow, not previously mapped in CO, is associated with a previously known chain of H$_2$ emission and several HH objects, and is referred here as the IC348-SMM2 outflow. The characteristics and energetic content of the two outflows are consistent with them being excited by Class 0 sources.  4. At low velocities ($3 < |\\Delta V| < 5$ km s$^{-1}$), the CO emission presents additional features associated with the SMM2 outflow and (probably) the interaction of the cluster and the dense gas. In addition, we find a compact outflow associated with with a mm/submm object (SMM3) also detected at 24 $\\mu$m. The small size and low energy content of this new outflow suggests that it is powered by a source of rather low luminosity.  5. The ongoing star formation activity in IC348-SW1 is producing a significant stellar density enhancement, similar to others previously identified in the halo of the IC 348 cluster from NIR observations. From an estimate of the production rate and dispersal timescale of aggregates like IC348-SW1, we conclude that similar episodes of star formation could have originated the substructure seen in the IC 348 cluster."
    },
    "0606/astro-ph0606359_arXiv.txt": {
        "abstract": "In recent years, gravitational lensing has been used as a means to detect substructure in galaxy-sized halos, via anomalous flux ratios in quadruply-imaged lenses.  In addition to causing anomalous flux ratios, substructure may also perturb the positions of lensed images at observable levels.  In this paper, we numerically investigate the scale of such astrometric perturbations using realistic models of substructure distributions. Substructure distributions that project clumps near the Einstein radius of the lens result in perturbations that are the least degenerate with the best-fit smooth macromodel, with residuals at the milliarcsecond scale.  Degeneracies between the  center of the lens potential and astrometric perturbations suggest that milliarcsecond constraints on the center of  the lensing potential boost the observed astrometric perturbations by an order of magnitude compared to leaving the center of the lens as a free parameter.  In addition, we  discuss  methods of substructure detection via astrometric perturbations that  avoid full lens modeling in favor of local image observables and also discuss modeling of systems with luminous satellites to constrain the  masses of those satellites. ",
        "introduction": "The cold dark matter (CDM) paradigm predicts that the dark matter (DM) halos hosting most galaxies contain a large number of low-mass, compact subhalos within their virialized regions.  These subclumps, collectively referred to as substructure, may or may not contain luminous stellar tracers.  Since the observed number of dwarf galaxy satellites in the Local Group falls short of the expected number of subhalos by more than an order of magnitude \\citep{klypin_etal99,moore_etal99}, then either the CDM model is incorrect or dwarf galaxies are biased tracers of DM in galaxy-scale halos \\citep[e.g.,][]{spergel_steinhardt00,hannestad_scherrer00,hu_etal00,bullock_etal00,benson_etal02,somerville02,stoehr02,nagai_kravtsov05}.  If the CDM paradigm is correct, then the paucity of optical counterparts of substructures leaves us with few avenues for detecting and investigating DM substructure.  One possibility for studying dark substructure around local galaxies is through its perturbative effect on kinematically cold systems like tidal streams \\citep[e.g.,][]{mayer_etal02} or galactic disks \\citep[e.g.,][]{toth_ostriker92,benson_etal04}. These tests probe only the closest and most massive substructures, and are only applicable in very nearby galaxies. As a result, this method provides little information to address fundamental questions such as the amplitude of the matter power spectrum on subgalactic scales or the detailed properties of dark matter itself.  Another promising method of study is through the gravitational lensing effects of substructure, referred to as substructure lensing. Much of the previous work in this field has focused on the effects of substructure on image magnifications and fluxes \\citep{mao_schneider98,metcalf_madau01,dalal_kochanek02,kochanek_dalal03, rozo_etal06} or on time delays \\citep{morgan_etal06}.  However, substructure  can perturb the deflection angle, $\\alpha$, of lensed images as well.  Image positions in lensed systems may not be subject to the same foregrounds  -- such as dust absorption --  as flux anomalies. Further, since the astrometric perturbation, $\\delta\\alpha$, is a steeper function of subhalo mass than flux perturbations, it may provide a qualitatively distinct probe of substructure properties.  A study of astrometric perturbations is particularly timely given the possibility of submilliarcsecond resolution observations of strong lenses.  For example, \\citet{biggs_etal04} have presented observations of a four image jet source system, CLASS B0128+437, in which milliarcseond astrometric perturbations may have been detected. Previous studies of astrometric perturbations by substructures have come to conflicting  conclusions about their overall size and probability.    \\citet{metcalf_madau01} use lensing simulations  of random realizations of substructure in regions near images to  suggest that in order to change image positions by a few tens of milliarcseconds (mas), there must be subclumps with masses   $\\gtrsim 10^8 M_{\\sun}$  in Milky Way sized halos that are well aligned with the images they perturb.  This alignment is likely to be rare in CDM, although the probability would  increase in systems where the source is elongated in a jet  \\citep[][]{metcalf02}.  On the other hand, \\citet{chiba02} tests the size of astrometric  perturbations in B1422+231 with a model of CDM subhalos as point masses   and finds deflections of 10 to 20 mas, using subclumps with masses greater than $\\sim 2 \\times 10^{8} h^{-1} M_{\\sun}$.  Additional studies of astrometric perturbations by single perturbers have focused on the detection of substructure via distortion of a finite source \\citep{inoue_chiba03,inoue_chiba05b,inoue_chiba05}.  In addition, \\citet{pen_mao05} studied the effect of multiple lens planes on the rotation of lensed images.  In this paper, we estimate the amplitude of astrometric perturbations produced  by substructures using realistic substructure models  and test the feasibility of  observing such perturbations by comparison with the image positions given by a  best-fit lens macromodel. While \\citet{metcalf_madau01} suggest that single clumps need to be well aligned  with the images in order to produce observable perturbations, measurable effects may be possible in more general scenarios, given the collective effects of entire substructure distributions.  In addition, the kinds of perturbations in these scenarios may be different.  For example, while the scenario of a nearby subhalo that perturbs only a single image ensures that degeneracies with best-fit smooth macromodels are small, large numbers of more  distant substructure may affect multiple images and exhibit large degeneracies  with the macromodel.  In addition to using detailed lens modeling, we discuss  model-independent methods  for identifying substructures, proposing a method using systems with multiply  imaged jets.  Substructure lensing may also be used for comparisons between luminous satellites and dark subhalos.  For example,  we can address one of the open questions regarding CDM substructure:  why are some subhalos dark and some luminous?  There have been many  proposed mechanisms to explain which subhalos have stars and which  do not. One possibility is that the efficiency of star  formation diminishes with decreasing halo mass, perhaps due to  photoionization squelching \\citep{bullock_etal01}.  Another suggestion has  been that luminous galaxies form only in the highest mass halos  ($M>10^9 M_\\odot$), and over time tidal stripping reduces subhalo  masses to the low values ($M\\sim 10^8 M_\\odot$) inferred for the  smallest local dwarfs \\citep{kravtsov_etal04a}.  Yet another possibility is  that the low masses inferred for local dwarfs are in fact mistaken,  and instead are much larger, $M\\gtrsim 10^9 M_\\odot$  \\citep{stoehr_etal02,stoehr_etal03}. \\citet{kazantzidis_etal04} find that subhalos do not experience the significant mass redistribution in their centers required to embed satellites in massive subhalos, and studies stellar kinematics in the dwarfs also dispute the existence of massive subhalos \\citep[e.g.,][]{wilkinson_etal04}.   Here, we propose directly  measuring the masses associated with luminous satellites projected near lensed  images, the results of which would direcly test the \\citet{stoehr_etal02} hypothesis.  The layout of the paper is as follows.  We begin by describing the host halo and substructure models we will be employing in our analysis in \\S\\ref{sec:models}.  The statistical properties of the astrometric perturbations in our models are presented in  \\S\\ref{sec:obsres}, and their degeneracy with macro-lens model parameters is discussed in \\S\\ref{sec:modres}.  We discuss  model-independent, astrometric signatures of substructure lensing in  \\S\\ref{sec:jets}.  Finally, the effect of luminous substructures on astrometric  perturbations and lens modeling is discussed in \\S\\ref{sec:lumsats}, and we present a summary of  our results and conclusions  in section \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Astrometric perturbations from substructure lensing offers a new avenue   for probing distributions of DM substructure in galactic halos.  Previous   work in this field has mostly probed the effects of single perturbers on   image positions \\citep{metcalf_madau01,inoue_chiba03,inoue_chiba05b,inoue_chiba05}, though \\citet{chiba02} investigated the lens system B1422+231 using a simple model in which CDM substructures were modeled as point masses -- finding that typical astrometric perturbations are of order 10 to 20 mas.  Since,  observations of  the strong lens system B0128+437 with submilliarcsecond resolution have found astrometric residuals to the best fit models to be slightly smaller at roughly 5 to 10 mas. Such residuals, however, are large enough to be significant relative to measurement errors \\citep{biggs_etal04}.  In light of these finding,  the question of how large astrometric perturbations by CDM substructures  are is a very timely one.  In this paper, we use realistic models of substructure in order to create populations  of artificial lens and estimate the residuals between modeled and observed image positions.   Substructure distributions were drawn from DM simulations and semi-analytic models in  order to bracket the expected abundance of dark matter substructures in the inner regions of galaxy size halos.  Our results can be summarized as follows:  \\begin{itemize}  \\item[1.]In general, astrometric perturbations from substructure are partly degenerate  with the smooth macromodel.  Intrinsic image position perturbations from substructure are typically of order $\\gtrsim$ 10 mas, even when few to no substructures are projected near the center of the lens. Modeled residuals are significantly smaller by about an order of magnitude.  \\item[2.]Substructure distributions in which clumps project near the Einstein radius of the lens appear to be the least degenerate with the macromodel parameters and, therefore, cause the largest residuals between the observed and modeled image positions.  When the center of the lens potential is fixed, substructure models with significant amounts substructure projected near the Einstein radius have modeled residuals of order a few milliarcseconds.  \\item[3.] Substructure astrometric perturbations appear to be coherent on scales comparable to the extent of typical radio jets. This coherence significantly enhances the degree to which substructure perturbations are degenerate with a change in the center of the lensing potential.  Allowing the center of the potential to be fit as a parameter results in modeled residuals that are an order of magnitude smaller than would be the case if submilliarcsecond constraints of the center of the lensing potential were obtained.  \\item[4.]  Tests of additional substructure realizations suggest that $\\sim$10 mas may represent the largest residuals for any realistic substructure distributions where the center of the lens potential is fixed.  These substructure realizations include nearby, massive dark subhalos and more concentrated halos.  \\item[5.]  Local lens modeling of jet sources can in principle detect substructure perturbations provided there exist substructures within the Einstein radius of the lens in consideration and the jet sources are intrinsically bent.  \\item[6.] Using gravitational lensing, it is possible to place constraints on the masses of luminous substructures of observed gravitational lenses through detailed lens modeling.  These constraints may be used to test proposed mechanisms for star formation in DM subhalos.  \\end{itemize}  This work suggests that lens modeling may provide an avenue for detecting substructure through astrometric perturbations.  Substructure distributions which project subclumps near the Einstein radius of the lens halo cause the greatest perturbations;  however, substructure catalogs from simulations and semi-analytic models suggest that the scale of the perturbations will be small, on the order of a few milliarcsecond.  However, unlike anomalous flux ratios between lensed images which have been thoroughly investigated in previous works \\citep[e.g.,][]{mao_schneider98,metcalf_madau01,dalal_kochanek02,kochanek_dalal03, rozo_etal06}, a complete accounting of the effects of astrometric perturbations and  their observational implications has not yet been done;  for example, observational biases remain unexplored. The results of future studies will offer a  better understanding of submilliarcsecond resolution observations of strong lens systems."
    },
    "0606/astro-ph0606503_arXiv.txt": {
        "abstract": "We designed, built, and calibrated a new spectropolarimeter for the HiVIS spectrograph (R$\\sim$ 12000-49000) on the AEOS telescope.  We also did a polarization calibration of the telescope and instrument.  We will introduce the design and use of the spectropolarimeter as well as a new data reduction package we have developed, then discuss the polarization calibration of the spectropolarimeter and the AEOS telescope.  We used observations of unpolarized standard stars at many pointings to measure the telescope induced polarization and compare it with a Zemax model.  The telescope induces polarization of 1-6\\% with a strong variation with wavelength and pointing, consistent with the altitude and azimuth variation expected.  We then used scattered sunlight as a linearly polarized source to measure the telescopes spectropolarimetric response to linearly polarized light.  We then made an all-sky map of the telescope's polarization response to calibrate future spectropolarimetry. ",
        "introduction": "We will describe the high-resolution visible spectrograph (HiVIS) and the design of the spectropolarimeter.  We will then show some examples of new data reduction packages we have developed, and discuss the observations of the unpolarized and polarized sources.  We will combine these measurements into a polarization calibration model for the instrument and telescope and discuss the polarization calibration methods as they can be applied to future observations with the HiVIS spectropolarimeter. ",
        "conclusions": "This was the first test of the HiVIS spectropolarimeter over broad wavelengths.  We have illustrated a method for calibrating an alt-az coud\\'e  telescope with many moving mirrors for induced polarization, efficiency of polarization detection, and rotation of the plane of polarization.  Measurements of unpolarized sources and linearly polarized sources ($\\sim1000$ spectra) with known position angles at many pointings were necessary to obtain good altitude-azimuth coverage.  The AEOS telescope shows a maximum of 6\\% induced polarization at 640nm and high altitude with values falling with wavelength and altitude to near 1\\%.  We constructed an all-sky map of the telescope-induced polarization and can use it to calibrate future observations.  We hope to do more observations of unpolarized and polarized sources to improve the resolution of the calibration sky maps.           \\clearpage"
    },
    "0606/astro-ph0606029_arXiv.txt": {
        "abstract": "We discuss the effects of convective overshooting in the PMS evolution of intermediate mass stars, by analysing in detail the early evolution towards the main sequence of a 2~$M_\\odot$ stellar model. These effects can be extremely important in the end of the PMS, when the abundances in CNO elements approach the equilibrium in the centre. We provide a possible physical explanation on why a moderate amount of overshooting produces, as the star approaches the ZAMS, an extra loop in the evolutionary tracks on the HR diagram.  An interesting feature is that there is a very well defined amount of overshooting (for a given stellar mass and chemical composition) beyond which a loop is produced. For smaller amounts of overshooting such a loop does not take place and the evolutionary tracks are similar to the ones obtained by \\cite{iben65}. The amount of overshooting needed to produce the loop decreases with stellar mass.  We discuss the underlining physical reasons for the behaviour predicted by the evolution models and argue that it provides a crucial observational test for convective overshooting in the core of intermediate mass stars. ",
        "introduction": "\\label{sec:intro}  Convection and convective overshooting can play a fundamental role in the evolution of a star, by being efficient mixing processes. From the coupling of convection with nuclear reactions, significant changes on the evolution may be expected when convective overshoot is present, regardless of using different prescriptions for convection and/or overshooting. Several authors have addressed the effect of including overshoot in modelling the evolution of main and post-main sequence stars \\cite[e.g.][ and references therein]{stothers90,mowlavi94}. But the problem of how much is the amount of overshooting that has to be included is still far from solved \\cite[see][]{renzini87}. Several authors modelled core overshooting in stars \\cite[e.g.][]{roxburgh78,bressan81,langer86,xiong85,xiong86}, arriving at different results concerning the extent of the convective penetration into the radiative zone, from negligible to quite important. Another open issue is how overshoot affects the local temperature stratification \\cite[e.g.][]{canuto97}. But as far as evolution is concerned the extent of overshooting is the key aspect \\citep{deupree00,ribas00}.  There are several indications, both theoretical and observational, that support the existence of overshooting in the convective core of intermediate and high mass stars. The comparisons of mass and radii of eclipsing binaries with theoretical models suggests not only the existence of a certain amount of overshooting, but also its dependence on stellar mass \\citep{ribas00}. On the other hand theoretical predictions of the apsidal motion rate for very eccentric binary orbits can be compatible with observed values, only if overshooting is included in the models \\citep{claret93}. Moreover 2D hydrodynamic simulations performed by \\cite{deupree00} for stars with masses from 1.2 to 20~$M_\\odot$ predict convective core overshooting for all models.  The effect of overshoot on stellar evolution has already been extensively discussed in the literature, but its effect on models of pre-main sequence stars has never been discussed taking the full network of nuclear reactions into account. The overall effect to be expected is a slight delay in the evolution, redefining the age at which the star arrives at the zero age main sequence (ZAMS). As the zero point for age of young stars is in itself an open problem, such a global effect of overshoot on the PMS has been ignored.  The mixing at most of the overshoot layer at the border of a convective core in main sequence stars is expected to be efficient \\cite[e.g.][]{browning04}. This should also be the case for the end of the PMS evolution, as the time scales for mixing due to convection are still much smaller than the time scales for nuclear reactions. The uncertainties on the thermodynamic stratification within an overshoot layer are not expected to be of major relevance for determining the effects on the early evolution, but the extent of the overshoot layer and its effect on mixing the stellar material at that location are bound to be decisive.  In a previous work \\citep{marques04} we studied the effects of the time step and convective overshooting in the calibration of the young binary EK~Cep. It has been shown that the time step adopted for the PMS evolution is a key numerical aspect to adequately reproduce the expected behaviour in this rapid phase of evolution. When the time step is properly defined for calculating the evolution of intermediate mass stars with overshoot the evolutionary tracks near the end of the PMS display an extra loop. This loop seems to be present in the output of previous works \\cite[e.g.][]{siess00} although it has never been discussed in the literature, as for most cases an inadequate time step in the integration procedure would mask the loop as being a numerical artifact. Consequently its nature and physical validity has never been analysed or discussed in the literature.  Given that an adequate definition of the time step for the PMS evolution together with the most recent set of physics produces the loop near the ZAMS (when there is a significant amount of overshooting), it has lead us to perform a detailed analysis on the underlying reasons why such a loop is predicted by an evolutionary code. In particular the goal of this work is to establish if such an effect should be expected in real stars. Instead of dismissing it as a numerical artifact, as implicitly done by other authors, we have investigated if it could be consistent with the physical picture that the new state-of-the-art physics predicts (in particular the network of nuclear reactions). Here we provide such an analysis showing that there may be a physical reason for such a behaviour. In spite of the difficulties of modelling convection and convective overshoot, our analysis seems to indicate that the loop in the HR diagram at the end of the PMS evolution could happen in real stars.  We start by discussing how overshoot is introduced in a one dimensional evolutionary code and move on to show evolutionary tracks without and with overshoot for a particular case of a 2~$M_\\odot$ star. In order to illustrate the physics behind the existence of the loop at the end of the PMS we then discuss the limit case of a specific value of the extent of the overshoot and the evolution of the chemical abundances that produce the loop. The paper ends by addressing how different input physics and stellar parameters change the limit value of the overshoot producing the loop and what observational evidence could be expected in order to identify a star undergoing such a phase in its PMS evolution. ",
        "conclusions": "We have shown that a moderate amount of overshooting in evolution models of young intermediate mass stars can cause a loop in the evolutionary tracks on the HR diagram just before the ZAMS. This is a consequence of using a more precise numerical scheme for calculating the time-step together with the best up-to-date physics. The loop corresponds, for a 2~$M_{\\odot}$ star, to a drop of about 35\\% in luminosity and about 1000~K in effective temperature; the loop lasts about 1.5~Myrs. A detailed description on why the full nuclear network predicts in the models such a loop has been presented. This analysis indicates that such a behaviour may be expected if the standard formulations for core overshoot in stellar models are representative of the average role convective overshooting has in the early evolution of intermediate mass stars. The major effect of the existence of the loop is on the age of the star at the ZAMS, being higher by as much as 20\\% relative to the evolution without the loop.  The observational identification of loop-stars can provide definite tests on some of the key aspects of the physics that determine the evolution in early and main sequence intermediate mass stars. In particular the existence and extent of overshoot at the core of these stars and the reaction rates for the branches in the CNO cycle responsible for producing the loop. Such a combination of precise observational constraints make the existence of loop-stars an interesting and worth pursuing observational opportunity to test the modelling of stellar structure and evolution at these early stages. The observational challenge is the short duration of the loop; less than about 2~Myrs. The loop is even shorter for higher stellar masses, for which the amount of overshooting needed to produce a loop becomes smaller and therefore more likely. Further analysis on how overshooting and nuclear reactions are coupled in the higher mass models is required in order to identify if the loop can be expected in this regime.  With the forthcoming asteroseismic space missions \\cite[e.g.][]{baglin03} and ongoing ground based campaigns for the seismic study of stars across the HR diagram, there will be a large set of young stars whose seismic properties will be determined with sufficiently high precision to be possible to use the seismic analysis we have discussed here. There are in particular campaigns planned for the seismic observation of young clusters whose ages are estimated to be around the expected values for the existence of loop stars (from about 5 to 12~Myrs). Consequently, the possibility of confirming the existence of loop stars may be possible in the not so distant future. The implications of finding a few such stars in a young cluster will be of great relevance for the modelling of the early evolution of intermediate mass stars."
    },
    "0606/astro-ph0606721_arXiv.txt": {
        "abstract": "The optical/near-infrared (OIR) region of the spectra of low-mass X-ray binaries appears to lie at the intersection of a variety of different emission processes. In this paper we present quasi-simultaneous OIR--X-ray observations of 33 XBs in an attempt to estimate the contributions of various emission processes in these sources, as a function of X-ray state and luminosity. A global correlation is found between OIR and X-ray luminosity for low-mass black hole candidate XBs (BHXBs) in the hard X-ray state, of the form $L_{\\rm OIR}\\propto L_{\\rm X}^{0.6}$. This correlation holds over 8 orders of magnitude in $L_{\\rm X}$ and includes data from BHXBs in quiescence and at large distances (LMC and M31). A similar correlation is found in low-mass neutron star XBs (NSXBs) in the hard state. For BHXBs in the soft state, all the near-infrared (NIR) and some of the optical emission is suppressed below the correlation, a behaviour indicative of the jet switching off/on in transition to/from the soft state. We compare these relations to theoretical models of a number of emission processes. We find that X-ray reprocessing in the disc and emission from the jets both predict a slope close to 0.6 for BHXBs, and both contribute to the OIR in BHXBs in the hard state, the jets producing $\\sim90$ percent of the NIR emission at high luminosities. X-ray reprocessing dominates the OIR in NSXBs in the hard state, with possible contributions from the jets (only at high luminosity) and the viscously heated disc. We also show that the optically thick jet spectrum of BHXBs extends to near the $K$-band. OIR spectral energy distributions of 15 BHXBs help us to confirm these interpretations. We present a prediction of the $L_{\\rm OIR}$--$L_{\\rm X}$ behaviour of a BHXB outburst that enters the soft state, where the peak $L_{\\rm OIR}$ in the hard state rise is greater than in the hard state decline (the well known hysteretical behaviour). In addition, it is possible to estimate the X-ray, OIR and radio luminosity and the mass accretion rate in the hard state quasi-simultaneously, from observations of just one of these wavebands, since they are all linked through correlations. Finally, we have discovered that the nature of the compact object, the mass of the companion and the distance/reddening can be constrained by quasi-simultaneous OIR and X-ray luminosities. ",
        "introduction": "X-ray binaries are binary systems in which a compact object -- either a black hole or a neutron star -- accretes matter from a companion star. The origin of the emission from these sources is known in some wavebands and not so well established in others. Radio emission is produced by the synchrotron process in collimated outflows \\citep*[see e.g.][]{hjelet95}, whereas X-ray emission could originate e.g. directly from the hot inner accretion disc, from a Comptonising corona, from an advective flow, or from the jets \\citep*{pout98,naraet98,mccoet00,market01,brocet04}.  The optical/near-infrared (optical/NIR; OIR) is perhaps the waveband for which the dominating emission processes in XBs are least known. OIR emission has been extensively studied in outburst and quiescence (for reviews see \\citealt{vanpet95} and \\citealt{charco06}; see also \\citealt*{chenet97}). The complex variety of spectral, timing and luminosity properties of the OIR emission indicates that many processes may be contributing, each depending on a number of factors. In high mass X-ray binaries (HMXBs), the OIR light is largely dominated by the massive companion star in the system \\citep*{vandet72,trevet80} with occasional additional contributions, for example from the reprocessing of X-rays. For low-mass neutron star X-ray binaries (NSXBs), there is strong evidence for a central X-ray source illuminating a disc that reprocesses the light to OIR wavelengths \\citep*[see e.g.][]{mcclet79}.  The optical emission of a low-mass black hole candidate X-ray binary (BHXB) is generally thought to arise in the outer accretion disc as the result of X-ray reprocessing \\citep*{cunn76,vrtiet90}, much like the NSXBs \\citep*[e.g.][]{vanpet95}. Indeed, timing and spectral analysis in many cases has led to this conclusion \\citep*[e.g.][]{wagnet91,callet95,obriet02,hyneet02a,hyne05}. However, reprocessed X-rays are often misleadingly \\emph{assumed} to dominate the OIR light in BHXBs. Some observations of BHXBs point towards alternative physical processes contributing (and sometimes dominating) the OIR emission \\citep[e.g.][]{homaet05}.  Intrinsic thermal emission from the viscously heated outer accretion disc is expected to contribute significant light in the optical, through UV to X-ray wavelengths \\citep*{shaksu73,franet02}. OIR behaviour has revealed this process to play a role in some BHXBs \\citep*[e.g.][]{kuul98,soriet99,brocet01a,brocet01b,homaet05}. Thermal emission from the companion star is observed in some low-mass XBs in quiescence \\citep*[e.g.][]{oke77,bail92,oroset96,greeet01,mikoet05}. In the last decade evidence has been mounting for the flat optically thick spectrum of the jets to extend from the radio to the OIR regime (\\citealt*{hanet92,fend01,corbet02,market03,chatet03}; \\citealt{brocet04,buxtet04,homaet05}; for a review see \\citealt{fend06}). Behaviour that is not consistent with intrinsic disc or reprocessed emission has in the past been attributed to e.g. magnetic loop reconnection in the outer disc \\citep*[e.g.][]{zuriet03}, emission from a magnetically dominated compact corona \\citep*[e.g.][]{merlet00} or emission from an advective region \\citep[e.g.][]{shahet03}.  In BHXBs (both transient and persistent), properties of the emission in all wavebands are often related to changes in the X-ray spectrum. The two main X-ray spectral states are the \\emph{hard} (or \\emph{low/hard}) state, which is characterised by a hard power-law spectrum and strong variability and the \\emph{soft} (or \\emph{high/soft}; \\emph{thermal--dominant}) state, where a thermal spectrum dominates with a power-law contribution (see \\citealt*{mcclet06} for a review of X-ray states; see also \\citealt*{homaet01,fendet04,homabe05}). The low luminosity `quiescent' state is likely to be an extension to the hard state (\\citealt*{nara96,esinet97,mcclet03,fendet03}; \\citealt{fendet04,gallet06}) but currently this is not universally accepted. Here, we treat quiescence as an extension to the hard state but we also show how our results differ when the quiescent data are removed. We hereafter class `optical' and `NIR' emission as that seen in the $BVRI$ ($\\sim4400-7900\\AA$) and $JHK$ ($\\sim1.25-2.22\\mu m$) wavebands, respectively.   \\subsection{Towards a Unified Model for the OIR Behaviour in BHXBs}  Power-law correlations between OIR and X-ray luminosities are naturally expected from a number of emission processes. \\cite{vanpet94} showed that the optical luminosity of an X-ray reprocessing accretion disc varies as $L_{\\rm OPT}\\propto T^2 \\propto L_{\\rm X}^{0.5}a$, where $T$ is the temperature and $a$ is the orbital separation of the system, and that this correlation has been observed in a selection of low-mass XBs. $L_{\\rm OIR}$--$L_{\\rm X}$ correlations are also expected when the OIR originates in the viscously heated disc as both X-ray and OIR are linked through the mass accretion rate (see Section 3.2).  In addition, OIR--X-ray correlations can be predicted if the OIR emission originates in the jets. Models of steady, compact jets demonstrate that the total jet power is related to the radio luminosity as $L_{\\rm radio}\\propto L_{\\rm jet}^{1.4}$ \\citep{blanko79,falcbi96,market01}. It was shown that the jet power is linearly proportional to the mass accretion rate in NSXBs and BHXBs in the hard state \\citep*{falcbi96,miglfe06,kordet06} and the X-ray luminosity scales as $L_{\\rm X}\\propto$ \\.m and $L_{\\rm X}\\propto$ \\.m$^2$ for radiatively efficient and inefficient objects, respectively \\citep[e.g][]{shaksu73,narayi95,kordet06}. The accretion in hard state BHXBs is found to be \\emph{radiatively inefficient} (the majority of the liberated gravitational potential is carried in the flow and not radiated locally), where jet-dominated states can exist, whereas in NSXBs, the accretion is \\emph{radiatively efficient}, and jet-dominated states are unlikely to exist \\citep[see also][]{fendet03}. We therefore have:  \\smallskip BHXBs: $L_{\\rm radio}\\propto L_{\\rm jet}^{1.4}\\propto $ \\.m$^{1.4}\\propto L_{\\rm X}^{0.7}$  \\smallskip NSXBs: $L_{\\rm radio}\\propto L_{\\rm jet}^{1.4}\\propto $ \\.m$^{1.4}\\propto L_{\\rm X}^{1.4}$ \\newline\\newline The correlation for BHXBs has been observed \\citep*{corbet03,gallet03} and very recently, \\cite*{miglfe06} have applied this technique to NSXBs and found $L_{\\rm radio}\\propto L_{\\rm x}^{\\ge1.4}$; which is also consistent with the above NSXB model. If the optically thick jet spectrum is indeed flat from the radio regime to OIR, we can expect the following correlations:  \\smallskip BHXBs: $L_{\\rm OIR}\\propto L_{\\rm radio}\\propto L_{\\rm X}^{0.7}$  \\smallskip NSXBs: $L_{\\rm OIR}\\propto L_{\\rm radio}\\propto L_{\\rm X}^{1.4}$ \\newline\\newline \\cite{homaet05} discovered a correlation between the quasi-simultaneous NIR (which was shown to originate in the jets) and X-ray fluxes for GX 339--4 in the hard state, with a slope $F_{\\rm NIR}\\propto F_{\\rm X}^{0.53\\pm 0.02}$ (3--100 keV). To date, no other sources have been tested for jet OIR emission using OIR--X-ray correlations.  It is now becoming clear that this profitable but simple technique of analysing the dependence of OIR and X-ray luminosities over many orders of magnitude, may prove fruitful for the understanding of the emission mechanisms involved. The unification of jet--X-ray state activity is now underway; a steady jet exists in the hard state, which is accelerated as the X-ray spectrum softens, and is finally quenched as it passes the `jet line' into the soft state \\citep{fendet04}. A unification (if one exists) of the origins of OIR light from BHXBs in different spectral and luminosity states is desired to understand the behaviour of these systems. Furthermore, a measure of the level of OIR emission from jets may be used to constrain jet power estimates. ",
        "conclusions": "We have collected a wealth of OIR and X-ray fluxes from 33 XBs in order to identify the mechanisms responsible for the OIR emission. A strong correlation between quasi-simultaneous OIR and X-ray luminosities has been discovered for BHXBs in the hard state: $L_{\\rm OIR}\\propto L_{\\rm X}^{0.61\\pm0.02}$. This correlation holds over 8 orders of magnitude in $L_{\\rm X}$ and includes data from BHXBs in quiescence and at large distances (LMC and M31), which were \\citep[until recently; see][]{gallet06} unattainable for radio--X-ray correlations in XBs. All the NIR (and some of the optical) BHXB luminosities are suppressed in the soft state; a behaviour indicative of synchrotron emission from the jets at high luminosities in the hard state. A similar correlation is found for NSXBs in the hard state: $L_{\\rm OIR}\\propto L_{\\rm X}^{0.63\\pm0.04}$, which holds over 7 orders of magnitude in $L_{\\rm X}$. At a given X-ray luminosity, a NSXB is typically 20 times fainter in OIR than a BHXB.  Comparing the hard state OIR data to radio data of \\cite{gallet03}, we find that the radio--OIR jet spectrum in BHXBs is $\\sim$ flat ($F_{\\nu}=$ constant) at a given $L_{\\rm X}$ \\citep[see also][]{fend01}. From this and OIR SEDs of 15 BHXBs, which show an IR excess in a number of systems, we deduce that the optically thick--optically thin turnover in the jet spectrum is likely to be close to the $K$-band for BHXBs. In comparison, the turnover probably lies further into the IR for NSXBs \\citep[see][]{miglet06}.  By comparing the observed OIR--X-ray relations with those expected from models of a number of emission processes, we are able to constrain the mean OIR contributions of these processes for XBs. Table 6 summarises the results. We find from the level of soft state quenching in BHXBs that the jets are contributing $\\sim$90 percent of the NIR emission at high luminosities in the hard state. The optical BHXB data could have a jet contribution between zero and 76 percent but the optical SEDs show a thermal spectrum indicating X-ray reprocessing in the disc dominates in this regime. In BHXBs, ambiguity arises from the fact that the slope of the expected OIR--X-ray relations from the jets and X-ray reprocessing are essentially indistinguishable. Emission from the viscously heated disc may contribute at low luminosities in BHXBs, but cannot account for the observed correlations. In the NSXBs, the correlations can be explained by the X-ray reprocessing model alone (agreeing with the current thinking), but the jets may play a role at high luminosities, especially in the NIR. The exact contributions of the emission processes are likely to be sensitive to many individual parameters, such as the size of the accretion disc and the shape of the jet spectrum (e.g. the OIR luminosity of the jets is very sensitive to the slope of the optically thick jet spectrum).   \\begin{table} \\small \\caption{The OIR emission processes that can describe the empirical OIR--X-ray relations and SEDs.} \\begin{tabular}{p{1.8cm}p{0.8cm}p{0.9cm}p{0.7cm}p{0.8cm}p{0.9cm}} \\hline Sample&X-ray state&X-ray reprocessing&Jet emission&Viscous disc&Intrinsic companion\\\\ \\hline NSXBs; OPT&hard&$\\surd$&$\\times$&$\\surd$&$\\times$\\\\ NSXBs; NIR&hard&$\\surd$&$\\surd$&$\\times$&$\\times$\\\\ BHXBs; OPT&hard&$\\surd$&$\\surd$&$\\times$&$\\times$\\\\ BHXBs; NIR&hard&$\\times$&$\\surd$&$\\times$&$\\times$\\\\ BHXBs; OIR&soft&$\\surd$&$\\times$&$\\surd$&$\\times$\\\\ HMXBs; OIR&all &$\\times$&$\\times$&$\\times$&$\\surd$\\\\ \\hline \\end{tabular} \\normalsize \\end{table}  The SEDs show a non-thermal component ($\\alpha < 0$) in most sources in the hard state, and thermal emission ($\\alpha > 0$) is present, probably in all sources in both hard and soft states. The soft state could originate from a combination of X-ray reprocessing (as the OIR--X-ray relations suggest) and the viscous disc \\citep[e.g.][]{homaet05}. The SEDs of many BHXBs are redder at low luminosities, which seems to be due to both a cooler disc blackbody and a higher fractional contribution from the jets. BHXBs may be jet-dominated at low luminosities (for more on jet-dominated states, see \\citealt*{fend01,falcet04}; \\citealt{fendet04,gallet05,kordet06}). We have also made a prediction of the `$L_{\\rm OIR}$--$L_{\\rm X}$ path' of a BHXB for a typical outburst, based on a hysteresis effect whereby the hard state rise reaches a higher luminosity than the hard state decline, if the source enters the soft state. The data obtained in this paper appear to agree with the prediction, with some inevitable scatter.  Since the X-ray, OIR and radio luminosities and the mass accretion rate are all linked through correlations in the hard state, it is possible to estimate the quasi-simultaneous values of all of these parameters, given the value of just one, by e.g. daily monitoring of the X-ray or OIR fluxes (which is currently being done for some sources by the \\emph{RXTE ASM} and by ground-based telescopes). In addition, we have discovered a potentially powerful tool to complement current techniques, that can constrain the nature of the compact object, the mass of the companion and the distance/reddening towards an XB, given only the quasi-simultaneous X-ray and OIR luminosities. Data from BHXBs (in the hard and soft states), NSXBs and HMXBs lie in different areas of the $L_{\\rm X}$--$L_{\\rm OIR}$ diagram, with small areas of overlap. The tool is most useful for e.g. faint sources with poor timing analysis or at large distances, i.e. extragalactic XBs and new transients.  Further work that could constrain the emission process contributions include identifying linear polarimetry perpendicular to the jet axis (a diagnostic of optically thin synchrotron emission), emission line equivalent width analysis \\citep[in particular testing for the Baldwin effect; e.g.][]{mushfe84} and analysing $L_{\\rm OIR}$--$L_{\\rm X}$ relations in the intermediate and very high X-ray states and in ULXs and SMBHs.   \\vspace{5mm} \\emph{Acknowledgements}. We would like to thank Erik Kuulkers for providing the A0620-00 data \\citep{kuul98} and Elena Gallo for the radio and X-ray data of \\cite{gallet03}. We thank the referee for many useful comments that significantly improved this work. Results provided by the ASM/RXTE teams at MIT and at the RXTE SOF and GOF at NASA's GSFC. This paper uses data taken with UKIRT and the Liverpool Telescope, the latter of which is funded via EU, PPARC and JMU grants."
    },
    "0606/astro-ph0606517_arXiv.txt": {
        "abstract": "{ Our aim is to confirm the nature of the long period radial velocity measurements for $\\beta$ Gem first found by Hatzes \\& Cochran (1993). We present precise stellar radial velocity measurements for the K giant star $\\beta$ Gem spanning over 25 years. An examination of the Ca II K emission, spectral line shapes from high resolution data ($R$ = 210,000), and Hipparcos photometry  was also made to discern the true nature of the long period radial velocity variations. The radial velocity  data show that the long period, low amplitude radial velocity  variations found by Hatzes \\& Cochran (1993) are long-lived and coherent. Furthermore, the Ca II K emission, spectral line bisectors, and Hipparcos photometry  show no significant variations of these quantities with the radial velocity period. An orbital solution assuming a stellar mass of 1.7 $M_\\odot$ yields  a period, $P$ = 589.6 days, a minimum mass of 2.3  $M_{Jupiter}$, and a semi-major axis, $a$ = 1.6 AU. The orbit is nearly circular ($e$ = 0.02). The  data presented here confirm the planetary companion hypothesis suggested by Hatzes \\& Cochran (1993). $\\beta$ Gem is the sixth intermediate mass star shown to host a sub-stellar companion and suggests that planet-formation around stars much more massive than the sun may common. ",
        "introduction": "\\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{hatzes_fig1.eps}} \\caption{Radial velocity measurements for $\\beta$ Gem from the 6 data sets: CFHT (crosses), DAO (diamonds), McD-2.1m (circles), McD-cs21 (inverted triangles), McD-MOPS (squares), and TOPS (triangles). } \\label{orbit} \\end{figure*}   Long-period, low amplitude radial velocity (RV) variations were reported in three K giant stars, including $\\beta$ Gem ($=$ Pollux $=$ HR 2990 $=$ HD 62509) by Hatzes \\& Cochran 1993 (hereafter HC93). One hypothesis for the RV variations was the presence of planetary companions. In the case of $\\beta$ Gem an orbital solution yielded a minimum mass $m$ sin $i$ = 2.9 $M_{Jupiter}$ (assuming a stellar mass of $M$ = 2.8 $M_\\odot$), semi-major axis, $a$ = 1.9 AU and eccentricity, $e$ = 0.12. Hatzes \\& Cochran noted that ``...it would seem that planetary companions around K giants have been detected.\" However, since RV variations of comparable periods were also found for $\\alpha$ Boo and $\\alpha$ Tau, HC93 were cautious in the interpretation: ``...it seems improbable that all three would have companions with similar masses and periods unless planet formation around the progenitors to K giants was an ubiquitous phenomenon.\" Indeed, up until that time extrasolar planet RV searches were yielding no detections, even around main-sequence stars, so it seemed odd that K giant stars would produce an abundance of sub-stellar companions. This, and the fact that the expected rotation periods of K giant stars were  comparable to the detected long period RV variations made rotational modulation by stellar surface structure a viable alternative.   Subsequent work by  Larson et al. (1993b) confirmed the long period RV variations of $\\beta$ Gem  with a revised period of 585 days. That work also showed that the Ca II $\\lambda$8662 equivalent width varied at the few percent level with the same period as the RV variations. This seemed to support  rotational modulation  by surface structure as a cause of the RV variations, although the Ca II equivalent width variations were marginal. The false alarm probability for the signal was about 1\\% and Larson et al commented: ``However, because of the weakness of the signal, $K$ $=$ 0.583 $\\pm$ 0.19 m{\\AA}, this signal needs confirmation.\"  Since the discovery of RV variations in $\\beta$ Gem, planetary companions have been established in several K giant stars. The planetary companion to $\\iota$ Dra (Frink et al. 2002) was largely accepted because of the eccentric orbit, a shape in the RV curve that is difficult to produce with rotational modulation. Companions to HD 47536 (Setiawan et al. 2003), HD 11977 (Setiawan et al. 2005), and HD 13189 (Hatzes et al. 2005) were established by the absence of Ca II H \\& K emission and the lack of  spectral line bisector variations with the same period as the RV.  Over a decade since their discovery the nature of the  long period variations in $\\beta$ Gem is  still unknown. For these reasons we continued to monitor this star  with precise stellar radial velocity measurements. These new measurements along with previous ones span over 25 years in time. In Section \\S3 we show that the RV variations in this star first reported in 1993 are still present with the same period and amplitude. In Section \\S4 we make a careful analysis of ancillary data that were available to us. These are Ca II K emission, Hipparcos photometry, and spectral line bisectors measured using very high spectral resolution data. Our analysis demonstrates that there are no other forms of variability with the same period as the RV variations. This confirms the planet hypothesis for this star first reported by HC93. ",
        "conclusions": "Our analysis of the radial velocity measurements for $\\beta$~Gem   show that the long period variations found by HC93 and confirmed by Larson et al. (1993a) are long-lived and coherent. These RV variations have not changed in period, amplitude, or  phase over the past 25 years. A careful examination of the Ca II K emission, spectral line shapes, and Hipparcos photometry reveals no convincing variation with the 590-day RV period. If the RV variations were due to stellar surface structure or stellar oscillations then it is difficult to reconcile the RV variations with a lack of spectral or photometric variability. Of course, we cannot entirely exclude that an exotic form of stellar oscillations could cause the RV variations. For example, toroidal modes have all of their atmospheric motion in the horizontal direction. These would produce no photometric or Ca II emission variations. However, these modes can produce line profile variations if the star is viewed from an intermediate inclination (Osaki 1986). Our bisector measurements exclude this possibility. The most likely and logical explanation for the RV variations is that they are indeed due to a planetary companion with minimum mass of 2.3 $M_\\odot$ at an orbital distance of 1.6 AU. These data confirm the planet hypothesis for the long period RV variations first proposed by HC93.  If the 135-d period found in the Hipparcos photometry indeed represents the rotational period, then this can be used to estimate the stellar inclination. HC93 measured a projected rotational velocity for $\\beta$ Gem of 1.6 $\\pm$ 0.9 km\\,s$^{-1}$. A radius of 8.8 $R_\\odot$ yields an equatorial rotational velocity of 3.3  km\\,s$^{-1}$. This yields sin $i$ $=$ 0.48 $\\pm$ 0.3. Assuming an alignment of rotational and orbital axes results in a true companion mass of 2.9 -- 12.8 $M_{Jupiter}$.  We can check if the possible photometric variations that are detected could  result from rotation in spite of the low activity level of this star by comparing it with models.  Given the period and velocity amplitude of the RV variations radial pulsations would produce a change in stellar radius of about 10\\%. This can be excluded by the lack of large photometric variations. Non-radial pulsations are ruled out by lack of bisector variations, and significant numbers of starspots would be surprising on such an inactive star (e.g., $F_X/F_X(\\odot) \\approx 0.02$; Rutten et al 1991).  The average solar spot coverage is $f_S \\sim$ 0.001 ($\\langle R_M \\rangle = 50, f_S \\approx 2.2\\times 10^{-5} R_M$; Cox 2000).  If $f_S$ scales with $F_X$, then for $\\beta$ Gem $f_S \\sim 2 \\times 10^{-5}$, far too small to yield the photometric variability seen by Hipparcos ($\\sigma_V \\approx 0.0033$).  A more likely possibility is microvariability due to the stellar granulation -- specifically, due to the finite number of (variable) convective granules on the stellar surface.  Ludwig (2006) estimates that the fractional flux RMS $\\sigma_F/F$ due to granulation is given by (combining their equations 56 and 59):  $\\sigma_F/F = 0.4 N^{-0.5} (\\delta I_{\\rm RMS}/I) f(a)$ \\\\ where $N$ is the number of granules on the stellar surface, $\\delta I_{\\rm RMS}/I$ is the fractional RMS intensity variation due to a granule at $\\mu=1$, and $f(a)$ is a slowly varying function of the linear limb-darkening coefficient $a$.  Freytag et al. (2002) use 3-D hydrodynamic  models to show that the size of a typical granule scales with the pressure scale height as  $x_{\\rm gran}/R_\\ast \\approx 0.0025 (T_\\ast/T\\odot)(R_\\ast/R_\\odot) (M_\\ast/M_\\odot)^{-1}$.  Thus, for $\\beta$ Gem $x_{\\rm gran}/R_\\ast \\approx 0.015$, implying $N \\sim 8900$ on the visible stellar surface.  Then, if we adopt the solar value of  $\\delta I_{\\rm RMS}/I$ = 0.18, and take $\\epsilon$ = 0.8, $f(a) \\approx 1.3$, and $\\sigma_F/F \\sim 0.001$.  While this a factor of $\\approx$3 smaller than $\\sigma_V$, we feel the two are sufficiently close (given the many approximations involved) to suggest it is quite probable that the  photometric variation is due to granulation-induced inhomogeneities combined with rotation.   $\\beta$ Gem with an estimated mass of 1.7 $M_\\odot$ is the sixth star of intermediate mass known to host an extrasolar planet. Table 6 lists those stars in the mass range 1.8 -- 5  $M_\\odot$ known to host giant planets ($\\iota$ Dra has an estimated mass of 1.05 $M_\\odot$).  Although the number of intermediate mass stars hosting giant planets is small, these already show some interesting characteristics. First, all would qualify as ``super planets\" having   masses much greater than a few  Jupiter masses.  Possibly, more massive stars have more massive protoplanetary disks   which could result in more massive planets. Second, the semi-major axes for most are around 2 AU. This might raise concerns that we are seeing some other phenomena such as rotation and not evidence for planetary companions; however, we believe this is not the case for two reasons: 1) For these stars the RV variations were not accompanied by other forms of variations which excluded rotational modulation or pulsations as a cause. 2) The derived orbital eccentricities span a wide range ($e$ = 0.01 -- 0.40). If the RV variations were due to stellar rotation or pulsations, then we would not expect similar shapes in the RV curves and not the wide variety that is provided by Keplerian motion.   Furthermore, not all K giant stars show long period RV variations. In a sample of 62 K giants D\\\"ollinger et al. (2006) found evidence for long period RV variations in at most 15\\% of  stars.   Most of the stars in Table 6 have low metallicities. Rice \\& Armitage (2005) have argued that stars with lower metallicities take longer to form planets. These would have had little time to migrate before the disks were dispersed leaving the giant planets near the snow-line of $\\approx$ 2--4 AU. However, $\\beta$ Gem and $\\gamma$ Cep have metallicities considerably higher than the solar value. Alternatively, radiation pressure from the more massive and thus hotter star may have dispersed the disks before the giant planet had time to migrate.  However, it is dangerous to draw conclusions based on such a small sample. The sample of exoplanets around intermediate stars  must be increased by at least an order of magnitude before we can discern the true distribution of semi-major axes and companion masses. Further discoveries of giant planets around intermediate stars may hold important clues for planet formation.  \\begin{table*} \\begin{center} \\begin{tabular}{cccccccc} Star & $M_{star}$  &  $R_{star}$ & M$_{planet}$$\\times$ sin $i$ & $a$ & $P$ &  $e$ & [Fe/H]   \\\\ \\hline HD 11977$^1$       & 1.9   & 10.2 & 6.5  & 1.9  & 1420  & 0.40 & $-$0.14  \\\\ HD 47536$^2$       & 2.0   & 21.3 & 5.0  & 2.0  & 712   & 0.20 & $-$0.61    \\\\ HD 13189$^3$       & 3.5   & -    & 14   & 1.8  & 471   & 0.27 & $-$0.59 \\\\ HD 104985$^4$      & 1.6   & 11   & 6.3  & 0.78 & 198   & 0.03 &  $-$0.35 \\\\ $\\gamma$ Cep$^5$   & 1.6   & 4.7  & 1.7  & 2.3  & 920   & 0.12 & $+$0.18 \\\\ $\\beta$ Gem    & 1.7   & 8.4  & 2.3  & 2.4  & 590   & 0.02 &  $+$0.19 \\\\ \\hline \\end{tabular} \\caption{Sub-stellar Companions to Intermediate Mass Stars ($^1$ Setiawan et al. 2005, $^2$ Setiawan et al. 2003, $^3$ Hatzes et al. 2005, $^4$Sato et al. 2003, $^5$Hatzes et al. 2003)} \\end{center} \\end{table*} -----------------------------------------------------------------------"
    },
    "0606/astro-ph0606667_arXiv.txt": {
        "abstract": "We present two-dimensional (2D) mapping of the stellar velocity field within the inner 5\\arcsec\\ of six nearby active galaxies, using spectra obtained with the Integral Field Unit of the GMOS instrument at the Gemini North telescope. The sampling of the observations is 0\\farcs 2, corresponding at the galaxies to spatial extents ranging from 10 to 30\\,pc. The spatial resolution range from 20 to about 180\\,pc, and the observed field of view covers a few hundred parsecs around the nuclei. The Calcium\\,II triplet absorption features at $\\approx$ 8500\\AA\\ were used to measure the stellar radial velocities and velocity dispersions. The radial velocity fields are dominated by rotation in all galaxies. A simple kinematical model assuming a purely rotating system with circular orbits in a plane was fitted to the radial velocity data. The turnover of the rotation curve is at only $\\approx 50$\\,pc for NGC\\,4051 and between 200 and 700\\,pc for the other 5 galaxies. The velocity dispersion ($\\sigma$) maps show the largest values ($100\\ge \\sigma \\ge 150 $ km\\,s$^{-1}$) at the centre. In the cases of NGC\\,2273 and NGC\\,3227, there is a decrease to $\\sigma \\approx 70-80$ km\\,s$^{-1}$ at $\\approx 200-300$\\,pc from the nucleus, delineating partial rings of low $\\sigma$ values. A similar broken ring seems to be present at $\\approx 400$\\,pc from the nucleus also in NGC\\,4593. We interpret these low $\\sigma$ rings as traces of recently formed stars that partially keep the cold kinematics of the original gas from which they have formed. In NGC\\,3516 there is a decrease of $\\sigma$ outwards with the steepest gradient observed along the direction of the galaxy major axis, where $\\sigma$ reaches $\\approx 80-90$ km\\,s$^{-1}$ at $\\approx 400$\\,pc from the nucleus.  The main novelty of the present work is the unprecedented spatial resolution reached by a 2D study of stellar kinematics of Seyfert galaxies using an IFU. The few similar IFU studies available in the literature for Seyfert galaxies have a much poorer spatial resolution and/or are restricted to the study of emission line kinematics. ",
        "introduction": "\\mbox{}  It is now widely accepted that active galactic nuclei (AGN) are powered by the accretion of material onto a central supermassive black-hole (SMBH). The present paradigm for the evolution of galaxies is that all galaxies which form bulges also form a SMBH at their centres (\\citealt{FM00}; \\citealt{geb00}). In this scenario, the active galaxies are those in which the SMBH is presently accreting material from its surroundings.  A problem still under investigation is the mechanisms by which the material is dragged to the nuclear regions to feed the AGN, as the galactic gas must loose almost all of its angular momentum in order to reach the central few parsecs. Signatures that such feeding is occurring include the frequent occurrence of recent star formation around AGN, which implies the existence of a gas reservoir close to the AGN \\citep{sch99, cid01, sto00, boi00, cid05, sto05}. Besides the signatures of young stars observed in spectra of AGN, kiloparsec scale kinematic studies \\citep{nw95,nw96} have suggested that Seyfert galaxies have lower mean mass-to-light ratios than normal galaxies, which also could be due to a younger near-nuclear stellar population \\citep{oli95,oli99}.  The presence of young stars near the nucleus seems also to be the explanation for the results of recent studies investigating the stellar kinematics on scales of hundred of parsecs \\citep{ems01,gar99,mar03}. These studies have found a central drop in the stellar velocity dispersion (hereafter $\\sigma$-drop) in a few galaxies with Seyfert nuclei. \\citet{woz03} presented simulations showing that a new generation of stars formed at the centre of the galaxy from cold material would create such a  drop which would remain visible for hundreds of Myrs. Nevertheless, a decrease in $\\sigma$ towards the nucleus has been recently observed also in late-type non-active galaxies \\citet{gan06}, supporting its link to recent star-formation, but not necessarily to nuclear activity in galaxies.  Most studies available in the literature on kinematics of Seyfert galaxies are based on long-slit observations, which are restricted to only one axis of the host galaxies. In order to properly probe the galactic gravitational potential, as well as to investigate the nature and extent of the $\\sigma$-drops, it is necessary to cover a 2D region, what we have done in the present work using the Integral Field Unit (IFU) of the Gemini Multi-object Spectrograph (GMOS).  The power of IFU observations to map the large scale kinematics of galaxies has been evidenced in recent studies (e.g. \\citealt{gan06}). For AGN, only few such studies are available (e.g. \\citealt{ems06} for NGC\\,1068). Most AGN kinematic studies have focused on gas emission-lines, as it is much easier to obtain kinematical data from emission lines than from stellar absorption features. Nevertheless, the gas in the nuclear region of AGNs is subject to non-gravitational effects, such as winds and jets, and in order to probe the gravitational potential, it is necessary to measure the stellar kinematics.   In this work, we use the IFU-GMOS to map the stellar kinematics in the inner few hundred parsecs of 6 of the closest Seyfert galaxies. The combination of the spatial resolution at the galaxies reached by the observations (down to a few tens of parsecs), 2D coverage, stellar kinematics, and a Seyfert sample is unique in the literature.  This paper is organized as follows. In Section 2 we present the criteria used to select the sample and discuss the relevant information for each galaxy. In Section 3 we describe the observations and reductions, in section 4 we present the methods used to analyze the data and in Section 5 we report and discuss our results. In section 6 we present a summary of the results and our conclusions. ",
        "conclusions": "\\mbox{}  In this work we have obtained 2D maps of the stellar kinematics of the inner few hundred parsecs of 6 nearby Seyfert galaxies at sub-arcsecond angular resolution, corresponding to spatial resolutions ranging from 30 to 180 parsecs at the galaxies.  The stellar velocity field is dominated by rotation, well represented by a simple model where the stars follow plane circular orbits under a Plummer potential. The residuals between measured and modeled radial velocities are always smaller than $\\pm\\,25$\\,km\\,s$^{-1}$.  In most cases, the turnover of the rotation curve seems to occur within or close to the edges of the observed field, at radial distances ranging from $\\sim 50$\\,pc for NGC\\,4051 to $\\sim 500-700$\\,pc for NGC\\,3516 and NGC\\,4593. Only for NGC\\,4941, for which our observations reach only as far as $\\approx 200$\\,pc from the nucleus, the turnover was not observed within the IFU FOV.  The case of NGC\\,4051 is particularly interesting because the turnover is at only $\\sim 50$\\,pc from the nucleus, suggesting that the stellar motions are dominated by a highly concentrated gravitational potential. Indeed, the scale length of the Plummer potential obtained for this galaxy is the smallest of the the sample, $A = 31$ pc. Adopting $\\sigma = 85$\\,km\\,s$^{-1}$ for the nuclear velocity dispersion (the mean value of innermost usable pixels), and the \\citet{tre02} relation, we infer a black hole mass of $4.5 \\times 10^6$ M$_\\odot$, which implies a sphere of influence of only $2.5$\\,pc radius. These values are consistent with those by \\citet{pet04}, who obtained a black hole mass of $1.9\\times 10^6$ M$_\\odot$ corresponding to a sphere of influence of $1.2$\\,pc radius. Such concentrated potential is unlikely to be responsible for the small potential-scale size we have found. A more likely interpretation is that the bulge is itself compact and concentrated. Indeed, unlike most Seyferts, this galaxy is late type (Scd), with a small bulge effective radius.  The velocity dispersion maps are featureless for NGC\\,4051 and NGC\\,4941. For the other galaxies the velocity dispersion maps show higher values at the centre and smaller values at $200-400$\\,pc from the nucleus. In the case of NGC\\,3516 the lowest $\\sigma$ values are observed towards the border of the IFU field, at location close to the kinematic major axis. For NGC\\,2273, the spatial distribution of lower $\\sigma$ values defines a ring-like morphology which is co-spatial with partial nuclear ring structures seen in line emission images and color maps of previous works. For NGC\\,3227, the velocity dispersion map shows a very good correspondence between the region of low $\\sigma$ values -- which seems to delineates a partial ring structure -- and regions of high $^{12}$CO (2-1) emission and blue $J-K$ colours observed in previous works. A partial ring of low $\\sigma$ values is also hinted in NGC\\,4593. As for the nuclear $\\sigma$-drops found by previous authors (e.g. \\citealt{ems01,gar99,mar03}) the circumnuclear $\\sigma$-drops we found in the present study can be interpreted as regions of higher (past or present) gas concentration, harboring younger stars which still preserve the lower velocity dispersion of the original gas from which they have formed."
    },
    "0606/astro-ph0606384_arXiv.txt": {
        "abstract": "{ We report the discovery of \\cog, a rich cluster of galaxies in the Zone of Avoidance, identified in public images from the XMM-\\textit{Newton} archive. We present the main properties of this cluster using the XMM-\\textit{Newton} X-ray data, along with new optical spectroscopic and photometric observations. \\cog\\ is located at $z = 0.271\\pm 0.001$, as derived from the optical spectrum of the brightest member galaxy. Such redshift agrees with a determination from the X-ray spectrum ($z$ = 0.263$^{+0.012}_{-0.010}$), in which an intense emission line is matched to the rest wavelength of the Fe K$\\alpha$ complex. Its intracluster medium has a plasma temperature of $4.92^{+0.50}_{-0.48}$\\,keV, sub-solar abundance of $0.38\\pm 0.12\\,Z_{\\odot}$, and a bolometric luminosity of $3.2\\times 10^{44}$ erg\\,s$^{-1}$. A density contrast $\\delta = 2500$ is reached in a radius of $0.5\\,h_{70}^{-1}\\,$Mpc, and the corresponding enclosed mass is $1.5 \\times 10^{14} M_{\\odot}$. Optical images show an enhancement of $g^\\prime - i^\\prime > 2.5$ galaxies around the central galaxy, as expected if these were cluster members. The central object is a luminous E-type galaxy, which is displaced $\\sim 40\\,h_{70}^{-1}\\,$kpc from the cluster X-ray centre. In addition, it has a neighbouring arc-like feature ($\\sim 22\\arcsec$ or $90\\,h_{70}^{-1}\\,$kpc from it), probably due to strong gravitational lensing. The discovery of Cl~2334+48 emphasises the remarkable capability of the XMM-\\textit{Newton} to reveal new clusters of galaxies in the Zone of Avoidance. ",
        "introduction": "\\label{sec:intro}  Rich clusters of galaxies have been successfully used as tracers of large-scale structure formation and evolution, which has allowed setting constraints on various cosmological parameters \\citep[see, e.g.,][]{Bardeen86,Henry00,Rosati02}. In the past, the majority of rich clusters were first identified in the optical and later observed in X-rays. However, with the advent of large and deep X-ray surveys, X-ray observations have become one of the most useful techniques for discovering clusters of galaxies, especially for intermediate and high-redshift systems \\citep[e.g.][]{Gioia94,Rosati98,Adami00}. On the other hand, X-rays may also be useful for detecting clusters near the Galactic plane, where the increasing number of stars and extinction makes the optical identification of background galaxies difficult \\citep[see, e.g.,][and references therein]{Kraan00}.  \\begin{figure*} \\centering \\includegraphics[width=17.5cm]{5681fig1.eps} \\caption[]{Dust map [$E(B-V)$ map from \\citet{Schlegel98} in logarithmic grey scale] superposed with all known clusters and groups found in NED with $|b| < 20^{\\circ}$ and $20^{\\circ} < l < 200^{\\circ}$. \\cog\\ is emphasised with a larger symbol near $l = 110^{\\circ}$,~$b = -12^{\\circ}$. \\label{fig:ZoA}} \\end{figure*}  The mapping of the large-scale structures hidden by the Galaxy has profound cosmological implications, as can be exemplified by the discovery of the Great Attractor at $(l, b) \\approx (320^{\\circ}, 0^{\\circ})$ \\citep{Lynden88, Kolatt95}. Indeed, the mass distribution derived from the large-scale distribution of galaxies and from the peculiar velocity field of nearby galaxies and/or clusters is still a matter of debate \\citep[see, e.g.,][]{Tonry00,Hudson04,Mieske05}. It is still not clear whether or not the bulk flow of galaxies is due mainly to the Great Attractor, or if there is a large contribution from the Shapley supercluster or other more distant large-scale structures \\citep{Nagayama06,Proust06,Radburn06}.  A systematic search for clusters of galaxies in the Zone of Avoidance, $|b| < 20^{\\circ}$, was carried out by \\citet{Ebeling02}, where they made use of the ROSAT All Sky Survey Bright Source Catalog \\citep[RASS-BSC][]{Ebeling96,Voges99} coupled with optical and near infrared follow-ups. A major difficulty with this approach is the low-energy band of RASS (0.1--2.4\\,keV) in which X-ray is strongly affected by the Galactic absorption, notably when $N_{\\rm H} \\ga 10^{21}\\,$cm$^{-2}$.  In this work we report the discovery of \\cog, a rich cluster of galaxies with $z$ = 0.271, projected onto the Galactic plane and identified in the XMM-\\textit{Newton} data archive. The source is described in Sect. \\ref{sec:source}. The main properties of this object were determined through the analysis of XMM-\\textit{Newton} X-ray data (described in Sect.~\\ref{sec:obs}), complemented by optical data (Sect.~\\ref{sec:optical}). The X-ray analysis is shown in Sect.~\\ref{sec:x-ray_analise}, followed by a mass determination in Sect.~\\ref{sec:massdet}. The optical data analysis is described in Sect.~\\ref{sec:OptObs}. For distances and luminosities, we use a $\\Lambda$CDM cosmology with $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_{0} = 70\\,h_{70}$km\\,s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{discussion}  In this paper we report on the discovery of Cl~2334+48, a cluster of galaxies in the Zone of Avoidance, from XMM-\\textit{Newton} X-ray images. New photometric and spectroscopic observations of the system are presented.  The main results of our analysis of the X-ray and optical data can be summarised as follows.  \\begin{itemize} \\item \\cog\\ is at a redshift of $z = 0.271 \\pm 0.001$, based on the optical spectrum of the brightest cluster member. The redshift derived from the X-ray spectral fitting is $z = 0.263^{+0.012}_{-0.010}$.  \\item The X-ray emission is consistent with an isothermal ICM with $kT = 4.92^{+0.50}_{-0.48}$\\,keV, gas metallicity of $0.38\\pm0.12$\\,$Z_{\\odot}$, and bolometric luminosity of $3.2\\times 10^{44}$ erg\\,cm$^{-2}$\\,s$^{-1}$.  \\item There is no evidence of a strong local $N_{\\rm H}$ absorption in the ICM.  \\item It has, within $800\\,h_{70}^{-1}$kpc, a gas mass of $3.2 \\times 10^{13}M_{\\odot}$ or $3.5 \\times 10^{13}M_{\\odot}$, a gas fraction of 0.09 or 0.16, and a dynamical mass of $3.8 \\times 10^{14}M_{\\odot}$ or $2.3 \\times 10^{14}M_{\\odot}$ if either a S\\'ersic or a $\\beta$-model, respectively, is considered for the gas and total mass cumulative profiles. The total virial mass of the cluster within the virial radius of $1.7\\,h_{70}^{-1}\\,$Mpc was determined to be $M_{200} = 5\\times 10^{14}M_{\\odot}$.  \\item There is a significant overdensity of galaxies around the brightest cluster object, within the inner $65\\arcsec$ ($\\sim 270\\,h_{70}^{-1}\\,$kpc). Most objects are red, with $g^\\prime - i^\\prime > 2.5$, as expected if these are members of the cluster.  \\item The surface brightness of the central galaxy corresponds to a $r^{1/4}$ law, indicating that the object is a normal elliptical and not a cD galaxy.  \\item An arc-like structure, probably due to strong lensing, is seen in both the $g^\\prime$ and $i^\\prime$ optical images. This feature is located at $\\sim 22\\arcsec$ ($90\\,h_{70}^{-1}\\,$kpc) from the central galaxy of the cluster. Better images are needed to obtain a secure mass determination of the cluster from the gravitational lensing map. \\end{itemize}  Although Cl 2334+48 was detected in X-ray in the past (see Sect.~\\ref{sec:source}), its nature was unknown. At a Galactic latitude of $-12^\\circ$, an optical identification of the cluster, alone, would be difficult. Due to the high sensitivity and good spatial resolution of the XMM-\\textit{Newton} EPIC cameras, such a cluster was easily detected, and we were able to derive its main X-ray properties, as well as its redshift.  The X-ray emission is almost spherically symmetric and the temperature profile has the usual shape (although the error bars are large) observed in relaxed clusters (Figures~\\ref{fig:VibStab} and \\ref{fig:tempprofile}). The X-ray derived properties based on hydrostatic equilibrium follow the scaling relations of nearby regular clusters. However, the X-ray peak emission does not coincide with the brightest member galaxy by an offset of $\\sim 10\\arcsec$ ($40\\,h_{70}^{-1}\\,$kpc; see Fig.~\\ref{fig:VibStab}). If this displacement is real, it indicates that this cluster is not quite relaxed and some dynamical event (possibly a sub-cluster merging we may surmise) took place a few Gyrs ago. The central galaxy is not a cD, owing to the lack of an extended stellar envelope, which may be interpreted as a sign of the relative youth of this cluster."
    },
    "0606/astro-ph0606451_arXiv.txt": {
        "abstract": "We investigate encounters between giant molecular clouds (GMCs) and star clusters. We propose a single expression for the energy gain of a cluster due to an encounter with a GMC, valid for all encounter distances and GMC properties. This relation is verified with $N$-body simulations of cluster-GMC encounters, where the GMC is represented by a moving analytical potential. Excellent agreement is found between the simulations and the analytical work for fractional energy gains of $\\Delta E/|E_0|<10$, where $|E_0|$ is the initial total cluster energy. The fractional mass loss from the cluster scales with the fractional energy gain as $(\\Delta M/M_0)=f(\\Delta E/|E_0|)$, where $f\\simeq0.25$.  This is because a fraction $1-f$ of the injected energy goes to the velocities of escaping stars, that are higher than the escape velocity. We therefore suggest that the disruption time of clusters, $\\tdis$, is best defined as the time needed to bring the cluster mass to zero, instead of the time needed to inject the initial cluster energy. We derive an expression for $\\tdis$ based on the mass loss from the simulations, taking into account the effect of gravitational focusing by the GMC. Assuming spatially homogeneous distributions of clusters and GMCs with a relative velocity dispersion of \\vcn, we find that clusters loose most of their mass in relatively close encounters with high relative velocities ($\\sim2\\,\\vcn$). The disruption time depends on the cluster mass ($\\Mc$) and  half-mass radius ($\\rh$) as $\\tdis~=~2.0\\,S\\,\\left(\\Mc/10^4\\,\\msun\\right)\\,\\left(3.75\\,\\mbox{pc}/\\rh\\right)^3{\\mbox{Gyr}}$, with $S\\equiv1$ for the solar neighbourhood and $S$ scales with the surface density of individual GMCs ($\\sigman$) and the global GMC density ($\\rhon$) as $S\\propto(\\sigman\\rhon)^{-1}$.  Combined with the observed relation between \\rh\\ and \\Mc, i.e. $\\rh \\propto \\Mcl^{\\lambda}$, \\tdis\\ depends on \\Mc\\ as $\\tdis\\propto\\Mc^\\gamma$. The index $\\gamma$ is then defined as $\\gamma=1-3\\lambda$.  The observed shallow relation between cluster radius and mass (e.g. $\\lambda\\simeq0.1$), makes the value of the index $\\gamma=0.7$ similar to that found from observations and from simulations of clusters dissolving in tidal fields ($\\gamma\\simeq0.62$).  The constant of 2.0 Gyr, which is the disruption time of a $10^4\\,\\msun$ cluster in the solar neighbourhood, is about a factor of 3.5 shorter than found from earlier simulations of clusters dissolving under the combined effect of galactic tidal field and stellar evolution. It is somewhat higher than the observationally determined value of 1.3 Gyr. It suggests, however, that the combined effect of tidal field and encounters with GMCs can explain the lack of old open clusters in the solar neighbourhood. GMC encounters can also explain the (very) short disruption time that was observed for star clusters in the central region of M51, since there $\\rhon$ is an order of magnitude higher than in the solar neighbourhood. ",
        "introduction": "\\label{sec:introduction}  Star clusters are subjected to various disruptive effects, which prevent low mass star clusters ($\\Mcl<\\,10^4\\,\\msun$) from surviving for a Hubble time. From observations this was first noted by \\citet{1958RA......5..507O} and later by \\citet{1971A&A....13..309W}, who showed that there is a lack of open star clusters older than a few Gyr in the solar neighbourhood. Later, more quantitative results were obtained for the life time of star clusters (e.g. \\citealt{1988PASP..100..576H} and \\citealt{2006MNRAS.366..295D} for clusters in the LMC; \\citealt{2003MNRAS.338..717B} for the SMC and M33; \\citealt{2005A&A...441..117L} and \\citealt{2006A&A...445..545P} for the solar neighbourhood and \\citealt{2005A&A...441..949G} for the central region of M51). From the theoretical side, the evolution and disruption of star clusters has been the subject of many studies using a variety of techniques to solve the gravitational $N$-body problem (e.g. \\citealt*{1972ApJ...176L..51O, 1990ApJ...351..121C, 1997ApJ...474..223G, 2000ApJ...535..759T, 2003MNRAS.340..227B}).  \\citet*{2005A&A...429..173L} compared observationally determined disruption times in four galaxies to the corresponding times following from the $N$-body simulations of clusters in tidal fields of \\citet{2003MNRAS.340..227B} and \\citet{1998A&A...337..363P,2002ApJ...565..265P}. They found that, based on the results of the simulations, the disruption time ($\\tdis$) should depend on the cluster mass $\\Mc$ as $\\tdis \\propto [\\Mcl/\\ln(\\Mcl)]^{0.75} \\propto \\Mcl^{0.62}$. The index (0.62), which is referred to as $\\gamma$, was also determined from the observed age and mass distributions of different cluster populations \\citep{2003MNRAS.338..717B} and the mean value from the observations is $\\bar{\\gamma}=0.62\\pm0.06$. The value of $\\tdis$ scales with the disruption time of a $10^4\\,\\msun$ cluster as $\\tdis=t_4\\,(\\Mcl/10^4\\,\\msun)^\\gamma$. Good agreement for $t_4$ between simulations and observations was found only for the star clusters in the SMC. In the solar neighbourhood, the observed $t_4$ for open clusters ($t_4=1.3\\,$Gyr, see \\citealt{2005A&A...441..117L}) is a factor of 5 lower than expected from the disruption time due to the galactic tidal field, combined with a realistic stellar mass function and with stellar evolution ($t_4 = 6.9\\,$Gyr, see \\citealt{2003MNRAS.340..227B}). An even larger disagreement of almost an order of magnitude was found for $\\tdis$ of clusters in the central region of M51 \\citep{2005A&A...441..949G}.  The fact that the observed disruption times are shorter could be a result of time-dependent external perturbations, such as spiral arm passages and encounters with giant molecular clouds (GMCs), which were ignored in the simulations of \\citet{2003MNRAS.340..227B}. \\citet{1987MNRAS.224..193T} and \\citet{1991MmSAI..62..909T} showed that a single encounter with a massive GMC could disrupt a star cluster of a few hundred $\\Msun$ and \\citet*{gieles06a} showed that in the solar neighbourhood passages of spiral arms contribute significantly to the disruption of open clusters.  In this study we investigate the disruption of clusters by encounters with GMCs. We want to quantify whether GMCs can explain the difference between the observed values of $t_4$ and the ones following from the simulations without the time-dependent external effects. Moreover, would simulations including encounters with GMCs preserve the numerical value of the index $\\gamma\\simeq0.62$, where the observations and previous simulations already agree?   The paper is organised as follows: In \\S~\\ref{sec:initial} we introduce the initial conditions for the clusters and GMCs and define the parameters involved in a cluster-GMC encounter. In \\S~\\ref{sec:analytical} we give analytical calculations of the energy gain for a cluster due to various types of encounters. These analytical formulae for the energy gain are verified with $N$-body simulations and compared to the mass loss of the cluster in \\S~\\ref{sec:nbody}. The results of the $N$-body simulations are used in \\S~\\ref{sec:disruption} to derive an expression for the cluster disruption time in environments with different GMC densities. The conclusions and discussion are presented in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have derive an expression for the cluster disruption time ($\\tdis$) due to encounters with GMCs, based on the results of $N$-body simulations of different types of encounters between a cluster and a GMC. The expression for \\tdis\\ (Eq.~\\ref{eq:tdis2}) allows us to estimate the disruptive effect of GMCs in different environments and to compare it to other effects such as the galactic tidal field, stellar evolution and spiral arm shocks. From Eq.~\\ref{eq:tdis2} we see that the disruption time for a $10^4\\,\\msun$ cluster in the solar neighbourhood is $t_4 = 2.0\\,$Gyr. This is a factor of 3.5 shorter than what was found by \\citet{2003MNRAS.340..227B}  for the combined effect of a realistic mass function, stellar evolution and the galactic tidal field ($t_4= 6.9\\,$Gyr). Note that the index $\\gamma$ is comparable, if we take the upper bound value of $\\lambda=0.10\\pm0.03$, since \\citet{2003MNRAS.340..227B} found $\\tdis\\propto\\left[\\Mcl/\\ln(\\Mcl)\\right]^{0.75}$. \\citet{2005A&A...429..173L} showed that $\\left[\\Mcl/\\ln(\\Mcl)\\right]^{0.75}\\propto \\Mcl^{0.62}$. The same mass dependence was also found from the age and mass distribution of cluster samples in different galaxies by \\citet{2003MNRAS.338..717B} and from the age distribution of clusters in the solar neighbourhood by \\citet{2005A&A...441..117L}. For this last sample they found a shorter scaling value of $t_4=1.3\\,$Gyr. Eq.~\\ref{eq:tdis2} combined with the observed relation between $\\rh$ and $\\Mcl$ (Eq.~\\ref{eq:rm}) suggests that the combined mass loss by GMC encounters and evaporation in a tidal field preserves the index $\\gamma\\simeq0.62$. A similar scaling between $\\tdis$ and $\\Mcl$ was found for the disruption time by spiral arm passages. The combination of all these effects and a comparison to the observed age distribution of open clusters in the solar neighbourhood will be discussed in more detail by \\citet{lamers06}.  The result of Eq.~\\ref{eq:tdis2} can easily be applied to environments/galaxies with different molecular gas densities ($\\sigmamol$). The mid-plane density ($\\rhon$) of GMCs scales with $\\sigmamol$ as $\\rhon=\\sigmamol/2h$, where $h$ is the vertical scale length of molecular gas, which is around $100\\,$pc in the solar neighbourhood \\citep{1987ApJ...322..706D}. Assuming that $h$ is the same in M51, and therefore $\\rhon=\\sigmamol/200\\,$pc, we can predict $t_4$ due to GMCs encounters once we know $\\sigmamol$ in M51 and compare this to the value of $t_4$ for clusters in M51 \\citep{2005A&A...441..949G}. \\citet{1993A&A...274..123G} find $\\sigmamol=90\\,\\msun\\,$pc$^{-2}$ for the central region of M51, which is 14 times higher than in the solar neighbourhood \\citep{1987ApJ...319..730S}. This implies $t_4\\simeq2100\\,{\\rm Myr}/14=150\\,$Myr. Note that this value is an upper limit, since we have assumed that the surface density of individual GMCs ($\\sigman$) is independent of $\\sigmamol$, while this could be higher in environments with higher $\\sigmamol$. In addition, the values of $f$ and $g$ in Eq.~\\ref{eq:tdis} will probably be higher in the centre of M51 (see \\S~\\ref{subsec:dedm} and \\S~\\ref{sec:disruption}), which will make \\tdis\\ shorter.  Also, we have assumed the same density profile as for the Galactic open cluster, e.g. $W_0 = 3$. A density profile of $W_0=7$, as found for the Galactic globular clusters, would reduce \\tdis\\  by a factor of two (Eq.~\\ref{sec:disruption}).  The resulting value of $t_4$, however, is very close to the value of $t_4=100-200\\,$Myr, determined observationally by \\citet{2005A&A...441..949G} from the age and mass distribution of star clusters in M51. Based on this work we think that we can explain the short disruption times of clusters in the centre of M51, if GMC encounters are the dominant disruption effect."
    },
    "0606/astro-ph0606498_arXiv.txt": {
        "abstract": "{We report the detection of 3 additional Wolf-Rayet stars in the young cluster Westerlund 1. They were selected as emission-line star candidates based on 1\\,$\\mu$m narrow-band imaging of the cluster carried out at OPD/LNA (Brazil), and then  confirmed as Wolf-Rayet stars by K-band spectroscopy performed at the 4.1\\,m {\\em SOAR} telescope (Chile). Together with previous works, this increases the population of Wolf-Rayet stars detected in the cluster to 22 members. Moreover, it is presented for the first time a K-band spectrum of the Luminous Blue Variable W243, which apparently implies in a higher temperature than that derived from optical spectra taken in 2003. The WC9 star WR-F was also observed, showing clear evidence of dust emission in the K-band. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606337_arXiv.txt": {
        "abstract": "A detailed imaging analysis of the globular cluster (GC) system of the Sombrero galaxy (NGC 4594) has been accomplished using a six-image mosaic from the Hubble Space Telescope Advanced Camera for Surveys.  The quality of the data is such that contamination by foreground stars and background galaxies is negligible for all but the faintest 5\\% of the GC luminosity function (GCLF).  This enables the study of an effectively pure sample of 659 GCs until $\\sim 2$ mags fainter than the turnover magnitude, which occurs at M$^{TOM}_{V} = -7.60\\pm0.06$ for an assumed m$-$M$=29.77$.  Two GC metallicity subpopulations are easily distinguishable, with the metal-poor subpopulation exhibiting a smaller intrinsic dispersion in color compared to the metal-rich subpopulation.  Three new discoveries include:  (1) A metal-poor GC color-magnitude trend.  (2) Confirmation that the metal-rich GCs are $\\sim17\\%$ smaller than the metal-poor ones for small projected galactocentric radii (less than $\\sim2$ arcmin).  However, the median half-light radii of the two subpopulations become identical at $\\sim$ 3 arcmin from the center.  This is most easily explained if the size difference is the result of projection effects.  (3) The brightest (M$_V < -9.0$) members of the GC system show a size-magnitude upturn where the average GC size increases with increasing luminosity.  Evidence is presented that supports an intrinsic origin for this feature rather than a being result from accreted dwarf elliptical nuclei.  In addition, the metal-rich GCs show a shallower positive size-magnitude trend, similar to what is found in previous studies of young star clusters. ",
        "introduction": "Extragalactic globular cluster (GC) systems provide an observational probe of galaxy formation and evolution.  GCs may have formed even before galaxies were fully assembled, as well as during subsequent star formation episodes \\citep[see references in ][]{bs06}.  Since the Milky Way serves as a standard against which other GC systems are compared, it is particularly important to study GC systems in other spiral galaxies to establish whether the Milky Way GC system is typical.  The study of spiral GC systems is complicated by their generally smaller GC populations and uncertainties in the internal extinction conditions. Nonetheless, data from the Hubble Space Telescope's (HST) Wide-Field Planetary Camera 2 (WFPC2) have helped establish the properties of GC systems in several spiral galaxies, including detailed studies of individual galaxies \\citep{kea99,lfb01,bh01,l00} as well as limited-scale surveys \\citep{fbl01,l02,gea03,cwl04}.  While these efforts illustrated the important role of space-based telescopes in extragalactic GC studies, large GC samples have been difficult to obtain for all but the most densely populated systems due to the small field of view of the WFPC2.  In contrast, the HST Advanced Camera for Surveys (ACS) has higher sensitivity and doubles the areal coverage. It is thus a powerful tool for global studies of GC systems, as demonstrated by a number of ACS surveys of such systems in early-type galaxies \\citep{pgea05,sea05B,hea05}.  \\begin{figure*} \\plotone{f1_color.eps} \\caption{Sombrero galaxy image mosaic from the Hubble Space Telescope Advanced Camera for Surveys.  The cross and circle symbols illustrate the locations of globular clusters belonging to metal-poor and metal-rich subpopulations, respectively.  Those objects found within the projected dust lanes were not included in the analysis to avoid uncertain extinction conditions and confusion with HII regions.  The region covered by a HST WFPC2 pointing used in an earlier Sombrero GC system study by \\citet{lfb01} is shown.\\label{figsomb}} \\end{figure*}  Figure~\\ref{figsomb} shows NGC 4594, the ``Sombrero galaxy'', which is a nearby ($9.0\\pm0.1$~Mpc or m$-$M$=29.77\\pm0.03$, see $\\S\\ref{lfsection}$ for derivation), nearly edge-on (i=84\\degr), spiral galaxy (type Sa) with a populous GC system.  The Sombrero is the most massive member of a small galaxy group.  While its classification as a spiral galaxy is questionable \\citep[e.g. ][ call it an S0]{rz04}, the Sombrero does appear to be an intermediate case between elliptical and late-type spiral galaxies, therefore a unique system to analyze.  The GC systems of both spiral and elliptical galaxies generally possess at least two GC subpopulations characterized as metal-poor (blue) and metal-rich (red) GCs \\citep{bs06}.  Due to the notorious metallicity-age degeneracy \\citep{w94}, ages and metallicities cannot be independently constrained by broad-band colors alone.  Ages derived from Keck spectra of 14 Sombrero GCs, suggest that both subpopulations are $\\sim 11-12$ Gyr old \\citep{lea02}, as is the case in the Milky Way GC system.  \\citet{hea02} used VLT spectra to derive mean metallicities of [Fe/H]$_{blue}=-1.7$ and [Fe/H]$_{red}=-0.7$ for the Sombrero's two GC subpopulations.  A ground-based, wide-field imaging study (R$_{gc} < $ 20\\arcmin) of the Sombrero by \\citet{mea02} showed that the red GCs are distributed spherically and suggested their association with the galaxy's bulge rather than its disk.  \\citet{rz04} used a $36\\arcmin\\times36\\arcmin$ wide-field, ground-based image to show that the GC surface density drops to zero at a projected radius of $\\sim19\\arcmin$ ($\\sim50$~kpc).  They estimated a total GC population of 1900 and calculated a specific frequency, S$_N=2.1\\pm0.3$, where S$_N = $N $\\times10^{0.4(M_V + 15)}$ \\citep{hv81}.  \\citet{lfb01} found that Sombrero's red GCs are, on average, $\\sim$30\\% smaller in size than its blue GCs. ",
        "conclusions": "A six-pointing mosaic from the HST Advanced Camera for Surveys was used to study the globular cluster system of the Sombrero galaxy.  An unprecedentedly deep ($\\sim95\\%$ of GCLF and $\\sim2$ mags fainter than the turnover magnitude), virtually contamination-free sample, allowed an accurate GC luminosity function (TOM) turnover magnitude to be determined:  m$_V^{TOM} = 22.17\\pm0.06$ or M$_V^{TOM} = -7.60\\pm0.06$, assuming m$-$M$=29.77\\pm0.03$.  The GCLF was found to be well fit by a t$_5$ to 2 magnitudes below the TOM (Figure~\\ref{figlfv}).  The dispersion of the t$_5$ fit is $\\sigma^{ACS}_{t5} = 1.08\\pm0.06$.  While the TOM is slightly fainter than a universal TOM, M$^{TOM}_{V} = -7.35$ \\citep{ric03,jea03}, it is likely consistent within the typical observational errors.  No variation with R$_{gc}$ is found in the Sombrero GCLF TOM, nor is a significant difference between the subpopulation TOMs detected.  The GC color distribution is clearly bi-modal (Figures~\\ref{figrawcmd} and \\ref{figgoodcmd}) with blue and red peaks at B$-$R $= 1.15$ and $1.44$ or [Fe/H] = $-1.38$ and $-0.49$, respectively (see Table~\\ref{tabcolor}).  Subpopulation peak colors are consistent with the expected values for a galaxy with a luminosity like the Sombrero, according to the relation between the GC subpopulation color peaks and host galaxy luminosity \\citep{sbf04}.  It was found that the color dispersion in the metal-poor subpopulation is smaller compared to the metal-rich subpopulation.  This is in agreement with the results of Harris et al. (2006), Peng et al. (2006), and Strader et al. (2006).  An estimated 53\\% of this GC sample belongs to the blue subpopulation.  This proportion contrasts to the results from ground-based observations where a larger proportion of 59\\% was found from data that covers the entire spatial distribution of the GC system \\citep{rz04}.  Thus the difference is likely caused by the decreasing ratio of red to blue GCs with galactocentric radius (Figure~\\ref{figradicolor}).  No color-R$_{gc}$ trend is observed within either subpopulation.  The galaxy color profile tends to be redder than metal-rich GCs and only begins to overlap this subpopulation at R$_{gc}>2\\arcmin$.  The best-fitting de~Vaucouleurs radial profile of the Sombrero GC system from \\citet{rz04} is steeper than the one found here.  However, their surface density values from within the ACS mosaic region are consistent with the ACS surface densities (Figure~\\ref{figsurfdensdeva}).  This observation is consistent with there being a core of constant number density in the inner GC system, confirming observations made in an earlier Sombrero GC system study \\citep{bh92}.  The number of red GCs normalized to the bulge luminosity, the bulge specific frequency, is found to be larger than all spiral galaxies for which this value has been calculated.  This suggests the rate of metal-rich GC formation was enhanced and/or GC destruction was less effective in the Sombrero galaxy relative to less massive spirals.  A color-magnitude trend, a so-called ``blue tilt'', was discovered in the metal-poor GC subpopulation that is likely a metallicity-mass trend (Figure~\\ref{figgoodcmd}).  The metal-rich subpopulation shows no strong evidence for an analogous trend.  The Sombrero provides the first example of this trend in a spiral galaxy and a galaxy found in a low-density galaxy environment.  The blue tilt extends from the brightest GCs to at least the GCLF TOM, in agreement with the findings of \\citet{sea05B}.  A shallower trend, in terms of metallicity-mass proportionality, is found here compared to examples in massive ellipticals found in high-density galaxy environments \\citep{sea05B,hea05}.  This difference and the absence of a blue tilt in NGC~4472 likely indicates the gradient of the color-magnitude trend depends on the specific properties of the host galaxy.  It is shown that GC sizes correlate with R$_{gc}$ (Figure~\\ref{figsizerad}).  From the extended radial coverage provided by the Sombrero ACS mosaic, the metal-rich GCs are found to have a steeper size-R$_{gc}$ gradient compared to the metal-poor subpopulation, thus causing the mean subpopulation sizes to become identical at larger R$_{gc}$ (Figure~\\ref{figcolsizerad}).  These observations are consistent with projection effects being largely responsible for the observed GC subpopulation size difference \\citep{lb03}, as opposed to an intrinsic feature of their respective metallicities \\citep{j04}.  A size-magnitude upturn was discovered among the brightest GCs (M$_V<-9.0$), where the sizes increase as a function of increasing luminosity (Figure~\\ref{figmagsiz}).  The accretion of dwarf elliptical nuclei into the Sombrero GC system provides a plausible scenario for this trend.  However, an analogous trend among young massive clusters \\citet{kpjb05} and a noticeable deficiency of luminous GCs (M$_V<-10.0$) with smaller than average sizes (Figure~\\ref{figmagsiz2}; also see figure~6~in Mieske et al. 2006), supports an intrinsic origin for the size-magnitude upturn.  An intrinsic size-magnitude upturn suggests the most massive GCs are at the low-mass end of a continuous trend that includes more massive systems such as dwarf nuclei and UCDs.  A shallower positive size-magnitude trend among metal-rich GCs is found at all magnitudes.  The luminosity-size proportionality, R$_{hl} \\propto L^{0.07\\pm0.03}$, is consistent with the same fits to young star clusters of the merger galaxy NGC 3256 \\citep{zea99} and 18 nearby spiral galaxies \\citep{l04}.  It is unclear whether the consistency is a reflection of a real trend or the result of a completeness effect."
    },
    "0606/astro-ph0606101_arXiv.txt": {
        "abstract": "We explore the possibility that extended disks, such as that recently discovered in M31, are the result of a single dwarf ($10^9$--$10^{10}M_\\odot$) satellite merger. We conduct N-body simulations of dwarf NFW halos with embedded spheriodal stellar components on co-planar, prograde orbits in a M31-like host galaxy. As the orbit decays due to dynamical friction and the system is disrupted, the stellar particles relax to form an extended, exponential disk-like structure that spans the radial range $30$--$200$ kpc. The disk scale-length $R_d$ correlates with the initial extent of the stellar component within the satellite halo: the more embedded the stars, the smaller the resulting disk scale-length. If the progenitors start on circular orbits, the kinematics of the stars that make up the extended disk have an average rotational motion that is $ 30 - 50$\\,km/s lower than the host's circular velocity.  For dwarf galaxies moving on highly eccentric orbits ($e \\simeq 0.7$), the stellar debris exhibits a much lower rotational velocity. Our results imply that extended galactic disks might be a generic feature of the hierarchical formation of spiral galaxies such as M31 and the Milky Way. ",
        "introduction": "\\label{sec:int}  Within the context of the hierarchical formation scenario, galaxies such as M31 are expected to have assimilated 100--500 smaller mass systems over a Hubble time (e.g.~Moore et al~1999). Today, stars from these systems are expected to be found in the disk, bulge and stellar halo (e.g Abadi et al. 2003, 2006, Governato et al. 2004 and references therein). The fossil signatures of these remnants are a veritable treasure trove of information about the galaxy assembly process and the nature of the merging subunits.  Detailed investigations of M31 (Ibata et al. 2001, 2004, 2005, Ferguson et al.  2002, 2005, McConnachie et al.  2003, 2004, Zucker et al.  2004, Fardal et al. 2006, Font et al.  2006, Guhathakurta et al. 2006, Kalirai et al. 2006, Brown et al. 2006, Chapman et al. 2006) reveal a large number of fossil substructures.  Intriguingly, many of these have disk kinematics and are distributed in a gigantic flattened structure surrounding the high-surface brightness inner disk (Ibata et al.~2005).  The observations are consistent with this extended structure being a rotating disk-like system whose midplane is coincident with that of the inner disk, but which extends $\\gta 80$ kpc from the center of M31.  The rotational velocities of the constituent stars lag the M31 disk rotation by $\\sim 30$\\,km/s (Ibata et al.~2005).  The mass of the extended disk is $\\sim 10\\%$ of the inner disk and its stellar surface density profile of the structure is consistent with a exponential decline although there is some uncertainty about the value of the radial scale length (cf. Ibata et al~2005; Kalirai et al.~2006).  Here, we explore whether the extended disk could be the remnant of a dwarf galaxy. ",
        "conclusions": ""
    },
    "0606/astro-ph0606271_arXiv.txt": {
        "abstract": "We report the results of a spectroscopic study of the Bo\\\"{o}tes (Boo) dwarf spheroidal (dSph) galaxy carried out with the WIYN telescope and the Hydra multifiber spectrograph. Radial velocities have been measured for 58 Boo candidate stars selected to have magnitudes and colors consistent with its red and asymptotic giant branches.  Within the 13\\arcmin~half-light radius, seven members of Boo yield a systemic velocity of $V_{r}=95.6\\pm3.4$ km s$^{-1}$ and a velocity dispersion of $\\sigma_{\\rm o}=6.6\\pm2.3$ km s$^{-1}$.  This implies a mass on the order of $1\\times10^7$ M$_{\\sun}$, similar to the inferred masses of other Galactic dSphs. Adopting a total Boo luminosity of $L=1.8\\times10^4$ L$_{\\sun}$ to $8.6\\times10^4$ L$_{\\sun}$ implies $M/L\\sim610$ to 130, making Boo, the most distorted known Milky Way dwarf galaxy, potentially also the darkest.  From the spectra of Boo member stars we estimate its metallicity to be [Fe/H] $\\sim-2.5$, which would make it the most metal poor dSph known to date. ",
        "introduction": "Since the early work of \\citet{Aaronson1983} on the Draco system, dwarf spheroidal (dSph) galaxies in the Local Group have been suspected to be heavily dark matter (DM) dominated.  Analyses of dSph internal dynamics under the assumption of virial equilibrium suggest mass contents far exceeding those inferred from their luminosities, with central mass-to-light ratios ($M/L$) ranging from a few to about a hundred (in solar units).  Such $M/L$ imply that dSphs have the largest DM fraction of all galaxy types in the universe.  Yet, despite the wide range of inferred $M/L$ for dSphs, their central velocity dispersions ($\\sigma_{\\rm o}$) and half-light radii ($r_{\\rm h}$) seem to be remarkably similar, $\\sim7-10$ km s$^{-1}$ and $\\sim200$ pc respectively.  Using these values, typical total dSph masses of a few times $10^7$ M$_{\\sun}$ are inferred (e.g., \\citealt{Mateo1998})\\footnote{Recent kinematical studies of dSphs have attempted to derive their mass content by modeling the shape of the velocity dispersion profile (e.g., \\citealt{Kleyna2002}; \\citealt{Walker2006a,Walker2006b}), but these calculations do not take into account likely tidal effects in the outskirts of dSphs, as pointed out by \\citet{Lokas2005,M05,M06,Westfall2006,Sohn2006}.}.  This mass conspiracy seemingly extends to even the very low end of the galaxy luminosity scale.  An early kinematical survey of the dSph in Ursa Major (UMa; \\citealt{Kleyna2005}), the faintest and most diffuse of the known Milky Way (MW) dwarf satellites (\\citealt[][hereafter {W05}]{W05}), found $\\sigma_{\\rm o}$ = 9.3$^{+11.7}_{-1.2}$ km s$^{-1}$, which, coupled with UMa's $r_{\\rm h}$ of 250 pc, also results in a mass of $\\sim10^7$ M$_{\\sun}$.  Recently, a new Galactic dSph candidate has been found in the constellation of Bo\\\"{o}tes (\\citealt[][hereafter {B06b}]{B06b}) during a search for halo substructure using Sloan Digital Sky Survey (SDSS; \\citealt{Abazajian2005}) data.  Along with UMa, Bo\\\"{o}tes (Boo) is one of the faintest MW satellites found, having (at a 60 kpc distance) an absolute magnitude of $M_{V}=-5.8$ (but maybe brighter; see \\S 3.2).  In addition, the Boo dSph exhibits the most irregular density contours of any Galactic dSph (\\citeauthor*{B06b}), which suggests that the satellite may be undergoing tidal disruption.  In this Letter, we present the results of a spectroscopic survey of Boo candidate stars. We have identified its kinematical signature, which allows us to estimate its systemic velocity, mass and central density, as well as its $M/L$.  We also use the spectral data to make a rough measurement of the metallicity of Boo giant stars. ",
        "conclusions": "We spectroscopically survey the Boo dSph and derive both its systemic velocity ($95.6\\pm3.4$ km s$^{-1}$) and central velocity dispersion ($6.6\\pm2.3$ km s$^{-1}$), which yields a mass of $M_{\\rm tot} = 1.1^{+1.3}_{-0.5}\\times10^7$ M$_{\\sun}$. This mass is similar to that of the other dSph galaxies and puts Boo squarely on the ``same mass-just different luminosities\" trend identified by \\citet{Mateo1998}, despite the fact that Boo is one of the faintest known Galactic satellites.  The dynamical mass derived for Boo, taken at face value, implies that this is possibly also the darkest dSph known to date.  If other systems of similar luminosity have the same mass, the current pace of discovery of these systems (three in the past year; \\citeauthor*{W05}; \\citealt{Zucker_CVn2006}; \\citeauthor*{B06b}) --- and in only the approximately 20\\% of the sky covered by the SDSS DR5 --- will help alleviate the current order of magnitude or two deficit of known Galactic satellites compared to that predicted by $\\Lambda$CDM simulations, albeit only for one part of the mass spectrum exhibiting the apparent ``missing satellites\" shortfall (e.g., \\citealt{Klypin1999}).  While Boo seems to be the most DM dominated dwarf, it is, at the same time, the Galactic satellite with the most distorted known morphology (\\citeauthor*{B06b}) and possibly most dramatic increase in velocity dispersion with radius.  This implies that Boo may be among the most disrupted Galactic dwarfs, to the extent that it even lacks a proper core (\\citeauthor*{B06b}).  In fact a puzzling correlation is now emerging between the DM fraction of a dSph and its morphology, wherein the faintest and most distorted systems (the UMi, UMa and Boo dSphs) seem also to present the largest central $M/L$'s.  If the distorted contours are a response to the influence of Galactic tides, and if all dSphs indeed have a similar current total mass and density, then why is it that tides seem to affect preferentially the least luminous systems? Perhaps these faintest systems represent the \\citet{Kroupa1997} regime in which disruption has proceeded to the point where the central velocity dispersions are inflated by tides, artificially increasing the derived $M/L$. But even if so, then it is curious that the current $\\sigma_{\\rm o}$ for Boo would be inflated to just such a value that the derived (but artificial) dynamical mass still participates in the ``same mass\" conspiracy of dSph galaxies.  We appreciate useful discussions with Gregory Sivakoff, Jeffrey Crane and Allyson Polak.  We gratefully acknowledge support by NSF grant AST-0307851, NASA/JPL contract 1228235, the Virginia Space Grant Consortium and Frank Levinson through the Celerity Foundation. D.L.N. is also supported by the ARCS Foundation and the Green Bank Telescope Student Support Program. P.M.F. is supported by NASA GSRP and UVa dissertation fellowships."
    },
    "0606/astro-ph0606047_arXiv.txt": {
        "abstract": "In light of the results from the WMAP three-year sky survey, we study an inflationary model based on a single-field polynomial potential, with up to quartic terms in the inflaton field. Our analysis is performed in the context of the Randall-Sundrum II braneworld theory, and we consider both the high-energy and low-energy (i.e. the standard cosmology case)  limits of the theory. We examine the parameter space of the model, which leads to both large-field and small-field inflationary type solutions. We conclude that small-field inflation, for a potential with a negative mass square term, is in general favored by current bounds on the tensor-to-scalar perturbation ratio $r_s$. ",
        "introduction": "Inflation, originally introduced to solve the initial condition problems of standard cosmology (SC)~\\cite{Guth:1980zm_plus}, is at present the favorite paradigm for explaining both the Cosmic Microwave Background Radiation (CMBR) temperature anisotropies and the initial conditions for structure formation. Indeed, it is during this epoch of accelerated expansion that the primordial density perturbations, which act as seeds for large-structure formation in the universe, are generated through the amplification of quantum fluctuations of the field(s)~\\cite{Mukhanov:1981xt_plus}. In the simplest inflationary models, the energy density of the universe is dominated by a single scalar field $\\phi$, the so-called inflaton, that slowly rolls down its self-interaction potential.  The recent publication of the three year results of the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{Hinshaw:2006ia_plus,Spergel:2006hy} continues to support the standard inflationary predictions. WMAP data provides no indication of any significant deviations from gaussianity and adiabaticity of the CMBR power spectrum and suggests that the universe is spatially flat to within the limits of observational accuracy. Moreover, it allows for very accurate constraints on the spectral index; the WMAP three-year data (WMAP3) yield~\\cite{Spergel:2006hy} \\begin{align} n_s = 0.951^{+0.015}_{-0.019}~, \\end{align} at $68\\%$ confidence level (CL), for vanishing running and no tensor modes. This result is in agreement with the one previously obtained in Ref.~\\cite{Sanchez:2005pi}, $n_s = 0.954 \\pm 0.023$, using  a joint analysis of the power spectrum of galaxy clustering measured from the final two-degree field galaxy redshift survey (2dfGRS) and a pre-WMAP3 compilation of measurements of both the temperature power spectrum and the temperature-polarization cross-correlation of the CMBR. A remarkable feature of these results is that the Harrison-Zel'dovich scale-invariant spectrum seems to be ruled out at around $3\\sigma$ level\\footnote{See Ref.~\\cite{Kinney:2006qm} for an analysis where the scale-invariant spectrum is consistent with WMAP3 data.}.  Another feature of WMAP3 data is the evidence for a significant running of the scalar spectral index~\\cite{Spergel:2006hy}, \\begin{align} \\alpha_s = -0.102^{+0.050}_{-0.043}~, \\end{align} at $68\\%$ CL and considering tensor perturbations\\footnote{If tensor perturbations are not taken into account slightly less negative values are obtained.}. This evidence was already present in the WMAP1 analysis \\cite{Spergel:2003cb_plus} and persists when large scale structure data is included~\\cite{Spergel:2006hy}, although it is diluted by the addition of Lyman-$\\alpha$ forest data~\\cite{Seljak:2004xh_plus,Lewis:2006ma,Seljak:2006bg}. We should note however that zero running is at about $2\\sigma$ from the central value and hence WMAP3 does not demand a nonvanishing running. In fact, vanishing running is still a good fit to the CMBR data, and the improvement of the $\\chi^2$ is small ($\\Delta \\chi^2 = -3$), if running is allowed. Notice also that if both tensor modes and running are taken into account, the WMAP team obtained $n_s= 1.21^{+0.13}_{-0.16}$ as the best fit value for the scalar tilt.   On the theoretical side, motivated by developments in string/M-theory, there has been considerable interest in the so-called braneworld constructions, where the matter fields are trapped in a lower dimensional brane, while (in the simplest models) only gravity can propagate into the bulk. Of special interest for cosmology is the Randall-Sundrum II (RSII) construction~\\cite{Randall:1999vf}, consisting of a single brane with positive tension embedded in a 5-dimensional bulk with a negative cosmological constant (anti-de Sitter spacetime). A remarkable feature of the RSII brane cosmology (BC) is the modification of the expansion rate of the universe, $H=\\dot{a}/a$, before the nucleosynthesis era~\\cite{Binetruy:1999hy_plus}. While in standard cosmology the expansion rate scales with the energy density $\\rho$ as $H\\propto\\sqrt{\\rho}$, this dependence becomes $H\\propto\\rho$ at very high energies in brane cosmology. This behavior may have important consequences on early universe phenomena such as inflation~\\cite{Maartens:1999hf,Bento:2001hu_plus} and the generation of the baryon asymmetry~\\cite{Mazumdar:2000gj_plus}.  In this paper we study the RSII braneworld inflationary period in the light of WMAP3 results, for a single-field power-law inflaton potential of the form~\\cite{Hodges:1989dw} \\begin{align} V(\\phi)= V_0 + \\dfrac{1}{2}\\, s\\, m^2\\, \\phi^2 + \\dfrac{1}{3} \\, m\\, \\widetilde{g} \\,\\phi^3 + \\dfrac{1}{4}\\, \\widetilde{\\lambda}\\, \\phi^4~.\\label{eq:potphi} \\end{align} Here $s=1$ (unbroken symmetry) or $s=-1$ (broken symmetry) and $m>0$. The parameters $\\widetilde{\\lambda}$ and $\\widetilde{g}$ are dimensionless constants; $\\widetilde{\\lambda}$ must be positive in order to ensure the stability of the potential, while $\\widetilde{g}$ can have either sign.  This potential has been studied previously in the literature in the SC context~\\cite{Hodges:1989dw,Cirigliano:2004yh,Boyanovsky:2005pw,deVega:2006hb} and covers a wide class of inflationary scenarios: for both spontaneously broken and unbroken (negative and positive mass square term, respectively) potentials one can obtain small and large-field inflation (hereafter also called new and chaotic inflationary solutions, respectively). Monomial potentials have already been widely studied both in standard cosmology (for a recent update following WMAP3 data see Refs. \\cite{Kinney:2006qm,Alabidi:2006qa_plus}) and in the braneworld context~\\cite{Maartens:1999hf,Bento:2001hu_plus,Tsujikawa:2003zd_plus}; these can be regarded as particular cases of the one considered here. It is also worth noticing that potentials of the type we are considering, Eq.~(\\ref{eq:potphi}), can be obtained as effective potentials in some supergravity models (for a recent study see Ref.~\\cite{Ibe:2006fs}).  The paper is organized as follows. After discussing in Section~\\ref{sec:potential} the main features of the inflaton potential, we briefly review in Section~\\ref{sec:BraneInflation} the slow-roll inflation formalism in the braneworld context. In Section~\\ref{sec:results} we present and discuss our results. Some final comments and conclusions are given in Section~\\ref{sec:conclusion}.  \\begin{figure*}[t] \\includegraphics[width=5.9cm]{Fig1a.eps} \\includegraphics[width=5.9cm]{Fig1b.eps} \\includegraphics[width=5.9cm]{Fig1c.eps}\\\\ \\caption{The potential of Eq.~(\\ref{eq:pot_dless}) for the cases discussed in the text. The full (dashed) lines correspond to large (small) field inflation. The curves correspond to different values of the asymmetry parameter $h$: (a) $h=-0.2,0$, (b) $h=-0.8,0$ and (c) $h=-1.15,-\\sqrt{9/8},-1.02$. We will only consider inflation occurring on the black branches of the potential (cf. Table~\\ref{table:pot}).} \\label{fig:potential} \\end{figure*} \\begin{table*}[t!] \\begin{tabular}{c c c c c c} \\hline \\hline &s & $h$ range & $\\varphi$ range & Inflation type & Case\\\\ \\hline &$-1$ & $h \\leq 0$ & $0 \\lesssim \\varphi < \\varphi_+$ & Small-field (a)&  A\\\\ &$-1$ & $h \\leq 0$ & $\\varphi > \\varphi_+$ & Large-field (a)&  B\\\\ &$+1$ & $-1 \\leq h \\leq 0$ & $\\varphi > 0$ & Large-field (b)& C\\\\ &$+1$ & $-\\sqrt{9/8} \\leq h \\leq -1 $ & $0 < \\varphi \\lesssim \\varphi_-$ & Large-field (c)& D\\\\ &$+1$ & $h \\leq -\\sqrt{9/8}$ & $\\varphi_- \\lesssim \\varphi < \\varphi_+$ & Small-field (c)& E\\\\ &$+1$ & $h \\leq -\\sqrt{9/8}$ & $\\varphi > \\varphi_+$ & Large-field (c)& F\\\\ \\hline \\hline \\end{tabular} \\caption{Cases A-F correspond to different ranges of the potential parameters and field values, see Eqs. (\\ref{eq:pot_dless}) and (\\ref{eq:hdef}), leading to large and small-field inflation. Labels (a)-(c) correspond to the cases shown in Fig.~\\ref{fig:potential}.} \\label{table:pot} \\end{table*} ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have explored inflationary solutions for a single-field polynomial potential, with up to quartic terms in the inflaton field, in light of the results from the WMAP three-year sky survey. Our analysis is done in the framework of the Randall-Sundrum II braneworld theory, and we have examined both the high-energy and standard cosmology regimes. As in the case of SC, the model displays large-field and small-field inflationary type solutions compatible with the observational data.  For a potential with a negative mass square term, cases A and B of Table I, we conclude that small-field inflation is in general favored by current bounds on $r_s$. We also get the bound $\\widetilde{\\lambda} \\lesssim 10^{-13}$ on the quartic coupling of the inflaton potential. The case with positive mass square term and $\\vert h\\vert\\leq 1$ (large-field solutions), which corresponds to case C of Table I, is also very much constrained by the $r_s$ bound and, in particular, this type of solutions  will be excluded if, as expected, $r_s < 0.2$. In this case we get the following bounds on the parameters of the inflationary potential: $\\widetilde{\\lambda} \\lesssim 10^{-12}$ and $|\\widetilde{g}| \\lesssim 10^{-6}$.  We have also made an estimate of the reheating temperature for this model and we have found that, for $M_5 > 10^{16}$~GeV and $N_\\star \\leq 60$, one has $T_{\\textrm{reh}}^{BC} \\simeq T_{\\textrm{reh}}^{SC} \\lesssim 10^{13}-10^{14}$ GeV, whereas for $M_5 < 10^{16}$~GeV, one gets the relation $T_{\\textrm{reh}}^{BC} \\sim 10^2\\, T_{\\textrm{reh}}^{SC}{M_5}/{M_P}\\,$. At this point, it is worth noting that in supergravity models the presence of the gravitino could lead to serious cosmological problems unless the reheating temperature is sufficiently low. In particular, in gravity-mediated supersymmetry breaking models, the gravitino mass is expected in the range $m_{3/2} \\sim \\mathcal{O}(10^{2-4})$~GeV. Such a gravitino is most likely unstable and its decays could destroy the light elements synthesized during BBN. On the other hand, in gauge-mediated supersymmetry breaking models, the gravitino can be the lightest supersymmetric particle, and thus stable, with a mass $m_{3/2} \\lesssim \\mathcal{O}(10)$~GeV. In this case, its contribution to the present cosmic density can be excessive. In either case, this leads to stringent constraints on $T_{\\textrm{reh}}$ after inflation~\\cite{Kawasaki:2006hm}. In the braneworld scenario, one expects such bounds to translate into additional constraints on the 5D fundamental Planck mass $M_5$.  Finally, let us comment on the running of the scalar index $\\alpha_s$. One should notice that it is not possible to obtain large values for $\\alpha_s$, either in SC or the high-energy BC scenarios, for polynomial potentials of the form given in Eq.~(\\ref{eq:potphi}). However, this is still consistent with observation, since the evidence for running is still weak and it can evaporate as more data becomes available. In fact, the addition of Lyman-$\\alpha$ forest data reduces the need for a negative value of $\\alpha_s$~\\cite{Seljak:2004xh_plus,Lewis:2006ma,Seljak:2006bg}. Recently Easther $\\&$ Peiris~\\cite{Easther:2006tv} have analyzed the implications of a large running spectral index for single-field slow-roll inflation in SC, using the inflationary flow equations~\\cite{Hoffman:2000ue_plus,Liddle:2003py} and retaining all terms in $\\alpha_s$ up to quadratic order in the slow-roll parameters. They found that, for all parameter choices consistent with a large negative running, inflation lasts less than 30 e-folds after CMBR scales leave the horizon. They concluded that a definitive observation of a large negative running would imply that inflationary stage requires multiple fields or the breakdown of slow-roll. The underlying conclusions of their work can presumably be carried over to the BC case, as the flow equations are quite insensitive to the expansion dynamics~\\cite{Ramirez:2004fb}. However, one should notice, as shown in~\\cite{Liddle:2003py}, that the flow equation approach does not directly incorporate inflationary dynamics. In fact, one can determine analytically the set of inflationary potentials which correspond to solutions of the truncated flow equations. Hence, conclusions obtained from an analysis based on the flow equations should be considered with some care, since it is certainly possible to find potential shapes that are not represented by the set of inflationary potentials for which the formalism is valid. An example of a model where it is possible to obtain a considerably large running was explored in Ref.~\\cite{Ballesteros:2005eg}."
    },
    "0606/astro-ph0606682_arXiv.txt": {
        "abstract": "It is now accepted that accretion onto classical T Tauri stars is controlled by the stellar magnetosphere, yet to date most accretion models have assumed that their magnetic fields are dipolar.  By considering a simple steady state accretion model with both dipolar and complex magnetic fields we find a correlation between mass accretion rate and stellar mass of the form $\\dot{M} \\propto M_{\\ast}^{\\alpha}$, with our results consistent within observed scatter.  For any particular stellar mass there can be several orders of magnitude difference in the mass accretion rate, with accretion filling factors of a few percent.  We demonstrate that the field geometry has a significant effect in controlling the location and distribution of hot spots, formed on the stellar surface from the high velocity impact of accreting material.  We find that hot spots are often at mid to low latitudes, in contrast to what is expected for accretion to dipolar fields, and that particularly for higher mass stars, the accretion flow is predominantly carried by open field lines. ",
        "introduction": "Classical T Tauri stars (CTTSs) are young, low mass, pre-main sequence stars that are actively accreting from a surrounding disc which is the eventual birth-place of planets. \\citet{uch84} suggested that the magnetic field of a CTTS disrupts the inner disc.  In the early 1990s several magnetospheric accretion models were developed (\\citealt{kon91}; \\citealt{col93}; \\citealt{shu94}) where material is lifted from the disc plane and is channelled along dipolar magnetic field lines onto the star,  terminating in a shock at the photosphere. In an idealised model of a CTTS's magnetic field there are closed field lines close to the star that contain the X-ray emitting corona, whilst at larger radii, there are closed field lines which thread the circumstellar disc.  It is along this latter set of field lines that accretion may proceed.  There are also regions of open field which carry outflows in the form of a wind, and in some cases, as large collimated bipolar jets.  Magnetospheric accretion models assume that CTTSs possess magnetic fields that are strong enough to disrupt the disc at a distance of a few stellar radii.  Such strong fields have been detected in a number of systems using a variety of techniques.  Average surface fields of 1-3kG have been detected most successfully by exploiting the Zeeman effect, both through Zeeman broadening (e.g. \\citealt*{joh99b}) and from the circular polarisation of lines which are sensitive to the presence of a magnetic field (e.g. \\citealt{joh99a}; \\citealt{sym05}; \\citealt*{dao06}).  Field detections have also been made from the increase in line equivalent width (\\citealt*{bas92}; \\citealt{gue99}) and also from electron cyclotron maser emission, a coherent emission process from mildly relativistic electrons trapped inside flux tubes close to the star \\citep{smi03}.  The mean magnetic field strengths detected so far appear to be roughly constant across all stars \\citep{val04}.  Traditionally magnetospheric accretion models have assumed the CTTSs have dipolar magnetic fields.  Dipole fields (or inclined dipole fields) have been successively used to explain some of the observations of CTTSs (e.g. the photopolarimetric variability of AA Tau, \\citealt{osu05}), but fail to account for others.  \\citet{val04} present magnetic field measurements for a number of stars, and despite detecting strong average surface fields from Zeeman broadening, often measurements of the longitudinal (line-of-sight) field component (obtained from photospheric lines) are consistent with no net circular polarisation.  This can be interpreted as there being many regions of opposite polarity on the stellar surface, giving rise to oppositely polarised signals which cancel each other out giving a net polarisation signal of zero.  This suggests that CTTSs have magnetic fields which are highly complex, particularly close to the stellar surface; however, as \\citet{val04} point out, as the higher order multi-pole field components will drop off quickly with distance from the star, the dipole component may still remain dominant at the inner edge of the disc.  Also their measurements of the circular polarisation of the HeI 5876{\\AA} emission line (believed to form in the base of accretion columns) are well fitted by a simple model of a single magnetic spot on the surface of the star, suggesting that the accreting field may be well ordered, despite the surface field being complex.  The fractional surface area of a CTTS which is covered in hot spots, the accretion filling factor $f_{acc}$, is inferred from observations to be small; typically of order one percent (\\citealt{muz03}; \\citealt{cal04}; \\citealt{val04}; \\citealt{sym05}; \\citealt{muz05}).  Dipolar magnetic field models predict accretion filling factors which are too large.  This, combined with the polarisation results, led \\citet{joh02} to generalise the Shu X-wind model \\citep{shu94} to include multipolar, rather than dipolar, magnetic fields.  With the assumption that the average surface field strength does not vary much from star to star the generalised Shu X-wind model predicts a correlation between the stellar and accretion parameters of the form $R_{\\ast}^2f_{acc} \\propto (M_{\\ast}\\dot{M}P_{rot})^{1/2}$, a prediction that matches observations reasonably well.  In this paper we present a model of the accretion process using both dipolar and complex magnetic fields.  We apply our model to a large sample of pre-main sequence stars obtained from the Chandra Orion Ultradeep Project (COUP; \\citealt{get05}), in order to test if our model can reproduce the observed correlation between mass accretion rate and stellar mass.  An increase in  $\\dot{M}$ with $M_{\\ast}$ was originally noted by \\citet{reb00} and subsequently by  \\citet{whi01} and \\citet{reb02}.  The correlation was then found to extend to very low mass objects and accreting brown dwarfs by \\citet{whi03} and \\citet{muz03}, and to the higher mass, intermediate mass T Tauri stars, by \\citet{cal04}.  Further low mass data has recently been added by \\citet{nat04}, \\citet*{moh05} and by \\citet{muz05} who obtain a correlation of the form $\\dot{M} \\propto M_{\\ast}^{2.1}$, with as much as three orders of magnitude scatter in the measured mass accretion rate at any particular stellar mass.  However, \\citet{cal04} point out that due to a bias against the detection of higher mass stars with lower mass accretion rates, the power may be less than 2.1.  Further data for accreting stars in the $\\rho-$Ophiuchus star forming region has recently been added by \\citet*{nat06}.  The physical origin of the correlation between $M_{\\ast}$ and $\\dot{M}$, and the large scatter in measured $\\dot{M}$ values, is not clear; however several ideas have been put forward.  First, increased X-ray emission in higher mass T Tauri stars (\\citealt{pre05}; \\citealt{jar06}) may cause an increase in disc ionisation, leading to a more efficient magnetorotational instability and therefore a higher mass accretion rate \\citep{cal04}.  Second, \\citet{pad05} argue that the correlation $\\dot{M} \\propto M_{\\ast}^{2}$ arises from Bondi-Hoyle accretion, with the star-disc system gathering mass as it moves through the parent cloud.  In their model the observed scatter in $\\dot{M}$ arises from variations in stellar velocities, gas densities and sound speeds. \\citet{moh05} provide a detailed discussion of both of these suggestions. Third, \\citet{ale06} suggest that the correlation may arise from variations in the disc initial conditions combined with the resulting viscous evolution of the disc.  In their model they assume that the initial disc mass scales linearly with the stellar mass, $M_d \\propto M_{\\ast}$, which, upon making this assumption, eventually leads them to the conclusion that brown dwarfs (the lowest mass accretors) should have discs which are larger than higher mass accretors.  However, if it is the case that the initial disc mass increases more steeply with stellar mass, $M_d \\propto M_{\\ast}^2$, then the stellar mass - accretion rate correlation can be reproduced with smaller brown dwarf discs of low mass (of order one Jupiter mass).  Thus the \\citet{ale06} suggestion, if correct, will soon be directly verifiable by observations.  Fourth, \\citet{nat06} suggest that the large scatter in the correlation between $\\dot{M}$ and $M_{\\ast}$ may arise from the influence of close companion stars, or by time variable accretion. It should however be noted that \\citet{cla06} take a more conservative view by demonstrating that a steep correlation between $\\dot{M}$ and $M_{\\ast}$ may arise as a consequence of detection/selection limitations, and as such $\\dot{M} \\propto M_{\\ast}^{2}$ is perhaps not a true representation of the correlation between $\\dot{M}$ and $M_{\\ast}$.  In \\S2 we describe how magnetic fields are extrapolated from observed surface magnetograms. In \\S3 we consider accretion onto an aligned, and then a tilted dipole field, to develop a simple steady state accretion model and to investigate how tilting the field affects the mass accretion rate.  In \\S4 these ideas are extended by considering magnetic fields with a realistic degree of complexity and we apply our accretion model to study the correlation between mass accretion rate and stellar mass, whilst \\S5 contains our conclusions. ",
        "conclusions": "By considering accretion to both dipolar and complex magnetic fields we have constructed a steady state isothermal accretion model were material leaves the disc at low subsonic speeds, but arrives at the star at large supersonic speeds.  We find in-fall velocities of a few hundred kms$^{-1}$ are possible which is consistent with measurements of the redshifted absorption components of inverse P-Cygni profiles (e.g. \\citealt{edw94}).  We found that for accretion along aligned and perpendicular closed dipole field lines that there was little difference in in-fall speeds between the two cases.  However, for tilted dipole fields in general (rather than just considering a single closed field line loop) there are many open and closed field lines threading the disc providing different paths that material could flow along from the disc to the star. The path that a particular field line takes through the star's gravitational potential combined with the strength of magnetic field and how that changes along the field line path, are all important in determining whether or not a field line can support transonic accretion.  At low accretion flow temperatures open field lines are typically not able to support transonic accretion, but at higher accretion flow temperatures they can accrete transonically and consequently we see an increase in the mass accretion rate. We only find a factor of two to three difference in $\\dot{M}$ values for accretion to tilted dipolar magnetic fields suggesting that the geometry of the field itself is not as significant as the stellar parameters (which control the position of corotation) in controlling the mass accretion rate.  The magnetic field geometry is crucial in controlling the location and distribution of hot spots on the stellar surface.  For accretion to complex magnetic fields we find that hot spots can span a range of latitudes and longitudes, and are often at low latitudes towards the star's equator.  We find that the accretion filling factors (the fractional surface area of a star covered in hot spots) are small and typically around 2.5\\%, but they can vary from less than 1\\% to over 4\\% and rarely to larger values.  This is consistent with observations which suggest small accretion filling factors and the inference of hot spots at various latitudes \\citep{val04}.  For accretion with complex magnetic fields there is a distribution of in-fall speeds, which arises from variations in the size and shape of accreting field lines.  The resulting effect that this will have on line profiles will be addressed in future work. In our model most of the accretion occurs from the corotation radius, but at some azimuths the disc extends closer to the star meaning that a small fraction of field lines are accreting material at lower velocity.  Lower mass stars, with their lower surface gravities, typically have larger coronae which would extend out to corotation \\citep{jar06}.  Therefore the lower mass stars are accreting along a mixture of open and closed field lines from the corotation radius. In contrast higher mass stars, with their higher surface gravities, have small compact coronae, and so the star is actively accreting along mainly open field lines from the corotation radius.  However, such open field lines do not thread the disc at all azimuths, with some accretion instead occuring along field lines which are much closer to the stellar surface.  This gives rise to a small peak at low velocity in the in-fall velocity distribution.  However, this represents only a small fraction of all accreting field lines which do not contribute significantly to the resulting mass accretion rate.  Finally we applied our accretion model to stars from the COUP dataset which have estimates of the stellar parameters and measurements of the equivalent width of the CaII 8542{\\AA} line, which is seen in emission for accreting stars.  For the complex magnetic fields we calculated mass accretion rates and accretion filling factors as a function of stellar mass.  The observed stellar mass - accretion rate correlation is $\\dot{M}\\propto M_{\\ast}^2$ (\\citealt{muz03}; \\citealt{cal04}; \\citealt{moh05}; \\citealt{muz05}), however, this may be strongly influenced by detection/selection effects \\citep{cla06}. By only considering observational data across the range in mass provided by the COUP sample of accreting stars, the observed correlation becomes $\\dot{M}\\propto M_{\\ast}^{1.4}$.  Our steady state isothermal model gives an exponent of 1.1 for the LQ Hya-like magnetic field and 1.2 for the AB Dor-like field with similar results for both the high and low coronal temperatures, with the caveat that the observed correlation may be less than 1.4 due to a strong bias against the detection of higher mass stars with lower mass accretion rates \\citep{cal04}.  It may be the case that an exponent of 1.2 compared to the observed 1.4, represents the best value that can be achieved with a steady state isothermal accretion model.  \\citet{jar06} have used this model to reproduce the observed increase in X-ray emission measure with stellar mass \\citep{pre05}. However, they find that when using the complex magnetic fields presented here (extrapolated from surface magnetograms of the young main-sequence stars AB Dor and LQ Hya) they slightly under estimate the emission measure-mass correlation.  When they use dipolar magnetic fields, which represent the most extended stellar field, they slightly over estimated the correlation.  This suggests that T Tauri stars have magnetic fields which are more extended than those of young main sequence stars, but are more compact than purely dipolar fields.  Also, our model has not been tested across a large enough range in mass to make the comparison with the observed correlation of $\\dot{M} \\propto M_{\\ast}^2$.  However, it already compares well with the alternative suggestion of \\cite{cla06} that the $\\dot{M}-M_{\\ast}$ correlation is not as steep with $\\dot{M} \\propto M_{\\ast}^{1.35}$.  In future we will extend our model to consider accretion from the lowest mass brown dwarfs, up to intermediate mass T Tauri stars. This will require estimates of, in particular, rotation periods and coronal temperatures.  Our accretion model reproduces a similar amount of scatter in calculated $\\dot{M}$ values compared with observations.  This can be attributed to different sets stellar parameters changing the structure of the accreting field.  The stellar parameters control the location of the corotation radius, whilst it is the position of corotation relative to the natural coronal extent, which determines whether or not accretion occurs predominantly along the open field.   However, for the moment we are restricted to using surface magnetograms of young main sequence stars which may not necessarily represent the true fields of T Tauri stars.  It could be possible that variations in magnetic field geometry from star to star are responsible for the observed large scatter in the mass accretion rate at any particular stellar mass.  If there is a large difference in the structure of the accreting field from star to star this could lead to large variations in mass accretion rates onto different stars.  The spectropolarimeter at the Canada France Hawaii telescope, ESPaDOnS: Echelle SpectroPolarimetric Device for the Observation of Stars (\\citealt{pet03}; \\citealt{don04}), will in the near future allow the reconstruction of the magnetic field topology of CTTSs from Zeeman-Doppler imaging.  This will allow a number of open questions to be addressed about the nature of the magnetic fields of T Tauri stars and allow us to test if our accretion model can indeed reproduce the stellar mass - accretion rate correlation."
    },
    "0606/astro-ph0606361_arXiv.txt": {
        "abstract": "We report on the results of the first 3D SPH simulation of massive, gravitationally unstable protoplanetary disks with radiative transfer. We adopt a flux-limited diffusion scheme justified by the high opacity of most of the disk. The optically thin surface of the disk cools as a blackbody. The disks grow slowly in mass starting from a Toomre-stable initial condition to the point at which they become marginally unstable. We find that gravitationally bound clumps with masses close to a Jupiter mass can arise. Fragmentation appears to be driven by vertical convective-like motions capable of transporting the heat from the disk midplane to its surface on a timescale of only about 40 years at 10 AU. A larger or smaller cooling efficiency of the disk at the optically thin surface can promote or stifle fragmentation by affecting the vertical temperature profile, which determines whether convection can happen or not, and by regulating the accretion flow from optically thin regions towards overdense regions. We also find that the chances of fragmentation increase for a higher mean molecular weight, $\\mu$, since compressional heating is reduced. Around a star with mass $1 M_{\\odot}$ only disks with $\\mu \\ge 2.4$, as expected for gas with a metallicity comparable to solar or higher, fragment. This suggests that disk instability, like core-accretion, should be more effective in forming gas giants at higher gas metallicities, consistent with the observed correlation between metallicity of the planet-hosting stars and frequency of planets. ",
        "introduction": "Long-lived clumps can form in a massive gravitational unstable protoplanetary disk if the gas is isothermal (Mayer et al. 2002) or cools on a timescale comparable to the local orbital time (Rice et al. 2003a,b, Mayer et al. 2004a, Rice et al. 2005, Mayer et al. 2005, Mejia et al. 2005). Fragmentation under such favourable conditions is a numerically robust result obtained with SPH codes as well as grid-based codes with both static and adaptive meshes (Durisen et al. 2005, Mayer et al., in prep.). The main open question is whether the required short cooling times can be obtained in the disks. Shock heating along spiral arms is intense during the phase of strong gravitational instability (Pickett et al. 2000, 2003; Mayer et al. 2004b) and will tend to erase clumps exactly when they have the best chance to collapse. Boss (2002a,2002b,2004) has used a grid-based code which solves radiation transport with the  diffusion approximation. He found that the disk cools rapidly through convection instead of radiation and forms long lasting clumps. Cai et al. (2006) adopt flux-limited diffusion in their cylindrical grid code. They find long cooling times, no evidence of convection and no fragmentation. Their method is very different from that of Boss (2004). While Boss does not use a flux-limiter but embeds the disk in a thermal bath at a constant temperature in the range 40-50 K,  Cai et al add an Eddington-like atmosphere on top of the optically thick layer, explicitly matching the fluxes at the boundary. Nelson et al. (2000) also did not find fragmentation in a 2D SPH simulation to which a plane-parallel atmosphere radiating as a blackbody was added assuming that rapid vertical cooling was achieved through convection. Here we present the first 3D SPH simulations of protoplanetary disks with flux-limited diffusion. ",
        "conclusions": "We have shown that disks can fragment even when radiative transfer is included. However, fragmentation requires a disk somewhat more massive than in previous works relying on simple equations of states or parameterizations of the cooling time., $M > 0.12 M_{\\odot}$ as opposed to $\\sim 0.1 M_{\\odot}$ (Mayer et al. 2004a,b). It is strongly dependent on the efficiency of cooling at the optically thin boundary and  on the molecular weight. The newly discovered dependence  on molecular weight is intriguing since it establishes for the first time a direct link between the disk instability mechanism and the metallicity of the gas.  Such a link is straightforward in the competing model, core-accretion (Pollack et al. 1996). Hence the formation of giant planets by disk instability should be favoured around more metal rich stars, in line with the observed correlation (Fischer \\& Valenti 2005). A more metal rich disk will also have a higher cooling rate in the optically thin region (Cai et al. 2006), which also appears to promote gravitational instability, whilst the associated increase of opacity in the optically thick regions will not affect the outcome based on our results.   \\medskip \\medskip {\\centering \\epsfxsize=8.5truecm \\epsfbox{Fig3b.ps} \\figcaption[mbar1.ps]{\\label{fig:asymptotic} \\small{Color coded logarithmic temperature plots of a region of size about 4.5 AU centered on one of the gravitationally bound clumps in the run with $\\mu=2.7$ and EDA=40 degrees. Left: a smoothed temperature map is shown (the temperature range goes from 10 K (dark blue) to 250 K (dark red) at time $t=1248$ years showing the cold, low density regions adjacent to the hot gas in the spiral shock and clump. Right: temperature map of the gas particles surrounding the same clump about half an orbit earlier with the particles that will accrete by $t=1248$ years marked in green. The brighter the color the higher the temperature. Most of the accreted mass appears to come from the colder regions outside the spiral shock.}}} \\medskip     The radiation physics in the simulations presented here is still simplified. In particular, the disk edge could be heated as a result of irradiation from the central star and/or neighboring bright stars, erasing the steep temperature gradient necessary to drive convection. On the other end, when an overdensity forms dust might accumulate at its location owing to the action of gas drag (Haghighipour \\& Boss,. 2003, Durisen et al. 2005), thus increasing locally the molecular weight and enhancing the chances of fragmentation. In the latter case the clumps would be automatically enriched in heavy elements relative to the surrounding disk, as seen in the gas giants of the Solar System (Saumon \\& Guillot 2004) and predicted by recent models of core accretion (Alibert et al. 2005).  There are situations in which fragmentation may be achieved more easily than suggested here. Indeed a new set of simulations being completed as we write indicates that if the mass of the central star is lowered to $M=0.5 M_{\\odot}$, the disk fragments for masses $< 0.1 M_{\\odot}$ for a molecular weight as small as $\\mu = 2$. In this case a lower shear allows a higher pressure front to be tolerated from growing overdensities. Therefore with disk instability formation of giant planets around M dwarfs is likely (see also Boss 2006). We also find that if the disk does not grow slowly in mass but starts with $Q_{min}=1.4$, as it might happen in response to a sudden external perturbation, fragmentation occurs with $\\mu=2$ and edge detection angles of up to 60 degrees. On the other end, previous work  shows that long lasting, time-dependent perturbations such as those induced by the presence of a binary companion might quench the fragmentation by increasing the heating due to spiral shocks (Mayer et al. 2005).  In summary, it seems that an ultimate answer requires that all the details of the radiation physics in the disks and in their surrounding environment are taken into account together with the details of the disk formation process.   \\bigskip  L.M. thanks Richard Durisen, Willy Benz, Annie Mejia, Willy Kley, Yann Alibert and Laure Fauchet for fruitful discussions. TQ acknowledges support from the Nasa Astrobiology Institute. Simulations were performed on LeMieux at the Pittsburgh Supercomputing Center and on Zbox at the University of Zurich."
    },
    "0606/astro-ph0606157_arXiv.txt": {
        "abstract": "Recent observations reveal galaxies in the early Universe ($2<z<6.4$) with large reservoirs of molecular gas and extreme star formation rates. For a very large range of sources, a tight relationship exists between star formation rate and the luminosity of the HCN J=1-0 spectral line, but sources at redshifts of z$\\sim$2 and beyond do not follow this trend. The deficit in HCN is conventionally explained by an excess of infrared (IR) radiation due to active galactic nuclei (AGN). We show in this letter not only that the presence of AGN cannot account for the excess of IR over molecular luminosity, but also that the observed abundance of HCN is in fact consistent with a population of stars forming from near-primordial gas. ",
        "introduction": "Recent work (Gao \\& Solomon 2004, Wu et al. 2005) indicates that there exists a strong correlation between the luminosity of molecular lines \\--- particularly the HCN J=1-0 transition\\--- and the star formation rate derived from infrared luminosity. Impressively, these relations apply to individual star forming regions in the Milky Way and also to external systems across ten orders of magnitude in luminosity. However, the three highest redshift systems in the sample used by Gao \\& Solomon (at $z\\sim 2$) do not fit this trend, but have luminosities in HCN lower than would be expected given their star formation rate. A detection of the HCN J=5-4 line in APM 08279+5255 at a redshift z=3.91 has been reported with $L_{FIR}$/$L_{HCN}$ similar to that seen in other high-z sources (Wagg et al. 2005). Limits on HCN luminosities obtained for systems at much higher redshifts ($z \\sim 6$) (Carilli et al. 2002) indicate that these systems also do not fit the low-z relation.  The conventional explanation for these discrepancies is to attribute the observed infrared excess to AGN activity. The $z\\sim 2$ systems are, indeed, quasar hosts, presenting us with the problem of distinguishing IR radiation associated with star formation and that due to the AGN. Detailed modeling of the spectral energy distribution (SED) of one such source, the Cloverleaf quasar (Wei\\ss~ et al. 2005), has shown that there exists a warm dust component which is believed to be associated with star-forming activity. In this letter we test the hypothesis that hidden AGN activity is responsible for the deviation of these high-z sources from the observed relationship between infrared and HCN luminosity. By using chemical models of high-mass star formation, we in fact find that AGN activity should be accompanied by an \\emph{increase} in the relative abundance of HCN. Furthermore, we find that models of star formation from near-primordial gas, enriched only by the ejecta from population III stars, can account for the observed HCN abundance; hence we show that there is no `lack' of HCN at high redshift. ",
        "conclusions": "We have argued that the observed differences between $L_{\\mathrm{IR}}$/$L_{\\mathrm{HCN}}$ at low and high redshifts cannot be due to the presence of AGN enhancement of the infrared luminosity. Such activity, we show, would increase the abundance of HCN in the observed systems. Detailed chemical models of such systems show promise in distinguishing between AGN enhanced and purely star forming activity through the ratio of HCN with species which are not enhanced by AGN. Examples of such species may include CS and $\\mathrm{SO_2}$ (both of which are enhanced in galactic hot cores (Hatchell et al. 1998) and are expected to be favoured in nitrogen-deficient gas (Lintott et al. 2005).  We explain the observations instead by hypothesizing that stars are forming from near-primordial material. Such material is believed to be deficient in nitrogen compared to the solar neighbourhood.  Jiminez and Haiman (2006) discuss four observational constraints on the nature of star formation at redshifts between 3 and 4. Specifically, they explain significant UV emission in Ly-break galaxies, Ly-alpha emission isolated from ionizing continuum radiation, a population of sources with strong Ly-alpha emission and strong HeII lines by proposing a significant fraction of star formation from primordial gas. These results imply that mixing of metals is inefficient within galaxies, and so solar abundance ratios may not have been reached at a redshift of z=2.  Such models present a challenge to existing models of chemical enrichment, which assume that the first population II stars form at much higher redshifts. Simulations (e.g. Abel et al. 2002) which suggest that population III stars will be massive assume that such stars will form at redshifts much greater than 4; it may be that the predictions of yields from such stars may need to be revised."
    },
    "0606/astro-ph0606227_arXiv.txt": {
        "abstract": "The Fisher matrix formalism has in recent times become the standard method for predicting the precision with which various cosmological parameters can be extracted from future data.  This approach is fast, and generally returns accurate estimates for the parameter errors when the individual parameter likelihoods approximate a Gaussian distribution.  However, where Gaussianity is not respected (due, for instance, to strong parameter degeneracies), the Fisher matrix formalism loses its reliability.  In this paper, we compare the results of the Fisher matrix approach with those from Monte Carlo simulations.  The latter method is based on the publicly available CosmoMC code, but uses synthetic realisations of data sets anticipated for future experiments.  We focus on prospective cosmic microwave background (CMB) data from the Planck satellite, with or without CMB lensing information, and its implications for a minimal cosmological scenario with eight parameters and an extended model with eleven parameters. We show that in many cases, the projected sensitivities from the Fisher matrix and the Monte Carlo methods differ significantly, particularly in models with many parameters. Sensitivities to the neutrino mass and the dark matter fraction are especially susceptible to change. ",
        "introduction": "In the last few years precision measurements of the cosmic microwave background (CMB), large scale structure (LSS) of galaxies, and distant type Ia supernovae (SNIa) have helped to establish a new standard model of cosmology. In this model, the geometry is flat so that $\\Omega_{\\rm total} = 1$, and the total energy density is made up of matter ($\\Omega_m \\sim 0.3$), and dark energy ($\\Omega_\\Lambda \\sim 0.7$, with equation of state $w \\equiv P/\\rho\\simeq -1$). With only a few free parameters this model provides an excellent fit to all current observations. Furthermore, each of these parameters is very tightly constrained by the observational data \\cite{Riess:1998cb,Perlmutter:1998np,Astier:2005qq,Spergel:2003cb,Spergel:2006hy,bib:sdss1}.  However, many other parameters can plausibly have a physical influence on the cosmological data, even if their presence has not yet been detected. Such parameters are, for example, a running spectral index of the primordial power spectrum, an equation of state for the dark energy which differs from $P=-\\rho$, and non-zero neutrino masses (e.g., \\cite{Hannestad:2005ey,Lesgourgues:2006nd}). Indeed, neutrinos are already known to have non-zero masses from oscillation experiments, where strong evidence points to a mass in excess of $\\sim 0.05$ eV for the heaviest mass eigenstate (see e.g. \\cite{Fogli:2006qg}). Given the sensitivities of future cosmological probes, even such a small mass will very likely be measured (e.g., \\cite{Hu:1997mj,Hannestad:2002cn,Kaplinghat:2003bh,Lesgourgues:2004ps,Lesgourgues:2005yv,% Hannestad:2006as,Wang:2005vr,Takada:2005si}). Thus, for future experiments such as the Planck satellite% \\footnote{ESA home page for the Planck project: {\\tt http://astro.estec.esa.nl/SA-general/Projects/Planck/} and Planck-HFI web site: {\\tt http://www.planck.fr/}} and beyond, it will be necessary to include the neutrino mass in the data analysis.  When performing a parameter error forecast for future experiments, it is customary to use the Fisher matrix formalism in which the formal error bar on a given parameter can be estimated from the derivatives of the observables with respect to the model parameters around the best fit point \\cite{Tegmark:1996bz,Eisenstein:1998hr}. However, this approach can only give a reasonable estimation when the likelihood of the cosmological parameters approximates a multivariate Gaussian function of the cosmological parameters.  In practice, significant departures from Gaussianity arise because of parameter degeneracies (in other words, because specific combinations of parameters are poorly constrained by the data), or because a parameter is defined in a finite interval, and its probability distribution does not fall to zero at one of the boundaries.  In this paper, we explore an alternative, simulation-based approach to derive projected sensitivities on the cosmological parameters for future experiments. This approach utilises synthetic data which emulate the expected instrumental and observational characteristics of the experiment under consideration. Using modern stochastic optimisation methods such as Markov Chain Monte Carlo or Importance Sampling \\cite{Lewis:2002ah}, the projected parameter errors can be extracted from the synthetic data set in the same way that parameters are extracted from existing data. Since this technique makes no assumption about the Gaussianity or otherwise of the parameter probabilities, we expect it to be much more reliable than the conventional Fisher matrix forecast.  This simulation-based forecast technique is in principle applicable to any one cosmological probe or combination of probes of interest, provided that one can reliably synthesise the data set given some underlying cosmological model and the instrumental noise and sensitivity. In this work, we focus on the example of the Planck satellite, to be launched in 2007 or 2008 by ESA for measuring with unprecedented sensitivitiy the CMB temperature and polarisation anisotropies on the full sky \\cite{:2006uk} .  The paper is structured as follows. In section \\ref{sec:CMB_data_model}, we summarise the statistical properties of the CMB data, with or without lensing information, and illustrate how synthetic data can be generated for some given fiducial model and instrumental characteristics. Sections \\ref{sec:Mock} and \\ref{sec:Fisher} outline, respectively, the simulation-based Monte Carlo Markov Chain method and the Fisher matrix formalism for forecasting cosmological parameter errors. Section \\ref{sec:minimal} contains a detailed comparison of the results obtained using  these two techniques for a minimal cosmological model with eight independent parameters.  The analysis is extended in section \\ref{sec:non-minimal} to a more complicated model with eleven  parameters, a case in which the differences between the Monte Carlo and Fisher matrix results are exacerbated.  We provide our conclusions in section\\ \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have studied error forecasts on cosmological parameters from the Planck satellite using two different methods. The first is the conventional Fisher matrix analysis in which the second derivative of the parameter likelihood function at the best fit point is used to calculate formal 1$\\sigma$ errors on the parameters, as well as the parameter correlation matrix.  The second is to use Markov Chain Monte Carlo methods such as CosmoMC on synthetic data sets.  This is the same method normally used to extract parameters from present data.  The Monte Carlo method has many advantages over the Fisher matrix approach.  While the Fisher matrix uses only information at the best fit point and assumes the likelihood function to be Gaussian with respect to the model parameters,  the use of Monte Carlo methods in conjunction with synthetic data maps out the true likelihood function for the given model realisation.  In this paper, we have shown that the likelihood function can be highly non-Gaussian, particularly with respect to the neutrino mass and the dark matter density, and, as a result, the CosmoMC analysis can give results that are very different from its Fisher matrix counterpart.  For prospective Planck data without lensing extraction and assuming a simple eight-parameter model, the difference in the projected errors can be as large as a factor of two or more for the said parameters. This indicates that in such cases the Fisher matrix method does not provide a reliable estimation. Including additional data such as CMB lensing information breaks some of the parameter degeneracies, and makes the likelihood more Gaussian.  The two methods thus become more compatible.  On the other hand, adding more cosmological parameters (all of which are physically motivated) leads to new parameter degeneracies, and generally worsens the agreement between the two forecast methods.  For the eleven-parameter model studied here, we find a difference for the neutrino mass fraction of 45\\% between the two methods at the 68\\% level, even with the inclusion of CMB lensing. The conclusion is that for some parameters, even with the very high precision of future data the likelihood function will not be sufficiently Gaussian to yield a reliable estimate of the precision with which the parameter can be measured using the Fisher matrix approach.  Given that Monte Carlo analysis of simulated data sets is computationally feasible with present computers, we propose that future error forecast analyses should employ this method rather than the Fisher matrix analysis. This will have the added advantage that the same parameter extraction pipeline can be used on real data as it becomes available. The present work only includes CMB data simulated to mimic Planck. However, the method can be easily generalised to include other data sets such as future weak lensing and baryon acoustic oscillation surveys."
    },
    "0606/astro-ph0606011_arXiv.txt": {
        "abstract": "Some gamma-ray burst (GRB) afterglow light curves show significant variability, which often includes episodes of rebrightening. Such temporal variability had been attributed in several cases to large fluctuations in the external density, or density \"bumps\". Here we carefully examine the effect of a sharp increase in the external density on the afterglow light curve by considering, for the first time, a full treatment of both the hydrodynamic evolution and the radiation in this scenario. To this end we develop a semi-analytic model for the light curve and carry out several elaborate numerical simulations using a one dimensional hydrodynamic code together with a synchrotron radiation code. Two spherically symmetric cases are explored in detail -- a density jump in a uniform external medium, and a wind termination shock. The effect of density clumps is also constrained. Contrary to previous works, we find that even a very sharp (modeled as a step function) and large (by a factor of $a \\gg 1$) increase in the external density does not produce sharp features in the light curve, and cannot account for significant temporal variability in GRB afterglows. For a wind termination shock, the light curve smoothly transitions between the asymptotic power laws over about one decade in time, and there is no rebrightening in the optical or X-rays that could serve as a clear observational signature. For a sharp jump in a uniform density profile we find that the maximal deviation $\\Delta \\alpha_{\\rm max}$ of the temporal decay index $\\alpha$ from its asymptotic value (at early and late times), is bounded (e.g, $\\Delta \\alpha_{\\rm max} < 0.4$ for $a = 10$); $\\Delta \\alpha_{\\rm max}$ slowly increases with $a$, converging to $\\Delta \\alpha_{\\rm max} \\approx 1$ at very large $a$ values. Therefore, no optical rebrightening is expected in this case as well. In the X-rays, while the asymptotic flux is unaffected by the density jump, the fluctuations in $\\alpha$ are found to be comparable to those in the optical. Finally, we discuss the implications of our results for the origin of the observed fluctuations in several GRB afterglows. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) are produced by a relativistic outflow from a compact source. The outflow sweeps up the external medium and drives a strong relativistic shock into it. Eventually the outflow is decelerated (by $pdV$ work across the contact discontinuity that separates the ejecta and the shocked external medium) and most of the kinetic energy is transferred to the shocked external medium (for recent reviews see \\citealt{Piran05,Meszaros06}). The shocked external medium produces long lived afterglow emission that is detected in the X-rays, optical, and radio for days, weeks, and months, respectively, after the GRB.  The afterglow emission is thought to be predominantly synchrotron radiation. This is supported both by the broad band spectrum and by the detection of linear polarization at the level of a few percent in the optical (or near infrared) afterglow of several GRBs \\citep[see][and references therein]{Covino03}. Inverse-Compton scattering of the synchrotron photons might dominate the observed flux in the X-rays in some cases \\citep{PK00,SE01,Harrison01}.  In the pre-{\\it Swift} era the best monitoring of GRB afterglow light curves was, by far, in the optical. Most afterglow light curves showed a smooth power law decay \\citep{Stanek99,LS03,Gorosabel06}, and often also a smooth achromatic transition to a steeper power law decay that is attributed to the outflow being collimated into a narrow jet \\citep{Rhoads99,SPH99}. It has been argued \\citep{WL00} that the smoothness of the afterglow light curve is directly related to (and thus enables to constrain) the smoothness of the external density field. Nevertheless, some optical afterglows have shown significant temporal variability, with strong deviations from the more typical smooth power law behavior. The best examples are GRBs 021004 \\citep{Pandey02,Fox03,Bersier03} and 030329 \\citep{Matheson03,Sato03,Uemura04,Lipkin04}.  Possible explanations for such temporal variability in GRB afterglow light curves include variations in the external density \\citep{WL00,Lazzati02,NPG03}, or in the energy of the afterglow shock. The latter includes energy injection by ``refreshed shocks'' -- slower shells of ejecta that catch up with the afterglow shock on long time scales \\citep{RM98,KP00a,SM00,R-RMR01,GNP03} or a ``patchy shell'' -- angular inhomogeneities within the outflow \\citep{KP00b,NPG03,HP03,NO04}. Another possible cause for variability in the afterglow light curve, although it is expected to be quite rare, is microlensing by an intervening star in a galaxy that happens to be close to our line of sight. GRB~000301C exhibited an achromatic bump in its optical to NIR light curves that peaked after $\\sim 4\\;$days \\citep{Sagar00,Berger00} which had been interpreted as such a microlensing event \\citep{GLS00,GGL01,BH05}, although other interpretations such as a bump in the external density have also been suggested \\citep{Berger00}.  Recent observations by {\\it Swift} have found flares in the early X-ray afterglows of many GRBs \\citep{Burrows05,Nousek06,Falcone06,Obrien06} which are probably due to late time activity of the central source \\citep{Nousek06,Zhang06}. Early optical variability also appears to be more common than previously thought \\citep[e.g.,][]{Stanek06}, although it is not yet clear if it is caused by similar mechanisms as the late time optical variability that had been detected before {\\it Swift}.  The most natural forms of variations in the external density are either clumps on top of a smooth background density distribution, or a global abrupt change in density with radius. The latter can be, e.g., the termination shock of the wind from the massive star progenitor of a long-soft GRB \\citep{Wijers01}. Such a stellar wind environment may have a richer structure and can include an abrupt increase in density with radius at the contact discontinuity between shocked wind from two different evolutionary stages of the progenitor star, as well as clumps that are formed due to Rayleigh-Taylor instability \\citep{R-R05,EGDM06}\\footnote{ \\citet{R-R05} find that the clump formation may also involve the Vishniac instability, and once such clumps are formed they stand a reasonable chance to survive until the time of the core collapse of the progenitor star.}. Density clumps with a mild density contrast may also be formed due to turbulence in the external medium. Furthermore, the external density profile is expected to vary between different progenitor models \\citep{FRY06}.  The variability that had been observed in optical afterglows was attributed to density clumps in the external medium \\citep{Lazzati02,NPG03}. The expected observational signatures of the afterglow shock running into the wind termination shock of the massive star progenitor have also been considered \\citep{Wijers01,R-R01,R-R05,EGDM06,PW06}. In all these cases it has been argued that there would be a clear observational signature in the form of a rebrightening in the afterglow light curve, before approaching the new shallower decay slope corresponding to the uniform density of the shocked wind.  Here we revisit the effect of density fluctuations on the afterglow light curve by solving in detail the case of a spherically symmetric external density with a single density jump (by a factor of $a > 1$) at some radius $R_0$, while the density at smaller and larger radii is a (generally different) power law in radius. This is done by constructing a semi-analytic model for the observed flux due to synchrotron emission at different power law segments of the spectrum and by carrying out numerical simulations. The semi-analytic model takes into account the effect of the reverse shock on the hydrodynamics and on the emissivity, as well as the effect of the spherical geometry on the arrival time of photons to the observer. Being semi-analytic, however, this model uses some approximations for the hydrodynamic evolution and the resulting radiation. Therefore, we also perform numerical simulations in which the light curves are calculated using a hydrodynamic+radiation numerical code. This code self-consistently calculates the radiation and evolves the electron distribution in every fluid element, which corresponds to a computational cell of the one dimensional Lagrangian hydrodynamic code. The results of this code are used to obtain light curves in cases of special interest and near spectral break frequencies, as well as to verify the quality of the semi-analytic model, which is found to agree well with the numerical results. These calculations are much more accurate than those presented in previous works, and our results are significantly different. In all cases we find a very smooth transition to the new asymptotic power law, with no rebrightening for an initially decaying light curve.  In \\S \\ref{sec:spherical} we develop the semi-analytic model. First, in \\S \\ref{sec:hydro} a simple analytic model is constructed for the hydrodynamics, which agrees very well with our numerical results. Then, in \\S \\ref{sec:rad} we construct a semi-analytic model for the observed flux density. Specific case studies (\\S \\ref{sec: case study}) are then analyzed in detail, for a wind termination shock (\\S \\ref{sec:wind}) and for a spherical density jump in a uniform medium (\\S \\ref{sec:uniform}). The light curves for these cases are also calculated using the numerical code (described in Appendix \\ref{sec:code}). The effect of proximity to a break frequency around the time of the density jump is investigated in \\S \\ref{sec:nu_break}, and our main conclusions are found to remain valid also in the vicinity of the break frequencies. \\S \\ref{sec:clump} is devoted to a discussion of the expected observational signatures of density clumps on top of a smooth underlying external density distribution. Such density clumps are found to have a very weak observational signature which would be very hard to detect. In \\S \\ref{sec:GRBs} we discuss the implications of our results for the origin of the observed fluctuations in several GRB afterglows. Our conclusions are discussed in \\S \\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis}  We have presented a semi-analytic model for the light-curve resulting from synchrotron emission by a spherical relativistic blast-wave that propagates into a power law external density profile with a single sharp density jump at some radius $R_0$. Our solution is general enough to include a transition in the density power-law index ($k$) at $R_0$, but is limited to cases in which the blast-wave remains relativistic after it encounters the density jump. This model has been used to explore in detail two density profiles that are most relevant to GRB afterglows: a wind termination shock, and a sharp density jump between two regions of uniform density. The latter results are also used to constrain the signature of density clumps in a uniform medium. We have also carried out detailed numerical simulations for several of the cases which we have studied in detail. These numerical results serve three purposes. First, they are used in order to obtain more accurate light curves for the cases which are of special interest.  Second, they are appropriate for calculating the light curves in the vicinity of a spectral break frequency (\\S \\ref{sec:nu_break}).  Third, they serve as a test for the quality of our simple semi-analytic model, which is found to give a very good qualitative description and reasonable quantitative description of the light curve.  Our main result is that density jumps do not produce sharp features in the light curve, regardless of their density contrast! The results of our specific case studies are as follows.\\vspace{0.1 in}\\linebreak A wind termination shock: \\begin{itemize} \\item The light curve shows a smooth transition, which lasts for about one decade in time, between the asymptotic power-law behavior at $T < T_0$ and at $T \\gg T_0$.  \\item There is no rebrightening or any other sharp feature that can be used as a clear observational signature.  \\item Above the cooling frequency, $\\nu_c$, there is no change in the asymptotic value of the temporal index, $\\alpha$, and it only fluctuates with a small amplitude ($\\Delta \\alpha \\approx 0.1$).  \\item The only observable signatures of a wind termination shock are a smooth break, with an increase of $\\Delta \\alpha = 0.5$ in $\\alpha$, below $\\nu_c$ and a transition in the temporal evolution of $\\nu_c$.   \\end{itemize} A density jump between two uniform density regions: \\begin{itemize} \\item The light curve shows a smooth transition between the two asymptotic power laws (at $T < T_0$ and $T \\gg T_0$).  \\item The transition time increases with the density contrast, $a$, and is about $\\sim 10aT_0$.  \\item The maximal deviation, $\\Delta \\alpha_{\\rm max}$, of the temporal index $\\alpha$ from its asymptotic value (at $T < T_0$ and $T \\gg T_0$) is small. For example, $\\Delta \\alpha_{\\rm max}(a=10) < 0.4$; $\\Delta \\alpha_{\\rm max}$ depends weakly on $a$ and approaches $\\approx 1$ at very large $a$ values.  Therefore, a density jump cannot produce an optical rebrightening when $\\nu_{\\rm optical} > \\nu_m$.  \\item The light curve fluctuates also above $\\nu_c$ (typically including the X-ray band). While the asymptotic flux (at $T \\gg T_0$) above $\\nu_c$ is unaffected by the density jump, the fluctuations in $\\alpha$ are {\\it comparable} to those below $\\nu_c$. \\end{itemize} An overdense clump on top of a uniform density background: \\begin{itemize} \\item Only a fairly large clump ($l_{\\rm cl} \\gtrsim R/\\gamma$) with a sufficiently large density contrast ($a \\gtrsim 10-10^2$) produces a significant fluctuation in the light curve.  \\item The effect of a clump on the light curve is significantly larger when it is located along the line of sight, than at an angle of $\\sim 1/\\gamma$ from the line of sight.  \\item The signature of a clump is most apparent at $\\nu < \\nu_m < \\nu_c$.  \\item Above $\\nu_m$ a clump can actually cause a small dip in the light curve, while below $\\nu_m$ it causes a (larger) bump.  \\end{itemize} For a spherical density jump our conclusions are based on accurate results, while in the case of a density clump we obtain only an approximate upper limit for its effect on the light-curve.  Therefore, our results for density clumps should be taken only as rough guidelines. Our main results remain valid also when the observed frequency is close to a spectral break frequency around the time of the density jump.  Our conclusions are very different from those of previous works, which predicted a significant optical rebrightening, and rather sharp features in the afterglow light curve. The main cause for this difference is our careful consideration of both the effect of the reverse shock on the dynamics (which we find cannot be neglected even when $a \\sim 2$), and the arrival time of the photons to the observer from different parts of the emitting regions. Both of these effects tend to smoothen the light curve significantly.  Finally, we considered the implications of our results for the origin of the fluctuations in the highly variable light curves of four GRBs (3 long-soft GRBs and one short-hard GRB). We find that density variations are unlikely to be the source of the fluctuations in any of these bursts.    We thank Enrico Ramirez-Ruiz, Avishay Gal-Yam, Eran Ofek, Brad Cenko and Pawan Kumar for useful comments. This research was supported by a senior research fellowship from the Sherman Fairchild Foundation (E. N.) and by US Department of Energy under contract number DE-AC03-76SF00515 (J. G.)."
    },
    "0606/gr-qc0606056_arXiv.txt": {
        "abstract": "We study cosmological perturbations in the brane models with an induced Einstein-Hilbert term on a brane. We consider an inflaton confined to a de Sitter brane in a five-dimensional Minkowski spacetime. Inflaton fluctuations excite Kaluza-Klein modes of bulk metric perturbations with mass $m^2 = -2(2\\ell-1) (\\ell +1) H^2$ and $m^2 = -2\\ell(2\\ell+3) H^2$ where $\\ell$ is an integer. There are two branches ($\\pm$ branches) of solutions for the background spacetime. In the $+$ branch, which includes the self-accelerating universe, a resonance appears for a mode with $m^2 = 2 H^2$ due to a spin-0 perturbation with $m^2 = 2H^2$. The self-accelerating universe has a distinct feature because there is also a helicity-0 mode of spin-2 perturbations with $m^2 = 2H^2$. In the $-$ branch, which can be thought as the Randall-Sundrum type brane-world with the high energy quantum corrections, there is no resonance. At high energies, we analytically confirm that four-dimensional Einstein gravity is recovered, which is related to the disappearance of van Dam-Veltman-Zakharov discontinuity in de Sitter spacetime. On sufficiently small scales, we confirm that the lineariaed gravity on the brane is well described by the Brans-Dicke theory with $\\omega=3Hr_c$ in $-$ branch and $\\omega = -3H r_c$ in $+$ branch, respectively, which confirms the existence of the ghost in $+$ branch. We also study large scale perturbations. In $+$ branch, the resonance induces a non-trivial anisotropic stress on the brane via the projection of Weyl tensor in the bulk, but no instability is shown to exist on the brane. ",
        "introduction": "There has been tremendous interest over the last several years in the brane-world scenario where we are assumed to be living on a four-dimensional hypersurface (brane) in a higher-dimensional spacetime (bulk) \\cite{Roy_review}. The simplest example of this scenario is proposed by Randall and Sundrum where there is a brane embedded in a five-dimensional anti-de Sitter (AdS) spacetime \\cite{RS99}. They showed that the four-dimensional gravity is recovered at low energies, because gravity is confined to a single positive-tension brane even if the extra dimension is not compact. Therefore, in order to see the deviation from the conventional four-dimensional gravity, it is necessary to investigate the phenomena occurred at high energies.  One of the most promising tools to extract the information of the extra-dimension is the primordial density fluctuations generated in the period of inflation in the early universe. In the inflation model where inflation is driven by an inflaton field confined to the brane, the amplitude of the curvature perturbation is calculated in the extremely slow-roll limit where the coupling between inflaton field fluctuation and bulk perturbations can be neglected \\cite{MWBH}. This work is proceeded to go beyond the zero-th order slow-roll approximation by solving the bulk metric perturbations classically \\cite{KLMW,KMW}. It is shown that we could have significant effects from the back-reaction due to the coupling to five-dimensional gravitational perturbations. On the other hand, it was pointed out that the localized matter on a brane can induce gravity on the brane via quantum loop corrections at high energies \\cite{induced_gravity}. This induced gravity can act as the ultra-violet cut-off for the inflaton perturbations.  Based on the induced gravity scenario, Dvali, Gabadadze, and Porrati (DGP) proposed a brane-world model with induced gravity \\cite{DGP} in which the four-dimensional brane is embedded in a five-dimensional Minkowski spacetime. (For a review of the phenomenology of DGP model, see \\cite{Lue_review}.) In the presence of the induced gravity, according to the embedding of the brane in the bulk, there appear two branches of background solutions. In $+$ branch, the size of the four-dimensional hypersurface becomes minimum at the brane while, in $-$ branch, it becomes maximum at the brane. Inflation models in this scenario were studied in \\cite{inducegw, Papan, Zhang}.  In this paper, we study the behavior of five-dimensional metric perturbations excited by the inflaton perturbations confined to the brane in DGP model. The solutions in $-$ branch are regarded as the high energy limit of the Randall-Sundrum model including the curvature corrections on a brane, depending on the parameters. Then our analysis can be used to discuss the generation of primordial fluctuations in these models. The solutions in $+$ branch are largely used to explain the present cosmic acceleration without introducing dark energy in the DGP model \\cite{DGP_FRW}, but it turns out that they are plagued by the ghost instability \\cite{Nicolis_Rattazzi,Luty_Porrati_Rattazzi,Koyama_ghost,Gorbunov_Koyama_Sibiryakov, Charmousis}. However, the study of the gravitational property of this model in the presence of matter and cosmic expansion is still very limited \\cite{Lue_Starkman, KM, Kaloper} so it is also worth considering the solutions in this branch.  The structure of the rest of the paper is as follows. In section II,  we derive the basic equations for the five-dimensional metric perturbation and the boundary condition imposed at the brane and then summarize the solution for the vacuum brane obtained in \\cite{Koyama_ghost}. In section III, we provide the solutions for curvature perturbations on the brane in the presence of the scalar field on the brane. Next we discuss the appearance and disappearance of van Dam-Veltman-Zakharov discontinuity in section IV. In section V we show the comparison of the behavior of perturbations on small scales with the Brans-Dicke theory. The solutions for perturbations on large scales are discussed in section VI. Section VII is devoted to conclusions. Technical details of the calculation of the $m^2 = 2H^2$ mode and the brief summary of the cosmological perturbation in four-dimensional Brans-Dicke gravity are presented in the Appendix A and B, respectively. ",
        "conclusions": "In this paper, we studied inflaton perturbations confined to a de Sitter brane with induced gravity in a five-dimensional Minkowski spacetime.  For a vacuum brane, the spin-0 mode appears as a discrete bulk mode with $m^2 = 2H^2$ in the $+$ branch, while in the $-$ branch there are no normalizable solutions for the spin-0 modes. Since there is another discrete bulk mode (helicity-0 mode of spin-2 perturbation) with mass $m^2 = H^2 (3 Hr_c -1) (Hr_c)^{-2}$ in the $+$ branch, there is a resonance between the spin-0 mode and the helicity-0 mode of spin-2 perturbation for $Hr_c =1$. In this paper we introduced inflaton perturbations on a brane. Then an infinite ladder of discrete modes with $m^2 = -2(2\\ell-1)(\\ell+1)H^2$ and $m^2 = -2\\ell (2\\ell +3)H^2$ are excited. Since there is a mode with $m^2 = 2H^2$ regardless of the value of $Hr_c$, in the $+$ branch the resonance inevitably appears. We obtained the solutions for the curvature perturbation on the brane ${\\cal{R}}_{(b)}$ and studied their behavior.  At high energies, $H r_c \\to \\infty$, we have confirmed that, in both branches, the four-dimensional general relativity solutions are recovered from the infinite sum of the modes. This results remind us that in the case of massive gravity, if we introduce a cosmological constant $\\Lambda_{(4)}$, the vDVZ discontinuity disappears if the limits $m \\to 0$ is taken for fixed $\\Lambda_{(4)}$ \\cite{Porrati:2000cp, Higuchi,Kogan:2000uy}. However, this argument holds only for $Hr_c \\to \\infty$. For a finite $Hr_c$, if we do $not$ fix $-k\\eta$ and take large enough $-k\\eta$, we can no longer perform the summation of the infinite ladder of the modes and the theory deviate from four-dimensional general relativity. In fact, on small scales, for a fixed $Hr_c$ we can see that the junction condition becomes the same as the gravitational field equations in Brans-Dicke theory where the brane bending mode acts as the BD scalar.  Then, in order to identify the BD parameter for the effective BD theory, we compared the solutions for Brans-Dicke scalar field and the brane bending mode. We have confirmed that, on sufficiently small scales, gravity is well described by the Brans-Dicke theory with $\\omega = -3Hr_c$ in $+$ branch and $\\omega = 3Hr_c$ in $-$ branch. This is consistent with the previous work \\cite{Lue_Starkman, KM} and the existence of the ghost in $+$ branch is confirmed.  We also studied the large scales perturbations. In $-$ branch, the solutions agree with four-dimensional GR with a modified gravitational constant. In $+$ branch, there is an additional contribution to ${\\cal R}_{(b)}$ due to the resonance, which diverges on large scales for $Hr_c=1$. We identified this contribution as the effect of the Weyl anisotropic stress and showed that the anisotropic stress itself does not diverge. In fact, we can find a suitable gauge where all metric perturbations remain small, so the resonance does not lead to a gravitational instability on the brane.  We make comments on future applications of our results. In $-$ branch, if $r_c \\ll H^{-1}$, small scale perturbations $a/k \\ll r_c$ can be described by the four-dimensional Brans-Dicke theory.  If the perturbations approaches to $r_c$, the gravity becomes five-dimensional and we expect significant effects from the coupling to five-dimensional metric perturbations \\cite{KMW}. Unlike the Randall-Sundrum model where the small scales perturbations are always coupled to five-dimensional perturbations and a quantum vacuum state is hard to be specified \\cite{KMW, K_Mennim_W, YK}, we can specify a vacuum state without ambiguity based on the four-dimensional BD theory on sufficiently small scales. Then we can estimate the effect of the coupling to five-dimensional gravity without ambiguity.  Finally we should emphasize that our analysis is limited to linear perturbations. It has been shown that we need to take into account the non-linearity of the brane bending modes before the gravity becomes non-linear. Actually, we can recover four-dimensional general relativity due to this non-linearity \\cite{Deffayet_vDVZdis,Lue_string,Gruzinov:2001hp,Giannakis:2001jg, Porrati:2002cp,Park:2003fc,Tanaka_DGP}. Ref.~\\cite{Lue_Starkman} suggested that this non-linear scale becomes very small at high energies by the factor $1 - \\epsilon 2 Hr_c$, which is the same factor seen in Eq.~(\\ref{bendcurve}). The non-linear interactions of the brane bending mode in de Sitter spacetime deserve further study."
    },
    "0606/astro-ph0606425.txt": {
        "abstract": "The spectra obtained above 100 MeV by the EGRET experiment aboard the Compton Gamma-Ray Observatory for a handful of gamma-ray bursts has given no indication of any spectral attenuation that might preclude detection of bursts at higher energies.  With the discovery of optical afterglows and counterparts to bursts in the last few years, enabling the determination of significant redshifts for these sources, it is anticipated that profound spectral attenuation will arise in the Gamma-Ray Large Area Space Telescope (GLAST) energy band of 30 MeV--300 GeV for many if not most bursts. This paper explores time-dependent expectations for burst spectral properties in the EGRET/GLAST band, focusing on how attenuation of photons by pair creation internal to the source generates distinctive spectral signatures.  The energy of spectral breaks and the associated spectral indices provide valuable information that constrains the bulk Lorentz factor of the GRB outflow at a given time.  Moreover, the distinct temporal behavior that is present for internal attenuation is easily distinguished from extrinsic absorption due to intervening cosmic background fields.  These characteristics define palpable observational goals for both spaced-based hard gamma-ray experiments such as GLAST, and ground-based \\v{C}erenkov telescopes, and strongly impact the observability of bursts above 300 MeV. ",
        "introduction": "\\label{sec:intro}  High energy gamma-rays have been observed for six gamma-ray bursts by the EGRET experiment on the Compton Gamma-Ray Observatory (CGRO). Most conspicuous among these observations is the emission of an 18 GeV photon by the GRB940217 burst (Hurley, et al. 1994).  An additional case of interest is provided by the so-called MILAGRITO burst (Atkins et al. 2000; Atkins et al. 2003), GRB 970417, with its uncorroborated detection of \\teq{\\sim 3\\sigma} significance in the TeV band. Taking into account EGRET's field of view, its detections indicate that emission in the 1 MeV--10 GeV range is probably common among bursts, if not universal (see Dingus 1995 for a discussion of EGRET high energy burst statistics). One implication of GRB observability at energies around or above 1 MeV is that, at these energies, spectral attenuation by two-photon pair production (\\teq{\\gamma\\gamma\\to e^+e^-}) is absent in the source.  From this fact, Schmidt (1978) deduced that a typical burst had to be closer than a few kpc, if it produced quasi-isotropic radiation.  In the aftermath of BATSE's revelation (e.g. Meegan et al. 1996) that most if not all long bursts are at cosmological distances, Krolik \\& Pier (1991) and Fenimore et al. (1992) proposed that GRB photon angular distributions are highly beamed, being produced by a relativistically-moving plasma, a suggestion that has become an underpinning of the GRB paradigm.  This can dramatically reduce \\teq{\\taupp} below the \\teq{\\sim 10^{11}-10^{12}} values realized for isotropic photons, and blueshift spectral attenuation turnovers above the observed spectral range.  Determinations of the bulk Lorentz factor \\teq{\\Gamma} of the GRB medium have mostly concentrated (e.g. Epstein 1985; Krolik \\& Pier 1991; Baring 1993) on cases where the angular extent of the source was of the order of \\teq{1/\\Gamma}.  These calculations generally assume an infinite power-law burst spectrum, and deduce (e.g. Woods \\& Loeb 1995; Baring \\& Harding 1997b) that gamma-ray transparency up to the maximum energy detected by EGRET requires \\teq{\\Gamma\\gtrsim 100}--\\teq{10^3} for cosmological bursts.  Since BeppoSax spawned the age of precise determination of GRB redshifts, such bounds can now be refined for the BeppoSax, HETE and Swift databases; the discovery of high redshift bursts (e.g. see Cummings et al. 2005; Kawai et al. 2005; Berger et al. 2006; Haislip et al. 2006, for Swift bursts) might suggest even higher estimates of bulk Lorentz factors in those sources.  While the power-law source spectrum assumption is expedient, the spectral curvature seen in virtually all GRBs by BATSE (Band et al. 1993; Preece et al. 2000) is expected to play an important role in reducing the opacity for super-GeV and TeV band emission from these sources (Baring \\& Harding 1997a). Such curvature is patently evident in 200 keV--2 MeV spectra of some EGRET-detected bursts (e.g. Schaefer, et al. 1992), and its prevalence is indicated by the generally steep EGRET spectra for bursts (e.g. Hurley, et al. 1994; Schneid, et al. 1992; Sommer et al. 1994). This curvature profoundly impacts \\teq{\\gamma\\gamma} opacity determinations, yielding a huge dearth of target soft photons in comparison with power-law spectra extrapolated down to the classic X-ray band. Hence, when considering emergent photons of dimensionless energies \\teq{\\erg_{\\gamma}\\gtrsim \\Gamma^2/\\eb}, where \\teq{\\eb} is the dimensionless energy (in units of \\teq{m_ec^2}) of the spectral break in the BATSE band, the target photon energy is below \\teq{\\eb m_ec^2} so that accurate pair opacity calculations mandate detailed treatment of the spectral curvature observed in bursts.  This paper enunciates the principal properties of pair production opacity that couple to spectral shape in the BATSE/EGRET energy range, embellishing upon and focusing the work of Baring \\& Harding (1997a), so as to identify possible observational diagnostics for the Gamma-Ray Large Area Space Telescope (GLAST; {\\tt http://glast.gsfc.nasa.gov/}) mission and ground-based experiments such as the MILAGRO, MAGIC, HESS and VERITAS atmospheric \\v{C}erenkov telescopes. Moreover, it helps define GRB science drivers for future initiatives in the hard gamma-ray band above 100 MeV. These spectral signatures at energies where significant opacity is realized are clearly distinguishable from absorption by background radiation fields, and in particular via their time-dependent evolution. The fact that turnovers produced by constant Lorentz factor evolution can be easily discriminated from those generated by decelerating expansions is compelling.  It may also be possible to elicit details of the adiabaticity or otherwise of the post-fireball expansion. Accordingly, hard \\teq{\\gamma}-ray telescopes can, in principal, provide powerful probes on such evolution of bulk motions in bright bursts, which can then be used to generate useful bounds on the explosion energy, a key piece of information elucidating the nature of the central engine for bursts.  Before discussing these evolutionary effects and the potential for observational diagnostics in Sections~\\ref{sec:evolution_effect} and~\\ref{sec:diagnostics}, the paper presents refinements in the formalism for pair production opacities in Section~\\ref{sec:pairprod}. ",
        "conclusions": "\\label{sec:conclusion}  This paper outlines the interesting potential the GLAST mission and future hard gamma-ray detectors can offer for probing the relativistic nature of outflow in gamma-ray burst prompt emission regions. The diagnostics highlighted here use a single, straightforward piece of physics, namely that of electron-positron pair creation, and are only dependent on model assumptions to second order.   Spectral attenuation by this process, while not evident in EGRET data for a handful of bright, hard bursts, should emerge from the GLAST database, unless the bulk Lorentz factors in the burst emission zones are higher than presently argued, i.e. greater than around \\teq{\\Gamma\\sim 10^3}.  Then, attenuation by interaction of source photons with those of the intervening cosmic background field will produce attenuation turnovers at above 30--50 GeV, whose energies are independent of the GRB flux at a given time.  In such cases, bursts can be used to explore the evolution of the background field with redshift by detectors with sufficient sensitivity in this band.  For the equally interesting possibility of lower \\teq{\\Gamma}, pair creation turnovers spawned by photons internal to the burst, will become prominent in the GLAST energy window.  In such cases, the representative calculations in Section~\\ref{sec:evolution_effect} clearly illustrate how the LAT instrument on GLAST might potentially discriminate between evolutionary scenarios with constant \\teq{\\Gamma}, and those where \\teq{\\Gamma} declines with time, should Nature be agreeable and offer bright, hard bursts with turnovers below 300 MeV--1 GeV.  A more sensitive, next-generation hard gamma-ray experiment might additionally be capable of distinguishing between adiabatic and radiative evolutionary scenarios: due to their different \\teq{\\Gamma (t)} dependence, their turnover energy/flux loci trace somewhat different tracks over time.  However, this probe will be more reliable for flat spectrum bursts and require relatively sharp turnovers, or correlated information from the MeV-band break evolution. While the GBM on GLAST, in conjunction with the LAT, while accumulate a host of data on the MeV-band break evolution, next-generation telescopes most likely will be needed to differentiate the adiabatic and radiative turnover tracks. Should distinction between these two keystone classes of blast wave deceleration be possible, the pair creation probes offer constraints on the explosion energy generated by the central engine.  Note that while the results presented focus on simple broken power-law source spectra, they are readily extended to apply to multiple components.  They encompass a variety of possibilities, and provide a template for burst pair production attenuation studies in the soon-to-be realized GLAST era, and for future initiatives down the line."
    },
    "0606/astro-ph0606541_arXiv.txt": {
        "abstract": "We provide classifications for all 143 million non-repeat photometric objects in the Third Data Release of the Sloan Digital Sky Survey (SDSS) using decision trees trained on 477,068 objects with SDSS spectroscopic data. We demonstrate that these star/galaxy classifications are expected to be reliable for approximately 22 million objects with $r \\lesssim 20$. The general machine learning environment Data-to-Knowledge and supercomputing resources enabled extensive investigation of the decision tree parameter space. This work presents the first public release of objects classified in this way for an  entire SDSS data release. The objects are classified as either \\textit{galaxy}, \\textit{star} or \\textit{nsng} (neither star nor galaxy), with an associated probability for each class. To demonstrate how to effectively make use of these classifications, we perform several important tests. First, we detail selection criteria within the probability space defined by the three classes to extract samples of stars and galaxies to a given completeness and efficiency. Second, we investigate the efficacy of the classifications and the effect of extrapolating from the spectroscopic regime by performing blind tests on objects in the SDSS, 2dF Galaxy Redshift and 2dF QSO Redshift (2QZ) surveys for which spectra are available. We find that, for a sample giving a maximal combined completeness and efficiency, the completeness values are 98.9\\% and 91.6\\% and the corresponding efficiencies are 98.3\\% and 93.5\\% for galaxies and stars, respectively,  in the SDSS.  We also test our star-galaxy classification by studying the inverse of this sample, finding quasars in the nsng sample with a completeness and efficiency of 88.5\\% and 94.5\\% from the SDSS and 94.7\\% and 87.4\\% from the 2QZ. Given the photometric limits of our spectroscopic training data, we effectively begin to extrapolate past our star-galaxy training set at $r \\sim 18$. By comparing the number counts of our training sample with the classified sources, however, we find that our efficiencies appear to remain robust to $r \\sim 20$.  As a result, we expect our classifications to be accurate for 900,000 galaxies and 6.7 million stars, and remain robust via extrapolation for a total of 8.0 million galaxies and 13.9 million stars. The latter sample should prove useful in constructing follow-up spectroscopic surveys. Characterizing the success of our classifications fainter than $r \\sim 20$ will require fainter spectroscopic data. ",
        "introduction": "\\label{Sec: Intro}  It is well-known that classification of objects in astronomical survey data provides a fundamental characterization of the dataset, and forms a vital step in understanding the ensemble properties of the dataset and the properties of individual objects. In this paper, we apply automated machine learning algorithms to classify the 143 million non-repeat objects in the third data release \\citep[DR3;][]{abazajian:dr3} of the Sloan Digital Sky Survey (SDSS\\footnote{\\url{http://www.sdss.org}}; York et al. 2000). The SDSS is used because the combination of the quality and accuracy of the survey CCD photometry, the approximately concurrent spectroscopy performed with the same telescope, and the large number of objects, is ideally suited to the methods we employ in this work.  As our primary goal is to perform reliable star-galaxy separation, we have used three classes: \\textit{galaxy}, \\textit{star}, and \\textit{nsng} (neither star nor galaxy).  The use of this third class is novel, and is included to not only cleanly separate out those objects that clearly are neither stars nor galaxies (e.g., quasars), but to also improve the extrapolation of our classifications well beyond the  nominal magnitude limits of our current training sample. The latter benefit arises since star-galaxy separation becomes less accurate at fainter magnitudes due to source confusion. In this case, our algorithm can assign sources to the third category when a clear star-galaxy classification is not feasible, thereby minimizing contamination for lower signal-to-noise sources.  In this paper, the learning algorithm we employ is the axis-parallel decision tree. A decision tree is trained on the subset of data for which spectra are available and is subsequently used to classify the much larger sample for which there is just photometry. This is the first public release of objects classified using this method for the entire SDSS DR3. A natural question is ``Why provide classification for all objects, most of which are near the photometric limits of the survey, where classifications are less robust?'' We have three reasons for classifying the entire SDSS DR3. First, given the speed of applying our classifications, we have found that it is considerably easier to handle large datasets when no artificial cuts are used. With the growth in the number of large astronomical surveys that are accessible from a virtual observatory, we feel this point is becoming increasingly more important, as it greatly simplifies any data federation issues a fellow scientist may need to perform if they wish to use our classifications. Second, classifying all objects allows others with fainter training data to verify the efficacy of our classifications. Finally, since our classifications are probabilistic in nature, by providing the classifications for all objects, we allow others to make their own probability cuts, perhaps using combinations of classes, to perform their own science. For example, our approach simplifies the task of performing follow-up spectroscopy, as a user of our fully classified catalog has an easier task in characterizing any selection effects.  Decision tree learning algorithms have been in existence since the 1970s, see, e.g., \\citet{brei:dt} or \\citet{quinlan:dt}. They are able to classify objects into discrete or continuous categories, scale well to large datasets, are fairly robust to inequitably-sampled training sets, and can cope with values from individual objects that happen to be bad or irrelevant for the training targets being considered. Axis-parallel decision trees, particularly smaller ones, are also easily understood, providing clear criteria for the splitting of the data set into constituent classes or values. This is in contrast to, for example, artificial neural networks \\citep[e.g.,][]{ball:ann}, where the resulting weights are basically a \\textit{black box}, or oblique decision trees \\citep[e.g.,][]{murthy:odt}, where linear combinations of parameters, or hyper-planes, are used to subdivide the data. Given the size of our dataset, however, our trees are generally rather large, making it difficult to fully interpret the classification rules, but one can in principle follow the splitting of the population. While oblique trees may be less transparant in their operation, they do have the benefit of generally producing smaller trees. Most astronomical results using neural networks and decision trees show that the two methods give comparable results---an example is \\citet{bazell:ensembles}, described below.  A number of previous efforts have used decision trees to classify objects in astronomical surveys. \\citet{salzberg:dt} use decision trees to separate stars and cosmic ray hits in Hubble Space telescope images and are able to do so with an accuracy of over 95\\%. Their algorithm, known as OC1, is an oblique decision tree. \\citet{owens:dt} use a version of OC1 to classify galaxy morphologies in the European Southern Observatory--Uppsala surface photometry catalog \\citep{lauberts:esouppsala} and show their results to be comparable to those of \\citet{storrielombardi:ann}, who use artificial neural networks on the same data.  \\citet{white:first} construct a radio-selected sample of quasars using the VLA-FIRST survey \\citep{becker:first} and the optical APM-POSS-I quasar survey of \\citet{mcmahon:apmqso}. They show that using decision trees improves on simpler methods and enables quasars to be selected with 85\\% efficiency. They also show the completeness and efficiency of the selection as a function of threshold probability from the decision tree that the object is a quasar. \\citet{bazell:ensembles} compare Naive Bayes, decision tree and neural network classifiers for morphological galaxy classification and show that the latter two are comparable.  \\citet{suchkov:dt} take advantage of the much improved data, both in quality and quantity of objects, available from the SDSS. They apply the oblique decision tree classifier ClassX, based on OC1, to the SDSS DR2 \\citep{abazajian:dr2} to classify objects into star, red star (type M or later), AGN photometric redshift and galaxy photometric redshift bins. Three representative samples, each of $10^5$ photometric objects, one of which extends 2 magnitudes fainter than the spectroscopy, are also investigated. They give a table of percentage correct classifications from which one can calculate completeness and efficiency.  Many studies, too numerous to list here, use other machine learning techniques to classify astronomical survey objects, including star-galaxy separation and quasar identification. These previous works, however, do not provide spectroscopically-trained classifications with probabilities on the scale of the 143 million sources given here for the SDSS DR3. As part of our work, we make publicly available\\footnote{\\url{http://quasar.astro.uiuc.edu/rml}} the full set of our 143 million object classifications for the SDSS DR3 and probability cuts to select galaxies and stars for a range of completenesses and efficiencies.  The rest of the paper is as follows. In \\S \\ref{Sec: Data} we describe the datasets used in the construction and testing of the catalog. \\S \\ref{Sec: Algorithms} describes the decision trees, their optimization using supercomputing resources, and the probability cuts to create subsamples of particular objects. \\S \\ref{Sec: Results} then describes the optimal learning parameters found and the resulting datasets. In \\S \\ref{Sec: Discussion} we discuss the efficacy of our star-galaxy separation, the extrapolation of our results to fainter magnitude, and new projects we have undertaken to extend this work. We conclude in \\S \\ref{Sec: Conclusions}. ",
        "conclusions": "\\label{Sec: Conclusions}  In this paper we have classified all of the 142,705,734 non-repeat (primary) objects in the Third Data Release of the Sloan Digital Sky Survey (SDSS DR3). The classifications were determined by using machine learning, the algorithm of choice being the axis-parallel decision tree. The algorithm was trained on 477,068 objects for which spectroscopic classifications were available within DR3. This training data consists of 361,490 galaxies, 62,333 stars, 49,545 quasars, and 3,700 unknown objects. Collaboration with domain experts at the National Center for Supercomputing Applications (NCSA) and the use of the general machine learning environment Data-to-Knowledge combined with supercomputing resources enabled extensive investigation of the decision tree parameter space and the associated datasets. To our knowledge, this level of investigation has not previously been published in the astronomy literature.  This is the first public release of objects classified in this way for the whole survey. The objects are classified as either galaxy, star or nsng, and we provide an associated probability for each class. These three probabilities always sum to unity for every object classified. By merely assigning the classification with the highest probability, we find that our full classification sample contains 38,022,597 galaxies, 57,369,754 stars and 47,026,955 nsng in the magnitude range $0 < r < 40$, with 589 outside this range and 285,839 (0.2\\%) ambiguous due to equal-highest probabilities for object type. The catalog is available for download at \\url{http://quasar.astro.uiuc.edu/rml}.  A major issue with this method of classification is the inevitable extrapolation from the spectroscopic regime. We investigate this by examining the apparent magnitude number counts for all sources and by performing several pseudo-blind or fully blind tests by using sources spectroscopically identified in the SDSS, 2dF Galaxy Redshift, and 2dF QSO Redshift (2QZ) surveys. We find that for a sample giving the optimal completeness and efficiency, the completeness values are 98.9\\%, 91.6\\%, and 88.5\\% for galaxies, stars, and nsngs in the SDSS, and 94.7\\% for quasars, which are classified as nsng, in the 2QZ. The corresponding efficiencies are 98.3\\%, 93.5\\%, 94.5\\%, and 87.4\\%. The number count distributions suggest that the classifications are reliable for $r \\lesssim 20$, giving 7,978,356 galaxies, 832,993 nsng, and 13,945,621 stars. As we have not applied any restrictions to the classification catalog, such as limiting by photometric error or detection flags, the number of reliable sources will be lower; however, we provide these classifications to eliminate any opaque selection effects from our classification catalog, which simplifies the task of using these classifications to supplement additional analyses, such as as defining a target sample for follow-up spectroscopy.  The assignment of probabilities to each object allows one to investigate the completeness and efficiency of the classifications as a function of these probabilities. While ratios of the probabilities were investigated, we find that the samples with an optimal completeness and efficiency are yielded by the simple cuts $p_g \\geq 0.49$, $p_n \\geq 0.54$ and $p_s \\geq 0.46$. Other values close to 0.5 also give very similar results. The full data describing the completeness and efficiency as a function of these three threshold probabilities is made available with the catalog.  We feel that the application of machine learning algorithms to large, astronomical surveys is a rich area of research. We are currently augmenting our training data with both additional wavelengths and fainter spectroscopic identifications. To improve our classification results, we are also performing additional tests to determine the optimal parameters for decision trees built from these data. Given the efficacy of our approach, we plan to increase the number of training classes used in our analysis, first by using just the SDSS, and later by using additional photometric and spectroscopic data. Finally, we are also exploring the efficacy of more powerful algorithms, such as instance-based classification, to tackle the fundamental goal of characterizing sources by the fraction of their energy that is derived from fusion and accretion. These results will presented in subsequent papers in this series."
    },
    "0606/astro-ph0606588_arXiv.txt": {
        "abstract": "The evidence for a rotation of the $\\epsilon$ Eridani debris disc is examined. Data at 850 $\\micron$ wavelength were previously obtained using the Submillimetre Common User Bolometer Array (SCUBA) over periods in 1997-1998 and 2000-2002.  By $\\chi^2$ fitting after shift and rotation operations, images from these two epochs were compared to recover proper motion and orbital motion of the disc. The same procedures were then performed on simulated images to estimate the accuracy of the results.  Minima in the $\\chi^2$ plots indicate a motion of the disc of approximately 0.6$''$ per year in the direction of the star's proper motion. This underestimates the true value of 1$''$ per year, implying that some of the structure in the disc region is not associated with $\\epsilon$ Eridani, originating instead from background galaxies. From the $\\chi^2$ fitting for orbital motion, a counterclockwise rotation rate of $\\sim$2.75$^{\\circ}$ per year is deduced. Comparisons with simulated data in which the disc is not rotating show that noise and background galaxies result in approximately Gaussian fluctuations with a standard deviation $\\pm$1.5$^{\\circ}$ per year. Thus counterclockwise rotation of disc features is supported at approximately a 2-$\\sigma$ level, after a 4-year time difference. This rate is faster than the Keplerian rate of 0.65$^{\\circ}$ per year for features at $\\approx$65~AU from the star, suggesting their motion is tracking a planet inside the dust ring.  Future observations with SCUBA-2 can rule out no rotation of the $\\epsilon$ Eridani dust clumps with $\\sim$4$\\sigma$ confidence. Assuming a rate of about 2.75$^{\\circ}$ per year, the rotation of the features after a 10-year period could be shown to be $\\geq$1$^{\\circ}$ per year at the 3$\\sigma$ level. ",
        "introduction": "\\subsection{$\\epsilon$ Eridani: A Special Case}  Debris discs around nearby stars represent extra-solar analogues of the Kuiper Belt. It is in this context that studying the $\\epsilon$ Eridani dust ring is of particular interest since it is of spectral type K2V, not too dissimilar to the Sun, but with a much younger age of 0.8 Gyr \\citep{Song 2000, Di Folco 2004} and thus represents an analogue to the young solar system. It is generally thought that protoplanetary discs evolve into debris discs after planets have formed, a process that is completed within 10-100 Myr after star birth \\citep{Schutz 2004,Holland 1998}. \\citet{Greaves 2004} have noted that there is an apparent lack of overlap between stars with debris discs detectable at submillimetre wavelengths and those with radial velocity planet detections. There are only a few stars with positive detections for both an infrared dust-excess and a radial velocity planet \\citep{Dominik 1998,Beichman 2005}. While this is likely to reflect the small number of stars searched for both phenomena, $\\epsilon$~Eridani remains unique in having resolved structure within the debris ring \\citep{Greaves 1998, Greaves 2005} plus an inner gas giant planet detected by Doppler wobble \\citep{Hatzes 2000}. Further evidence for this planet comes from the forced offset of the centre of the ring \\citep{Greaves 2005}, consistent with a forced eccentricity of the dust particles \\citep{Wyatt 1999}. The ring structure also appears perturbed as if by a more distant gas giant, but this body has not yet been imaged, with a maximum mass constraint at around the dust ring radius of approximately 5 M$_J$ \\citep{Macintosh 2003}.  \\subsection{Planet Hunting by Tracking of Disc Features}  An infrared excess around $\\epsilon$ Eridani was first detected during photometric measurements by the Infrared Astronomical Satellite (IRAS) \\citep{Aumann 1988}, soon after the initial discovery of an IR excess around Vega \\citep{Aumann 1984}. Subsequent submillimetre observations at 850 $\\micron$ with the Submillimetre Common User Bolometer Array (SCUBA) by \\citet{Greaves 1998} resolved the disc and showed a ring-like structure. The faint ring surrounding $\\epsilon$ Eridani has not yet been successfully imaged at optical wavelengths \\citep{Proffitt 2004}.  Further SCUBA observations made in 2000-2002 \\citep{Greaves 2005} allow us to study the motions of the substructure in the ring over a 5 year period. If clump features are tracking the orbital motion of a planet within the ring, they should appear to rotate faster than the Keplerian rate at the ring radius; this would be an unambiguous signature of a planet with the forced period providing a measure of its orbital semi-major axis.  Clumps will appear in a disc where planet formation has taken place because the planets will migrate outwards via angular momentum exchange with the remaining planetesimals. As the planets migrate outwards, the associated gravitational resonances also move out, thus trapping dust particles outside the planet's orbit into mean motion resonances \\citep{Wyatt 1999,Wyatt 2003}. The Kuiper Belt in the Solar System probably formed in a similar way, where the material is thought to have been pushed outwards by a mean motion resonance associated with Neptune as it migrated further away from the Sun \\citep{Gomes 2003, Levison 2003,Malhotra 1995}.  Circumstellar dust, trapped into resonances by an outwardly migrating planet \\citep{Wyatt 1999, Wyatt 2003, Quillen 2002}, is expected to be observable at sub-millimetre wavelengths. Whilst the Vega and Fomalhaut disc asymmetries have been modelled \\citep{Wyatt 2003, Wyatt 2002}, the $\\epsilon$ Eridani dust ring offers a unique prospect to directly track motion of clumps over time, because of its favourable inclination ($\\sim$25$^{\\circ}$ from face-on to the observer) and the fact that this is one of the closest stars to the Sun at a distance of 3.22 pc.  Models of the distribution of the $\\epsilon$ Eridani clumps \\citep{Quillen 2002,Ozernoy 2000} suggest the perturbing planet lies $\\approx$40-60 AU from the star, thus the clumps are expected to orbit with a period of $\\approx$280--520 years. These rotation rates of 0.7--1.3$^{\\circ}$ per year, corresponding to particular resonant periods \\citep{Quillen 2002,Ozernoy 2000}, require long times to detect.  Here we present a preliminary analysis over 5~years of data, making use of image fitting to recover statistically any small rotation present. ",
        "conclusions": "We have identified a motion of the $\\epsilon$ Eridani disc consistent with features tracking in the same direction as the proper motion of the star. The best fitting proper motion $\\mu$ = 0.6$''\\pm$0.4$''$ per year is most probably less than the star's proper motion of 1$''$ per year, confirming that some of the features in the vicinity of the disc are likely to be background galaxies. The measured rotation rate of 2.75$^{\\circ}$ per year counterclockwise is in the same direction but nearly twice as large as that suggested by \\citet{Greaves 2005}; it could also be an underestimate as stationary galaxies tend to bias the rotation solution below the real amplitude. Comparisons with simulated data show that the measured rotation is significant at the $\\sim$2$\\sigma$ level, and a future image with SCUBA-2 in 2007 could rule out no or very slow rotation at the $\\sim$4$\\sigma$ level. The technique of $\\chi^2$ fitting is inherently limited by trying to find one solution that fits both stationary galaxies and moving ring features, hence a method that identifies individual clumps or a method capable of using the proper motion to distinguish the foreground and background images will be needed to measure the true rotation rate.  Radial velocity detections of a planet are restricted to finding planets out to only a few AU from the star. In the case of $\\epsilon$ Eridani, finding a planet out to tens of AU via radial velocity measurements will be extremely difficult since the period of the planet is deduced to be $>$100 years, and the magnitude of the Doppler wobble is reduced by the nearly face-on inclination ($\\sim$25$^{\\circ}$) of the disc. However, this viewing angle also means that $\\epsilon$ Eridani provides a unique opportunity to track the motion of the substructure within the disc. This is the first ever analysis attempting to track clumps rotating in a dusty disc. Future observations with SCUBA-2 can confirm the rotation of clumps in resonance with a planet at tens of AU, while imaging with submillimeter interferometers such as SMA and ALMA could radically shorten the timescale to identify such motions."
    },
    "0606/astro-ph0606631_arXiv.txt": {
        "abstract": "We report the detection of interstellar hydrogen deuteride (HD) toward the supernova remnant IC443, and the tentative detection of HD toward the Herbig Haro objects HH54 and HH7 and the star forming region GGD37 (Cepheus A West).  Our detections are based upon spectral line mapping observations of the R(3) and R(4) rotational lines of HD, at rest wavelengths of 28.502 and 23.034$\\,\\mu$m respectively, obtained using the Infrared Spectrograph onboard the {\\it Spitzer Space Telescope}.  The HD R(4)/R(3) line intensity ratio promises to be a valuable probe of the gas pressure in regions where it can be observed. The derived HD/H$_2$ abundance ratios are $(1.19^{+0.35}_{-0.24}) \\times 10^{-5}$, $(1.80^{+0.54}_{-0.32}) \\times 10^{-5}$, and $(1.41^{+0.46}_{-0.33}) \\times 10^{-5}$ respectively (68.3$\\%$ confidence limits, based upon statistical errors alone) for IC443 (clump C), HH54, and HH7.  If HD is the only significant reservoir of  gas-phase deuterium in these sources, the inferred HD/H$_2$ ratios are all consistent with a gas-phase elemental abundance $[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} \\sim 7.5 \\times 10^{-6}$, a factor 2 -- 3 below the values obtained previously from observations of atomic deuterium in the local bubble and the Galactic halo.   However, similarly low gas-phase deuterium abundances have been inferred previously for molecular gas clouds in the Orion region, and in atomic clouds along sight-lines within the Galactic disk to stars more distant than 500~pc from the Sun. ",
        "introduction": "With a cosmic abundance $> 10^{-5}$, deuterium is more common than all but a half-dozen heavy elements. The initial deuterium abundance was set by primordial nucleosynthesis (e.g.\\ Schramm \\& Turner 1998), roughly $\\rm 100 \\, s$ after the Big Bang, and has been reduced by stellar nuclear reactions over the subsequent 13.7 Gyr. The primordial deuterium abundance revealed from intergalactic absorption components (e.g.\\ Kirkman et al.\\ 2003, who obtained [D/[H]=$2.78^{+0.44}_{-0.38} \\times 10^{-5}$ toward Q1243+3047) is in good agreement with models for primordial nucleosynthesis, given the baryon densities derived independently from an analysis of the cosmic microwave background (e.g.\\ Spergel et al.\\ 2003). Thus, the distribution of deuterium in the Galaxy probes both stellar processing and the degree to which material is mixed efficiently within the interstellar medium. The abundance of interstellar deuterium has been measured primarily by ultraviolet absorption line observations of atomic deuterium within diffuse clouds (reviewed by Moos et al.\\ 2002). Recent measurements of the Galactic deuterium abundance reveal significant variations (over the range $\\sim 0.7 - 2.2 \\times 10^{-5}$) that are inconsistent with current predictions from Galactic chemical evolution models (e.g.\\ Wood et al.\\ 2004, Friedman et al.\\ 2006, and references therein); understanding the nature of these inconsistencies is critical to our understanding of chemical evolution within the Galaxy.  While deuterium has been extensively studied in atomic clouds, the deuterium abundance in molecular clouds is less certain. In diffuse molecular clouds, interstellar hydrogen deuteride (HD) has been widely observed by means of ultraviolet spectroscopy -- starting over 30 years ago with the {\\it Copernicus} satellite (reviewed by Spitzer \\& Jenkins 1975) -- and is the most abundant deuterium-bearing molecule. However, deriving the deuterium abundance from the HD column in diffuse clouds is difficult for several reasons: the chemistry is complex, HD contains only a trace amount of D, and the HD abundance is sensitive to the density and UV field in the cloud (Lacour et al.\\ 2005; Le~Petit et al.\\ 2002).  In dense molecular clouds, by contrast, the observations are much fewer. Prior to the results reported here, detections of hydrogen deuteride emissions had been reported toward just one cloud: the Orion Molecular Cloud (Wright et al.\\ 1999; Howat et al.\\ 2002; Bertoldi et al.\\ 1999; the latter's detection of HD from warm shocked gas having been anticipated theoretically by Timmermann 1996). In addition, HD R(0) line absorption has been detected toward two far-IR continuum sources: Sgr B2 (Polehampton et al.\\ 2002) and W49 (Caux et al.\\ 2002). HD is not easily detected by infrared emission-line spectroscopy; its small dipole moment leads to relatively weak emission and its large rotational constant places its pure rotational transitions at infrared frequencies that are inaccessible from ground-based telescopes. On the other hand, a large number of other deuterium-bearing molecules are readily detected with ground-based radio telescopes. Many molecules, such as methanol, show high levels of deuterium fractionation, in which the deuterated/non-deuterated abundance ratio can exceed the cosmic deuterium abundance by more than four orders of magnitude; several doubly- and even triply-deuterated species have been observed (Parise et al.\\ 2004).  In this paper, we report a detection of HD in shocked molecular gas within the supernova remnant IC443, clump C, and tentative detections in three other sources -- the Herbig-Haro objects HH54 and HH7, and the star-forming region GGD37 (Cep A W) -- all from observations with the Infrared Spectrometer (IRS) on the {\\it Spitzer Space Telescope}.  The observational results are discussed in \\S2, below, and the implied gas density and HD abundance derived in \\S3.  The implications of the inferred HD abundance are discussed in \\S 4. ",
        "conclusions": "The signal-to-noise ratio is significantly better toward IC443C, HH54FS, and HH7, the only cases for which R(4) is detected at the $5 \\, \\sigma$ level.  The best fit gas pressures in these sources, found to lie in the range $1.6 \\times 10^7 - 1.0 \\times 10^8 \\, \\rm cm^{-3}\\, K$, correspond to gas densities in the range $4 \\times 10^4$ to $2.5 \\times 10^5 \\, \\rm cm^{-3}$ for the warm gas component.  These values are in reasonable agreement with those derived previously for these regions (see discussion of HH7 and HH54 by Neufeld et al.\\ 2006; and of IC443 by Snell et al.\\ 2005.)  The HD R(4)/R(3) intensity ratio promises to be a valuable probe of the gas pressure in regions where it can be observed.  Our analysis yields best estimates of $(1.19^{+0.35}_{-0.24}) \\times 10^{-5}$, $(1.80^{+0.54}_{-0.32}) \\times 10^{-5}$, and $(1.41^{+0.46}_{-0.33}) \\times 10^{-5}$ (68$\\%$ confidence intervals) respectively for the HD abundance relative to H$_2$ in IC443C, HH54FS, and HH7, but all three estimates are consistent (within the 68$\\%$ confidence intervals) with $n({\\rm HD})/n({\\rm H}_2) = 1.5 \\times 10^{-5}$.  If HD accounts for all the deuterium in the gas phase\\footnote{In interpreting their observations of Orion, Bertoldi et al.\\ argued that significant destruction of HD occurred in the shocked gas that they observed, as the result of the reaction $\\rm H + HD \\rightarrow D + H_2$ (Timmermann 1996), and applied a correction in deriving the gas-phase elemental abundance of HD.  In the shocked regions that we have observed, however, the inferred shock velocities ($\\sim 10 - 20\\, \\rm km \\,s^{-1}$; Neufeld et al.\\ 2006) are significantly smaller than those for the Orion shock; thus the resultant atomic H abundance is expected to be much smaller and the extent of HD destruction negligible.}, this would imply $[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = 7.5 \\times 10^{-6}$, a value which lies a factor 2 -- 3 below those inferred from atomic D absorption line observations of the local bubble ($[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = (1.52 \\pm 0.08)\\times 10^{-5}$; Moos et al.\\ 2002), atomic D absorption line observations of the halo ($[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = (2.2 \\pm 0.7) \\times 10^{-5}$; Sembach et al.\\ 2004), and the expected abundance from primordial nucleosynthesis ($[n_{\\rm D}/n_{\\rm H}]= (2.62^{+0.18}_{-0.20}) \\times 10^{-5}$; Spergel et al.\\ 2003, given the cosmological parameters determined by the Wilkinson Microwave Anisotropy Probe).  However, it is comparable to those derived for the Orion Molecular Cloud from a claimed detection of HD R(5) toward Orion Peak 1 ($[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = (7.6 \\pm 2.9) \\times 10^{-6}$; Bertoldi et al.\\ 1999) and from observations of HD R(0) emission from the Orion Bar ($[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = (1.0 \\pm 0.3) \\times 10^{-5}$; Wright et al.\\ 1999). It is also remarkably consistent with the deuterium abundances observed recently in atomic clouds along long sight-lines within the Galactic disk toward stars further than 500~pc from the Sun and with atomic H column densities in excess of 10$^{20.5}\\,\\rm cm^{-2}$, which yield $[n_{\\rm D}/n_{\\rm H}]_{\\rm gas} = (8.5 \\pm 1.0) \\times 10^{-6}$ (Hoopes et al.\\ 2003; Wood et al.\\ 2004).  The explanation of these variations has been a subject of considerable debate (Friedman et al.\\ 2006).  H\\'ebrard and Moos (2003) have argued that the lower values inferred for high-$N_{\\rm H}$ sight-lines to distant stars represent the typical Galactic elemental abundance for deuterium, requiring a substantial degree of destruction via astration (i.e.\\ in stars), while Draine (2004) has argued for the importance of depletion in reducing the gas-phase deuterium abundance, a possibility originally raised by Jura (1982).  In this picture, material within the local bubble is anomalous because it has suffered recent grain destruction in shocks driven by stellar winds or supernovae. An observed correlation between the gas-phase abundances of deuterium and the refractory element titanium (Prochaska et al.\\ 2005), and the possible detection of deuterated polycyclic aromatic hydrocarbons (Peeters et al.\\ 2004), have been interpreted as supporting the importance of depletion. While our determination of the HD abundance in dense shocked clouds does not directly settle this debate, it does suggest that the gas-phase deuterium abundance in dense clouds differs little from that inferred for typical diffuse clouds along high-$N_{\\rm H}$ sight-lines to distant stars within the Galactic disk."
    },
    "0606/astro-ph0606407_arXiv.txt": {
        "abstract": "We report a discovery of proto-cluster candidates around high redshift radio galaxies at $z\\sim2.5$ on the basis of clear statistical excess of colour-selected galaxies around them seen in the deep near-infrared imaging data obtained with CISCO on Subaru Telescope. We have observed six targets, all at similar redshifts at $z\\sim2.5$, and our data reach to $J=23.5$, $H=22.6$ and $K=21.8$ (5$\\sigma$) and cover a $1.6' \\times 1.6'$ field centered on each radio galaxy. We apply colour cuts in $JHKs$ in order to exclusively search for galaxies located at high redshifts, $z>2$.  Over the magnitude range of $19.5<K<21.5$ we see a significant excess of red galaxies with $J-K>2.3$ by a factor of two around the combined radio galaxies fields compared to those found in the general field of GOODS South. The excess of galaxies around the radio galaxies fields becomes more than factor of three around $19.5<K<20.5$ when the two-colours cuts are applied with $JHKs$. Such overdensity of the colour-selected galaxies suggest that those fields tend to host high density regions at high redshifts, although there seems to be the variety of the density of the colour-selected galaxies in each field. In particular, two radio galaxies fields out of the six observed fields show very strong density excess and these are likely to be proto-clusters associated to the radio galaxies which would evolve into rich clusters of galaxies dominated by old passively evolving galaxies. ",
        "introduction": "\\label{sec:intro}  From many previous studies based on fundamental planes and colour-magnitude (CM) relations in clusters up to $z\\sim 1.3$, it appears that major star formation in the brightest cluster galaxies was virtually completed by $z\\sim 2$ (e.g., \\citealp{bow92}, \\citealp{ell97}, \\citealp{kod98}, \\citealp{sta98}, \\citealp{van98}, \\citealp{bla03}). Therefore, one expects to find proto-clusters dominated by already passively evolving galaxies even beyond $z\\sim 2$. Searching for such ancestors of present-day rich clusters of galaxies is not an easy task, however, since it requires a large volume to survey in order to find such rare objects. Radio galaxies can be viewed as good markers of such high density regions at high redshifts, since they are among the most massive galaxies at any redshift ($M_{\\rm star}>10^{11} M_{\\odot}$, \\citealp{roc04}). In fact, about half of powerful radio galaxies at intermediate redshift inhabit rich environments (\\citealp{hil91}, \\citealp{bor06}). The hypothesis that radio galaxies point us towards proto-clusters at high redshift has been further supported by the recent finding of a correlation between AGN activity and bulge mass (Magorrian relation, \\citealp{geb00}). Given this relation, it is natural to imagine that massive cD galaxies sitting in deep potential well of clusters tend to host massive powerful AGN and hence they are viewed as radio galaxies at high redshift.  Narrow-band surveys of Ly$\\alpha$ emitters around high-$z$ radio galaxies have been intensively conducted by the Leiden group (e.g., \\citealp{kur00}, \\citealp{ven02}) over many years who discovered a large number of emitters at the same redshifts of the radio galaxies, and a large fraction of those emitters have been spectroscopically confirmed (e.g., \\citealp{kur00}, \\citealp{ven02}). These overdense regions of star forming galaxies are embedded in large scale structures at high redshifts and are likely to evolve into massive systems at the present day, such as clusters of galaxies.  Alternative approach is to find evolved populations at longer wavelength. This is complementary to the emitter surveys. \\cite{bes03} found an excess of red galaxies with $R-K>4$ around radio-loud AGNs at $z\\sim1.6$. \\cite{hal98} and \\cite{kod03} also found statistical overdensities of red galaxies around radio-loud AGNs at $1\\lsim z\\lsim1.5$ in near-infrared surveys. \\cite{wol03} also report overdensities of red galaxies in optical--NIR colours around 13 QSO's at $2\\lsim z\\lsim3$. These studies indicate that there is also an excess of {\\it evolved} population around the high-$z$ radio galaxies, as expected if they are embedded in proto-clusters.  We take the second approach and extend the previous analyses to higher redshift based on the near-infrared (NIR) imaging. And in this {\\it Letter}, we report a discovery of the overdensity of NIR-selected galaxies in the fields of six radio galaxies at $z\\sim2.5$.  The Vega-referred magnitude system is used throughout the paper. ",
        "conclusions": "\\label{sec:summary} We searched for proto-clusters around the radio galaxies at $z\\sim2.5$ based on the deep $JHK$-bands imaging taken with Subaru telescope. In order to largely subtract foreground/background contamination and have high contrast of proto-cluster member candidates associated to the radio galaxies, we applied several colour selections; (1) a widely used single colour cut of $J-K>2.3$ and (2) our new two-colours cut with $JHK$-bands.  In both cases, we see a clear excess in the number densities of colour-selected galaxies in the observed radio galaxies fields compared to the general field GOODS-South. The overdensity in the six combined radio galaxies fields is about factor of two for the galaxies with $J-K>2.3$ and is factor of three for those selected by the two-colours cut, although the variety of the overdensity in each field is also seen. In particular, two radio galaxies fields, namely 4C~$-$00.62 and 4C~+23.56, show the highest overdensities of the colour-selected galaxies of more than factor of five. Therefore, we are likely to be witnessing the assembly and formation of proto-clusters with many evolved galaxies at $z\\sim2.5$ associated to the radio galaxies."
    },
    "0606/hep-ph0606256_arXiv.txt": {
        "abstract": "We consider the nonlinear dynamics of inflaton fragmentation during and after preheating in the simplest model of chaotic inflation. While the earlier regime of parametric resonant particle production and the later turbulent regime of interacting fields evolving towards equilibrium are well identified and understood, the short intermediate stage of violent nonlinear dynamics remains less explored.  Lattice simulations of fully nonlinear preheating dynamics show specific features of this intermediate stage: occupation numbers of the scalar particles are peaked, scalar fields become significantly non-gaussian and the field dynamics become chaotic and irreversible. Visualization of the field dynamics in configuration space reveals that nonlinear interactions generate non-gaussian inflaton inhomogeneities with very fast growing amplitudes. The peaks of the inflaton inhomogeneities coincide with the peaks of the scalar field(s) produced by parametric resonance.  When the inflaton peaks reach their maxima, they stop growing and begin to expand. The subsequent dynamics is determined by expansion and superposition of the scalar waves originating from the peaks.  Multiple wave superposition results in phase mixing and turbulent wave dynamics.  Thus, the short intermediate stage is defined by the formation, expansion and collision of bubble-like field inhomogeneities associated with the peaks of the original gaussian field. This process is qualitatively similar to the bubble-like inflaton fragmentation that occurs during tachyonic preheating after hybrid or new inflation. ",
        "introduction": "The origin of matter in the universe from a decaying inflaton field in the process of (p)reheating is a basic feature of all realistic inflationary models. If four dimensional effective field theory is sufficient (for specifics of reheating in string theory inflation see e.g. \\cite{KY}) this process is described by the non-equilibrium QFT of particle creation and thermalization. In chaotic inflation this particle creation typically involves a period of parametric resonance, when occupation numbers of Bose particles rapidly become exponentially large \\cite{KLS}. In this case the QFT is well approximated by classical field theory, which can be investigated in detail with classical lattice simulations \\cite{lattice}.  A full treatment of the quantum field theory of the non-linear stages of preheating in controllable models in the limit of high occupation numbers is in agreement with the classical approximation \\cite{berges}.    The simplest possible inflationary potential contains a massive inflation field $V=\\frac{1}{2} m^2 \\phi^2$ and the simplest preheating model involves a  coupling of the inflaton to another field $ \\frac{1}{2}g^2 \\phi^2 \\chi^2$. The regime of parametric resonant particle production is  understood analytically  \\cite{KLS}. Backreaction of inhomogeneous fluctuations quickly brings the system of interacting scalar fields to a strongly  non-linear regime characterized by very high occupation numbers. The turbulent regime of interacting classical scalar field waves was studied in detail in numerical simulations \\cite{lattice,MT,FK}, most of which are based on the LATTICEASY code \\cite{FT}, and in the $\\lambda \\phi^4$ model even analytically with the kinetic theory of Kolmogorov-type turbulence \\cite{MT}. The least understood stage of (p)reheating is the short, violent transition from linear preheating to the turbulent stage, which shows  anomalies in the momentum space picture, and in the departure from gaussian statistics \\cite{FK}.  Hybrid inflation is another very important class of inflationary models.  At first glance preheating in hybrid inflation, which contains a symmetry breaking mechanism in the Higgs field sector, has a very different character than in chaotic inflation.  Preheating in hybrid inflation occurs via tachyonic preheating \\cite{GBFKLT}, in which a tachyonic instability of the homogeneous modes drives the production of field fluctuations.  In hybrid inflation, the decay of the homogeneous fields leads to fast non-linear growth of scalar field lumps associated with the peaks of the initial (quantum) fluctuations. The lumps then build up, expand and superpose in a random manner to form turbulent, interacting scalar waves \\cite{GBFKLT,lumps}.  A similar picture emeres in new inflation preheating \\cite{new}.  Like parametric resonance, tachyonic preheating can be interpreted via the reciprocal picture of copious particle production far away from thermal equilibrium, and consequent cascades of energy through interacting, excited modes.    In this paper we investigate in detail the structure of preheating after chaotic inflation in configuration space. We find that the intermediate stage between linear preheating and turbulence proceeds via non-linear growth, expansion, and superposition of large value field bubbles, similar to what we earlier observed in hybrid inflation. Thus, the bubble-like intermediate structure of the non-linear fields is another consistent feature of preheating. However, the details of the non-linear dynamics are different: while the bubbles in preheating after hybrid inflation initially occur as isolated patches with large spaces in between them, the bubbles that appear in preheating in the model being considered here appear more densely throughout the space and persist for some time as a pattern of standing waves before they begin spreading and colliding.  In section \\ref{model} we describe the model we are considering and review the basic nature of preheating in this model, focusing on the (well studied) behavior of the fields in momentum space. In section \\ref{results} we discuss different diagnostics of the system of interacting scalar fields  in configuration space and describe the fully non-linear dynamics of inflaton fragmentation. In the concluding section we discuss some of the implications of these results, in particular, for the generation of gravitational waves from preheating and baryo/leptogenesis from preheating. ",
        "conclusions": "It is convenient to split the process of transition from the homogeneous inflaton condensate to the radiation of randomly moving waves into four stages: The first is exponentially rapid growth of small inhomogeneities that emerge from vacuum fluctuations. During this stage fluctuations are linear and the fields are gaussian random fields. The second stage is violent backreaction and rescattering of waves, with non-gaussian, nonlinear, nonthermal fluctuations. The third stage is Kolmogorov turbulence. During this period the fluctuations are again (nearly) gaussian and energy gradually cascades towards high momentum modes. Finally, there is thermalization.  In this paper we used lattice simulations to investigate the second stage of violent field restructuring. We considered the model (\\ref{potential}), in which preheating occurs through parametric resonance, and examined the evolution of the fields in configuration space. The picture that emerged is similar in many ways to what we observed earlier for tachyonic preheating in hybrid inflation.  For instance in the F-term inflation example where the non-linear potential around the bifurcation point is $V(\\sigma)=-\\frac{\\lambda}{3}v \\phi^3 +\\frac{\\lambda}{4} \\phi^4+\\frac{\\lambda}{12}v^4$,  small initial random fluctuations of $\\phi$ are amplified by the non-linear $\\phi^3$ term. The peaks of the gaussian field thus begin to grow very fast relatively to the surrounding regions of $\\phi$. Along those same lines, we found here that oscillations of the $\\chi$ and $\\phi$ fields grow initially at the locations of the peaks in the initial $\\chi$ field. For the model (\\ref{potential}) these growing peaks form a pattern of standing waves that persists throughout the linear regime and then begin to spread and overlap as rescattering becomes important.  The bubble-like structure we see here (and in hybrid inflation) has nothing to do with first-order phase transitions, but just with the initial structure of vacuum fluctuations plus non-linear dynamics \\footnote{Similar physics operates in a different context with gravitational instability in an expanding universe, when growth of initial matter inhomogeneities results in the cosmic web pattern of large scale structure.  Clusters (voids) form from the high positive (negative) peaks of the initial gaussian field, originated from vacuum fluctuations during inflation. See e.g. J.R.~Bond, L.~Kofman, D.~Pogosian, ``How filaments are woven into the cosmic web''.  Nature 380:603-606,1996.}.  Growth of the individual peaks results in the build-up of the scalar field gradients.  Subsequent evolution is defined by the expansion of the bubbles.  The superposition of many almost spherically expanding bubbles leads to decoherence and turbulent motion of the scalar waves.  The main lesson we have learned from this work is that the preheating stage of linear fluctuations and the turbulent stage of interacting waves are divided by a short, violent stage of non-linear formation and collision of bubble-like large value field regions.  Eventually the fields reach thermal equilibrium characterized only by the temperature. Does that mean that all traces of inflaton fragmentation history are erased?  There are potential tracers of the non-linear stage of preheating related to out-of-equilibrium processes. For instance, people have discussed realizations of baryogenesis at the electroweak scale via tachyonic preheating after hybrid inflation \\cite{GGKS}, and this process is ultimately related to the bubble-like lumps of the Higgs field that form during tachyonic preheating \\cite{GGG}. Since we now see that fragmentation through bubbles can also occur in chaotic inflation, baryogenesis via out-of-equilibrium bubbles can also be extended to these models.  There is another, potentially observable consequence of the non-linear ``bubble'' stage of inflaton fragmentation. Lumps of the scalar fields correspond to large (order of unity) energy density inhomogeneities at the scale of those bubbles, $R$. Collisions of bubbles generate gravitational waves.  The fraction of the total energy at the time of preheating converted into gravitational waves is significant. We estimate it is of the order of \\begin{equation}\\label{fraction} \\frac{\\rho_{gw}}{\\rho_{rad}} \\simeq (RH)^2 \\ , \\end{equation} where $1/H$ is the Hubble radius. This corresponds to a present-day fraction of energy density $\\Omega_{GW} \\sim 10^{-5}(RH)^2$. The way to understand formula (\\ref{fraction})  is the following: The energy converted into gravitational waves from the collision of two black holes is of the order of the black hole masses. If the mass of lumps of size $R$ is a fraction $f$ of a black hole of the same size, then the fraction of energy converted to gravitational waves from two lumps colliding is $f$. Scalar field lumps at the Hubble scale would form black holes, so in our case $f=(RH)^2$.  The present-day frequency of this gravitational radiation is \\begin{equation}\\label{frequency} f\\simeq \\frac{10^7 Gev}{M} Hz \\ , \\end{equation} where $M=V^{1/4}$ is the energy scale of inflation with the potential $V$.  For the chaotic inflation model considered in this paper the size of the bubbles is $R \\sim few/m$ and at the time they begin colliding $H \\sim m/100$, so that the fraction of energy converted into gravitational waves is of the order $10^{-3}-10^{-4}$.  This figure is in agreement with the numerical calculations of gravitational wave radiation from preheating after chaotic inflation \\cite{GW}.  For chaotic inflation with $M$ at the GUT scale the frequency (\\ref{frequency}) is too short and not observable.  Gravitational waves continue to be generated during the turbulent stage and even during equilibrium due to thermal fluctuations, but with a smaller amplitude. It is a subject of further investigation if they can be observed.  The most promising possibility for observations is, however, generation of gravity waves from low energy hybrid inflation, where $f$ can much much smaller.    \\bigskip  We would like to thank Marco Peloso for useful discussions. G.F. was supported by NSF grant PHY-0456631. L.K. wa supported by NSERC and CIAR. G.F. would like to thank the Canadian Institute for Theoretical Astrophysics for their hospitality during part of this work"
    },
    "0606/hep-th0606162_arXiv.txt": {
        "abstract": "We consider quasinormal modes with complex energies from the point of view of the theory of quasi-exactly solvable (QES) models. We demonstrate that it is possible to find new potentials which admit exactly solvable or QES quasinormal modes by suitable complexification of parameters defining the QES potentials. Particularly, we obtain one QES and four exactly solvable potentials out of the five one-dimensional QES systems based on the $sl(2)$ algebra. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606139_arXiv.txt": {
        "abstract": "In this paper we discuss the optical and X-ray spectral properties of the sources detected in a single 200\\,ks {\\it Chandra} pointing in the Groth-Westphal Strip region. A wealth of optical photometric and spectroscopic data are available in this field providing optical identifications and redshift determinations for the X-ray population. The optical photometry and spectroscopy used here are primarily from the DEEP2 survey with additional redshifts obtained from the literature. These are complemented with the deeper ($r\\approx26$\\,mag) multi-waveband data ($ugriz$) from the Canada France  Hawaii Legacy Survey to estimate photometric redshifts and to optically identify sources fainter than the DEEP2 magnitude limit ($R_{AB}\\approx24.5$\\,mag). We focus our study on the 2-10\\,keV selected sample comprising 97 sources to the limit $\\approx  8 \\times 10^{-16}\\rm \\, erg \\, s^{-1} \\, cm^{-2}$, this being the most complete in terms of optical identification rate (86\\%) and redshift determination fraction (63\\%; both spectroscopic and photometric). We first construct the redshift distribution of the sample which shows a peak at $z\\approx1$. This is in broad agreement with models where less luminous AGNs evolve out to $z\\approx1$ with powerful QSOs peaking at higher redshift, $z\\approx2$. Evolution similar to that of broad-line QSOs applied to the entire AGN population (both type-I and II) does not fit the data. We also explore the {\\it observed} $N_H$ distribution of the sample and estimate a fraction of obscured AGN ($N_H \\rm > 10^{22} \\, cm^{-2}$) of $48\\pm9$ per cent. This is found to be consistent with both a  luminosity dependent {\\it intrinsic} $N_H$ distribution, where less luminous systems comprise a higher fraction of type-II AGNs and models with a fixed ratio 2:1 between type-I and II AGNs. We further compare our results with those obtained in deeper and shallower surveys. We argue that a luminosity dependent parametrisation of the intrinsic $N_H$ distribution is required to account for the fraction of obscured AGN observed in different samples over a wide range of fluxes. ",
        "introduction": "In the last few years the study of the diffuse X-ray background (XRB) has witnessed significant observational progress allowing detailed comparison with model predictions. The ultra-deep {\\it Chandra} surveys in particular, have demonstrated that most of the XRB, at both soft and hard energies, is resolved into discrete point sources (Brandt et al. 2001; Giaconni et al. 2002; Alexander et al. 2003), the vast majority of which are without doubt AGNs. To the first approximation, this finding has been a huge success for models that reproduce the spectral properties of the XRB under the zero order assumption that it originates in a combination of obscured and unobscured AGN (Comastri et al. 1995; Gilli, Salvati \\& Hasinger 2001). Under more careful examination however, a number of inconsistencies emerge. Firstly, luminous ($L_X>10^{44} \\rm \\, erg \\, s^{-1}$) heavily obscured type-II QSOs at $z\\approx1.5-2$, predicted in large numbers by the models, are scarce in the surveys above.  Secondly, the redshift peak of the X-ray population lies below $z=1$ in stark contrast with the model expectation of $z\\approx1.5-2$.  Although these inconsistencies suggest that some revision of the models is almost certainly required (Hasinger 2003), observational biases may complicate any interpretation. For example, about $\\approx25$ per cent of the sources in the {\\it  Chandra} Deep Fields (CDF) are optically faint, $R>24$\\,mag, rendering optical spectroscopy  difficult or even impossible with current technology (Rosati et al. 2002; Barger et al. 2003). Any information about the nature of these sources is therefore limited and they are proposed as best candidates for heavily obscured AGN (Alexander et al. 2001; Treister et al. 2004), likely to comprise a fraction the elusive population of high-$z$ type-II QSOs. Moreover, the small field-of-view of the CDFs (0.07\\,deg$^2$ each) makes them sensitive to cosmic variance further complicating interpretation of the derived redshift distribution.  Wide-area shallower ($\\approx 10^{-14} \\rm \\, erg \\, s^{-1} \\, cm^{-2}$) surveys are less affected by the observational biases above (e.g. Baldi et al. 2002; Kim et al. 2004; Georgantopoulos et al. 2004). These samples although of key importance, comprise a large fraction of unobscured AGNs that are not representative of the sources responsible for the spectral shape of the XRB ($\\Gamma=1.4$; e.g. Gruber et al. 1999).  The evidence above suggests that deep surveys with relatively wide field-of-view are essential to improve our understanding of the XRB. Observational programs in this direction are already well underway such as the {\\it XMM-Newton} Cosmic Evolution Survey [COSMOS; $\\rm 2\\,deg^2$ , $f_X ( \\rm 0.5 - 2 \\, keV) \\approx 5\\times 10^{-16} \\, erg \\, s^{-1} \\, cm^{-2}$] and the Extended {\\it Chandra} Deep Field South  [E-CDF-S; Lehmer et al. 2005; Virani, Treister \\& Urry 2006; $\\rm 0.3 \\, deg^2$, $f_X ( \\rm 0.5 - 2 \\, keV) \\approx  1.1 \\times 10^{-16} \\, erg \\, s^{-1} \\, cm^{-2}$]. In this paper we present  results on a single  200\\,ks {\\it Chandra} pointing, which is part of an on-going X-ray survey in the Extended Groth Strip region, which will eventually cover a total area of about 0.5\\,deg$^2$ to the depth above (200\\,ks per pointing). This sample, when completed, will be intermediate in terms of area coverage and depth to the CDFs and shallower wide-area surveys, minimising any observational biases affecting the ultra-deep fields and comprising a large fraction of obscured AGNs responsible of the XRB properties (Nandra et al. 2005). Moreover, the Extended Groth Strip is targeted by the largest space and ground-based facilities for  multiwavelength observations: (i) the DEEP and DEEP2 surveys provide  optical spectroscopy to the limit $R_{AB} \\approx 24$\\,mag, (ii) multiwaveband optical photometry to fainter magnitudes is underway as part of the Canada France Hawaii Telescope Legacy Survey, (iii) deep imaging and spectroscopy, independent from the programs above, has been performed by Steidel et al. (2003) in search for Lyman break galaxies, (iv) comprehensive imaging with HST/ACS has recently been completed, (v) {\\it Spitzer} mid-IR data are available, (vi) radio observations to sub-mJy levels have been obtained by Fomalont et al. (1991) with new much wider VLA observations recently completed, (vii) SCUBA has observed part of this field to the deep limits of the Canada-UK Deep Submillimetre survey (Webb et al. 2003). A combination of the X-ray observations with the mutliwavelength datasets above promises a breakthrough in the study of the evolution and large scale structure of AGNs as well as the connection between AGN activity and host galaxy formation.  This paper presents the optical and X-ray spectral properties of the sources detected in the first 200\\,ks {\\it Chandra} pointing observed as part of the Extended Groth Strip X-ray survey. This observation encompasses the original Groth-Westphal Strip region (Groth et al. 1994). In addition to studying the properties of the X-ray sources in the context of XRB models, our purpose is to demonstrate the power of the full $\\rm 0.5 \\, deg^2$ Extended Groth Strip {\\it Chandra} survey, when completed, for XRB studies. Throughout this paper we adopt $\\rm Ho = 70 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\rm \\Omega_{M} = 0.3$ and $\\rm \\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "This paper presents first results from an ongoing deep (200\\,ks per pointing) X-ray {\\it Chandra} survey in the Extended Groth Strip region. Data from the first of a total of 8 pointings in this field are used here. We analyse and discuss the optical and X-ray properties of the sample in the context of XRB population synthesis models.  We first construct the photometric and spectroscopic redshift distribution of the 2-10\\,keV selected sample to the limit $\\approx 8 \\times 10^{-15}\\rm \\, erg \\, s^{-1} \\, cm^{-2}$ which has a peak at $z\\approx1$ in agreement with previous studies. Although there is about 40 per cent redshift incompleteness we find that luminosity dependent density evolution, where lower luminosity systems peak at lower redshifts, is in fair agreement with the observations. Luminosity functions that assume evolution out to $z\\approx2$ for the entire AGN population (both type-I and II), similar to that of unobscured broad-line QSOs do not fit the data. The  luminosity dependent evolution supported by our data is consistent with scenarios suggesting that lower-luminosity systems (dominated by obscured AGNs) are associated with star-formation activity and peak at $z \\approx 1$ while, more powerful QSOs evolve out to higher-$z$.  We also explore the $N_H$ distribution of the sample which is consistent with either a fixed obscured AGN fraction of ${\\cal R} \\approx 2$ or the luminosity dependent $N_H$ distribution proposed by Ueda et al. (2003), where less luminous systems comprise a higher fraction of type-II AGNs. We also argue that such luminosity dependent parametrisation of the  $N_H$ distribution is essential to account for the fraction of obscured AGN observed in different samples over a wide range of flux limits.  The X-ray survey of the Extended Groth Strip region, when completed, will cover a total of about $\\rm 0.5 \\, deg^2$ at flux limits similar to those presented here. The final sample will be about 8 times larger, comprising a total of over 1000 sources. The wealth of follow-up optical  photometric and spectroscopic data available in this field (DEEP2, CFHTLS) will provide optical identifications as well as spectroscopic and photometric redshift estimates for a large fraction of the X-ray population, for example close to 90 per cent for the hard-band. The use of {\\it Spitzer} mid-infrared data will further increase the number of X-ray identifications and also has the potential to refine the photometric redshift estimates, particularly for the optically faint subsample. Such a high identification rate combined with photometric/spectroscopic redshifts and the complementary multi-wavelength data (infrared, radio, sub-mm) available for the EGS promise a major step forward in our understanding of the nature and the evolution of the AGN populations that make up the XRB."
    },
    "0606/astro-ph0606625_arXiv.txt": {
        "abstract": "Globular clusters are often assumed to be good tracers of major star formation episodes in their host galaxies. While observations over the past 2 decades have confirmed the presence of young objects with globular cluster-like properties in many galaxies, it is still not well understood exactly how the formation efficiency of bound star clusters relative to field stars and their mass spectrum depend on external factors. The cluster initial mass function typically appears to be consistent with a power-law with a slope $\\sim-2$, but most attempts to constrain any upper limit on the CIMF have been limited by size-of-sample effects. However, evidence is starting to accumulate for possible truncation of the cluster mass function.  It is tentatively suggested that the upper mass limit may currently be at $\\sim10^5$ M$_\\odot$ in the Milky Way disk, while there are indications that it is $\\sim5\\times10^5$ M$_\\odot$ in M51 and about $2\\times10^6$ M$_\\odot$ in the Antennae. Some extreme starbursts (e.g.\\ Arp 220, NGC~7252) are (or were) able to form clusters as massive as $\\sim10^7$ M$_\\odot$. The overall formation efficiency of star clusters (relative to field stars) in the Galactic bulge may not have been much different from that in the disk today, but was probably significantly higher for metal-poor GCs in halos. ",
        "introduction": "\\label{sec:intro}  Three or four decades ago, globular cluster (GC) formation was thought by many to be a phenomenon occurring only in the early Universe (e.g.\\ \\cite{pd68}). In the meantime, young, compact star clusters with masses in the range $10^4$--$10^6$ M$_{\\odot}$ have been found in many different galaxies (e.g.\\ \\cite{lar05,schweizer,whit01}), and there is a growing concensus that these ``Young Massive Clusters'' (YMCs) may well be young analogues of the old GCs associated almost universally with the spheroidal components of galaxies.  By implication, GCs have become potentially interesting not just as fossil left-overs from the early Universe, but more generally as test particles for studies of extragalactic stellar populations.  In keeping with this spirit, contemporary observing proposals or papers on GCs often include an introductory remark along the lines of: \\emph{``GCs are thought to be good tracers of the major star forming episodes in their host galaxies''}. YMCs/GCs may indeed trace star formation quite generally, but it is also clear that some caution must be exercised. For example, the number of GCs per field star varies from galaxy to galaxy (the classical ``specific frequency problem''), as well as between stellar populations within galaxies \\cite{hh02,ffg05}.  This must be due to differences in the formation efficiency of GCs relative to field stars, in the GC survival rates, or (more likely) some combination of the two.  Of the roughly 150 catalogued GCs in our own Galaxy \\cite{harris96}, about 2/3 are associated chemically, spatially and kinematically with the halo, while this is true for only $\\sim1$\\% of the stars. Conversely, some 90\\% of the stellar mass resides in the thin disk of our Galaxy, while no GCs are currently known to be associated with this component. Even though some quite massive clusters might be located in remote parts of the Galactic disk \\cite{clark05,figer06}, a simple scaling by stellar mass of the GCs in the bulge or halo would predict hundreds or thousands of GCs in the disk, seemingly at odds with the observations.  All this is of course well known, and is the reason why GCs are often assumed to trace only ``\\emph{major} star forming episodes''. The problem then remains to define when a star forming episode qualifies as ``major''. Perhaps GCs mainly trace spheroid formation \\cite{bs06}, but some stars currently residing in spheroids may originally have formed in disks. As an example, it is illustrative to consider the outcome of merging two Milky Way galaxies: The merger product would contain about 200 metal-poor and 100 metal-rich pre-existing GCs from the progenitor galaxies, assuming no GCs are destroyed.  The current gas mass in the Milky Way is about $0.5\\times10^{10}$ M$_{\\odot}$ \\cite{co96}, which is about half the mass of the bulge. Assuming that the merger would form GCs with the same efficiency as the bulge, using all the available gas, about 50 new metal-rich GCs would form (and survive).  The resulting GC population would then consist of three major sub-components: metal-poor, old GCs originating in the progenitor galaxy halos, moderately metal-rich GCs from the pre-existing bulges, and very metal-rich GCs formed in the merger. The estimate of the number of new GCs is obviously very crude. However, the main point here is that while the majority of the \\emph{stars} in the resulting spheroid (about 90\\%) would have formed in the disks already before the merger took place, their formation history might not be reflected in the GC system.  Mergers at higher $z$ were likely more gas-rich, and the discrepancy between the star- and GC formation histories may have been less extreme than in the example outlined above. Nevertheless, since GCs play such an important role in attempts to constrain the star formation histories of early-type galaxies, there is a clear need to also quantify the limitations better. ",
        "conclusions": "While research in extragalactic star clusters remains a very active field, we are still facing a number of important questions. The ubiquity of YMCs in external galaxies is now well established, but the apparent absence of a large population of clusters with masses in the range $10^5$--$10^6$ M$_\\odot$ in the Galactic disk remains somewhat of a puzzle. This may suggest an upper limit to the CIMF near $10^5$ M$_\\odot$ in the Milky Way, as noted already by van den Bergh \\& Lafontaine \\cite{vl84}, although studies of disk clusters are still hampered by our own location close to the Galactic plane.  However, the recent realization that the Milky Way is still forming clusters with masses near $10^5$ M$_\\odot$ hints that the CIMF in the Milky Way may not be very different from that in the LMC, the main difference being that the census of YMCs is more complete in the LMC.  There is some evidence for truncation at a somewhat higher mass in M51 ($\\sim5\\times10^5$ M$_\\odot$) and the Antennae ($\\sim2\\times10^6$ M$_\\odot$), while Arp 220, NGC~1316 and NGC~7252 host clusters as massive as $\\sim10^7$ M$_\\odot$. These galaxies also define a sequence of increasing star formation rate, suggesting that galaxies with higher SFRs are physically able to form more massive clusters. This may provide a hint as to why GC formation was common at high $z$, when SFRs and gas densities were generally higher."
    },
    "0606/astro-ph0606555_arXiv.txt": {
        "abstract": "We demonstrate that a transition from decelerated to accelerated cosmic expansion arises as a pure interaction phenomenon if pressureless dark matter is coupled to holographic dark energy whose infrared cutoff scale is set by the Hubble length. In a spatially flat universe the ratio of the energy densities of both components remains constant through this transition, while it is subject to slow variations for non-zero spatial curvature. The coincidence problem is dynamized and reformulated in terms of the interaction rate. An early matter era is recovered since for negligible interaction at high redshifts the dark energy itself behaves as matter.  A simple model for this dynamics is shown to fit the SN Ia data. The constant background energy density ratio simplifies the perturbation analysis which is characterized by non-adiabatic features. ",
        "introduction": "Nowadays an overwhelming, direct and indirect, observational evidence supports the idea  that the Universe is currently undergoing a phase of accelerated expansion. The most recent and precise confirmation of this rather unexpected feature (see, however \\cite{Turner97,Priester})  was provided by the 3rd year data of the WMAP mission \\cite{WMAP3}. Likewise, data from high--redshift supernovae type Ia \\cite{SNIa}, the cosmic microwave background radiation \\cite{cmb}, the large scale structure \\cite{lss}, the integrated Sachs--Wolfe effect \\cite{isw}, and weak lensing \\cite{weakl}, endorse it to the point that the discussion has now shifted to when the acceleration began and which might be the agent behind it. For recent reviews see \\cite{recentrev}.  According to our understanding on the basis of Einstein's gravity, more than $70$\\% of the cosmic substratum, dubbed  dark energy, must be endowed with a high negative pressure. Less than $30$\\% is found to be pressureless matter. Most of the latter is in the form of cold dark matter; only about $5$\\%  corresponds to ``normal\" baryonic matter. Currently, we neither know what the dark matter is made of nor do we understand the nature of dark energy.  Lacking a fundamental theory, most investigations in the field are phenomenological and rely on the assumption that these two unknown substances evolve independently, i.e., their energies are assumed to obey separate conservation laws. In particular, this implies that the dark matter energy density varies as $a^{-3}$, where $a$ is the scale factor of the Robertson-Walker metric. The behavior of the density of the dark energy is then entirely governed by its equation of state, the determination of which is a major subject of current observational cosmology. One should realize, however, that any coupling in the dark sector will change this situation. A coupling will modify the evolution history of the Universe. As one of the consequences, the energy density of the (interacting) dark matter will no longer evolve as $a^{-3}$. Ignoring a potentially existing interaction between dark matter and dark energy from the start, may result in a misled interpretation of the data concerning the dark energy equation of state. Das {\\em et al.} \\cite{Das} and Amendola {\\em et al.} \\cite{Luca} have shown that a measured phantom equation of state may be mimicked by an interaction, while the bare equation of state may well be of a non-phantom type. Further, models showing interaction fare well when contrasted with data from the cosmic microwave background \\cite{german1} and matter distribution at large scales \\cite{german2}. It can therefore be argued that the possibility of dark energy to be in interaction with dark matter must be taken seriously.  On the other hand, there exist limits for the strength of this interaction for various configurations \\cite{limits}. It is characteristic for these approaches that they admit a non-interacting limit. If the interaction is switched off, they continue to represent models of a mixture of dark matter and dark energy. The best known examples are models with a decaying cosmological ``constant\": if the decay is switched off, they reduce to the $\\Lambda$CDM model. An interaction here leads to (possibly important) corrections of the non-interacting configuration. E.g., for a given (not necessarily constant) negative equation of state parameter of the dark energy, the interaction manifests itself in third order in the redshift in the luminosity distance of supernovae type Ia \\cite{WDGRG}. It does not, however, (directly) influence the leading orders.  The qualitatively new aspect of the present approach is the circumstance that an interaction is crucial already in leading order. In this context the accelerated expansion itself is a consequence of the interaction. The coupling does not just lead to higher order corrections of a non-interacting reference model which by itself provides an accelerated expansion of the Universe. The non-interacting limit of the present approach reproduces an Einstein-de Sitter universe, i.e., there is no accelerated expansion if the interaction is switched off. This non-interacting limit  is supposed to characterize our Universe at high redshifts.    The data also suggest that the present Universe is nearly spatially flat \\cite{WMAP3}, \\cite{seljak}.  Many studies neglect the spatial curvature term and focus solely on the spatially flat case, thereby taking for granted that the spatial curvature term necessarily leads to trifling corrections. Several of the not so many works that retain that term, for the sake of generality, even seem to justify that it is of minor importance. On the other hand, thanks to the increasing observational precision also higher order corrections may become within reach in the not too distant future. The authors of \\cite{Caldwell} have demonstrated that the spatial curvature enters the luminosity distance of SNIa supernovae in third order in redshift (see also \\cite{WangGongSU}). It is also known that there are degeneracies between the curvature and the dark energy equation of state in the corresponding parameter space \\cite{Degeneracies}. The possible relevance of spatial curvature for the lowest multipoles in the cosmic microwave background radiation was discussed in \\cite{Abdalla}. These examples indicate that apart from general theoretical grounds, the observational situation will require the inclusion of spatial curvature, even if its contribution is small, at some level of precision. Further, as demonstrated by Ichikawa {\\it et al.} \\cite{Ichikawa}, depending on the parametrization of the equation of state parameter the present value of the spatial curvature can be as large as $0.2$. Therefore, the additional dynamics provided by the curvature, should certainly not be dismissed too quickly since it contains information which is needed to further restrict the cosmological parameter space. In this connection, Clarkson {\\it et al.} \\cite{Bassett} have shown that by excluding the spatial curvature when interpreting the empirical data one can incur  in  gross mistakes when reconstructing the equation of state of dark energy. We mention that an effective spatial curvature term also appears as the result of an averaging procedure within the ``macrosopic gravity\" approach \\cite{Zala} (for a further discussion of the effects of averaging on cosmological observations see, e.g., \\cite{alan}).  In this paper we consider pressureless dark matter in interaction with an unknown component which is supposed to describe dark energy. We neither specify the dark energy equation of state nor the interaction rate from the beginning. These two (generally time dependent) parameters will influence the ratio of the energy densities of dark matter and dark energy. The behavior of this ratio is crucial for  the ``conventional\" form of the ``coincidence problem\", namely: ``why are the matter and dark energy densities of precisely the same order today\". In principle, matter and dark energy redshift at different rates. We show, that there exists a preferred class of dark energy models for which the dynamics of the energy density ratio is entirely determined by the spatial curvature. For vanishing curvature the energy density ratio remains constant. These models are singled out by a dependence $\\rho_{X} \\propto H^{2}$, where $\\rho_{X}$ is the dark energy density and $H = \\dot{a}/a$ is the Hubble parameter. Exactly this dependence is characteristic for a certain type of dark energy models, inspired by the holographic principle \\cite{hooft,leonard}.  Holographic dark energy models must specify an infrared cutoff length scale \\cite{cohen}. The choice of this scale is presently a matter of debate. The most obvious choice, the Hubble length, seemed to be incompatible with an accelerated expansion of the Universe \\cite{Hsu} (see, however, \\cite{raul1,guberina}). This is why, starting with the work of Li \\cite{mli1}, many researchers have adopted the future event horizon as the cutoff scale as this choice allows for a sufficiently negative equation of state parameter and hence an accelerated expansion -see, e.g. \\cite{feventh}.  In a previous paper \\cite{DW}, we showed that a cutoff set by the Hubble length may well be compatible with an accelerated expansion provided  that the dark energy and dark matter do not evolve separately but interact, also non--gravitationally, with each other. In this setting, a negative equation of state parameter that gives rise to accelerated expansion arises as a direct consequence of the interaction. Here we put this feature in a broader context and  demonstrate that a constant or slowly varying (as the consequence of a non-vanishing spatial curvature) energy density ratio is compatible with a transition from decelerated to accelerated expansion under the condition of a growing interaction parameter. We show that in holographic dark energy models with a Hubble length cutoff a transition from decelerated to accelerated expansion is realized as a pure interaction effect. At high redshifts and for negligible interaction the dark energy equation of state approaches the equation of state for matter, such that a preceding matter era is naturally recovered.  In this context the coincidence problem can be rephrased as follows: ``why is the interaction rate between dark matter and dark energy of the order of the Hubble rate precisely at the present epoch?\" This reformulation allows us to address the coincidence problem as part of the interaction dynamics in the dark sector.  Different interaction rates will imply a different perturbation dynamics. We shall demonstrate that non-adiabatic features appear as a characteristic signature for a large class of interacting models. These effects may be used to discriminate between different models of the cosmic medium which share the same background dynamics.  The outline of this paper is as follows. Section \\ref{general} provides the general formalism for an interacting two-component fluid, where one of the fluids is pressureless. Then it focuses on the case for which the ratio of the energy densities of both fluids is constant or slowly varying. This will single out models for which the energy density of the second component is proportional to the square of the Hubble parameter. A realization of this dependence is provided by certain holographic dark energy models. The basic properties of these models are recalled in Section \\ref{holography}. Section \\ref{interaction} discusses in detail matter in interaction with holographic dark energy. It provides the conditions under which an interaction driven transition from decelerated to accelerated expansion is feasible. It also comments on the possibility of a weak time dependence of the saturation parameter of the holographic bound, which is usually assumed constant. Section \\ref{model} introduces a simple interacting model and compares it with the $\\Lambda$CDM model. The perturbation dynamics for a fluctuating interaction rate is discussed in \\ref{perturbations} with special emphasis on non-adiabatic aspects on large perturbation scales. Finally, section \\ref{discussion} summarizes our findings. ",
        "conclusions": "\\label{discussion} We have established a holographic dark energy model in which a negative equation of state parameter as well as the transition from decelerated to accelerated expansion arise from pure interaction phenomena. A remarkable property of the model is the fact that, in a spatially flat universe, the ratio of the energy densities of dark matter and dark energy remains constant during the transition. The non-interacting limit of this model is an Einstein-de Sitter universe which is supposed to describe our cosmos at high redshifts.  At first sight, the fact that $r$ was never large may seem at variance with the conventional scenario of cosmic structure formation as one may think that at early times the amount of dark matter may have been insufficient to  produce gravitational potential wells deep enough to lead to the condensation of the galaxies. However, this is not so; an earlier matter dominated phase is naturally recovered since for negligible interaction at high redshifts the equation of state of the dark energy is similar to that of non-relativistic matter. In the context of our approach the coincidence problem is dynamized and takes the form: ``why is the interaction rate that drives the accelerated expansion of the order of the Hubble rate precisely at the present epoch?\" Based on the assumption that the relevant interaction parameter grows with a power law in the scale factor (Eq.~(\\ref{ansgamma})), we have worked out a specific model for the transition from decelerated to accelerated expansion under the condition of a constant ratio of the energy densities of dark matter and dark energy. A preliminary analysis shows that this model fits the SNIa data not less well than the $\\Lambda$CDM model.  We have also shown that taking into account the spatial curvature term gives rise to an additional dynamics which implies a small (compared with the Hubble rate) change of the energy density ratio.  Furthermore, we discussed to what extent a slowly varying saturation parameter of the holographic bound may modify the cosmological dynamics. Besides, our approach can cope with a later transition to a new decelerated phase of expansion \\cite{carvalho}-something incompatible with holographic models whose infrared cutoff is set by the radius of the future event horizon.  The interaction between dark matter and dark energy introduces non-adiabatic features into the perturbation theory. For the special case of a constant decay rate of the dark energy, in the background equivalent to a generalized Chaplygin gas, we find, Eq. (\\ref{zeta}), that the large scale curvature perturbations on hypersurfaces of constant density (which are constant in the adiabatic case, in particular for the $\\Lambda$CDM model) vary with the power $3/2$ of the scale factor."
    },
    "0606/astro-ph0606763_arXiv.txt": {
        "abstract": "Even with the renaissance in gamma-ray burst (GRB) research fostered by the Swift satellite, few bursts have both contemporaneous observations at long wavelengths and exquisite observations at later times across the electromagnetic spectrum.  We present here contemporaneous imaging with the KAIT robotic optical telescope, dense optical sampling with Lulin, and supplemented with infrared data from PAIRITEL and radio to gamma-ray data from the literature.  For the first time, we can test the constancy of microphysical parameters in the internal-external shock paradigm and carefully trace the flow of energy from the GRB to the surrounding medium.  KAIT data taken $\\lessim 1$ minute after the start of GRB~051111 and coinciding with the fading gamma-ray tail of the prompt emission indicate a smooth re-injection of energy into the shock.  No color change is apparent in observations beginning $\\sim 1.5$ minutes after the GRB and lasting for the first hour after the burst. There are achromatic flux modulations about the best-fit model at late ($t\\approx 10^4$ s) times, possibly due to variations in the external density. We find that the host-galaxy extinction is well fit by a curve similar to that of the Small Magellanic Cloud.  Low visual extinction, $A_V\\approx 0.2$ mag, combined with high column densities determined from the X-ray and optical spectroscopy ($N_H> 10^{21}$ cm$^{-2}$), indicate a low dust-to-metals ratio and a possible over-abundance of the light metals. An apparent small ratio of total to selective extinction ($R_V\\approx 2$) argues against dust destruction by the GRB.  Time constancy of both the IR/optical/UV spectral energy distribution and the soft X-ray absorption suggests that the absorbing material is not local to the GRB. ",
        "introduction": "The {\\it Swift}~gamma-ray burst (GRB) satellite \\citep{gehrels04} continues to unleash a torrent of finely time- and energy-sampled photons arising from the bursts and their impact on the surrounding medium (i.e., the GRB afterglows). The X-ray Telescope (XRT) provides detailed light curves for 1-2 bursts and burst afterglows per week, allowing us to routinely view the practically uncharted first 100s to $\\sim 1$ day in the life of a GRB.  In the X-ray data, we see complex phenomena: unexpected rapid decays, also sometimes temporally flat and apparently re-energized afterglows, and massive X-ray flares \\citep[see, e.g.,][]{nousek06}.  Unfortunately, it has proven very challenging to complement these data with an early and rapid-cadence data set at longer wavelengths.  Because studies of GRB afterglows prior to {\\it Swift}~rely primarily on observations in the optical and longer wavelengths, and mostly at later times, contemporaneous observations are critical for connecting new-found insights to the large body of previous work.  Correlated, broadband observations of events which exhibit extremely energetic broadband emission from the radio to the gamma-rays (e.g., Figure \\ref{fig:broadband_spec}) allow us to study and potentially understand the new phenomenology and to pose answers to numerous open questions.  Optical observation of GRBs in the prompt phase and shortly there after are extremely rare.  Of $\\sim 50$ optical afterglows detected prior to {\\it Swift}, only one \\citep[GRB~990123,][]{akerlof99} was detected during the prompt phase and only a handful were detected in the first 10 min of the afterglow: GRB~021004 \\citep{fox03a}, GRB~021211 \\citep{fox03b,li2003b,vestrand04}, and GRB~030418 \\citep{rykoff04}. In the {\\it Swift}~era, the rate of early detections has jumped dramatically, thanks to rapidly communicated and tight GRB localizations from {\\it Swift}~and also due to the maturing system of ground-based robotic telescopes (e.g., $\\sim$15 early detections by ROTSE\\footnote{http://www.rotse.net} alone). The UV Optical Telescope (UVOT) on {\\it Swift}~has the potential to match this performance, particularly with a recent prioritization of early unfiltered observations of the very red afterglows. Perhaps most impressive, the RAPTOR experiment has detected two GRBs during the prompt phase: GRB~041219 \\citep{vestrand05} and GRB~050820A, \\citep{vestrand06}. The prompt and contemporaneous long wavelength emission of GRB~041219A was actually discovered at IR wavelengths \\citep{blake05} with the Peters Automated Infrared Imaging Telescope \\citep{bloom06}. Even though the Galactic extinction toward the GRB was large, GRB~041219A is the only prior burst with prompt, long-wavelength observations at multiple frequencies.  Here we present observations of GRB~051111 conducted with the robotic 0.76-m Katzman Automatic Imaging Telescope \\citep[KAIT;][]{fil01,li2003a,li2003b}, beginning with an unfiltered exposure 43.7 s after the GRB trigger from the Burst Alert Telescope (BAT) on {\\it Swift}~and just catching the tail of the prompt emission. We obtained color information in the form of $V$ and $I$-band observations beginning just 73.7 s after the trigger.  Supplemented with $B$, $V$, $R$, and $I$-band observations with the Lulin One-meter Telescope and $J$, $H$, and $K_s$-band observations with PAIRITEL taken hours after the burst and with other observations reported in the literature, we present an unrivaled and impressive broadband view of this afterglow and its early evolution.  The study of the early emission during and following a GRB is a critical step toward understanding the origin of the emission and the nature of the surrounding medium. It is widely accepted that GRBs are produced by a self-interacting relativistic outflow \\citep[``internal shocks,'' e.g.,][]{fmn96} which heats and shocks the surrounding medium (``external shocks'') to form a long-lasting afterglow at longer wavelengths \\citep{mnr97,snp99}. Transient emission from a ``reverse shock,'' which propagates backward toward the central engine in the shock frame, can also be produced.  Despite an early indication to the contrary from GRB~990123 \\citep[e.g.,][]{akerlof99}, the reverse shock does not appear to be a common feature in the early optical data \\citep[see, also,][]{mkp06}. Primarily due to the larger physical size, the external shocks are expected to generate smoother light curves, and the $\\sim 10$\\% of GRBs with smooth time histories may require only external shocks \\citep[e.g.,][]{mkp04}.  However, the growing sample of early afterglow observations militates against such a simple interpretation.  As mentioned above, the X-ray data \\citep[and also some optical data, e.g.,][]{fox03a,woz05} evidence shock refreshment \\citep{reesNmes98,snm00,rretal01} or continued central engine activity \\citep{reesNmes00,macfad01,rr04,leeNrr02}, which can persist for several hundreds of seconds or longer \\citep[e.g.,][]{nousek06}. While its physical size is still small compared to the timescale for synchrotron cooling, the external shock can cool rapidly. Energy can be efficiently radiated from the shock at this stage \\citep[e.g.,][]{spn98}, prior to the onset of an adiabatic evolution.  Depending on the breadth of the shock shell, which depends on the duration of the period of internal shocks generation, the internal shock emission can overlap with the external shock emission. Below, we use the exquisitely sampled data for GRB~051111 to disentangle the early-time emission components and also to study the dust and gas properties of the host galaxy. ",
        "conclusions": "We have presented a thorough analysis of the rich broadband data available for {\\it Swift} GRB~051111, with a focus on the early, multi-band optical data from KAIT and Lulin.  The optical data prior to $t\\approx 700$ s show a very flat decline, with little or no evidence for a color change as compared to the data after $t\\approx 700$ s.   The data at $t\\gtrsim 1$ minute to several hours after the GRB are well fit using a simple, modified external shock model with absorption by gas and dust.  The modelling entails energy injection or a rapidly increasing density profile with radius, prior to the break time.  At later times there are large ($\\gtrsim 30$\\%), possibly achromatic modulations in the optical, IR, and UV about the best-fit model. The modulations appear to be common in well-studied optical afterglows, yet their origin remains mysterious.  The increasing density model may allow for an external shock explanation of both the prompt gamma-ray and later emission.  Such an increase in density might be expected if a supernova (SN) occurred prior to the GRB, as in the ``supra-nova'' model \\citep{vietri99}.  However, the simultaneity of GRB and SN in the nearby, well-studied cases (GRB~060218/SN~2006aj, GRB~980425/SN~1998bw, GRB~031203/SN~2003lw) argues against this possibility.  On the other hand, slowly declining light curves in the X-ray band are common and are thought to be due to shock refreshment \\citep[e.g.,][]{nousek06}, and the GRB~051111 observations from KAIT show that this emission can also dominate the optical light curve at early times.  In fact, the optical light curve for this event, which shows a rapid early decline, followed by a levelling off and then a moderate decline typical of those found in the past and modelled with external shocks, appears quite similar to the ``canonical'' behavior observed in the X-ray band and reported by \\citet{nousek06}. Perhaps there is a canonical optical afterglow behavior, too.  (As a counter-point, the early optical behavior reported for GRBs 060206 \\citep{monfar06,stanek06} and 060210 \\citep{stanek06} appears quite different from that here.)  An important feature of the GRB~051111 optical afterglow is the lack of a turnover in the optical decay rate at early times.  We do not detect the peak in the synchrotron spectrum passing through the optical bands.  Consonant with a lower than average X-ray flux at late times, this leads to a low value for the fraction (assuming constant equipartition) of shock kinetic energy winding up in the synchrotron-emitting electrons, $\\epsilon_e \\lessim 0.3$\\%. Such a low value for $\\epsilon_e$ is uncommon, but not unheard of, in GRBs \\citep[see, e.g.,][]{pnk02}. The very low value may indicate a wider diversity than previously suspected in the microphysical parameters from GRB to GRB \\citep[see, also,][]{bkf03}. Aside from the low $\\epsilon_e$ value, the derived external shocks parameters are comparable to those previously found.  It is important to stress that whether we find $\\epsilon_e \\approx 0.3$\\% or $\\approx 3$\\% depends on the deconvolution of the early data into prompt internal shock and afterglow external shock components. Our favored smaller $\\epsilon_e$ arises if the synhrotron peak frequency has passed through the optical well before the break at $t\\approx 700$ s, in which case mild energy re-injection into the shock can explain the gradually decaying light curve. To get the larger $\\epsilon_e$ value, the light curve prior to $t\\sim 700$ s must be dominated by the GRB and not by the external shocks which dominate after $t\\sim 700$ s.  Otherwise, the early light curve would be rising rather than declining.  In any case, reverberations of the prompt emission reprocessed into optical radiation \\citep[e.g.,][]{vestrand06} appear to be required to explain the earliest few optical data points for GRB~051111. It does not appear to be necessary to invoke reverse-shock emission \\citep[e.g.,][]{akerlof99,li2003b}, which is expected to produce time decays more rapid than those observed.  The absorption at IR, optical, and UV wavelengths in the observer frame is well fit by an SMC extinction curve.   Little evidence for the 2175 \\AA~dust ``bump'' and an excellent fit of the SMC extinction profile are common features in optical GRB afterglow spectra \\citep[e.g.,][]{vrees04,jakob03,sNf04,watson06}.  Combined with the soft X-ray absorption measurement, there is an implied low dust-to-gas ratio and a possible overabundance of the light-metals relative to the Fe-group metals. This has also been observed for GRB~050401 \\citep{watson06}. The light metal overabundance works against a direct association with an SMC-like environment, because the SMC has a metalicity $\\sim 1/10$ times solar \\citep{pei92}.  The SED fitting also implies a low ratio of total to selective extinction, $R_V\\approx 2$. This is a clue that the absorbing medium exhibits unusual dust properties.  The work of \\citet{galama01} establishing large typical $N_H/A_V$ values has led to suspicions that GRBs could destroy dust out to distances $R\\approx 20$ pc \\citep[see also,][]{wNd00,fruchter01,dNh02,pNl02,perna03}. The small dust grains are preferentially destroyed, creating a flat (or ``gray'') extinction curve \\citep{galama01,sNf04,stratta05}.  A small $R_V$, however, suggests small dust grains are dominant, and our SED is very curved. These facts argue against the GRB playing a direct role in defining the extinction properties. The unchanging optical color implies that the absorbing column is not local to the GRB (i.e., $R\\gtrsim 0.1$ pc) and may be associated with a nearby giant molecular cloud or the GRB host galaxy. Time-resolved spectroscopy in the optical and X-ray bands is of utmost importance for answering these questions.  For GRB~051111, as for a few other {\\it Swift}~events for which ground-based observers have been fortunate enough to capture early data, there is an emerging complicated interplay between the GRB and the subsequent shocking of the external medium.  It is critical that more long-wavelength data be taken and published for other GRBs in order to facilitate modelling similar to that performed above. Only with such a disentangling of the competing emission processes can we hope to answer open questions regarding which shock components are truly the most important and what microphysical parameters define the shocks and characterize the transition from internal to external shocks."
    },
    "0606/astro-ph0606280_arXiv.txt": {
        "abstract": "{ We examine the conditions of the plasma along a sample of 'classical' Herbig-Haro (HH) jets located in the Orion and Vela star forming regions,  through combined optical-infrared spectral diagnostics. Our sample includes HH 111, HH 34, HH 83, HH 73, HH 24 C/E, HH 24 J, observed quasi-simultaneously and in the same manner at moderate spatial/spectral resolution. Once inter-calibrated, the obtained spectra cover a wide wavelength range from 0.6 - 2.5 $\\mu$m, including many transitions from regions of different excitation conditions. This allows us to probe the density and temperature stratification which characterises the cooling zones behind the shock fronts along the jet. From the line ratios  we derive  the variation of the visual extinction along the flow, the electron density and temperature ($n_e$ and $T_e$), the hydrogen ionisation fraction $x_e$, and the total density $n_H$ in the emission region of different lines. The knowledge of such parameters is essential for testing existing jet models and for planning follow-up high-angular resolution observations.  From the diagnostics of optical forbidden lines we find, on average, that in the examined jets, in the region of optical emission, $n_e$ varies between 50 cm$^{-3}$ and 3 10$^{3}$ cm$^{-3}$, $x_e$ ranges between 0.03 and 0.6, and the electron temperature $T_e$ is $\\sim$1.3 10$^{4}$ K in the HH 111 and HH 34 jets, while it appears to be higher (1.8 10$^{4}$ K on average) in the other examined jets. The electron density and temperature derived from \\fe\\, lines, turn out to be, respectively, higher and lower in comparison to those determined from optical lines, in agreement with the fact that the \\fe\\, lines arise in the more compressed gas located further from the shock front. An even denser component in the jets, with values of $n_e$ up to 10$^{6}$ cm$^{-3}$ is detected using the ratio of Calcium lines.  The derived physical parameters are used to estimate the depletion onto dust grains of Calcium and Iron with respect to solar abundances. This turns out to be  quite substantial, being between  70\\% and 0\\% for Ca and $\\sim$90\\% for Fe. This leads us to suggest that the weak shocks present in the beams are not capable of completely destroying the ambient dust grains, confirming previous theoretical studies. We then derive  the mass flux rates, $\\dot{M}_{\\rm jet}$, in the flows using two independent methods. Taking into account the filling factor of the emitting gas, $\\dot{M}_{\\rm jet}$ is on average 5 10$^{-8}$ M$_\\odot$ yr$^{-1}$. The associated linear momentum fluxes ($\\dot{P}_{\\rm jet} = v_{jet} \\dot{M}_{\\rm jet}$) are higher than, or of the same order as, those measured in the coaxial molecular flows, where present, suggesting that the flows are jet driven.  Finally, we discuss differences between jets in our sample. In general, we find that higher ionisation and electron temperatures are associated with  less dense jets. The comparison suggests that the shock mechanism exciting the knots along the flows has the same efficiency in all the examined objects, and the observed differences are consistent with the different densities, and hence cooling rates, found in the various flows. ",
        "introduction": "Spectral analysis is a powerful tool to investigate the nature and properties of astrophysical nebulae excited by shocks and/or energetic radiation. The diffuse matter in such objects produces a wealth of permitted and forbidden emission lines, whose excitation properties are reasonably well known. This in principle  allows us to go back from the observed line intensities and line  ratios to the physical conditions of the emitting material, which in turn provide essential information to test theoretical models and plan challenging observational programs, e.g. those involving interferometry. For quite a long time line diagnostics have been used mainly to determine the gas electron density ($n_e$) and electron  temperature ($T_e$) (see, e.g. Osterbrock 1994). The spectra of line emitting regions, however, contain much more information and the development of a number of new diagnostic techniques allow us to put, directly or indirectly, new important constraints on the gas physics. The study  of stellar jets is one of the fields that has benefitted enormously from progress in spectral techniques. These collimated flows appear right from the birth of a star and indeed are believed to play a central role in the star formation process itself (see, e.g., \\cite{eisloffel00}, \\cite{reipurth01}, \\cite{ray03}, \\cite{bacc04}). In fact, jets can help clear the circumstellar environment, thus setting a limit to the final central mass, and may even extract most of the excess angular momentum from the star/disk system (\\cite{bacc02}, \\cite{coffey04}, \\cite{woitas05}). Prototypical cases are the HH 34 and HH 111 jets in Orion (\\cite{reipurth01}), also studied in this work. All the observed features in stellar jets, such as bright knots along the flow and giant bow shocks, are believed to trace shocks that develop within these highly supersonic flows. Passing through the shock front, part of the bulk kinetic energy of the flow is turned into thermal motions, and the gas is suddenly heated to high temperatures (10$^5$ K or more), compressed, and partially ionised. This situation favours the collisional excitation of upper levels of atomic transitions. A number of radiative lines, whose type and strength are characteristic of shock excitation, are produced. These, in turn, radiatively cool down the gas to its original temperature. The whole process takes place in a limited region of space behind the shock front called the 'cooling zone', whose length may vary from 10$^{13}$ to about 10$^{15}$ cm, depending on the shock strength, the pre-shock conditions and the composition of the gas. The structure of radiatively cooling shocks has been widely studied in the past by several authors (see e.g. \\cite{hartigan87}, \\cite{hollenbach89}, \\cite{hartigan94}, \\cite{hollenbach97}, \\cite{flower03}, \\cite{hartigan03} and references therein). Within the cooling region individual physical quantities vary enormously, creating a stratification of excitation conditions that results in the production of different kinds of lines. In ground based observations at moderate spatial resolution, such as those presented here, the internal structure of the shock cooling zone remains spatially unresolved. Nevertheless, the stratification of such regions can be investigated spectrally by analysing the many different lines arising from different layers.  In the  optical and Near Infrared (NIR) ranges used in this work one observes forbidden and permitted  lines of abundant atomic and molecular species, such as H, He, O, S, N, Fe, C, Ca, H$_2$  (\\cite{reipurth01}). To extract more information from these lines, several attempts have been made to use novel techniques by various groups. B\\\"{o}hm et al. (1980), Brugel et al. (1981) were among the first to examine the potential of combining the observed line ratios using low spatial resolution spectra integrated over the shock cooling zone. Hirth et al. (1997) examined long-slit spectra to study the spatial and kinematical properties of the forbidden emission line regions and micro-jets of a large sample of T Tauri stars. Intensities of atomic lines  have been compared with the prediction of shock models in, e.g., Raga \\& B\\\"ohm (1986), Hartigan et al. (1987), and Hartigan et al. (1994), indirectly inferring quantities such as the hydrogen ionisation fraction $x_e$ and the  intensity of the magnetic field in the pre-shock region. A new simple technique to measure $x_e$ and $T_e$ from the ratio of optical  lines was first presented by Bacciotti, Chiuderi \\& Oliva (1995), and subsequently refined in Bacciotti \\& Eisl\\\"offel (1999) (hereafter BE99). The method, referred to as the `BE' technique, is based on the fact that the gas emitting forbidden lines is collisionally excited, but no assumption is made regarding the heating agent. This is clearly  an advantage if the results are used to validate a given thermal model. On the other hand, the method assumes that the emitting gas is at a single temperature, which is not true in the cooling region behind  a shock front. In this regard, one should consider that the results of the BE technique are relevant to the region behind the shocks in which the employed lines have their peak emission (see the discussion and diagrams in BE99). As we show in this paper, the stratification of temperature and densities present in the entire cooling zone can be traced using a larger sample of lines appropriate for different excitation conditions.   A number of jets have been analysed with the BE technique, leading to the finding that jets are only partially ionised, with average $x_e$ values between 0.01 and 0.6. This procedure is much easier to apply than a grid of shock models (although it is of more limited application), and thus well suited for the analysis of big datasets, as those provided by high angular resolution observations (\\cite{bacc_eisl_ray99}, \\cite{bacciotti02}). The  main advantage in determining $x_e$, independently from the gas heating mechanism, is that by combining its value with the derived electron density one can estimate the total density, n$_H$, a fundamental parameter of the jet that is critical in the various models.  Diagnostic techniques using both optical and NIR lines, on the other hand, provides important complementary information on jet parameters. In particular, the combination of optical and NIR \\fe\\, lines gives an independent tool to determine T$_e$ and n$_e$, which does not rely on the choice of elemental abundances, in the more compressed and cooler post-shock layer where these lines are excited (\\cite{nisini02}, \\cite{pesenti03}, \\cite{hartigan04}). In addition, the NIR H$_2$ lines provide a means to probe the molecular component of shock excited gas, which may give a significant contribution to the gas cooling in low velocity, magnetized shocks (\\cite{eislsmithdavis00}, \\cite{giannini04}).  Very recently, we have investigated the advantages of a combined optical and NIR spectral analysis of the HH 1 jet (\\cite{nisini05}) using a variety of diagnostic lines over a wide wavelength range (from 0.6 to 2.2 $\\mu$m). The adopted procedure turned out to be extremely powerful for constructing detailed physical maps of the stratified media in the beams of stellar jets. In this paper,  we analyse a large sample of 'classical' stellar jets,  located in the Orion and Vela star formation regions, with the same technique. We obtain their basic parameters, as well as important quantities derived from them, such as the mass flux rates. We analyse the variations of the parameters both behind each shock, i.e. depending on the tracers, and along the jet, i.e. as a function of the distance from the emitting source. We then discuss the differences/similarities among jets having different properties and ages. This information can in turn be used for the selection of suitable candidates for high angular resolution observations.   This paper is organized as follows: we describe our observations in Sect. 2 and briefly recall our diagnostic procedure in Sect. 3, adding a detailed discussion on the choice of the best set of elemental abundances, an issue that was not examined in our previous papers. In Sect. 4 we describe the results obtained for each target in the sample, and we discuss them in Sect. 5 together with the derivation of the depletion of refractory elements and the mass and linear momentum flux rates. Sect. 6 summarizes our findings. ",
        "conclusions": "The application of our combined optical/NIR diagnostics to a sample of protostellar jets allows us to define their physical structure, to infer important parameters related to their dynamics such as their mass and linear momentum fluxes, and to investigate the acceleration of associated molecular outflows, as well as the mechanism of dust reprocessing. For the HH 111 and HH 34 jets, which are detected both in the optical and in the near infrared, we obtain the following results:  \\begin{itemize}  \\item[-] In HH 34 the visual extinction, $A_{\\rm V}$, derived from \\fe\\, lines varies between $A_{\\rm V}$$\\sim$7.1 mag, at a distance of 3$''$ from the source, to 1.3 mag, from 5$''$ to 30$''$ from the source. Along HH 111 the visual extinction is $A_{\\rm V}$$\\sim$2 mag between 15$''$-30$''$ from the source  and negligible in the outer knots.  \\item[-] In both jets the  electron density from the \\s\\, lines varies between  0.5-3.4 10$^{3}$ cm$^{-3}$, while in the more compressed region traced by \\fe\\, lines we find higher electron density (n$_e$$\\sim$ 1-5 10$^{3}$ cm$^{-3}$). An even denser layer with electron densities up to 10$^{6}$ cm$^{-3}$ is traced by [\\ion{Ca}{ii}] lines. Such a density stratification is expected in the post-shocked gas of each unresolved knot.  \\item[-] A temperature stratification is also found: the temperature derived from optical S$^{+}$, N$^{+}$ and O lines is on average  1.3 - 1.4 10$^{4}$ K, while in the region of Iron emission the  temperature is always lower than 10$^{4}$ K.  \\item[-] In HH 111, which is a strong \\h\\, emitter, we also derive the physical conditions of the molecular component. These lines probably trace the lateral wings of the bow shocks associated with the knot working surfaces, where \\h\\, post-shock temperatures of 2000 - 3000 K are found.  \\item[-] The ionisation fraction derived through the BE technique demonstrates that the gas in the jet is only partially ionised (x$_e$ $\\sim$ 0.03$\\sim$0.3). The derived total density (n$_e$/x$_e$) is about 0.1-3 10$^{4}$ cm$^{-3}$ in the region of optical emission. If we assume the same ionisation fraction in the more compressed regions where \\fe\\, and [\\ion{Ca}{ii}] lines are excited, total densities up to 10$^{7}$ cm$^{-3}$ are obtained in the dense layers traced by these lines.  \\item[-] Estimates of the gas-phase abundance of carbon, calcium and iron are obtained. We find that carbon is not depleted with respect to solar values, while calcium and iron are strongly depleted (the depletion is $\\sim$ 87\\% for iron atoms and varies between 70\\% and 0\\% for calcium). On the other hand, the estimated depletion is lower than the one found in the Orion Cloud by Baldwin et al. (1991) and Esteban et al. (2004) both for Calcium and Iron. This result demonstrates that weak shocks only partially destroy dust in the jet and that some grains still survive. In fact a velocity of 100-400 km\\,s$^{-1}$, like that of a primary bow-shock, is required to completely destroy the dust (\\cite{draine03}).  \\item[-] Taking into account the filling factor of the gas in the knots, we derive mass flux rates of $\\sim$ 4-5 10$^{-8}$ M$_{\\odot}$\\,yr$^{-1}$. In HH 111 the flux of mass transported by the  molecular \\h\\, component is two orders of magnitude lower than the atomic one.  \\item[-] The estimated flux of linear momentum is on average $\\sim$ 10$^{-5}$ M$_{\\odot}$\\,yr$^{-1}$\\,km\\,s$^{-1}$. This is higher than, or comparable to, the flux of momentum measured for the associated CO molecular outflows, and indicates that jets can drive these molecular outflows.  \\end{itemize}   For the HH 83, HH 73 and HH 24 C/E jets we could only apply optical diagnostics, obtaining the following results:  \\begin{itemize} \\item[-] The electron densities inferred from the \\s\\, doublet are very low (n$_e$ $<$ 10$^{3}$ cm$^{-3}$). This explains why these HH jets are not detected in the NIR. In fact, the transition of \\fe\\, lines comes from levels which have a critical density higher than 10$^{4}$ cm$^{-3}$.  \\item[-] The ionisation fraction, on the other hand, is always higher in these jets (x$_e$ up to 0.6), consistent with the fact that for a given shock velocity a higher ionisation is produced when the shock propagates into a low density medium, and with the fact that recombination is slowed down at low electron densities. The latter can also be partially responsible for the higher temperatures in these objects (T$_e$ $\\sim$ 1.4 - 3 10$^{4}$ K), because the collisional cooling process is less efficient at low density. We suggest that the efficiency of the excitation mechanism is the same for all of the jets in our sample, and the differences in the physical conditions of the gas are due to the fact that the HH 111 and HH 34 jets are denser.  \\item[-] The total density is quite low (n$_H$ $\\sim$ 10$^{3}$ cm$^{-3}$). From these values we infer mass fluxes of $\\sim$ 8 10$^{-8}$ M$_{\\odot}$\\,yr$^{-1}$. Note that these values are upper limits because we are not taking into account the filling factor of the emitting gas for these jets.  \\item[-] The linear momentum fluxes are of $\\sim$ 1-4 10$^{-5}$ M$_{\\odot}$\\,yr$^{-1}$\\,km\\,s$^{-1}$. \\end{itemize}   The above results demonstrate the potential of our combined optical/NIR analysis to study the jet physical structure and represent a good basis to plan future higher spectral-spatial observations with VLT-CRIRES/UVES and AO or interferometric instruments. A higher spectral resolution, in fact, would allow us to analyse the kinematic structure of the jets, while with higher spatial resolution we could extend our analysis to the base of the jet, where it is launched and accelerated, testing the theoretical MHD models."
    },
    "0606/astro-ph0606249_arXiv.txt": {
        "abstract": "In the present work we study a dark energy model in which a non-linear scalar field (tachyon) with a Born-Infeld type of action is responsible for the observed cosmic acceleration. The potential of the tachyon is well-motivated since it comes from open string theory and the model is subjected to various cosmological constraints with data coming from supernovae as well as from microwave background radiation. Our analysis shows that in the particular model under study the tachyon can be an excellent candidate for dark energy in the universe, as the model agrees with a series of observational data and for a wide range of the parameters of the model. ",
        "introduction": "Perhaps the most dramatic discovery in modern cosmology is the observational fact that the present universe undergoes an accelerating phase. During the last ten years or so a remarkable progress in technology has allowed us to witness extraordinary precision measurements in cosmology. A plethora of observational data are now available, which show that we live in a flat universe that expands with an accelerating rate and that the dominant component in the energy budget of the universe is an unusual material, the nature of which still remains unknown. Identifying the origin and nature of dark energy is one of the great challenges in modern theoretical cosmology. The simplest candidate for dark energy is the cosmological constant, which corresponds to a perfect fluid with parameter $w=p/\\rho=-1$. However, over the years many other theoretical models have been proposed and studied. One class of such models is based on some modification of Einstein's gravity~\\cite{gde1,gde2} and one is talking about the so called geometrical dark energy models. Another class contains the dynamical dark energy models, in which a new dynamical field (almost certainly a scalar field) is coupled to gravity. In the second class one would find models called quintessence~\\cite{quintessence}, phantom~\\cite{phantom}, quintom~\\cite{quintom}, k-essence~\\cite{kessence}, tachyonic~\\cite{tachyonic} etc. A recent review on dark energy dynamics one can find in~\\cite{copeland}.  Since Sen's original publications~\\cite{sen}, much work has been done on tachyonic cosmology, concerning either inflation~\\cite{tachinfl} or dark energy~\\cite{tachyonic}. In the present work our aim is to test a particular tachyonic dark energy model against observational cosmological data that are available from supernovae and microwave background radiation. A very recent work~\\cite{liddle} is closely related to ours and hopefully the present analysis will be a complementary one, at least for the particular model that we shall study here. In~\\cite{liddle} the authors studied a wide range of different tachyonic models. In one of these models the potential for the tachyon emerges from a class of open string theories (both bosonic and superstrings) with explicit computation. That way string theory can be linked to cosmology and provide a natural candidate for playing the role of dark energy in the universe. This is exactly the potential that we shall be using in what follows.  The present work is organized as follows: There are four sections of which this introduction is the first. The theoretical framework is presented in section 2, while our analysis and results are presented in section 3. Finally we conclude in the fourth section. ",
        "conclusions": "In the present work we have studied a dark energy model with a tachyon playing the role of dark energy in the universe. The tachyon is non-linearly coupled to gravity, with a Born-Infeld type of action. Its equation of motion is different than that of a Klein-Gordon scalar field, which couple linearly to gravity. However, here also the homogeneous tachyon behaves like a perfect fluid with an energy density and pressure that are given by certain expressions in terms of the potential and the time derivative of the tachyon. The potential for the tachyon that we have used emerges from explicit computations in open string theories (bosonic string or superstrings), and therefore is well-motivated. We have tested the particular model against a series of cosmological constraints coming from observational data both from supernovae and microwave background radiation. Our detailed analysis shows that in the model under consideration the tachyon can be an excellent candidate for dark energy, as the model passes successfully all the aforementioned tests for a wide range of the parameters of the model."
    },
    "0606/astro-ph0606005_arXiv.txt": {
        "abstract": "It is shown that Fermi acceleration at an ultra-relativistic shock wave cannot operate on a particle for more than 1 {\\small 1/2} Fermi cycle (i.e., $\\rm u\\rightarrow d\\rightarrow u\\rightarrow d$) if the particle Larmor radius is much smaller than the coherence length of the magnetic field on both sides of the shock, as is usually assumed. This conclusion is shown to be in excellent agreement with recent numerical simulations. We thus argue that efficient Fermi acceleration at ultra-relativistic shock waves requires significant non-linear processing of the far upstream magnetic field with strong amplification of the small scale magnetic power. The streaming or transverse Weibel instabilities are likely to play a key r\\^ole in this respect. ",
        "introduction": "Fermi acceleration of charged particles bouncing back and forth a collisionless shock wave is at the heart of a variety of phenomena in high energy astrophysics. According to standard lore, this includes the acceleration of electrons at gamma-ray bursts internal/external relativistic shock waves, whose synchrotron light is interpreted as the prompt/afterglow radiation.  However the inner workings of Fermi acceleration at relativistic shock waves remain the subject of intense study and debate, even in the test particle limit. It has been argued that a universal energy spectral index $s\\simeq 2.2-2.3$ should be expected (e.g., Bednarz \\& Ostrowski 1998, Achterberg {\\it et al.} 2001, Lemoine \\& Pelletier 2003, Ellison \\& Double 2004, Keshet \\& Waxman 2005), which would agree nicely with the index that is inferred from gamma-ray bursts observations. Yet recent numerical simulations that include more realistic shock crossing conditions have indicated otherwise (e.g., Lemoine \\& Revenu 2006, Niemiec \\& Ostrowski 2006), so that the situation is presently rather confuse. Unfortunately, numerical Monte-Carlo simulations of Fermi acceleration, although they remain a powerful tool, do not shed light on the physical mechanisms at work.  In the present {\\em Letter}, we offer a new analytical discussion of Fermi acceleration in the ultra-relativistic regime. We rely on the observation that, under standard assumptions, the Larmor radius $r_{\\rm L}$ of freshly injected particles is much smaller than the coherence length $l_{\\rm coh}$ of the upstream magnetic field, and we integrate the equations of motion to first order in the quantity $r_{\\rm L}/l_{\\rm coh}$ (Section 2). We thus find that a particle cannot execute more than 1 1/2 $\\rm u\\rightarrow d\\rightarrow u$ cycles through the shock before escaping downstream, hence Fermi acceleration is mostly inoperative. This result is related to the very short return timescale in the ultra-relativistic regime: over its trajectory, the particle experiences a nearly coherent and mostly transverse magnetic field, hence the process is akin to superluminal acceleration in regular magnetic fields discussed by Begelman \\& Kirk (1990). We also show that our analytical predictions agree very well with numerical simulations. Finally, we argue that the (expected) non-linear processing of the magnetic field at ultra-relativistic shock waves and, in particular, the strong amplification of small-scale power, may be the agent of efficient Fermi acceleration (Section 3); we suggest new avenues of research in this direction. ",
        "conclusions": "Fermi acceleration is thus inefficient at ultra-relativistic shock waves if the Larmor radius $r_{\\rm L}$ of injected particles is much smaller than the coherence length $l_{\\rm coh}$ of the turbulent magnetic field on both sides of the shock. This does not mean that Fermi acceleration is bound to fail. In particular, if $r_{\\rm L}\\gg l_{\\rm coh}$, powerlaw spectra must emerge as the memory of the magnetic field direction (in the transverse plane) at shock crossing is erased during propagation in small scale turbulence (as we have checked numerically). The final value of the spectral index will depend on the transport properties of the particle in this small-scale turbulence. It is difficult to probe this regime using numerical simulations as integration timescales become large (the scattering timescale $\\gg r_{\\rm L}/c$).  On analytical grounds, one expects $s\\simeq 2.3$ if the small scale turbulence is isotropic downstream of the shock wave (Keshet \\& Waxman 2005); if it is anisotropic, one ought to expect a different value however.  Although Fermi acceleration could operate efficiently on high energy initial seed particles with $r_{\\rm L}\\gg l_{\\rm coh}$, the abundance of such particles is generally so low in realistic astrophysical shock wave environments that the injection efficiency would be extremely small (see Gallant \\& Achterberg 1999 for instance).  In the above context, the interpretation of the afterglow emission of $\\gamma-$ray bursts as the synchrotron radiation of electrons accelerated at the ultra-relativistic external shocks ($\\Gamma_{\\rm sh\\vert u}\\ga 100$) becomes particularly enlightening. The success of this model indeed requires both efficient Fermi acceleration as well as very significant amplification of the interstellar magnetic field (Gruzinov \\& Waxman 1999).  Hence it is tempting to tie these two facts together and to wonder whether this non-linear MHD processing could not be the agent of efficient Fermi acceleration.  One proposal discussed so far is the transverse Weibel instability which could possibly produce sufficiently strong magnetic fields on very small spatial scales \\\\ $\\sim 10^5\\,{\\rm cm}\\,(\\Gamma_{\\rm sh\\vert u}/10)^{-1/2}(n_e/1\\,{\\rm cm}^{-3})^{-1/2}$ (Medvedev \\& Loeb 1999); there is however ongoing debate on the lifetime and strength at saturation of the magnetic field (e.g., Wiersma \\& Achterberg 2004; Lyubarsky \\& Eichler 2006). Nonetheless, such a small-scale turbulence should result in powerlaw spectra of accelerated particles; albeit the value of the resulting spectral index is not known.  Recently, it has been suggested that the generalization of the streaming instability to the relativistic regime could amplify the magnetic field to the values required by $\\gamma-$ray bursts observations (Milosavljevic \\& Nakar 2006).  Due to the very short upstream return timescale $\\sim r_{\\rm L}/\\Gamma_{\\rm sh\\vert u}$, the particle can never stream too far ahead of the shock so that the turbulence is generated on small scales $\\sim 10^7-10^8\\,{\\rm cm}\\,\\ll r_{\\rm L}$ (Milosavljevic \\& Nakar 2006). Therefore, one naturally expects in this case too that Fermi acceleration would be efficient, here as well, one needs to understand the turbulence properties before conclusions can be drawn on the index $s$.  To summarize, we have shown in Section 2 that Fermi acceleration cannot operate successfully at ultra-relativistic shock waves if one assumes (somewhat na\\\"\\i vely) large-scale turbulence on both sides of the shock wave. The conclusions of the present discussion are thus more optimistic and open a wealth of new possibilities; in particular they suggest that the success of Fermi acceleration is intimately connected with the mechanism of magnetic field amplification in the shock vicinity. The comprehension of Fermi acceleration will eventually require understanding the generation of the magnetic field, deriving the properties of the turbulence as well as characterizing the transport of accelerated particles in this possibly anisotropic turbulence.\\smallskip  \\noindent{\\bf Note added:} while this work was being completed, a recent preprint by Niemiec {\\it et al.} (2006) appeared, reporting on Fermi acceleration with small-scale turbulence. Although their simulations are limited to $\\Gamma_{\\rm sh\\vert u}= 10$, these authors observe that the inclusion of small scale turbulence allows powerlaw spectra to emerge through Fermi acceleration, in good agreement with the above discussion."
    },
    "0606/astro-ph0606719_arXiv.txt": {
        "abstract": "The effect of gravitational microlensing on the determination of extragalactic distances using surface brightness fluctuations (SBF) technique is considered and a method to calculate SBF amplitudes in presence of microlensing is presented. With a simple approximation for the magnification power spectrum at low optical depth the correction to the SBF-based luminosity distance is calculated. The results suggest the effect can be safely neglected at present but may become important for SBF-based Hubble diagrams at luminosity distances of about $1\\,\\mathrm{Gpc}$ and beyond. ",
        "introduction": "For many centuries, measuring distances to celestial objects has been a central part of astronomy. Historically, building on the trigonometric parallax technique, astronomers have developed a long cosmic distance ladder, within which every step is supplied by observing more and more luminous objects further and further away, whose calibration is based on distances obtained at preceding steps.  About two decades ago a new method to determine extragalactic distances using surface brightness fluctuations (SBFs) was proposed by \\citet{tonryschneider}. The method relies on a simple but powerful idea that less numerous stellar populations show greater relative fluctuations in their total brightness as Poissonian fluctuations in the number of stars are more pronounced in them compared to more numerous ones. This allows one to estimate the true average number of stars in and therefore intrinsic luminosities of pixels in images of distant galaxies. Comparing these estimates with the observed flux then estimates the luminosity distance directly. Although the idea of the method was clear to some astronomers long ago (see comments in \\citealt{tonryschneider}), it was not until the mass introduction of linear panoramic detectors (CCDs) into astronomical observations, that the method could be usefully employed.  Since then, SBF technique has proven extremely fruitful in independent distance determination for galaxies and also in studies of their stellar populations (for review material, see \\citealt{blakeslee98} or \\citealt{blakeslee01}). Various ways for the calibration of surface brightness fluctuations -- both empirically and on the basis of stellar evolution simulations -- have been developed, allowing to reduce uncertainty in distance determination to the level of five to ten per cent, making it one of the most precise methods in extragalactic astronomy. The range of the method has also increased nearly ten-fold compared to original estimates, and the distances to galaxies at nearly 200 Mpc have been successfully measured using the {\\it Hubble Space Telescope} \\citep{jensen01}.  Since the major limiting factor in the implementation of the method is photon noise, the trend in SBF observations has been towards the infrared for at least a decade, as one gets more infrared photons at a given flux level and the SBF signal is dominated by late-type giant stars. In addition, IR surface brightness fluctuations turned to be very interesting for stellar populations studies (e.g., \\citealt{jensen03, mouhcinegonzalezliu}).  One of the external factors which can potentially contribute to the fluctuations of stellar magnitudes is gravitational microlensing by compact objects in the Universe. The presence of such population has been inferred from long-term photometrical monitoring of quasars \\citep{hawkinsnature, hawkins96}. The effect of such population on photometrical and spectroscopical properties of quasars has been considered previously \\citep{canizares82, dalcanton94} and useful observational constraints on the properties of such population have been obtained \\citep{schneider93}.  This effect has been proposed as a tool to detect dark objects in galaxy clusters by observing {\\it temporal} variations in the flux of a given pixel \\citep{liw, lewisibata, ourcluster}. However, microlensing also introduces extra scatter in flux variations {\\it among different pixels} and makes the distances inferred via SBF method systematically lower than the true ones. The correction due to microlensing increases with redshift dramatically, and the compact objects therefore have a potential to profoundly affect SBF-based Hubble diagrams at high $z$.  In the present paper, we investigate how strong this effect is and whether it represents a serious issue for the present-day SBF observations. Borrowing some technique from our previous work on galaxy clusters \\citep{ourcluster}, we show how microlensing can be taken into account for SBF studies, and then calculate relevant quantities under some general assumptions about potential microlensing population.  We start in Section 2 with an alternative derivation of the SBF amplitude which allows to take microlensing into account. Section 3 presents a simple summation approximation we use for the magnification spatial power spectrum and discusses a possible contribution to microlensing by an occasional intervening low surface brightness galaxy on the line of sight to the source. Then, in section 4 we first show that the interstellar correlations within the pixel can be neglected and then introduce and calculate the correction to SBF distances due to microlensing for simulated stellar populations. Section 5 discusses our results. ",
        "conclusions": "In the present paper we have considered the correction to distances inferred from the surface brightness fluctuations analysis, caused by a potential population of compact objects along the line of sight to the source galaxy. Using a simple approximation for microlensing magnification fluctuations we have studied how this effect can be taken into account and whether such accounting is needed in observations. Although we have considered microlensing as the main factor of individual star brightness fluctuations, the technique developed should be applicable to a wider range of geometrical variability.  By combining the results shown in figure~\\ref{qfig} with the upper limit of $\\sim 0.3$ on the total fraction of matter in the energy budget of the Universe, one can see that microlensing can for now be safely neglected when interpreting SBF-based distances. The present-day observational facilities limit the range of the method to less than 200 Mpc while the uncertainty of this method is presently between five and ten per cent. The finite-sample effect discussed at the end of Section 2 further reduces the magnitude of the lensing correction.  However, for the next generations of telescopes, this situation can -- at least, in principle -- change. The main factor limiting the accuracy of the surface brightness fluctuations method is the photon noise, which is of geometrical origin. Therefore, a naive expectation would be that increasing the diameter of the telescope $n$-fold would lead to a directly proportional increase in the range in terms of the comoving distance $c/H_0 r(z)$. Taken with the estimates given in this study, this implies that microlensing might just start to matter when space telescope diameters get closer to a ten-metre scale.  Alternatively, one might notice that $q$ turned to be of the same order as the total optical depth. Therefore, any large overdensity of compact microlenses -- such as those potentially associated with galaxy clusters -- would introduce noticeable change in the surface brightness fluctuations statistics and can thus be detected in this manner. However, the properties of strong lensing mapping need to be known in full detail for such project since local variations of the apparent luminosity distance are greatest near the critical curves of the mapping which yield galaxies most suitable for SBF studies. Temporal variability, therefore, seems to be a better tool to study compact dark matter in galaxy clusters \\citep{ourcluster}.  We conclude that microlensing by a cosmological population of compact objects is unlikely to significantly bias determination of distances using surface brightness fluctuations technique. Similarly, present accuracy of this method does not allow to put useful constraints on the properties of such population from the absence of systematic discrepancies between distances determined via SBF and other methods. The effect considered in this paper, however, should not be overlooked when the accuracy and range of the SBF method reach the level of a few per cent and a gigaparsec, respectively.  {\\it Acknowledgements.} AVT is supported by IPRS and IPA from the University of Sydney. We thank Santi Cassisi, Daniel Cordier and the BaSTI team for performing population synthesis simulations for our study. We thank the referee for their suggestion to extend the applications of this work and numerous useful comments."
    },
    "0606/astro-ph0606233_arXiv.txt": {
        "abstract": "{}{We report the theoretical spectral energy distributions (SEDs), Br$\\gamma$ line profiles, and visibilities for two scenarios that can explain the disk dissipation of active hot stars and account for the transition from the Be to the B spectroscopic phase.} {We use the SIMECA code to investigate two scenarios, the first one where the disk is formed by successive outbursts of the central star. A low-density region is developing above the star and slowly grows outward and forms a ring-like structure that will gradually excavate the disk. The second one has a slowly decreasing mass loss due for instance, to a decrease in the radiative force through an opacity change at the base of the photosphere, and may also be responsible for the vanishing of the circumstellar disk.} {We find that a clear signature of the disk dissipation following the ring scenario will be the disappearance of the high velocity tails in the emission lines and a nearly constant peak separation. Moreover, we found that, following the ring-like scenario, the visibilities must show an increasing second lobe, an increase in the value of the first zero, and assuming an unresolved central star, a first zero of the visibility curves that appends at shorter baselines as far as the disk has been excavated. We propose to use the AMBER instrument on the VLTI to probe whether the ring scenario is the one to rule the Be phenomenon.}{} ",
        "introduction": "According to advances in the study of stellar mass-loss phenomena since the discovery of Be stars, different kinds of models have attempted to describe their circumstellar envelope (CE) or disk formation. Although they coexist, we may attempt to divide the mass-ejection phenomena into two broad classes: 1) continuing mass ejections and 2) discrete mass ejections. Even though one or the other of these mass-loss events might be prevalent, somewhat different types of CE structures could be envisioned:\\par \\begin{figure}[t] \\begin{center} \\includegraphics[height=8.cm]{4690F001.epsf} \\end{center} \\caption{Isovelocity intensity maps computed with the SIMECA code, seen with an inclination angle of 45$\\degr$, and showing the effect of the disk kinematics (mainly Keplerian rotation) as a function of wavelength, 21636, 21641, 21646, 21651, 21656, 21661, 21666, 21671, and 21676 $\\AA$ across the Br$\\gamma$ line within a spectral bandwidth of 5 $\\AA$.  } \\label{visi_baseline_radius} \\end{figure}  1) \\textsc{Continuing mass ejections}. The most frequently suggested mechanisms of CE formation evoke continuing mass ejections. Struve (1931, 1942) imagined Be stars as secularly stable critical rotators that build a rotating extended envelope flattened towards the equator through equatorial ejection of mass(Limber (1964, 1967, 1969). Somewhat modified by the inclusion of an expansion velocity field, these CE structures were used by Marlborough (1969) and Poeckert \\& Marlborough (1987a, b) to predict observable spectroscopic, photometric, and polarimetric quantities. Thenceforth, several ad hoc flaring and flattened disc-shaped empirical CE models were explored in the literature to see if the CEs were outwardly accelerated, decelerated, and/or rotationally supported (Waters 1986; Waters et al. 1987; Kastner \\& Mazzali 1989; Hummel \\& Dachs 1992; Hummel 1994; Hanuschik 1995, 1996; Wood et al. 1997). Thomas \\& Doazan (1982) proposed an active spherical exophotospheric structural pattern, where the low-density outwardly accelerated hot chromospheric-coronal zone is followed by a cool and decelerated high density region. However, interferometric and polarimetric observations, as well as statistical inferences from spectroscopy, support some flattening of CE in Be stars (Arias et al. 2006; Quirrenbach et al. 1997; Tycner et al. 2005). In the frame of theoretical inferences of CE properties based on continuing mass ejections, the thinness of circumstellar discs in Be stars still remains an open issue (Ara\\'ujo \\& Freitas Pacheco 1989; Bjorkman \\& Cassinelli 1993; Chen et al. 1992; Owocki et al. 1996; Poe \\& Friend 1986; Porter 1997;  Stee \\& Ara\\'ujo 1994; Stee et al. 1995, 1998). Moreover, from physical first principles (Schmitz 1983) and the kinetic theory of gases applied to steady-state isothermal self-gravitating gaseous masses with cylindrical differential rotation, Rohrmann (1997) predicts bagel-shaped CE that are not reminiscent of flat discs.\\par 2) \\textsc{Discrete mass ejections}. In addition to these continuing mass ejections, there are discrete mass ejections. They can be divided roughly into two types: 1) small-scale to moderate ejections, detected either as sharp absorption features crossing the line profiles or as a sudden appearance of emission shoulders in the wings of some absorption lines and/or H$\\alpha$ line emission outbursts that imply ejecta with masses $\\Delta{\\rm M}\\!\\la$ $10^{-11}$M$_{\\odot}$ (cf. Floquet et al. 2000); 2) Large-scale or massive ejections. A clear example is $\\omega$~Ori, where a sudden light increase of nearly $\\Delta V\\!\\approx$ $-0.3$ mag was reported by Guinan \\& Hayes (1984). The polarimetric data and H$\\alpha$ spectroscopy of the event were interpreted by Brown \\& Wood (1992) as due to a discrete mass ejection equivalent to an instantaneous mass-loss rate $\\dot{M}\\!\\approx$ $2.5\\!\\times\\!10^{-8}M_{\\odot}$~yr$^{-1}$. The sudden H$\\alpha$ emission outbursts in $\\mu$~Cen on timescales ranging from 20 to 245 days were interpreted by Hanuschik et al. (1993) as being produced by discrete mass ejecta of $\\Delta M\\!\\approx$ $10^{-10}M_{\\odot}$. In the past decade, short- and long-lived photometric outbursts have been discovered (Cook et al. 1995; Hubert \\& Floquet 1998; Keller et al. 2002; Mennickent et al. 2002) from {\\sc macho}, {\\sc ogle}, and {\\sc hipparcos} data. They are currently called `flickers' or 'sharp' variations when the time scales of phenomena range from $t\\sim$ 50 to 200 days, and they imply light changes of $\\Delta V\\!\\la$ 0.2 mag. They are called `bumpers' or `humps' if $t\\ga$ 200 days to 3-4 years and $\\Delta V\\!\\ga$ 0.3 mag. Mennickent et al. (2002)  speculate that thermal instabilities may produce these photometric behaviors, while Moujtahid et al. (1999), Hubert et al. (2000), and Zorec et al. (2000a) have attempted to explain them in terms of sudden or discrete mass-ejections of $\\Delta M\\!\\approx\\! 10^{-10}M_{\\odot}$. These variations also summarize the irregular photometric behaviors, including the well-known B$\\rightleftharpoons$Be$\\rightleftharpoons$Be-shell phase transitions, observed in the past more or less sporadically (cf. Pavlovski et al. 1997, Percy et al. 1997, Moujtahid et al. 1998). Finally, we should not omit mentioning the far-UV C\\,{\\sc iv}, Si\\,{\\sc iii}, and Si\\,{\\sc iv} discrete line-absorption components in stars with $V\\!\\sin i\\ga$ 150 km~s$^{-1}$ (Grady et al. 1987, 1989), which were interpreted as due to wind instabilities, but may actually correspond to discrete mass-ejections.\\par According to these observational grounds for mass ejections, two scenarios can be sketched to roughly account for the CE variations and/or formation, to which the observed long-term Be star emission variations should be associated.\\par This paper is organized as follows. In Sect~2 we introduce the two possible scenarios for the mass ejection. In Sect.~3 we briefly present the SIMECA code and its main assumptions. Section 4 presents the {\\it ring scenario} and Sect.~5  treat the {\\it vanishing mass flux scenario}. The comparison of different observables (photometric, spectroscopic, and interferometric) is done in Sect.~6. Finally, in Sect.~7 we summarize and discuss our main results.\\par \\medskip ",
        "conclusions": "It is clear from this study that, in order to definitely favor the ring scenario of disk dissipation instead of the mass-loss vanishing, we must follow both spectroscopically and interferometrically a Be star that is showing a decrease in its IR excess, as well as in its emission line intensities.  A clear signature of the disk dissipation following the ring scenario could be:\\\\ -the disappearance of the high velocity tails in the emission line and a nearly constant peak separation;\\\\ - visibilities with an increasing second lobe, an increase in the value of the first zero, and, assuming an unresolved central star, a first zero at shorter baselines as a function of the disk excavation.\\\\  Thanks to the VLTI and its AMBER instrument, which combine high spatial resolution (2 mas for the 200m baseline) and a spectral resolution of 10000, the ``direct\" observation of a varying inner disk radius, certainly correlated to circumstellar outbursts, will be soon accessible. This kind of observation has already be done by Guilloteau et al. \\cite{guilloteau}  for GG Tau using  the IRAM interferometer.  We can see in Fig. \\ref{ggtau} that the real part of the 1.4 mm continuum visibility  as a function of the IRAM baseline length (dots with error bars) agrees very well with a model (plain line) where 90\\% of the flux is coming from a circumbinary ring and the rest from an extended disk. The global shape of the visibility curve is globally the same as for our pole-on ring models from Fig. \\ref{lines_visibilities_ring} even if we have plotted visibilities modulus (between 0 and 1) in our modeling whereas this is the real part of the visibility,  which can be negative, that is plotted in Guilloteau et al. \\citealp{guilloteau}."
    },
    "0606/astro-ph0606469_arXiv.txt": {
        "abstract": "This is the third in a series of papers that present observations and results for a sample of 76 ultra-steep-spectrum radio sources designed to find galaxies at high redshift. Here we present multi-frequency radio observations, from the Australia Telescope Compact Array, for a subset of 37 galaxies from the sample. Matched resolution observations at 2.3, 4.8 and 6.2~GHz are presented for all galaxies, with the $z<2$ galaxies additionally observed at 8.6 and 18~GHz. New angular size constraints are reported for 19 sources based on high resolution 4.8 and 6.2~GHz observations. Functional forms for the rest-frame spectral energy distributions are derived: 89\\% of the sample is well characterised by a single power law, whilst the remaining 11\\% show some flattening toward higher frequencies: not one source shows any evidence for high frequency steepening. We discuss the implications of this result in light of the empirical correlation between redshift and spectral index seen in flux limited samples of radio galaxies. Finally, a new physical mechanism to explain the redshift -- spectral index correlation is posited: extremely steep spectrum radio galaxies in the local universe usually reside at the centres of rich galaxy clusters. We argue that if a higher fraction of radio galaxies, as a function of redshift, are located in environments with densities similar to nearby rich clusters, then this could be a natural interpretation for the correlation. We briefly outline our plans to pursue this line of investigation. ",
        "introduction": "\\indent Until the late 1970s the search for high redshift radio galaxies (HzRGs) was restricted to optical spectroscopy at the position of a radio source whose optical host galaxy is undetected on, or at the detection limit of, the Palomar Observatory Sky Survey (POSS) plates (e.g. \\citealp{kri74,kri78,smi80,gun81,lai83,spin84,per84,spin85}). Indeed, the first galaxy discovered above redshift one was found in this way \\citep{spin82}.\\newline  The field was revolutionised when it was realised that the radio sources without optical counterparts on the POSS plates tended to have the steepest spectral indices $\\alpha$ (where $S_{\\nu}\\propto \\nu^{\\alpha}$) measured between 178 and 1415~MHz \\citep{tie79,blu79}. Since then most radio selected samples designed to find HzRGs have used some degree of spectral index selection. \\newline  Three explanations have been put forward in the literature to justify the trend for HzRGs to have steeper spectral indices; the so-called $z-\\alpha$ correlation. The first is based on the observation that the spectral energy distributions (SEDs) of radio galaxies typically steepen toward higher frequencies. For a fixed set of observing frequencies, a radio galaxy at higher redshift will be sampled at higher rest-frame frequencies where, on average, it will exhibit progressively steeper spectral indices: a k-correction. It has also been suggested that the SEDs of HzRGs will be intrinsically steeper thanks to enhanced inverse Compton (IC) losses of the relativistic electron population against the cosmic microwave background (CMB) photons whose energy density increases as $(1+z)^4$ \\citep{kro91}. These two explanations have been adopted in the literature over recent years as the main contributors to the correlation \\citep{cdb00sample,cdb02wish,ped03,coh04,jar04}. The third explanation holds that the $z-\\alpha$ correlation itself is indirect, reflecting an intrinsic correlation between radio luminosity and spectral index coupled to a Malmquist bias \\citep{blu99a}. Regardless of the underlying reason for the correlation, the steep-spectrum technique has enjoyed tremendous success over the years, with at least 30 radio galaxies above $z=3$ found to date using some form of spectral index culling.\\newline  In the first paper of this series (paper I; \\citealt{cdb04}), we defined a sample of 76 high redshift radio galaxy candidates selected between 843 and 1400~MHz on the basis of their ultra steep radio spectral index (USS; $\\alpha\\leq-1.3$). In the second paper of the series (paper II; \\citealt{cdb06}) we reported the results of optical spectroscopy for 52 sources in the sample which included 35 spectroscopic redshifts. The aim of this paper III is to better understand the radio spectral characteristics of our sample of USS radio galaxies, over a broad range of redshifts. To this end, we report multi-frequency radio observations from the Australia Telescope Compact Array (ATCA) and present rest-frame radio SEDs for 37 galaxies in the sample. \\newline  This paper is organised as follows. Some relevant synchrotron physics is given in \\S\\ref{synchphysics}. The subset of SUMSS-NVSS USS radio galaxies we observed for this paper is defined in \\S\\ref{sourceselection}. Multi-frequency ATCA observations and data analysis follow in \\S\\ref{ATCAobservations}. The results of the observational program are presented in \\S\\ref{ATCAresults}, which includes rest-frame radio SEDs. The remainder of the paper is discussion based. \\S\\ref{discussion} and \\S\\ref{zedalphadiscussion} focus on the implications of our results for existing interpretations of the redshift -- spectral index correlation. Penultimately, in \\S\\ref{neighbours}, we speculate on an alternative physical mechanism for this correlation and concluding remarks are given in \\S\\ref{summary}. Throughout this paper, we assume a flat $\\Lambda$-dominated cosmology with $H_0=71$~\\kms Mpc$^{-1}$ and $\\Omega_{\\Lambda}=0.73$ \\citep{spe03}. ",
        "conclusions": "\\label{summary} In this paper we have presented multi-frequency radio observations for a sample of 37 steep-spectrum radio galaxies. The sample was originally designed to find distant radio galaxies by exploiting a well-known correlation between redshift and radio spectral index. We determined rest-frame SEDs for the sources in our sample, which cover a broad range of redshifts from $0.1<z<4.0$. We have found that 34 (89\\%) of the SEDs are well-characterised by a single power law, with the remaining 4 (11\\%) sources showing evidence for spectral flattening toward higher frequencies. This result appears in stark contrast to complete samples of radio sources in which approximately 70\\% of SEDs progressively steepen toward high frequencies. We have begun investigating whether differences in the sample selection frequencies, flux density limits or spectral index thresholds are responsible for these differences. We also note that the Malmquist bias inherent in our sample currently prevents us from separating out the contribution from a luminosity-spectral index correlation. \\newline  Coupled to the rarity of USS sources in nearby GHz-selected samples, we are not persuaded that a k-correction plays a dominant role in our SUMSS-NVSS USS sample. Instead, we believe that USS HzRGs are intrinsically steeper than their nearby counterparts. Since the surface density of $z>3$ radio galaxies in our sample is up to three times higher than in lower-frequency-selected USS samples, and since the derived space density of $z>3$ radio galaxies in our sample is consistent with representing the majority of all powerful radio galaxies at high redshift (paper~II), we speculate that a mechanism other than a k-correction or luminosity-induced correlation is driving the $z-\\alpha$ correlation. \\newline  There is a well-known trend between the spectral index and the ambient surroundings of nearby, low-luminosity radio galaxies, in the sense that the steepest spectrum sources reside at the centre of the richest galaxy clusters. By analogy, the $z-\\alpha$ correlation could be driven by evolution in the richness of the environment of radio galaxies {\\textit{if}} a higher fraction of radio galaxies reside within cluster-like ambient densities as a function of redshift. Further investigations are now underway to test this hypothesis."
    },
    "0606/astro-ph0606143_arXiv.txt": {
        "abstract": "Over the last five decades, AGN studies have produced a number of spectacular examples of synergies and multifaceted approaches in astrophysics. The field of AGN research now spans the entire spectral range and covers more than twelve orders of magnitude in the spatial and temporal domains. The next generation of astrophysical facilities will open up new possibilities for AGN studies, especially in the areas of high-resolution and high-fidelity imaging and spectroscopy of nuclear regions in the X-ray, optical, and radio bands. These studies will address in detail a number of critical issues in AGN research such as processes in the immediate vicinity of supermassive black holes, physical conditions of broad-line and narrow-line regions, formation and evolution of accretion disks and relativistic outflows, and the connection between nuclear activity and galaxy evolution. ",
        "introduction": "Recent years have witnessed substantial progress in studies of active galactic nuclei (AGN).  Activity is observed in galaxies throughout the entire range of the electromagnetic spectrum~\\cite{lobanov:fab99}. It is manifested on spatial and temporal scales that span over 12 orders of magnitude, ranging from several Schwarzschild radii to megaparsecs and from minutes to millions of years~\\cite{lobanov:pet02}. On the surface, the activity of galactic nuclei comes in many faces, outlined by the uncomfortably large number of the AGN classes --- yet the underlying mechanism and physical conditions in the nuclear regions of all galaxies are probably more similar than  may appear at first glance~\\cite{lobanov:ant93,lobanov:gm98,lobanov:ver00}.  Last but not least, AGN are now detected at redshifts of 6 and beyond~\\cite{lobanov:fan+03}, and they may be closely connected to the formation and evolution of the large-scale structure~\\cite{lobanov:mb00} and to the epoch of reionization of the Universe~\\cite{lobanov:fan+02,lobanov:mad+04,lobanov:ric+04}.  There is growing acceptance of ubiquity of the nuclear activity in galaxies and its connection to the presence of supermassive black holes (SMBH) in the centres of all massive galaxies~\\cite{lobanov:and+04,lobanov:gh04,lobanov:ho99a,lobanov:ho99b}. Observations in the infrared~\\cite{lobanov:gc00} and sub-millimetre bands~\\cite{lobanov:mb00} indicate that almost every galaxy exhibits certain characteristics previously thought to be related only to a much narrower class of powerful AGN. This implies that virtually all massive galaxies may go through a strong AGN phase in the course of their cosmological evolution.  Historically, observational studies of AGN have been concentrated largely within six broadly defined areas.  \\emph{1.~Surveys} have been used effectively, first in the optical and radio domains and later throughout much of the electromagnetic spectrum, to find correlations between different types of AGN and to trace their cosmological evolution. The databases provided by the ROSAT, IRAS, ISO, SIRTF and VLA (FIRST~\\cite{lobanov:bwh95} and NVSS~\\cite{lobanov:con+98}) surveys have become the true cornerstones of AGN research, as well as the HST and Chandra deep field observations. These efforts will be taken further by the ongoing SDSS survey~\\cite{lobanov:aba+03} and the Spitzer surveys~\\cite{lobanov:fad+04}.  The Spitzer data have already yielded a fundamental result on the obscuration in AGN, indicating that as much as three quarters of all AGN are likely to be obscured at all redshifts~\\cite{lobanov:urr04}, which implies that the number of active galaxies may be much higher than was thought. AGN studies in the next decade will also benefit from multi-wavelength surveys such as the great observatories origins deep survey (GOODS~\\cite{lobanov:gia+04}) which addresses AGN structure and demographics~\\cite{lobanov:hec+04}, unification scheme and and AGN--SMBH co-evolution~\\cite{lobanov:gm98,lobanov:urr04}. A number of ground and space VLBI surveys undertaken in the radio in the centimetre~\\cite{lobanov:fc00,lobanov:fom+00,lobanov:hen+95,lobanov:hir+00,lobanov:kel+98,lobanov:lis03b,lobanov:pr88,lobanov:pol+95,lobanov:tay+94,lobanov:tay+96,lobanov:tha+95,lobanov:xu+95} and millimetre~\\cite{lobanov:lob+00,lobanov:ldp98,lobanov:rbb98} domains have revealed a wealth of information about morphology~\\cite{lobanov:hor+04,lobanov:kov+05,lobanov:zen+02}, kinematics~\\cite{lobanov:kel+04} and evolution of radio-emitting material in the nuclear regions and relativistic outflows in AGN on scales down to fractions of a parsec. VLBI surveys have also been used effectively for addressing a variety of general astrophysical problems including AGN evolution~\\cite{lobanov:sha+96} and population modelling~\\cite{lobanov:lis03,lobanov:ls02}, jet formation~\\cite{lobanov:lob+00}, fundamental astrophysical emission processes~\\cite{lobanov:kel02,lobanov:lm99}, and cosmology~\\cite{lobanov:gkf99}.   \\emph{2.~Morphological studies} of AGN have been done extensively in the radio~\\cite{lobanov:kas+93,lobanov:ohc89}, optical~\\cite{lobanov:per+99} and X-ray regimes~\\cite{lobanov:mar+01,lobanov:mar+02}, recovering large-scale structures produced by AGN.  Circumnuclear disks and tori have been revealed conclusively in recent optical, near-infrared and radio observations~\\cite{lobanov:jaf+04,lobanov:kbg03,lobanov:mun+03,lobanov:wei+04} (Fig.~\\ref{lobanov:fig1}). Most spectacularly, recent X-ray images obtained with Chandra have shown the shocks and ripples in the intergalactic medium excited by the central engine in NGC\\,1275~\\cite{lobanov:fab+03}.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.9\\textwidth,angle=0,bb=14 14 326 253,clip=true]{lobanovF1.eps} \\caption{Montage of the inner 250 pc of NGC 4151. The H$_2$ emission traces a torus, the H{\\,\\scriptsize I} absorption comes from a ring inside the torus. Ionized gas (black) is assumed to fill the torus inside the H{\\,\\scriptsize I} ring~\\cite{lobanov:mun+03}.} \\label{lobanov:fig1} \\end{center} \\end{figure}   \\emph{3.~Variability studies} of long- and short-term changes of continuum and line emission have enabled detailed investigations of non-stationary processes in the immediate vicinity of the SMBH and understanding better the physics of nuclear regions of AGN~\\cite{lobanov:ulr92}. Timescales probed by these studies range from hours to several decades~\\cite{lobanov:aah03,lobanov:bvf03,lobanov:pet02,lobanov:sha+01,lobanov:sta+04,lobanov:wam+97,lobanov:zhe96}. The variable continuum flux is believed to be responsible for ionizing the cloud material in the broad-line region (BLR).  Optical spectral line and continuum variability data have been combined to model the BLR~\\cite{lobanov:kas+00,lobanov:pet02,lobanov:pet+02}, which provided a reliable method for estimating the mass of the central black holes in a number of AGN~\\cite{lobanov:pet02}. In radio-quiet AGN, short-term variations of the continuum X-ray emission~\\cite{lobanov:are+05,lobanov:bri+04} are most likely connected with processes occurring the accretion disk at distances of $\\sim 10$--$100\\,R_\\mathrm{g}$ ($R_\\mathrm{g} = G\\,M_\\mathrm{bh}/c^2$ is the gravitational radius for a black hole of mass $M_\\mathrm{bh}$, where $G$ is the Newtonian gravitational constant) from the black hole~\\cite{lobanov:mdp93}. This conclusion is contested by multi-wavelength monitoring observations~\\cite{lobanov:die+98,lobanov:pet02} showing no apparent time delay between flux variations in the ultraviolet (UV) and optical continua, which indicates that instabilities in the accretion disk or random fluctuation in the accretion rate alone cannot explain the continuum variability~\\cite{lobanov:pet02}. In radio-loud AGN, continuum emission from the relativistic plasma in the jet dominates at all energies~\\cite{lobanov:umu97,lobanov:wor05}, swamping the X-ray emission associated with the accretion flow.  Hence, the continuum variability in radio-loud AGN may be related to both the jet and the instabilities of accretion flows~\\cite{lobanov:mdp93,lobanov:umu97} near the central engine.  Long-term variability of X-ray~\\cite{lobanov:mar+02,lobanov:min+03,lobanov:shi02} and radio emission~\\cite{lobanov:aah03} probably reflects instabilities developing in the disk~\\cite{lobanov:lr05} and shocks propagating in the relativistic jet~\\cite{lobanov:fer98}.  Ejections of new jet components have been reported to occur after minima in the X-ray light curve~\\cite{lobanov:mar+02}, suggesting that the jet is fed by the material falling onto the black hole from inner radii of the accretion disk.  No observational evidence has been reported so far for a link between optical/UV continuum variability and the compact radio jet.  Intraday variations observed in the radio~\\cite{lobanov:kra+03,lobanov:ww95}, optical~\\cite{lobanov:hw96}, IR~\\cite{lobanov:gup+04}, and X-ray domains~\\cite{lobanov:are+05} most likely reflect processes in the immediate vicinity of the central engine and in internal shocks developing in relativistic outflows~\\cite{lobanov:ssp98,lobanov:ssp99,lobanov:spa+01}.  In this respect, intraday variability in AGN may be similar to quasi-periodic oscillations observed in Galactic X-ray binaries. Interstellar scintillations are also likely to play a role in the intraday variations observed in the radio regime~\\cite{lobanov:big+02,lobanov:rkd02}.  \\emph{4.~Multifrequency campaigns}, albeit proven difficult to organize, have yielded accurate broad-band spectra of AGN~\\cite{lobanov:mar+94,lobanov:tur+99} and have enabled detailed investigations of physical processes governing the production of the non-thermal continuum emission from circumnuclear regions in AGN~\\cite{lobanov:har+01,lobanov:tur+00}.  Multifrequency data were used to determine the spectral energy distribution in AGN in the quiescent and flaring stages~\\cite{lobanov:fos+98,lobanov:mat04,lobanov:smu96} and to characterize the properties of the synchrotron emission in relativistic outflows~\\cite{lobanov:deh05,lobanov:lob98b}.  \\emph{5.~Spectroscopic studies} in the optical~\\cite{lobanov:bec+01,lobanov:gre+96,lobanov:mar+03,lobanov:whi+00} infrared~\\cite{lobanov:lut+00} radio~\\cite{lobanov:lob05b}, and increasingly also in the X-ray domain~\\cite{lobanov:fab+95,lobanov:fab+02,lobanov:mal+97,lobanov:nan+97a,lobanov:nan+97b,lobanov:ogl+01,lobanov:sam+98,lobanov:tan+95,lobanov:yaq+00} have provided measurements of the gas and stellar kinematics near the central black hole~\\cite{lobanov:zak+03}, in the broad and narrow line regions of galactic nuclei. These studies will be greatly enhanced with the systematic spectroscopy provided by the SDSS data and with the high-energy spectroscopy data from the existing and planned X-ray missions.   \\emph{6.~Dedicated monitoring programs} have been employed successfully virtually in all observational domains to study in detail the dynamics and physical conditions in the nuclear regions of a number of prominent AGN.  Structure of the nuclear regions has been imaged extensively with high resolution radio interferometry, uncovering the evolution of radio emitting plasma on linear scales from $\\approx 100R_\\mathrm{g}$~\\cite{lobanov:jbl99,lobanov:kri+02} to $\\approx 1$\\,kpc~\\cite{lobanov:wii01,lobanov:zen97,lobanov:zen+96}. Radio monitoring programs are now reaching timescales of several decades~\\cite{lobanov:cou98,lobanov:lz99,lobanov:lr05}. Radio frequency observations of maser emission have been used to probe the presence of accretion disks~\\cite{lobanov:her+99,lobanov:miy+95} and molecular tori~\\cite{lobanov:kb04} around putative black holes in the centres of AGN.  Extragalactic maser emission has provided direct evidence of interaction between the dense molecular material and the ionization cones~\\cite{lobanov:gal+96} or nuclear jets~\\cite{lobanov:cla+98,lobanov:pec+03}.  Properties of relativistic outflows have been connected to the physical conditions in the nuclear regions~\\cite{lobanov:kad+04,lobanov:lob98} and used for identifying possible binary systems of SMBH in AGN~\\cite{lobanov:ca04,lobanov:lr05,lobanov:rom+00}.   High-resolution radio observations of Seyfert galaxies have offered an opportunity to study the optical emission from the immediate vicinity of the central engine of AGN, on scales comparable to the largest extent of the BLR.  Optical observations of broad emission lines have been instrumental for understanding the structure and dynamics of the circumnuclear gas, and they have yielded an attractive concept of a disk-like morphology of the BLR~\\cite{lobanov:pop+04}. The detection and modelling of several double-peaked Balmer lines has further supported this idea. In most cases, the double-peaked emission lines can be explained with line emission that originates in the disk~\\cite{lobanov:eh03}.  There are, however, two very notable exceptions among the radio-quiet AGN. The core of the broad double-peaked emission lines in RX\\,J1042+1212~\\cite{lobanov:pmc96} and Ark\\,120~\\cite{lobanov:pop+01} is most likely to originate from an outflow-related component of the BLR (Fig.~\\ref{lobanov:fig2}). It is possible that this component is indeed related to an interaction between the relativistically moving plasma and high-velocity clouds situated in a sub-relativistic outflow/wind surrounding the jet~\\cite{lobanov:elv00,lobanov:mc97,lobanov:psk00}. In this case, the broad-line emitting region associated with the jet/outflow can be located at a significant distance from the nucleus, perhaps as large as $\\sim 1$\\,pc.  Verifying the existence of the outflow-related component of the BLR and testing the putative connection between the BLR and relativistic outflows in Seyfert galaxies can therefore be viewed as critical experiments that may advance our understanding of the nuclear activity in galaxies.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.45\\textwidth,bb=0 0 504 503,clip=true]{lobanovF2a.eps} \\includegraphics[width=0.45\\textwidth,bb=-94 37 691 805,clip=true]{lobanovF2b.eps} \\caption{{\\bf Left}: The observed H$_\\beta$ line of Ark\\,120 fitted by the two-component model (solid line): a disk (in the line wings) plus an outflow-related component of the BLR (the core of the line). The two components in the core correspond to the line emission generated in the approaching (left) and receding (right) part of a bipolar outflow.  {\\bf Right}: The scheme of the proposed BLR model for Ark\\,120~\\cite{lobanov:pop+01}.} \\label{lobanov:fig2} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0606/astro-ph0606375_arXiv.txt": {
        "abstract": "Many astrophysical sources of high energy emission, such as black hole magnetospheres, superstrongly magnetized neutron stars (magnetars), and probably relativistic jets in Active Galactic Nuclei and Gamma Ray Bursts involve relativistically magnetically dominated plasma. In such plasma the energy density of magnetic field greatly exceeds the thermal and the rest mass energy density of particles.  Therefore the magnetic field is the main reservoir of energy and its dissipation may power the bursting emission from these sources, in close analogy to Solar flares.  One of the principal dissipative instabilities that may lead to release of magnetic energy is the tearing instability.  In this paper we study, both analytically and numerically, the development of tearing instability in relativistically magnetically-dominated plasma using the framework of resistive magnetodynamics. We confirm and elucidate the previously obtained result on the growth rate of the tearing mode: the shortest growth time is the same as in the case of classical non-relativistic MHD, namely $\\tau = \\sqrt{\\tau_a \\tau_d}$ where $\\tau_a$ is the \\Alfven crossing time and $\\tau_d$ is the resistive time of a current layer. ",
        "introduction": "\\label{intro} Dissipation of magnetic energy may power high energy emission in a variety of relativistic astrophysical phenomena, e.g. pulsar wind nebulae \\cite{Cor,Usov,LyK,KiS}, jets of active galactic nuclei \\cite{RoL,JLT}, gamma-ray bursts \\cite{Dre,DrS}, and magnetars \\cite{ToD,Lyu03}.  Perhaps the most clear-cut case for magnetic dissipation is magnetars -- \\NSs with superstrong \\Bfs, sometimes in excess of quantum \\Bf $B_Q=m_e^2 c^3/\\hbar e = 4 \\times 10^{13}$G. (Observationally, magnetars are separated in two classes -- Anomalous X-ray Pulsars (AXPs) and the Soft Gamma-ray Repeaters (SGRs) -- both showing X-ray flares and quiescent X-ray emission.)  A number of evidence suggest that processes that lead to the production of magnetars X-ray flares (and possibly of the persistent emission) are similar to those operating in the Solar corona. The bursting activity of SGRs is strongly intermittent, showing a power law dependence of the number of flares on their energy, $dN/dE \\sim E^{\\alpha}$, with $\\alpha=1.66$ and log-normal distribution of waiting times between the flares \\cite{gogus}, both being similar to Solar flares.  Giant flares (GFs below), immense explosions releasing up to $10^{46}$ ergs in a fraction of a second, provide a number of independent evidence in favour of magnetic dissipation.  Two recently observed explosions, from SGR 1900+14 and SGR 1806-20, showed similar behavior leading to and following the GFs: months before the GF X-ray activity increased, spectrum hardened and spindown increased \\citep{turolla}. In the post-flare period, pulsed fraction and the spin-down rate have significantly decreased and the spectrum softened \\citep{woods01,rea05}.  All these effects are in agreement with the prediction of twisted magnetosphere model \\citep{tlk02}, with the twist increasing before the GF and decreasing during the GF brought about by reconnection \\cite{Lyu06}. Most importantly, observations of the December 27 GF from SGR 1806-20 show very short rise time $\\sim 0.25 $ msec \\citep{palmer05}. Such short rise times are inconsistent with crustal deformations and points to magnetospheric origin of GF \\cite{Lyu06}.  In analogy with Solar flares magnetic energy to be dissipated may build-up on long time scales and then be dissipated on very short times scales, possibly as short as \\Alfven crossing time.  For example, a slow plastic motion of the crust implants a twist (current) in the magnetosphere on a long time scale. At some point a global system of magnetospheric currents and sheared magnetic fields loses equilibrium and produces a flare. This instability may be dynamical (\\eg loss of magnetic equilibrium) and/or resistive (this work).  Though similar in underlying physical process, magnetar and Solar plasmas may differ considerably, since properties of plasma expected in magnetar magnetospheres are very different from the conventional Solar and laboratory plasmas.  The principal difference is that magnetar plasma is strongly (relativistically strongly) magnetized. It is convenient to introduce a parameter $\\sigma_m$ \\citep{kc84} as the ratio of the magnetic energy density $u_B=B^2/8\\pi$ to the total plasma energy density (including rest mass!): \\beq \\label{sigmadef} \\sigma_m={u_B\\over u_{p}} \\label{sig} \\ee Conventionally, in a non-relativistic Solar and laboratory plasmas, parameter $\\sigma_m$ is small, $\\sigma_m \\ll 1$. On the other hand, in magnetar magnetospheres it is expected that $\\sigma_m$ is very large, \\beq {\\om_B R_{NS} \\over c} \\left({m_e \\over m_p} \\right) \\sim 10^{13} \\leq \\sigma_m \\leq {\\om_B \\over \\Omega} \\left({m_e \\over m_p} \\right) \\sim 10^{16} \\ee (the upper limit here comes from electron-ion plasma density equal to the Goldreich-Julian density, the lower limit corresponds to plasma density $\\sim B/ (R_{NS} e)$ at which point currents drifting with near the speed of light create toroidal \\Bfs of the order of the poloidal).  Here $\\om_B = e B/m_e c$ is cyclotron frequency, $B$ is \\Bf at the \\NS surface, $R_{NS} $ is radius of \\NS and $\\Omega$ is rotational frequency.  Large expected value of $\\sigma_m$ (or small $1/\\sigma_m$) may be used as an expansion parameter in equations of relativistic magnetohydrodynamics (MHD). The zero order equations, the equations of {\\it magnetodynamics} (MD), describe the dynamics of magnetic field under the action of magnetic pressure and tension \\cite{Kom02}. Since to this order the rest mass density and thermodynamic pressure of plasma completely vanish, the only effect plasma has on this dynamics is due to its high conductivity which requires ${\\bf E} \\cdot {\\bf B}=0$ and $E^2<B^2$.  Alternatively, one can write the equations of magnetodynamics as Maxwell's equations closed with a suitable Ohm's law that gives the electric current as a function of the electromagnetic field only \\cite{Gru99}. This Ohm's law is derived from the condition of vanishing Lorentz force that explains the alternative name for this system, {\\it force-free degenerate electrodynamics} or simply {\\it force-free electrodynamics}.  The limit $\\sigma_m \\rightarrow \\infty$ is somewhat reminiscent of the subsonic incompressible hydrodynamic. In this limit the \\Alfven four-velocity is $u_A \\sim \\sqrt{2 \\sigma_m} \\rightarrow \\infty$, so that both cases are applicable when (four)-velocities are much smaller than the velocity of propagation of disturbances in a medium (\\Alfven and sound waves correspondingly).  In fact, this analogy turns out to be much deeper: as we discuss in this paper there are deep similarities in governing equations between magnetodynamics a small velocity limit of the conventional magnetohydrodynamics.   In order to address energy dissipation and eventually particle acceleration, one needs to go beyond the ideal approximation.  One possible approach is to take into account the dissipation of mostly force-free currents due to plasma resistivity.  Resistivity will result in the decay of currents supporting the magnetic field, which would influence the plasma dynamics.  Though plasma resistivity will be of the anomalous type, mediated by plasma turbulence and not by particle-particle collisions, as a first step we assume that plasma resistivity can be represented by macroscopic resistivity parameter $\\eta$.  As another simplification we neglect a possible back reaction of the heated plasma on the global dynamics. This may be justified if the dissipated energy is quickly radiated away.  The principal resistive instability in a conventional, non-relativistic plasma is a so-called tearing mode. It is one of the principle unstable resistive modes, which plays the main role in various TOKAMAK discharges like sawtooth oscillations and major disruptions \\citep{Kadomtsev75}, unsteady reconnection in Solar flares \\citep{Shivamoggi85,Aschwanden02} and Earth magnetotail \\citep{gca78}. Qualitatively, the most important property of the tearing mode is that a current layer of thickness $l$ may dissipate on time scales much shorter than it the resistive time scale $\\tau_d \\sim l^2/\\eta$. In addition, tearing mode may be an initial stage of the development of the (steady-state) reconnection layers.  Development of the tearing mode in a relativistic, strongly magnetized plasma ($\\sigma_m \\gg 1$) is the principal topic of this work. Initially, this problem has been considered by Lyutikov \\shortcite{Lyu03}, who found that in resistive magnetodynamics the current layers are unstable towards formation of resistive small-scale current sheets, that's to development of tearing mode.  He also found that the growth rate of tearing mode is intermediate between the short \\Alfven time scale $ \\tau_a \\sim l/v_a$ (which in a $\\sigma_m \\gg 1$ plasma is the light crossing time scale $\\tau_c \\sim l/c$) and a long resistive time scale $\\tau_d$: $\\tau \\sim (\\tau_d \\tau_a)^{1/2}$. Surprisingly, this is exactly the same expression as the one found in the non-relativistic framework of incompressible magnetohydrodynamics.  In this paper we uncover the deep underlying reasons for this coincidence. It turns out that slow, resistively-driven evolution of strongly magnetized plasma is described by a system of equations that is very similar to nonrelativistic MHD.  In addition to the analytical study we test the theory by means of numerical simulations.  In Section \\ref{basics} we describe the basic equations of resistive magnetodynamics.  In Section \\ref{equi-approximation} we consider the case of slow evolution near equilibrium and show that in this regime the MD equations can be reduced to the system that is very similar to nonrelativistic MHD.  Further reduction is possible when the equilibrium is characterized by zero magnetic tension. This is shown in Section \\ref{incomp-limit}. The analytic theory of the tearing instability is reviewed in Section \\ref{tearing}. Section \\ref{method} describes our numerical method and the results of numerical simulations are presented in Section \\ref{simulations}. ",
        "conclusions": "The original idea of this study was to verify the surprising analytical results of \\cite{Lyu03} on the tearing instability in magnetically dominated relativistic plasma by means of direct numerical simulations. We have carried out such simulations and they fully confirmed Lyutikov's conclusion on the growth rate in magnetically dominated regime -- it indeed coincides with that found earlier for the tearing instability in non-relativistic incompressible MHD, the shortest growth time being equal to the geometric mean of the \\Alfven and resistive timescales \\cite{Furth,Priest}.  Moreover, like in the non-relativistic case we observed formation of magnetic islands in the non-linear phase of the instability. This is accompanied by the appearance of noticeable O-type and X-type neutral points (or lines). The development of an X-point may be considered as a first step in setting up of reconnection layer where dissipated magnetic energy and magnetic flux are constantly replenished due to plasma inflow \\citep{syr81}.  Our results support the possibility of magnetar flares being magnetically driven reconnection-like events that occur in magnetar magnetospheres \\citep{Lyu03,Lyu06} in a similar fashion to flares and coronal mass ejections of Sun's magnetosphere. More complicated \\Bf geometries, e.g. those with nonvanishing magnetic tension, should also be subject to tearing instability on time scales intermediate between the resistive and \\Alfven time scales \\citep{Priest}. The growth rate typically scales as $\\tau \\sim \\tau_c^{\\alpha} \\tau_d^{1-\\alpha}$, where $0<\\alpha<1$ is some coefficient ($\\alpha=1/2$ in our case).  In the case of magnetar flares the observed flare rise-times range from fractions of a millisecond to $\\sim $ ten milliseconds, which is indeed intermediate between the \\Alfven crossing time, $\\sim 0.03 $ msec, and the resistive time, $\\sim$ seconds, of the magnetosphere \\cite{Lyu03}.  The short timescale of tearing mode is achieved entirely through the formation of a very narrow resistive sub-layer with very short diffusion time scale. However, it is also expected that the resistivity itself may be enhanced inside such sub-layers due to the development of plasma turbulence (anomalous resistivity).  Elucidating the properties of anomalous resistivity in such plasma is an important next step. Moreover, the collapse of current sheets and the eventual development in our simulations of the drift speeds exceeding the speed of light clearly indicates a breakdown of the MD approximation. In order to move forward one needs to put back the dynamical effects of matter which are bound to become important in this singular domain. The simplest way of doing this is to return to the model of relativistic MHD like in Lyutikov \\& Uzdensky \\shortcite{lu03} and Lyubarsky \\shortcite{Lyub05}.   The coincidence of the growth rates for the tearing instability in incompressible MHD and MD clearly suggests some hidden similarity between the evolution equations of these two systems. We have carried out further analysis of resistive MD and discovered that in the quasi-equilibrium limit, that is characterized by low drift speeds, its equations do indeed become similar to those of non-relativistic MHD. In particular, the mass-energy density of the magnetic field takes on the role the mass density of classic MHD whereas the drift velocity plays the role of the fluid velocity. {\\it So the dynamics of the system can be described as a flow of magnetic mass-energy under the action of magnetic pressure and tension.}  Moreover, the fact that the magnetic tension of the equilibrium current sheet subjected to tearing instability is zero allows further reduction of the MD equations which leads to almost the same system as incompressible MHD.  When, this system is utilized in the analysis of the tearing instability it generates exactly the same equations for the perturbation as in incompressible MHD. Thus, our analysis provides perfect mathematical explanation for the surprising results in Lyutikov \\shortcite{Lyu03} and improves our understanding of the dynamics of magnetically-dominated plasma.  Finally, we seem to have come up with a more suitable name for the system of equations describing the dynamics of relativistic magnetically dominated plasma. Instead of the currently used {\\it force-free degenerate electrodynamics} (FFDE) or simply {\\it force-free electrodynamics} (FFE) we propose to use {\\it magnetodynamics} (MD).  Let us summarize our arguments in favour of this new name. The current name originates from the early studies of stationary magnetospheres of neutron stars and black holes.  Besides the steady-state Maxwell's equations, the key equations used in those studies were the degeneracy condition, $\\spr{E}{B}=0$, and the condition of vanishing Lorentz force. The full system of the corresponding time-dependent equations was not known at the time and so were not known its connection with the relativistic MHD and the properties of it's waves.  All these have been discovered only very recently \\cite{Uchida,Gru99,Kom02}.  In Komissarov \\shortcite{Kom02} the dynamical equations of relativistic magnetically dominated plasma, equations (\\ref{MD1}-\\ref{MD4}), have been derived from the system of ideal relativistic MHD in the limit of vanishing rest mass density and pressure of matter.  This way of derivation immediately suggests simply to remove the {\\it hydro}-component from the word {\\it magneto-hydro-dynamics}, which obviously results in the name we propose here. Like in ideal MHD, in ideal MD the electric field also vanishes in the the fluid frame, or to be more precise in the frame moving with the drift velocity.  An inertial observer moving with this velocity would experience a pure magnetic field.  Thus, it is desirable to retain the emphasis on the magnetic component of the electromagnetic field, that is present in the name of MHD.  The key wave of MHD, the Alfv\\'en wave, is also present in MD. It propagates along the magnetic field lines with the speed of light \\cite{Kom02}. Like in MHD the electric field vector can be eliminated from the set of dependent variables of equations (\\ref{MD1}-\\ref{MD4}) (see Sec.\\ref{basics}).  Finally, as we have shown in this paper one can qualitatively describe the evolution of magnetically dominated plasma as a flow of magnetic energy under the action of Maxwell's stress. Apparently, we are not the only ones who are not particularly happy with the name force-free electrodynamics. Recently, A.Spitkovsky proposed to use the name {\\it force-free MHD} instead \\cite{Spi06}. This name is somewhat better but it is still not completely satisfactory. First of all, matter and the electromagnetic field appear in this name on equal terms, contrary to the relevant physical conditions. Secondly, the name MHD implies a somewhat different set of evolution equations. Thirdly, as we have seen, the resistive case can no longer be described as force-free. Finally, the name {\\it magnetodynamics} is made of just one word that makes it more esthetically pleasing."
    },
    "0606/astro-ph0606696_arXiv.txt": {
        "abstract": "We have used the deep, multi--wavelength images obtained with the Advanced Camera for Surveys (ACS) on the {\\it Hubble Space Telescope (HST)} by the Great Observatories Origins Deep Survey (GOODS) to identify $\\sim 4700$ Lyman-break galaxies (LBGs) at redshifts $2.5<z<5$, and 292 starburst galaxies at $z\\sim 1.2$. We present the results from a parametric analysis of the two-dimensional surface brightness profiles using a S\\'{e}rsic function, for the 1333 brightest LBGs with rest-frame M$_{1600\\AA}$ $\\leq -20.5$ AB magnitudes. We distinguish the various morphological types based on the S\\'{e}rsic index, $n$, which measures the profile shape. About 40\\% of LBGs at $z\\sim 3$ have light profiles close to exponential, as seen for disks, and only about 30\\% of the galaxies have the high central concentrations seen in spheroids. We also identify a significant fraction ($\\sim$ 30\\%) of galaxies with shallower than exponential profiles, which appear to have multiple cores or disturbed morphologies suggestive of close pairs or on-going mergers. The fraction of spheroid-like ($n>2.5$) LBGs decreases by about 15\\% from $z\\sim 5$ to 3. A comparison of LBGs with the starburst galaxies at $z\\sim 1.2$, shows that disk-like and merger morphologies are dominant at both redshifts, but the fraction of spheroid-like profiles is about 20\\% higher among LBGs. The ellipticity distribution for LBGs exhibits a pronounced skew towards high ellipticities ($\\epsilon >0.5$), which cannot be explained by galaxy morphologies similar to the local disks and spheroids viewed at random orientations. The peak of the distribution evolves toward lower ellipticities, from 0.7 at $z=4$ to $\\sim 0.5$ at $z=3$. The ellipticity distribution for the $z\\sim 1.2$ galaxies is similar to the relatively flat distribution seen among the present--day galaxies. The dominance of elongated morphologies among LBGs suggests that in a significant fraction of them we may be witnessing star-formation in clumps along gas-rich filaments, or the earliest gas-rich bars that encompass essentially the entire visible galaxy. Similar features are found to be ubiquitous in hydrodynamical simulations in which galaxy formation at high redshifts occurs in filamentary inflows of dynamically--cold gas within the dark matter halos, and involves gas-rich mergers. ",
        "introduction": "% Galaxy morphologies at various look-back times provide important insights into the physical process associated with galaxy assembly. In the local Universe, the morphologies seen among the galaxy population are well described by the Hubble sequence (Hubble 1936) and both the physical, and kinematic properties of galaxies are known to vary systematically with the Hubble type (Roberts \\& Haynes 1994). Quantitative morphological analysis of galaxies at {\\it HST} resolution ($\\sim$0.\\asec 1) suggest that the Hubble sequence was already in place by $z\\sim 1$ (Abraham et al. 1996; Schade et al. 1999; Brinchmann et al. 1998; Lilly et al. 1998; Marleau \\& Simard 1998; Im et al. 1999; van den Bergh 2001; Trujillo \\& Aguerri 2004; Conselice, Blackburne, \\& Papovich 2005). There is growing evidence that a significant fraction of the present--day spirals and ellipticals were formed beyond $z=1$, and very little size evolution has occurred after this epoch (Lilly et al. 1998; Simard et al. 1999; Im et al. 2002; Stanford et al. 2004; Ravindranath et al. 2004; Barden et al. 2005). Therefore, the assembly of the spheroid and disk components that define the morphological types appears to have occurred at much higher redshifts ($z>1$). Although based on small samples, the morphological analysis of galaxies with spectroscopic redshifts $1<z<3$, have shown that the regular Hubble types can be identified out to $\\sim z=1-1.5$. However, at $z>1.5$ the Hubble sequence is no longer clearly discernible, even in deep {\\it HST}/NICMOS infrared images which measure the rest-frame optical light. At $z>1.5$, the galaxies are often irregular and compact, and show tidal features, double nuclei, or disturbed morphology suggesting that mergers may be dominant at these redshifts (Papovich et al. 2005; Conselice, Blackburne, \\& Papovich 2005). Several earlier studies (Abraham et al. 1996; Glazebrook et al. 1994; Griffiths et al.  1994; Driver et al.  1995; Cowie, Hu, \\& Songaila 1995) have also emphasised the dominance of peculiar and irregular galaxies at the faintest magnitudes in optical {\\it HST} surveys.  For almost a decade, star-forming galaxies at higher redshifts ($z>2.5$) have been identified in deep, multi--wavelength surveys by using color criteria that are sensitive to the Lyman-break and Lyman-$\\alpha$ forest absorption features in their otherwise flat ultraviolet spectral continuum (Steidel et al. 1996; Giavalisco 2002; Giavalisco et al. 2004). The robustness and efficiency of the ``Lyman-break'' technique to identify $z>2.5$ galaxies has been established through extensive spectroscopic redshift determinations (Steidel et al. 2003). Although Lyman-break galaxies (LBGs) are among the largest samples of high-$z$ galaxies that have been identified, until recently, the high resolution {\\it HST} images required to study their morphological properties were available only for small samples of LBGs (Giavalisco, Steidel, \\& Macchetto 1996; Giavalisco et al. 1996; Lowenthal et al. 1997). The previous morphological analysis using 19 LBGs have suggested that they have a range of morphologies, and their surface brightness profiles in most cases show a core that can be represented by the $r^{1/4}$ profile seen for spheroids (Giavalisco, Steidel, \\& Macchetto 1996). About 90-95\\% of the stars are formed in these compact cores with sizes in the range 2.4$-$3.6 h$_{50}^{-1}$ kpc. Ferguson et al. (2004) used the {\\it HST}/ACS images from the initial three epochs of GOODS observations to measure the concentration index, ellipticity, and sizes of 386 LBGs at $z\\sim4$. Although the concentration index measures favor a significant fraction of spheroid-like systems among the LBGs at $z\\sim 4$, their ellipticity distributions seemed to suggest that extended disk-like morphology was more dominant. In a recent analysis, Lotz et al. (2005) used GOODS images to study morphology of 82 LBGs at $z\\sim 4$ with sizes $r_{p}$ $>$ 0.\\asec 3, where $r_{p}$ is the Petrosian radius. They measured the Gini coefficients, and the second--order moment of the brightest 20\\% of the galaxy light (M$_{20}$), and find that only 30\\% have bulge-like morphology and more than 50\\% of the galaxies have morphologies that are disk-like, minor mergers, or post--mergers.  In this paper, we use the surface-brightness profile shape, and ellipticity measurements to characterize the rest-frame UV morphologies among the high redshift, star-forming galaxies in the GOODS images. Our aim is to quantify the frequency of spheroid, disk, or merger-like morphology among LBGs at $z>2.5$ in a statistical manner, and to study the morphological evolution of actively star-forming galaxies through a comparison of LBGs at $z>2.5$ and starbursts sample at $z\\sim 1.2$. These redshifts are particularly interesting because they correspond to two widely separated epochs marked by lack of the Hubble type morphologies at $z=3$, to their appearance by $z=1.2$.  The central concentration of the light profiles (Abraham et al. 1996; Conselice et al. 2003), and axial ratios (Im et al. 1995; Odewahn et al. 1997; Alam \\& Ryden 2002) are among the commonly used measurements to characterize the morphological types among faint, distant galaxies. The interpretation of these quantities is based on the observed trends with galaxy morphology in the local Universe. Elliptical galaxies and bulges of spiral galaxies tend to have steep ($r^{1/4}$) light profiles, while disk-dominated galaxies have exponential or even shallower light profiles. The apparent axial ratio ($b/a$) or ellipticity ($\\epsilon = 1-b/a$), of a galaxy is related to its intrinsic three-dimensional structure, and elliptical and spiral galaxies are known to have different observed axial ratio distributions locally (Sandage, Freeman, \\& Stokes 1970; Lambas, Maddox, \\& Loveday 1992). The distribution for local ellipticals peaks around $\\epsilon=0.2$ and declines rapidly beyond $\\epsilon > 0.5$, while spiral galaxies show a relatively flat distribution from $\\epsilon=0.2$ to 0.7 and falls off at higher $\\epsilon$ (Lambas, Maddox, \\& Loveday 1992).  Most of the morphological studies of low redshift galaxies are done at rest-frame optical wavelengths, but the {\\it HST} optical images of the LBGs at $z>2.5$ sample the rest-frame UV wavelengths where the contribution from actively star forming regions dominates the light. It is known that morphological $k-$correction or the wavelength dependence of galaxy morphology can be significant for local galaxies (Kuchinski et al. 2000, 2001), and Giavalisco et al. (1996) have examined how this may influence the morphologies inferred for the high redshift galaxies. Galaxies tend to be classified as later Hubble types in the UV, because the redder bulge and bar components which are pre-dominantly made of evolved stellar populations and low-mass stars, become extremely faint with respect to the actively star-forming disk component. However, a comparison of the optical and near-infrared {\\it HST} images which sample the rest-frame UV and optical light respectively for LBGs at $z>2.5$, show very similar morphology implying that the morphological $k-$correction is negligible (Dickinson 2000; Papovich et al. 2003). Quantitative measurements of the internal color gradients also show that this effect is more severe at low redshifts ($z\\leq 1$) than at high redshifts (Papovich et al. 2005). Therefore, the UV morphologies derived for the LBGs are likely to be representative of the typical morphological mix among this population.  In the present study, we perform a parametric analysis of the galaxy surface brightness distribution, in order to measure the profile shape that quantifies the central concentration, and the ellipticity which allows to infer the intrinsic shape.  We use the deep, multi-wavelength {\\it HST}/ACS images for a large sample ($\\approx 4700$) of LBGs available from GOODS, with unprecedented high spatial resolution ($\\sim 700-800$ parsecs at $2.5< z <5.0$) which is essential for exploring the morphology of high redshift galaxies. We also make use of the Hubble Ultra Deep Field (HUDF) images to understand the selection effects, and use the UV images from GALEX to redshift local galaxies to high-$z$ for comparison with the morphology of LBGs. The UV morphology is dominated by the actively star-forming regions which are often clumpy and irregular, and cannot be well-described by a smooth analytical function. However, we use the S\\'{e}rsic function (S\\'{e}rsic 1968) to model the light profile because it is sensitive to the central concentration, and can help to broadly distinguish between a bulge-dominated and disk-dominated galaxy. Also, we obtain a measure of the ellipticity for the overall galaxy light distribution from the parametric model. We describe the observations, and sample selection in \\S\\ 2, and the morphological analysis is detailed in \\S\\ 3. The simulations that were performed in order to assess the selection biases and measurement errors are described in (\\S\\ 4). We discuss the observed morphological diversities among LBGs, and its implications for galaxy formation in \\S\\ 5, and \\S\\ 6. We adopt the cosmology defined by $H_0$ = 70 \\kms\\ Mpc$^{-1}$, $\\Omega_{M} = 0.3$, and $\\Omega_{\\lambda} = 0.7$ throughout this paper. All magnitudes are in units of AB magnitudes (Oke \\& Gunn 1983). ",
        "conclusions": ""
    },
    "0606/astro-ph0606544.txt": {
        "abstract": "We examine, from a theoretical viewpoint, how the physical parameters of \\HII\\ regions are controlled in both normal galaxies and in starburst environments. These parameters are the \\HII\\ region luminosity function, the time-dependent size, the covering fraction of molecular clouds, the pressure in the ionized gas and the ionization parameter. The factors which control them are the initial mass function of the exciting stars, the cluster mass function, the metallicity and the mean pressure in the surrounding interstellar medium. We investigate the sensitivity of the H$\\alpha$ luminosity to the IMF, and find that this can translate to more than a factor two variation in derived star formation rates. The molecular cloud dissipation timescale is estimated from a case study of M17 to be $\\sim1$~Myr for this object. Based upon \\HII\\ luminosity function fitting for nearby galaxies, we suggest that the \\HII\\ region  cluster mass function is fitted by a log-normal form peaking at $\\sim 100 M_{\\odot}$. The cluster mass function continues the stellar IMF to higher mass. The pressure in the \\HII\\ regions is controlled by the mechanical luminosity flux from the central cluster. Since this is closely related to the ionizing photon flux, we show that the ionization parameter is not a free variable, and that the diffuse ionized medium may be composed of many large, faint and old \\HII\\ regions. Finally, we derive theoretical probability distributions for the ionization parameter as a function of metallicity and compare these to those derived for SDSS galaxies. ",
        "introduction": "Introduction} In distant galaxies and starbursts, much of what we infer about the star formation rates, the physical conditions in their interstellar medium (ISM), and the chemical abundances (or ``metallicity\") in the ISM is determined by a study of their emission line spectra. However, we must be constantly aware that what we observe are the integrated properties of the population of \\HII\\ regions in these galaxies, and the integration extends over a wide range of \\HII\\ region ages, cluster masses, sizes, pressures, ionization parameters and potentially, over a wide range of metallicity as well.  In nearby disk galaxies the individual \\HII\\ regions can be spatially resolved and subjected to independent study, but the properties of these \\HII\\ regions are still intimately connected to the properties of their local ISM and to the nature and the state of evolution of their central cluster.  In both cases, it would be wise for us to attempt to understand how the physical parameters of the \\HII\\ regions and their emission line spectra are controlled by their cluster, their environment, and their metallicity. Surprisingly, there has been very little effort devoted to this problem. In general, the modelling of the emission line spectrum of distant galaxies has been treated as if it originated from a single \\HII\\ region of a given metallicity and ionization parameter. This ionization parameter is defined as either the ratio of the mean photon flux to the mean atom density, $q$ or in its dimensionless form as the ratio of mean photon density to mean atom density ${\\cal{U}}=q/c$. The excitation of the \\HII\\ region is very sensitive to both the IMF and the age distribution of the exciting stars. In the modelling hitherto, this is usually treated in two limits; either the instantaneous burst approximation or in the continuous star formation approximation in which stars are supposed to be born continually within the ionized region (see \\citet{Kewley01a}, and references therein).  Clearly, this modeling approach is rather crude, and does not provide insight into the dynamical and age evolution of the individual clusters and their associated \\HII\\ regions. These are important parameters because they determine the geometry of the gas with respect to the exciting stars. The geometry determines both the excitation of the \\HII\\ region and the radiation field density incident on photodissociation region (PDR) around the \\HII\\ region. This in its turn determines the characteristic dust temperature, and hence the wavelength of the peak in the FIR dust re-emission feature.  As discussed in \\citet{Dopita05}, hereinafter referred to as Paper I, for starburst galaxies dominated by their young stellar population and possessing an ISM of high density, the global SED can be adequately be described by the sum of the SEDs of its embedded \\HII\\ regions. This approximation requires that the radiative transfer problem can be approximated as a local, rather than as a global problem. Such a model is correct provided that the majority of the far-IR emission arises from absorption of the UV radiation field in a relatively thin dust layer identified with the classical photo-dissociation region (PDR). These typically have an optical depth corresponding to $A_{\\rm V} \\sim 3$, and a thickness $dR \\sim 300/n_{\\rm H}$ pc. In the molecular regions surrounding normal galactic  \\HII\\ regions, densities are typically $100 - 1000$ cm$^{-3}$, implying that much of the far-IR which they produce comes from a layer of parsec or sub-parsec dimensions. In starburst galaxies, interstellar pressures may range up to a factor of 100 higher than in normal disk galaxies, and these produce correspondingly thinner PDR zones. The assumption that the PDR is thin justify the use of models which treat the radiative transfer as either localized within or immediately surrounding the \\HII\\ regions.  A theory for the size distribution of \\HII\\ regions has been presented by \\citet{Oey97} and \\citet{Oey98}. According to these papers, the \\HII\\ regions expand and evolve as mass-loss and supernova energy driven bubbles for as long as their internal pressure exceeds the ambient pressure in the ISM. When this condition is no longer met, the \\HII\\ region is assumed to ``stall\".  The time taken to reach the stall radius turns out to be proportional to the stall radius, and the mass of the OB star cluster is key in determining the stall radius; low mass cluster \\HII\\ regions stall early and at small radius. The resultant size distribution therefore depends strongly upon the cluster mass distribution, which can in principle de derived from the observed \\HII\\ region H$\\alpha$ luminosity function. In the \\citet{Oey98} paper, this luminosity function was modelled on the assumption that the number of ionizing stars was a power-law, $N(N_*)dN_*=BN_*^{-\\beta}dN_*$. In a later paper, \\citet{Oey03} suggest that the differential nebular size distribution can be described by a power law of slope -4, with a flattening for sizes below $\\sim130$~pc, in contrast to the exponential size distribution usually adopted.  During the expansion phase the \\HII\\ region is treated as a bubble powered by the mechanical energy input of the mass-losing central stars, and following the classical \\citet{Castor75} theory.The radius $R$ and internal pressure,  $P$, of a mass-loss and/or supernova explosion pressurised \\HII\\ region are then given by  \\cite{Oey97, Oey98}; \\begin{equation} R = \\left(\\frac {250}{ 308\\pi} \\right)^{1/5}L_{\\rm mech}^{1/5}\\rho_0^{-1/5}t^{3/5} \\label{R} \\end{equation} and \\begin{equation} P = {\\frac{7}{(3850\\pi)^{2/5}}}L_{\\rm mech}^{2/5}\\rho_0^{3/5}t^{-4/5}. \\label{P} \\end{equation} Here, $L_{\\rm mech}$ is the instantaneous mechanical luminosity of from the central stars from all sources, $\\rho_0$ is the density  of the ambient medium and $t$ is the time. The particle density is given in terms of the ambient density by $n_0=\\rho_0/\\mu m_H$, and the ambient pressure $P_0=nkT_0$. Eliminating $t$ between equations \\ref{R} and \\ref{P}, we obtain the pressure in terms of the radius of the \\HII\\ region, assuming that the ionized plasma layer is thin compared with the total radius: \\begin{equation} P = {\\frac {7 } {(3850\\pi)^{2/5}}}\\left(\\frac{250}{ 308\\pi }\\right)^{4/15}{\\left(\\frac {L_{\\rm mech}}{ \\mu m_Hn}\\right)}^{2/3} {\\frac{\\mu m_H n_0}{ R^{4/3}}} \\label{P1} \\end{equation}   This theory emphasizes the important role of the pressure in the ISM in determining the radius of an individual \\HII\\ region, and its evolution with time. In particular, with higher ISM pressure, the \\HII\\ region stalls at smaller radius, and therefore the mean dust temperature in the atomic and molecular shell around the \\HII\\ region should be higher in high-pressure environments. This relationship was investigated in detail in  Paper I, but the models presented there were for only a single cluster mass and single set of chemical abundances. However, the dynamical evolution evolution was treated in a more sophisticated self-consistent manner which allowed for the time-dependence of both the mass-loss driven mechanical energy input of the central stars, and the energy input from supernova explosions.  Equation \\ref{R} shows that the size of an \\HII\\ region is correlated, albeit weakly, with the $L_{\\rm mech}$ of the cluster, that is to say, with cluster mass, metallicity and with the age of the cluster. The radius of the \\HII\\ region and the properties of the exciting cluster together determine the characteristic peak of the far-IR dust re-emission bump. Likewise, equation \\ref{P1} shows that the ionization parameter of the \\HII\\ region is determined not only by the flux of ionizing photons from the cluster (which is a function of metallicity, cluster mass and time; $N_*(Z, M_{clust}, t)$ but also by the radius of the \\HII\\ region and by $L_{\\rm mech}$.  In this paper, we use \\emph{Starburst 99 v.5} \\citep{Leitherer99} stellar synthesis models to investigate the effect of metallicity, the stellar initial mass function and the cluster mass function in controlling the excitation of the populations of \\HII\\ regions in both disk and starburst galaxies galaxies. In particular, we will investigate what constraints are placed by observations of H$\\alpha$ luminosity functions upon the cluster mass functions, and what are the probability distributions of \\HII\\ region ionization parameters in galaxies as a function of metallicity and pressure in the ISM.  In the third paper in this series, we will investigate what effect these have upon the emission line diagnostics of ensembles of \\HII\\ regions in galaxies with the view to assisting the analysis of the strong line emission spectra of distant unresolved galaxies in order to derive star formation rates, metallicities and ISM pressures. ",
        "conclusions": "\\subsection{Sensitivity of Star Formation Rates to the IMF} Because they are so important in drawing conclusions about the both evolution of galaxies and the global star formation rate throughout cosmic time, we have investigated the sensitivity of commonly-used star formation indicators to assumptions about the IMF. We have found that, even ignoring the possibility that the form of the upper stellar mass part of the IMF may be changed by the local environment (\\emph{e.g} truncated or top-heavy IMFs), the uncertainty in star formation rates introduced by assumptions about the form of the lower stellar mass part of the IMF is considerable. For both determinations based upon the H$\\alpha$ line, and upon the IR luminosity,  $L_{{\\rm IR}}$ errors of order a factor of two arise directly from plausible differences in the assumed form of the lower mass part of the IMF.  \\subsection{Cluster Mass Function} We have demonstrated the the luminosity function of \\HII\\ regions in galaxies can be fit better by a log-normal cluster mass distribution, rather than a simple power-law, as has been assumed hitherto. The  mean cluster mass is about $m_0 = 100M_{\\odot}$ and $\\sigma \\sim1.7$.  The probability of forming a stellar cluster increases just as the probability of forming a single star decreases, perhaps indicating that cluster formation can be considered as a continuation of the IMF to higher masses. This is probably the result of angular momentum constraints preventing the formation of massive single stars, so that a cluster of stars is formed instead. This conclusion is consistent with those of \\citet{Oey04}, based upon number counts of cluster stars in the SMC and the embedded cluster mass function inferred by \\citet{Lada03}. This result is also consistent with the idea that it is the density distribution in a turbulent ISM which determines the form of both the stellar and cluster IMF \\citep{Elmegreen02,Elmegreen04,Padoan02,Padoan05,Bonnell03}.   \\subsection{Ionization Parameter Distribution} We have demonstrated that the ionization parameter of \\HII\\ regions, hitherto treated as a free variable, is in fact determined largely by the instantaneous ratio of the ionizing photon flux to the mechanical energy flux of the central stars. Although this conclusion has been obtained using a simple 1-D spherical evolution model for the \\HII\\ region, it should remain valid in more complex geometries, since wherever un-ionized or molecular inclusions exist, these will be pressure confined by the hot shocked stellar wind gas (modulus the recoil momentum of the gas streaming off the ionization fronts), and this gas pressure is coupled with the overall size of the \\HII\\ region, which itself determines the dilution of the radiation field of the central stars. Because the ratio of the the ionizing photon flux to the mechanical energy flux of the central stars is a strong function of metallicity according to the stellar atmosphere and stellar wind theory of massive stars, the ionization parameter should also be quite sensitive to metallicity. This appears to be the case for the SDSS galaxies. For individual \\HII\\ regions there is also some indication of such a systematic variation of ionization parameter with metallicity (\\emph{c.f.} figure 6 of \\citet{Dopita00}). The ionization parameter is only weakly dependent upon pressure in the interstellar medium and upon cluster mass; $\\log q \\propto P_{0}^{-1/5}$ and $\\log q \\propto M_{\\rm cl}^{1/5}$, respectively.  Finally, we have predicted that there should exist a substantial fraction ($20-30$\\%) of the total H$\\alpha$ flux  coming from the faint, extended, old and evolved \\HII\\ regions with $\\log q \\le 7$cm s$^{-1}$, or $\\log{\\cal{U}} \\le -3.5$. This emission has the appropriate flux, and should have the appropriate line ratios to be identified with the diffuse ionized gas (DIG) component of the ISM in galaxies."
    },
    "0606/hep-ph0606107_arXiv.txt": {
        "abstract": "The axion is one of the best motivated candidates for particle dark matter. We study and update the constraints imposed by the recent CMB and LSS experiments on the mass of axions produced by the misalignment mechanism, as a function of both the inflationary scale and the reheating temperature. Under some particular although not unconventional assumptions, the axionic field induces too large isocurvature perturbations.  Specifically, for inflation taking place at intermediate energy scales, we derive some restrictive limits which can only be evaded by assuming an efficient reheating mechanism, with $\\trh>10^{11}$ GeV. Chaotic inflation with a quadratic potential is still compatible with the axion scenario, provided that the Peccei-Quinn scale $f_a$ is close to $10^{10}$ or $10^{11}$ GeV. Isocurvature bounds eliminate the possibility of a larger $f_a$ and a small misalignment angle. We find that isocurvature constraints on the axion scenario must be taken into account whenever the scale of inflation is above $10^{12}$ GeV; below this scale, axionic isocurvature modes are too small to be probed by current observations. ",
        "introduction": "The fact that the strong sector of the Standard Model conserves the discrete symmetries P and CP while the electroweak sector doesn't, also known as the \\textit{strong CP problem}, is considered a serious puzzle for modern particle physics~\\cite{reviews}.  The most elegant and compelling solution to this problem was proposed in 1977 by Peccei and Quinn~\\cite{PQ} with the introduction of a new $U(1)_{\\rm{PQ}}$ global symmetry at high energies. The Peccei-Quinn (PQ) scalar field $\\psi=\\rho/\\sqrt{2}\\,e^{i\\Theta}$ has a potential of the form \\be V(\\psi)= {\\lambda\\over4}\\left(\\rho^2-{f_a^2\\over N^2}\\right)^2 = V_{\\rm PQ} - \\half m_\\psi^2 \\rho^2 + {\\lambda\\over4} \\rho^4\\,, \\ee where $N$ is the number of degenerate QCD vacua associated with the color anomaly of the PQ symmetry. Spontaneous symmetry breaking (SSB) occurs when the energy density of the universe falls below $V_{\\rm PQ}$ and the field acquires a vacuum expectation value (VEV) $\\rho=f_a/N$. The axion~\\cite{axion} is then the Goldstone boson of the broken PQ symmetry, \\be a(\\vec x)=\\frac{f_a}{N}\\Theta(\\vec x)\\,.  \\ee Note, however, that SSB is effective only when the typical fluctuations on $\\delta\\rho$ are smaller than $f_a/N$. If either $T$ or $\\hinf$ are of order $f_a/N$ at reheating or during inflation, thermal~\\cite{Harari} or quantum~\\cite{LythStewart2} fluctuations (respectively) will modify the effective potential and restore the PQ symmetry, with a coherence length comparable to the Hubble scale at that time.  As the universe expands, its energy decreases to about $\\Lambda_{\\rm QCD} \\sim 200$~MeV, and non-perturbative instanton effects tilt the previously flat axion potential, explicitly breaking the residual symmetry~\\cite{witten}: \\be\\label{potential} V(a)\\simeq \\ma^2{f_a^2\\over N^2}(1-\\cos\\Theta)~. \\ee The axion field acquires a mass about the minimum of the potential that depends on the temperature in the vicinity of $T \\sim \\Lambda_{\\rm QCD}$ as~\\cite{Gross:1980br,Fox:2004kb} \\be \\label{mass_timedep}\\ma(T)\\simeq m_a(0) C \\left(\\frac{\\Lambda_{\\rm QCD}}{200\\rm{MeV}} \\right)^{\\frac{1}{2}}\\left(\\frac{\\Lambda_{\\rm{QCD}}}{T}\\right)^4\\,, \\ee where $C$ is a model-dependent factor calculated in Refs.\\ \\cite{Gross:1980br,Fox:2004kb} to be of the order of $C\\simeq 0.018$. The mass finally reaches its asymptotic zero-temperature value $m_a \\equiv m_a(0)$: \\begin{equation}\\label{mass} m_a = {\\sqrt z\\over1+z}\\,{f_\\pi m_\\pi\\over f_a} = 6.2\\,\\mu{\\rm eV}\\,\\left({10^{12}\\ {\\rm GeV}\\over f_a}\\right)\\,. \\end{equation} Here $z=m_u/m_d\\simeq0.56$ is the mass ratio of up to down quarks, while $m_\\pi$ and $f_\\pi$ are respectively the pion mass and decay constant. The field equation of motion is \\begin{equation} \\label{axion_field_eq} \\ddot{a}   + 3H \\dot{a} +V'(a)+ {1\\over R^2} \\nabla^2 a =0\\,, \\end{equation} where $V'(a)=\\partial V/\\partial a$, $\\nabla^2$ is the comoving laplacian, and $R$ is the scale factor of the universe. If the axion field is initially displaced from the minimum of the potential when it acquires a mass, it starts oscillating in the potential described by Eqs.~(\\ref{potential}), (\\ref{mass_timedep}), (\\ref{mass}).  The axion coupling to the rest of matter is inversely proportional to the symmetry breaking scale~\\cite{curr,PDG}: \\begin{equation}\\label{couplings} {\\cal L}_{\\rm int} = g_\\gamma\\,{\\alpha\\over\\pi}\\,{a(x)\\over f_a}\\, \\vec E\\cdot\\vec B +i\\,g_f\\,m_f\\,{a(x)\\over f_a}\\,\\bar f\\,\\gamma_5\\,f\\,, \\end{equation} where $\\alpha$ is the electromagnetic coupling constant, $\\vec E$ and $\\vec B$ are the electric and magnetic fields, $f$ is a generic fermion and for each species $g_i$ is a model-dependent coefficient of order one.  The main distinction between the various axion models comes from their coupling with electrons; for the ``hadronic'' models, such as the KSVZ model~\\cite{KSVZ}, one has $g_e=0$, while the tree-level coupling does not vanish for the non-hadronic DFSZ models~\\cite{DFSZ}. The other couplings are of the same order.  For example, $g_\\gamma=-0.36$ in the DFSZ model, while $g_\\gamma=0.97$ in the KSVZ model.  In the currently accepted \\textit{invisible axion} model, the scale $f_a$ is in principle arbitrary, well above the electroweak scale so that the axion coupling to matter is weak enough to pass undetected, for the moment. There are at present several experiments searching for the axion in the laboratory, like ADMX~\\cite{ADMX} and CAST~\\cite{CAST}, which have recently reported bounds on the axion coupling to matter~\\cite{bounds}.  Since in the large $f_a$ limit the axion remains effectively decoupled from other species, its fluctuations during inflation could induce isocurvature perturbations in the CMB anisotropy spectrum which would in principle be observable today~\\cite{Axenides:1983hj,Lindeaxion,SeckelTurner,hybrid,TurnerWilczek,LindeLyth,Lyth:1991ub,Shellard:1997mf}. In this paper, we review the consequences of up-to-date bounds on the allowed isocurvature fraction and CDM density for the axion mass.  The paper is organized as follows: in Section \\ref{production} we review the different ways in which axions could be produced, with a particular emphasis on the misalignment angle mechanism.  In Section \\ref{isocurvature} we study the induced isocurvature perturbations. In Section \\ref{window} we present the various constraints bounding the axionic window. Finally, our results are summarized and discussed in Section \\ref{discussion}. ",
        "conclusions": "In this paper, we reanalyzed the bounds on the cosmological scenario in which the cold dark matter is composed of axions plus some other component (like e.g. neutralinos), and the energy density of axions is produced by the misalignment mechanism at the time of the QCD transition. In this case, a fraction of the cosmological perturbations consists of isocurvature modes related to the quantum perturbations of the axion during inflation. We use this possible signature plus the contribution to dark matter energy density as a tracer of axions in the universe.  In that sense, this work is similar to that presented in \\cite{Sikivie:2006ir}. However, we make a stronger statement about the generation of isocurvature modes in the case in which the PQ symmetry is restored by quantum fluctuations during inflation, and we significantly improve previous bounds coming from the non-observation of isocurvature modes in the CMB, in particular, given the recent WMAP 3-year data.  A narrow window for the PQ scale $10^{10}\\,{\\rm GeV} < f_a < 10^{12}$~GeV still remains open, but we show that the main consequence of recent data on cosmological perturbations is to limit the possible energy scale of inflation in this context, in order not to have an excessively large contribution of isocurvature perturbations to the CMB anisotropies. In particular, we show that in the axion scenario, the energy scale of inflation cannot be in the range $8 \\times 10^{12} < M_{\\inf} < 5 \\times 10^{14}$~GeV, unless one of the two situations occurs: either reheating leads to a temperature $T_{\\rm rh}>f_a$, with an efficiency parameter $\\epsilon_{\\rm eff} \\geq 4 \\times 10^{-12}$, and there is another allowed region with $6 \\times 10^{10} \\epsilon_{\\rm eff}^{-1/2} \\,{\\rm GeV} < M_{\\rm inf} < 3 \\times 10^{16}$~GeV; or the misalignment angle is fine-tuned to very small values (this coincidence can be motivated by anthropic considerations) and the upper bound $M_{\\inf} < 10^{13}$~GeV can be slightly weakened (by at most one order of magnitude).  These bounds may be of interest taking into account particle physics experiments searching for the axion, which may help to put very stringent constraints on inflationary models.  Detecting the QCD axion could shed some light on the scale of inflation and possibly into the mechanism responsible for inflation.   In this respect, there is an intriguing possibility that the PVLAS experiment~\\cite{PVLAS} may have observed a pseudo-scalar particle coupled to photons.  The nature of this particle is yet to be decided, since its properties seem in conflict with present bounds on the axion coupling to matter~\\cite{ADMX,CAST}, see however~\\cite{Masso,Mendonca}."
    },
    "0606/astro-ph0606604_arXiv.txt": {
        "abstract": "The X-ray spectra of Anomalous X-ray Pulsars have long been fit by smooth, empirical models such as the sum of a black-body plus a power law.  These reproduce the $\\sim\\!0.5$ to 10\\,keV range well, but fail at lower and higher energies, grossly over-predicting the optical and under-predicting the hard X-ray emission.  A poorly constrained source of uncertainty in determining the true, intrinsic spectra, in particular at lower energies, is the amount of interstellar extinction.  In previous studies, extinction column densities with small statistical errors were derived as part of the fits of the spectra to simple continuum models.  Different choices of model, however, each produced statistically acceptable fits, but a wide range of columns.  Here, we attempt to measure the interstellar extinction in a model-independent way, using individual absorption edges of the elements O, Fe, Ne, Mg and Si in X-ray grating spectra taken with {\\em XMM-Newton}.  We find that our inferred equivalent hydrogen column density \\NH\\ for 4U 0142+61 is a factor of 1.4 lower than the typically quoted value from black-body plus power-law fits, and is now consistent with estimates based on the dust scattering halo and visual extinction.  For three other sources, we find column densities consistent with earlier estimates.  We use our measurements to recover the intrinsic spectra of the AXPs empirically, without making assumptions on what the intrinsic spectral shapes ought to be.  We find that the power-law components that dominate at higher energies do not extend below the thermal peak. ",
        "introduction": "The Anomalous X-ray Pulsars (AXPs) are a group of about six young, isolated neutron stars showing pulsations with periods of the order 10\\,s, whose high-energy luminosity vastly exceeds what is available in rotational energy losses. They are modeled as {\\em magnetars}, their energetics dominated by energy released from a decaying super-strong magnetic field, $\\sim\\!10^{14}\\,$G externally. See Woods \\& Thompson (2004) for a summary of current observations and their interpretation in the context of the magnetar model.  Since their discovery, the X-ray spectra of the AXPs have been fitted with simple continuum models, most commonly the sum of a black body, representing the peak of the spectrum around a few keV, and a power law to account for the emission at higher energies.  The overall spectra are relatively soft, as indicated by high power-law indices. In the fitting process, the interstellar extinction is estimated as well, by including a multiplicative term to take account of the absorption of all elements, parameterized by the hydrogen column density, \\NH.  While these smooth continuum spectral models reproduce the spectra of all AXP well in the range observed by satellites such as {\\em Chandra}, {\\em XMM} and {\\em ASCA} (roughly 0.5 to 10\\,keV), their extrapolation fails at both higher and lower energies.  At energies from tens of keV to $\\sim\\!1\\,$MeV, Kuiper et al.\\ (2004) and Den Hartog (2006) found from RXTE/HEXE and Integral observations that much harder power-law components are present, which dominate the total energetics.  In the optical and infrared, the first detections of AXPs made by Hulleman et al.\\ (2000, 2001) already showed that while the optical and infrared emission is only a tiny part of the energy budget, it is well below (by orders of magnitude) the extrapolation of the X-ray spectra dominated by the power law, yet well above the extrapolation of just a blackbody component.  Thus, the optical and infrared emission also appears to require a separate emission component.  More recently, the situation was complicated even further, when Wang et al.\\ (2006) found from mid-infrared observations from {\\em Spitzer} that at least one of the AXPs (4U 0142$+$61) shows evidence of a circumstellar dusty disc.  The above discrepancies led us to wonder whether the measurements of interstellar extinction were reliable.  These are particularly sensitive to the the low-energy ($<\\!500\\,$eV) part of the spectra (where the absorption is highest), and one would expect that they would depend strongly on the assumed model for the intrinsic emission. For instance, the above-mentioned model consisting of a black-body and a power-law component rises towards lower energies, while a model of, say, two black-body components, would turn over.  Thus, to match a given observation, the former would require a larger column density than the latter.  From previous studies (see Table~\\ref{nhlit}), it is indeed clear that the inferred column density \\NH\\ depends strongly on the assumed intrinsic model, with the differences in \\NH\\ for different models fitted to the same spectrum far exceeding the statistical uncertainty obtained for any given one.  Clearly, without knowing the intrinsic shape of the spectrum, one cannot measure the extinction accurately, and, conversely, without knowing the extinction independently, one can obtain only little information about the true intrinsic spectrum.  Another clue that the extinction is not estimated correctly, and hence that the models used for the intrinsic spectrum are incorrect, comes from variability studies.  AXP spectra have been seen to vary both from one epoch to another, and between phase bins (e.g., Woods et al.\\ 2004; Rea et al.\\ 2005), and, generally, different values of \\NH\\ are found for these different spectra (see Table~\\ref{nhlit}).  Although an intrinsic, variable contribution to the extinction is not impossible in light of the discovery of a possible debris disk around 4U~0142+61 (Wang et al.\\ 2006), it seems unlikely, especially on the time-scale of seconds.  Indeed, discrepancies in fitted parameters obtained from spectra taken with different instruments, which are often blamed on poor cross-calibration, could well be purely an artifact of the fitting process, with differences in sensitivities leading to differences in weight as a function of energy, and hence different values for the parameters, including the extinction.\\footnote{Another source of differences may be the use of different sets of cross-sections and abundances; see, e.g., Weisskopf et al.\\ (2004).}  \\begin{deluxetable}{lccccccl} \\tablecaption{Hydrogen column densities towards the AXPs inferred using different models.\\label{nhlit}} \\tablewidth{0pt} \\tabletypesize{\\footnotesize} \\tablehead{ \\colhead{Object} &\\colhead{Telescope/}& \\multicolumn{5}{c}{\\dotfill Model\\tablenotemark{a}\\dotfill} & \\colhead{Reference}\\\\ & \\colhead{instrument}& \\colhead{PL} & \\colhead{PL+BB} & \\colhead{BB+BB} & \\colhead{BR} & \\colhead{BR+BB} } \\startdata 4U 0142+61 & XMM/EPIC &       \\nodata & 0.96 &\\nodata & 0.82 &   0.68 & G\\\"ohler et al.\\ (2004) \\\\ & Chandra/HETGS &            1.43: &   0.88 &\\nodata & 0.92 &   0.69 & Juett et al.\\ (2002)\\\\ & ASCA/SIS,GIS &             \\nodata & 0.95 &\\nodata &\\nodata & 0.90 & White et al.\\ (1996)\\\\ 1E 2259+586&XMM/EPIC &        \\nodata & 0.94-1.10\\tablenotemark{b} & \\nodata &\\nodata &\\nodata & Woods et al.\\ (2004)\\\\ & Chandra/ACIS &             \\nodata & 0.93 &\\nodata &\\nodata &\\nodata & Patel et al.\\ (2001)\\\\ & ASCA/PSPC &                1.14: &   0.85 & 0.63 &   0.8 &   \\nodata & Rho \\& Petre (1997)\\\\ 1E 1048.1$-$5937&XMM/EPIC&    \\nodata & 0.95-1.10\\tablenotemark{b} & 0.55-0.67\\tablenotemark{b} &\\nodata &\\nodata  & Tiengo et al. (2005)\\\\ & BeppoSAX/LECS &            1.54: &   0.45 & 0.14 &  \\nodata &\\nodata & Oosterbroek et al.\\ (1998)\\\\ 1RXS J170849.0& BeppoSAX/LECS&\\nodata & 1.3-1.7\\tablenotemark{c} & \\nodata &\\nodata &\\nodata & Rea et al.\\ (2005)\\\\ \\multicolumn{1}{r}{$-$400910}& BeppoSAX/LECS& \\nodata & 1.1-1.6\\tablenotemark{c} &\\nodata &\\nodata &\\nodata & Rea et al.\\ (2003) \\enddata \\tablecomments{Column Densities are in units (10$^{22}$\\,cm$^{-2}$). For all models listed, the reduced $\\chi^2$ of the fits was less than 2 (those marked with colons are relatively poor fits).} \\tablenotetext{a}{Acronyms are PL: Power Law, BB: Black Body, BR: Bremsstahlung.} \\tablenotetext{b}{Ranges represent measurements made at different epochs.  One does not expect the column density to vary with time (see text).} \\tablenotetext{c}{Range represents measurements in different phase bins.  Again, one would not expect the column density to vary with rotational phase.} \\end{deluxetable}  From a physical perspective, it is not clear why there should be a power-law component in the spectra of AXPs.  Recent physical models, such as those presented by Lyutikov \\& Gavriil (2006) can account for a power-law like high-energy tail, but do not predict a soft part.  In these models, the blackbody surface flux is modified by interactions in the outer atmosphere or magnetosphere, for example through the inverse-Compton scattering of photons off high-energy particles.  This creates an extended high-energy tail but leaves the spectrum at low energies (the Rayleigh-Jeans side of the blackbody spectrum) unaffected.  Absorption in the interstellar medium of X-ray photons is primarily by the photo-electric capture by electrons in inner shells of metals and helium.  Since the creation of X-ray absorption models (Morrison \\& McCammon 1983; Balucinska-Church \\& McCammon 1992), great advances have been made in our understanding of absorption edge structure and energies (e.g., Juett et al.\\ 2003), largely driven by the improved resolution and sensitivity of X-ray observatories.  We are also beginning to understand more about interstellar abundances (e.g., Lodders 2003; but see \\S5 below).  With these advances, it is now possible to measure individual elemental absorbing columns to a source independently, and so recover the intrinsic spectrum without further assumptions.  In this paper, we measure the extinction to the four best-studied AXPs by analyzing the individual absorption edges present in the sensitivity range of XMM/RGS, and use this to derive intrinsic X-ray spectra for each AXP, as well as to estimate reddenings at optical and infrared wavelengths.  In \\S2, we present the data sets we use and their reduction.  We describe in \\S3 how we infer individual element column densities, and what atomic data we use.  In \\S4, we use Monte-Carlo simulations to estimate the uncertainties of our measurements. We present our results in \\S5, and show de-extincted X-ray spectra in \\S6. We continue by deriving optical extinctions for all sources in \\S7, and discussing the overall spectral energy distribution of the best-studied source, 4U 0142+61, in \\S8.  We briefly summarize our results and look forward to future work in \\S9. ",
        "conclusions": "We have attempted to measure the extinction to the AXPs without making assumptions about what their intrinsic spectral shapes might be.  With our resulting best estimates, we derived intrinsic spectra, which can be compared with each other as well as with predictions, such as those from simulations and semi-analytic models that are now being produced within the magnetar framework (e.g., Lyutikov \\& Gavriil 2006; R.\\ Fern\\'andez \\& C.\\ Thompson 2006, pers.\\ comm.).  Apart from these comparisons, future work could include extending our analysis to other sources (once better spectra become available), or improving the precision of our measurements using further {\\em XMM-Newton} observations (some already taken but not yet public) or {\\em Chandra} spectra (some available).  More interestingly, by measuring the run of reddening with distance along the line of sight to the AXPs (using field stars), our extinction estimates can be used to estimate distances and thus determine the intrinsic luminosities of the AXPs. It turns out that although the extinctions are not very well determined, the AXPs fall into areas of rapidly rising extinction associated with spiral arms, and so can be well localized (Durant \\& Van Kerkwijk, 2006).  Finally, our reddening estimates will place on much firmer footing inferences from further optical and infrared studies, such as could be used, e.g., to uncover the nature of the break seen between $V$ and $B$ in 4U 0142+61 (Hulleman et al.\\ 2004) or to measure the precise parameters of the possible debris disk around that source (Wang et al.\\ 2006)."
    },
    "0606/astro-ph0606118_arXiv.txt": {
        "abstract": "{ The feasibility of observing the \\ceo (3--2) spectral line in cold clouds with the APEX telescope has been tested.  As the line at 329.330 GHz lies in the wing of a strong atmospheric H$_2$O absorption it can be observed only at high altitude observatories. Using the three lowest rotational levels instead of only two helps to narrow down the physical properties of dark clouds and globules.  The centres of two \\ceo maxima in the high latitude low mass star forming region \\object{CG 12} were mapped in \\ceo (3--2) and the data were analyzed together with spectral line data from the SEST.  The \\TMB(3--2)/\\TMB(2--1) ratio in the northern \\ceo maximum, \\cgnp, is 0.8, and in the southern maximum, \\cgsp, $\\sim$2.  \\cgn is modelled as a 120\\arcsec \\ diameter (0.4pc) cold core with a mass of 27 \\Msun. A small size maximum with a narrow, 0.8 \\kmps, \\ceo (3--2) spectral line with a peak temperature of \\TMB $\\sim$ 11 K was detected in \\cgsp.  This maximum is modelled as a 60\\arcsec$-$80\\arcsec \\ diameter ($\\sim$0.2pc) hot (80 K $\\lesssim$ \\Tex $\\lesssim$ 200 K) $\\sim$1.6 \\Msun \\ clump.  The source lies on the axis of a highly collimated bipolar molecular outflow near its driving source. This is the first detection of such a compact, warm object in a low mass star forming region. ",
        "introduction": "The two lowest rotational transitions of \\twcop, \\thco and \\ceo have been used extensively for studies of dark clouds and globules. Also the $J=3-2$ lines of \\twco and \\thco can be routinely observed. In contrast, there are very few published observations of the $J=3-2$ transition. Besides the nine point \\ceo (3--2) map of S106 (Little et al. 1995) and mapping of the centre of the pre-stellar cloud core L1689B (\\cite{jessop-WT}) only pointed observations are available (e.g., J\\o rgensen et al. 2002, 2004; Sch\\\"oier et al (2002); Lee et al. 2003). The addition of the \\ceo (3-2) transition to the lower transitions would, however,  substantially improve the accuracy of physical parameters derived from \\ceo data as the relative line intensities depend heavily on the cloud temperature  and density.    The SEST telescope has been  used to map numerous southern molecular clouds in \\ceo (1--0) and (2--1).  The new 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 f\\\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}) provides an excellent opportunity to complement these data with \\ceo (3-2) observations.  The APEX beam size at 329 GHz, $\\sim$19\\arcsec, is about equal to  the SEST beam size at the \\ceo (2--1) frequency, 24\\arcsec. This simplifies the comparison of APEX and SEST data.    In this Letter we demonstrate the use of \\ceo (3-2), (2--1) and (1--0) observations for radiative transfer modelling in the case of \\object{Cometary Globule 12}, CG 12, which is a high-latitude star forming region at a distance of 630 pc (Williams et al. 1977). The head of the cloud contains a highly collimated bipolar outflow (White et al. 1993) and two compact \\ceo maxima separated by about 3' (Haikala \\& Olberg 2006, Paper I). The centres of these maxima were mapped in \\ceo (3--2). Based on these observations a simple model is derived for both cores. ",
        "conclusions": "\\label{sec:conclusions}  The addition of the \\ceo (3--2) transition to the \\ceo data on the two lower transitions restricts strongly the parameter space of possible cloud models.  The \\ceo (3--2) observations of \\cgn confirm the cold temperature of this core. The \\ceo (3--2) line profile varies more over the mapped region than the \\ceo (2--1) profile. The \\ceo (3--2) line is therefore probably better suited for studying the small-scale stucture of molecular clouds than the lower transitions of this molecule.  The detection of high, $\\sim$10 K, \\ceo (3--2) line temperatures in \\cgs was unexpected and forced us to reconsider the temperature and density structure of the material traced by \\ceo (2--1) and (1--0) presented in paper 1.  The small, hot clump lies projected on the edge of a dense core traced by \\dcoplp. It also lies on the axis of a highly collimated bipolar molecular outflow near its driving source. This suggests that the two are related. As the \\ceo lines from the hot spot are narrow and appear at the same radial velocity as the parent cloud, they probably do not originate in shocked gas. Neither is direct radiative heating from the strong IRAS point source 13547-3944 likely. The heating mechanism and the relation of the hot spot to the collimated outflow remain therefore unspecified.  Further studies of this phenomenom seem warranted as they may lead to a better understanding of the interaction between newly born stars and the surrounding molecular material.  Better spatial resolution can be achieved  by observing still higher \\ceo transitions which should be detectable in \\cgsp. The sizes of the CG 12 \\ceo maxima are well suited for the future ALMA interferometer"
    },
    "0606/astro-ph0606432_arXiv.txt": {
        "abstract": "We present the results of a program to acquire photometry for eighty-six late-M, L, and T dwarfs using the Infrared Array Camera (IRAC) on the {\\it Spitzer Space Telescope}. We examine the behavior of these cool dwarfs in various color-color and color-magnitude diagrams composed of near-IR and IRAC data. The T dwarfs exhibit the most distinctive positions in these diagrams. In M$_{5.8}$ versus [5.8]-[8.0], the IRAC data for T dwarfs are not monotonic in either magnitude or color, giving the clearest indication yet that the T dwarfs are not a one parameter family in \\teff. Because metallicity does not vary enough in the solar neighborhood to act as the second parameter, the most likely candidate then is {\\it gravity}, which in turn translates to {\\it mass}.  Among objects with similar spectral type, the range of mass suggested by our sample is about a factor of five ($\\sim$70 $M_{Jupiter}$ to $\\sim$15 $M_{Jupiter}$), with the less massive objects making up the younger members of the sample.  We also find the IRAC 4.5 \\mic fluxes to be lower than expected, from which we infer a stronger CO fundamental band at $\\sim$4.67 \\mic. This suggests that equilibrium CH$_4$/CO chemistry underestimates the abundance of CO in T dwarf atmospheres, confirming earlier results based on $M$-band observations from the ground.  In combining IRAC photometry with near-IR $JHK$ photometry and parallax data, we find the combination of $K_S$, IRAC 3.6 \\mic, and 4.5 \\mic bands to provide the best color-color discrimination for a wide range of M, L, and T dwarfs.  Also noteworthy is the M$_{K_S}$ versus $K_S$-[4.5] relation, which shows a smooth progression over spectral type and splits the M, L, and T types cleanly. ",
        "introduction": "The study of low-mass dwarfs has progressed enormously in the last decade: from the strictly theoretical beginnings for objects of sub-stellar mass (e.g. Kumar 1963; Grossman 1970; Nelson, Rappaport, \\& Joss 1985) to the actual discovery in the late-1980's and mid-1990's of objects in two new spectral classes later than M (Becklin \\& Zuckerman 1988; Nakajima et al. 1995). As a result of a large amount of observational effort over the past decade, the temperature sequence of dwarfs has been extended down to T $\\sim$ 700 K (Golimowski et al. 2004a). Several hundred L dwarfs and several dozen T dwarfs have now been identified, leading to well-defined L and T spectral sequences (Kirkpatrick 2005), and to a large body of accurate photometry and quantitative spectroscopy to characterize those spectral types (Leggett et al. 2001, 2002; Basri et al. 2000; Reid et al. 2002).  Also of importance, other groups have been able to obtain accurate trigonometric parallaxes for a large number of these very cool dwarfs (e.g., Dahn et al. 2002; Tinney et al. 2003; Vrba \\etal 2004) and to obtain dynamical masses for a significant sample of stars (e.g., Henry \\& McCarthy 1993; Henry et al. 1999; Lane \\etal 2001; Zapatero Osorio \\etal 2004; Bouy \\etal 2004; Close et al. 2005).  The observational progress was accompanied by an equally dramatic evolution in the sophistication of the theoretical models of cool dwarfs. The improved observational data spurred theoretical model efforts such as Allard et al. (1994) and Brett (1995), who were able to incorporate much more extensive molecular line lists and thus improve the fit to observations and extend the model spectra further into the infrared.  The model atmospheres have since rapidly evolved, incorporating ever-larger molecular line lists, dust formation and dust-settling, and improved treatment of pressure broadening, particularly for the alkali lines that dominate the optical and near-IR spectra of the L dwarfs. More recently, the models have begun to take into account non-equilibrium chemistry (e.g., Noll, Geballe, \\& Marley 1997; Saumon et al. 2000; Burrows, Marley, \\& Sharp 2000; Ackerman \\& Marley 2001; Allard et al. 2001; Lodders \\& Fegley 2002; Marley et al. 2002; Tsuji 2002; Saumon et al. 2003, 2006).  While much of the work on sub-stellar mass objects to date has been carried out in the optical and near-IR, there is great potential to gain even more insight into brown dwarf atmospheres by extending observations into the mid-IR regime.  The advantages include the peaking of the spectral energy distribution for L and T dwarfs between 1$-$4 \\mic and the presence of prominent molecular features such as CH$_4$, H$_2$O, and NH$_3$ (e.g. Burrows et al. 1997; 2001).  Ground-based infrared observations are hampered longward of $K$ band by telluric atmospheric lines, such as OH and H$_2$O, making the atmosphere increasingly opaque at these wavelengths. Thermal emission of the atmosphere and telescope structure also cause very high backgrounds for ground-based observations longward of $K$ band. Thus, space-based instruments with sensitivity in the mid-infrared are required for the study of sub-stellar mass objects in this wavelength regime. It was in this context that the Infrared Array Camera (IRAC) Guaranteed Time Observer (GTO) team decided that it would be important to obtain accurate mid-IR fluxes for a representative sample of M, L, and T dwarfs when the {\\it Spitzer Space Telescope} was launched.  We present in this paper the results of this survey.  In section 2, we describe the sample selection of M, L, and T dwarfs and their basic properties.  In section 3, we summarize the observing strategy used with {\\it Spitzer}/IRAC and the subsequent data reduction and extraction of the IRAC photometry.  In section 4, we present the basic color-color and color-magnitude diagrams for both IRAC and near-IR $JHK$ photometry and discuss trending with color and spectral type.  Finally, in section 5, we conclude with a few tentative interpretations of the observations. ",
        "conclusions": "As shown in the previous sections, the infrared colors of low mass stars and brown dwarfs are generally well-correlated with their spectral types. In turn, the spectral types of M, L, and T dwarfs are reasonably well-correlated with their effective temperatures, though to a lesser degree than for higher mass objects because molecular chemistry and particulate clouds compete aggressively in determining the detailed spectral energy distributions of cool dwarfs.  Theoretical models are beginning to be able to match the observed spectral properties of very cool dwarfs, and we have used those models to help interpret the IRAC photometry in Section 4.  While a detailed theoretical analysis of these new {\\it Spitzer}/IRAC data is beyond the scope of this paper, there are a few conclusions of a general nature that can be extracted at this preliminary stage. To do so, we have generated three color-magnitude plots (Figures 11--13) for most of the T dwarfs in our sample with parallaxes and superposed synthetic magnitudes and colors derived from theoretical spectra for an effective temperature (\\teff) range from 700 K to 1300 K in steps of 100 K at three gravities. These gravities are 10$^{4.5}$ cm s$^{-2}$ (green circles), 10$^5$ cm s$^{-2}$ (red triangles), and 10$^{5.5}$ cm s$^{-2}$ (blue squares).  Figure 11 depicts the absolute magnitude in the IRAC 3.6 \\mic bandpass versus the [3.6]-[4.5] color.  The legend on this figure indicates the T dwarfs included on this figure and they are numbered in order of spectral sub-type.  As Figure 11 demonstrates, despite the fact that the numbering is in order of increasing spectral sub-type, the 3.6 \\mic IRAC data are not monotonic in either magnitude or color.  That they are not in order in magnitude may in part be ascribed to multiplicity.  However, this cannot be the explanation for the majority of these T dwarfs.  Furthermore, the [3.6]-[4.5] colors too are non-monotonic in spectral type.  This is even more obvious in Figure 12, which portrays M$_{5.8}$ versus [5.8]-[8.0]. For example, SDS1346-0031 and 2MA0727+1710 only differ by half a spectral sub-type, but they have very different colors.  The non-monotonicity of these {\\it Spitzer}/IRAC data is the clearest indication yet that the T dwarfs are not a one-parameter family in \\teff, but that more than one parameter is influential in determining the spectroscopic type. Because the metallicity can not generically vary enough in the solar neighborhood to explain this, the extra parameter may be gravity.  This was the conclusion of Burrows et al. (2002) and Knapp et al. (2004) using different datasets at shorter wavelengths, and is confirmed here. For the T dwarfs, gravity will translate into mass (Burrows et al. 1997). From the comparison of the spread in the data in Figures 11 and 12 with the spread with gravity of the theoretical models, we conclude that a range of gravities of approximately a factor of five is represented in the extant T dwarf sample.  This can be converted into a range of masses from $\\sim$70 M$_{Jupiter}$ to $\\sim$15 M$_{Jupiter}$. The less massive objects would also be younger and a range of ages of about a factor of ten (0.2$-$0.3 to 10 Gyr) is indicated (Burrows et al. 1997; Burrows et al 2001).  The latter conclusion might be problematic, but such is indicated by our preliminary analysis.  Figure 13 is an HR diagram of IRAC M$_{4.5}$ versus the IRAC [4.5]-[5.8] color.  The non-monotonicity of the data seen in previous plots color survives in this plot as well.  However, as is abundantly clear in the figure, the theoretical models do not fit these data.  Because CO has a strong spectral feature at $\\sim$4.67 \\mic, one can interpret the discrepancy between theory and the IRAC data as an indication that equilibrium CH$_4$/CO chemistry underestimates the abundance of CO in T dwarf atmospheres. This conclusion was already reached by Golimowski et al. (2004a) using $M$-band measurements from the ground.   The magnitude of the discrepancy translates into an overestimate in the $\\sim$4--5 \\mic flux by factors of 1.5 to 3.0.  It should be noted that the opacities of water in the mid-infrared are still being studied and refined (Schwenke 2002). Continuing ambiguity in the H$_2$O opacities are a source of systematic theoretical error, with the product of the O abundance and the H$_2$O opacity directly connected to the goodness of fit at 5.8 and 8.0 \\mic for the later T dwarfs. Moreover, at the lowest \\teff represented here, the condensation of K into KCl and Na into Na$_2$S could introduce hazes with interesting optical depths.  Finally, the early T dwarf spectra show signs of silicate clouds (Stephens, Marley, \\& Noll 2001; Marley et al. 2002; Burrows et al. 2002) not included in the models we have presented.  Hence, the study of sub-stellar mass object atmospheres is still in its formative years.  Nevertheless, as Figure 14 implies, one can be encouraged that the basics are falling into place. Figure 14 compares a theoretical T dwarf spectral model with \\teff/g = 750 K/10$^{5}$ cm s$^{-2}$ (Burrows et al. 2002; Geballe et al. 2001) with the four IRAC fluxes measured for the T7.5 dwarf GJ 570D. This model was generated in 2002 for GJ 570D to fit its optical spectrum {\\it shortward} of 1.0 \\mic (Burrows et al. 2002).  A slight ($\\sim$40\\%) discrepancy in the 4.5-\\mic bandpass, attributable to a CO abundance excess in its atmosphere, is visible. Despite this, Figure 14, and Figure 7 in the related paper by Burrows et al. (2002), together represent an acceptable fit from 0.6 \\mic to $\\sim$8.0 \\mic and indicate that the IRAC data were successfully anticipated."
    },
    "0606/astro-ph0606268_arXiv.txt": {
        "abstract": "{}{ The aim of the paper is to understand the emission from the photon dominated regions in Cepheus~B, estimate the column densities of neutral carbon in bulk of the gas in Cepheus~B and to derive constraints on the factors which determine the abundance of neutral carbon relative to CO.} {This paper presents 15\\arcmin$\\times$15\\arcmin\\ fully sampled maps of \\CI\\ at 492~GHz and \\twCO\\ 4--3 observed with KOSMA at 1\\arcmin\\ resolution.  The new observations have been combined with the FCRAO \\twCO\\ 1--0, IRAM-30m \\thCO\\ 2--1 and \\CeiO\\ 1--0 data, and far-infrared continuum data from HIRES/IRAS. The KOSMA-$\\tau$ spherical PDR model has been used to understand the \\CI\\ and CO emission from the PDRs in Cepheus~B and to explain the observed variation of the relative abundances of both C$^0$ and CO.} {The emission from the PDR associated with Cepheus~B is primarily at $V_{\\rm LSR}$ between -14 and -11~\\kms. We estimate about 23\\% of the observed \\CII\\ emission from the molecular hotspot is due to the ionized gas in the \\HII\\ region. Over bulk of the material the C$^0$ column density does not change significantly, $(2.0\\pm1.4)\\,10^{17}$\\,cm$^{-2}$, although the CO column density changes by an order of magnitude.  The observed \\cbyco\\ abundance ratio varies between 0.06 and 4  in Cepheus B. We find an anti-correlation of the observed \\cbyco\\ abundance ratio with the observed hydrogen column density, which holds even when all previous observations providing \\cbyco\\ ratios are included. Here we show that this observed variation of C/CO abundance with total column density can be explained only by clumpy PDRs consisting of an ensemble of clumps. At high H$_2$ column densities high mass clumps, which exhibit low C/CO abundance, dominate, while at low column densities, low mass clumps with high C/CO abundance dominate.}{} ",
        "introduction": "Photon Dominated Regions (PDRs) are predominantly neutral, atomic and molecular regions where the physical and chemical processes are dominated by Far Ultraviolet (FUV) radiation ({\\em cf.} \\citet{sternberg1995,hollenbach1997}) and the major carbon-bearing species changes from C$^+$ through C$^0$ to CO. While FUV photons (6 eV$< h\\nu <13.6$~eV) are primarily responsible for the heating of the cloud surfaces via photoelectric effect on dust grains, at larger depths of about 10~mag into the clouds cosmic ray induced heating dominates. Owing  to the clumpy nature of the molecular clouds, contrary to the initial  ideas, PDRs are not strictly confined to the surface of the molecular  clouds. Rather the extended \\CII\\ and \\CI\\ emission observed from  allover the molecular clouds suggest formation of PDRs deep inside the  molecular clouds, at the surfaces of clumps inundated with FUV  photons escaping into the cloud. Further, the chemical abundances of several key species like e.g., CN and HCN depend to a large extent on the abundances of C$^0$ and C$^+$ in the gas phase \\citep{boger2005}. PDRs may therefore be seen as the most  ubiquitous phase of the ISM, observations of which together with proper  modeling provide insight into the physical conditions within the  molecular clouds.  On a more global scale, PDRs provide direct insight  into the chemical evolution of the ISM deluged with FUV photons and PDR models suggest radiation-induced feedback mechanisms  that may regulate star formation rates and the column density of gas  through giant molecular clouds.  To the south of the Cepheus OB3 association of early-type stars lies the giant Cepheus molecular cloud, at a distance of about 730~pc \\citep{blaauw1964}. Cepheus~B, with a size of 4~pc, is the hottest \\twCO\\ component \\citep{sargent1977,sargent1979} and is located near the northwestern edge of the cloud, adjacent to the \\HII\\ region S155 \\citep{sharpless1959}. The interface between the molecular cloud and the OB stars, is delineated by the optically visible \\HII\\ region  S155, whose very sharp edges clearly indicate the presence of ionization fronts bounding the dust/molecular cloud. The OB association itself consists of two subgroups of different ages, with the youngest lying closer to the molecular cloud \\citep{sargent1979}.  \\citet{felli1978} presented radio continuum observations which confirmed that the molecular cloud Cepheus~B is surrounded by an ionization front driven by the radiation field of the bright stars that are part of the youngest generation of the Cepheus~B OB3 association.  The structure of the diffuse ionized gas and its energy balance can be accounted for by the UV radiation coming from the brightest members of the OB3 association (viz., HD217086 and 217061). The physical association between the Cepheus~B molecular cloud and the S155 \\HII\\ region was also confirmed by  the H$_2$CO and recombination line observations of \\citet{panagia1981}.  \\citet{minchin1992} used a combination of low-J CO and far-infrared  (FIR) continuum observations to show that the Cepheus~B molecular cloud  terminates at its north-west edge with a hotspot. The hotspot lies  1\\arcmin\\ to the east of the peak in CO column density. The high temperature of  the hotspot and the associated FIR emission cannot be explained by only  considering the stars in the OB3 association. It was further proposed by  \\citet{minchin1992} and later confirmed using radio and near-infrared  (NIR) continuum observations by \\citet{testi1995} that the hotspot is  indeed energized by a highly reddened main sequence star (NIR-A) of  spectral type B1-B0.5, together with a small cluster of young stars.  In a first study with the KOSMA 3m telescope, \\citet{beuther2000} mapped the Cepheus B cloud in 2--1 and 3--2 transitions of $^{12}$CO, \\thCO, and \\CeiO\\  at 2\\arcmin\\ resolution and used PDR models of \\citet{stoerzer2000} to interpret the emission at four selected positions. The hotspot emission indicated the presence of shocks. Based on the derived local volume densities of $\\sim 2~10^4$~\\cmcub, and average volume density less than 10$^3$~\\cmcub, \\citet{beuther2000} concluded that Cepheus~B is highly clumped with clumps filling  only 2\\% to 4\\%  by volume of the cloud. In this work \\CI\\ at 492~GHz was detected from two selected positions and detection of \\CI\\ emission from deep inside the cloud was also put forward as an additional evidence for the clumpiness of Cepheus~B.   In clumpy or homogeneous PDR scenario both \\CI\\ and \\twCO\\ 4--3 emission arise from regions closer to the surfaces of the FUV irradiated clumps or molecular cloud surfaces respectively.  Thus the \\CI\\ and \\twCO\\ 4--3 emission trace similar regions.  While \\CI\\ emission necessarily implies photodissociation of molecular CO, the higher kinetic temperature (the upper energy level being $56$~K) required to excite the \\twCO\\ 4--3 line suggest origin of this emission from the outer surfaces of dense (since the critical density of the \\twCO\\ 4--3 transition for collisions with H$_2$ is $\\sim 10^5$~\\cmcub) clumps embedded in a UV radiation field.  Thus in order to trace the PDR we have observed largescale ($15'\\times15'$) fully-sampled maps of the \\CI\\ and \\twCO\\ 4--3  emission from the Cepheus~B molecular cloud. At a common resolution of 1\\arcmin, this dataset has been combined with FCRAO \\twCO\\ 1--0, IRAM-30m \\thCO\\ 2--1 and \\CeiO\\ 1--0 data, and FIR continuum data from HIRES/IRAS.  This allows for a detailed analysis of the PDRs in Cepheus B, since  observed transitions span a substantial range of critical densities and temperatures (Table~\\ref{tab_exc}).   \\begin{table}[hbt] \\caption{Parameters of the observed transitions used to study the PDRs in Cepheus~B. \\label{tab_exc}} \\begin{center} \\begin{tabular}{llll} \\hline \\hline Species & Transition & $E_{\\rm u}$/k & $n_{\\rm cr}$ \\\\ && (K) &  (\\cmcub)\\\\ \\hline \\CI\\ & $^3P_1\\rightarrow^3P_0$ & 23.6 & 1~10$^3$\\\\ \\twCO\\  & 1--0 & 5.6 & 4.8~10$^3$ \\\\ \\thCO\\  & 2--1 & 16.8 & 3.2~10$^4$ \\\\ \\twCO\\ &  4--3 & 56  & 2.6~10$^5$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We have presented the first largescale fully-sampled maps of \\CI\\ (492~GHz) and \\twCO\\ 4--3, of the edge-on PDR, Cepheus~B. The \\CI\\ emission does not delineate the PDR clearly, rather it traces the bulk emission of the more quiescent part of the molecular ridge extending to the east. Bulk of the PDR gas emits at velocities between -14 and -11~\\kms. The \\twCO\\ 4--3 emission is strongly dominated by the hotspot present toward the north-western tip of the molecular cloud. We observed \\twCO 4--3/1--0 intensity ratios in excess of $1.5$ towards the north-western edge of the cloud. Such ratios can only be explained in terms of temperature gradients along the line of sight.  Excitation temperature, derived from the observed \\twCO\\ 4--3/1--0 ratio, varies between 20~K (in the quiescent parts of the cloud) to  50~K (close to the hotspot). We conclude that \\CI\\ primarily traces the higher column density regions while \\twCO\\ 4--3 traces the higher temperature and density regions.  In Cepheus~B we find that N(C) remains nearly constant although N(CO) varies over one order of magnitude.  The observed \\cbyco\\ ratio varies between 0.06 and 4.0 and shows a direct anti-correlation with the hydrogen column density N(H$_2$).  As estimated from the FIR  continuum emission, the FUV radiation field in Cepheus~B varies between 25 \\& 1500 in Draine units, with an average value of 100 over most of the \\CI\\ emitting region. We have used the KOSMA-$\\tau$ PDR model to explain the observed intensity ratios and the observed \\cbyco\\ vs N(H$_2$) anti-correlation.  These PDR model calculations show that for $\\chi = 100$, the observed CI/\\thCO2--1 and CI/\\twCO4--3 ratios indicate local hydrogen densities between 10$^4$--10$^5$~\\cmcub\\ with even higher values at a few positions. The \\twCO\\ 4--3/1--0 ratio can be reproduced by the PDR models and the modelled clumps indeed exhibit outwardly increasing temperature gradients, as predicted from LTE analysis.  The observed global \\cbyco\\ vs \\nhtwo\\ anti-correlation and its slope cannot be reproduced by single spherical PDR models or by plane-parallel homogeneous slabs with or without inclination. We were able to explain the observed anti-correlation between \\cbyco\\ and N(H$_2$) considering an ensemble of clumps of the same mass and density (n=10$^5$~\\cmcub) for each observed position. We find that (i) the \\cbyco\\ ratio is characteristic of the dominant mass in the ensemble, clumps of smaller mass have higher \\cbyco\\ ratio and (ii) clumps of larger mass have to necessarily dominate emission along lines of sight with higher column density in order to reproduce the observed lower \\cbyco\\ ratio."
    },
    "0606/astro-ph0606742_arXiv.txt": {
        "abstract": "We consider the possibility that the ultra-high-energy cosmic ray flux has a small component of exotic particles which create showers much deeper in the atmosphere than ordinary hadronic primaries.  It is shown that applying the conventional AGASA/HiRes/Auger data analysis procedures to such exotic events results in large systematic biases in the energy spectrum measurement which may distort the shape of the measured spectrum near the expected Greisen-Zatsepin-Kuzmin (GZK) cutoff energy.  Sub-GZK exotic showers may be mis-reconstructed with much higher energies and mimic super-GZK events.  Alternatively, super-GZK exotic showers may elude detection by conventional fluorescence analysis techniques. ",
        "introduction": "The measurement by the AGASA ground array \\cite{Takeda:1998ps} of an unabated ultra-high-energy cosmic ray (UHECR) flux beyond the Greisen-Zatsepin-Kuzmin (GZK) energy of $\\sim 5\\times 10^{19}$ eV \\cite{Greisen:1966jv} \\cite{Zatsepin:1966} has provoked much speculation in the particle astrophysics community.  Meanwhile, the HiRes fluorescence experiment has reported evidence of the expected GZK suppression \\cite{Abbasi:2002s} \\cite{Abbasi:2005s}.  Explanations of the AGASA super-GZK events have mainly focused on mechanisms by which the energetic particles may evade interaction with the CMB.  If the UHECRs are ordinary hadronic particles, then their sources must be very close, within $\\sim$100 Mpc.  No local astrophysical sources have been compellingly identified, though interesting possibilities do exist \\cite{Farrar:2005me}.  In some models, local decays of heavy cosmological relic particles create the UHECR flux.  In this case, the flux of conventional particles is usually expected to be dominated by energetic photons from $\\pi^0$ decay.  These types of models are starting to become constrained by limits on the photon fraction of the UHECR flux \\cite{Gelmini:2005wu} \\cite{Risse:2005jr} \\cite{Risse:2005hi} \\cite{Rubtsov:2006tt}.   Alternatively if the UHECRs originate very far away, then in order to propagate, their interactions with the CMB must be greatly suppressed.  If, for example, Lorentz invariance is violated for particles with large boosts relative to our local cosmological rest frame, then the thresholds for the various interactions and hence the GZK feature can be shifted towards higher energies \\cite{Kirzhnits:1972sg} \\cite{Sato:1972wy} \\cite{Gonzalez-Mestres:1997mq} \\cite{Gonzalez-Mestres:1997zg} \\cite{Coleman:1998ti} \\cite{Stecker:2001vb}.  Other conventional particles with small interaction cross-sections include neutrinos and axions.  Neutrinos could induce hadronic air showers if their hadronic interaction cross-sections were greatly enhanced due to new physics at the TeV scale\\cite{Feng:2001ib} \\cite{Anchordoqui:2001cg} \\cite{Ahn:2003qn} \\cite{Anchordoqui:2005ey}.  Axion-like particles could propagate large distances before converting locally to photons which then shower in the atmosphere \\cite{Gorbunov:2001gc}.  Of course, one can always postulate the existence of a new species of particle with interaction processes tuned to solve the propagation problem and to still make detectable air showers.  For example, the GZK energy threshold for a heavy particle of mass $M_X$ is increased by a factor of $M_X/M_{proton}$.  To create detectable air showers which can mimic ordinary hadronic events at moderate zenith angles, the cross-section $\\sigma_{X-nucleon}$ would need to be in the range 10-1000 mb.  Specific models for such ``UHEcrons'' were analyzed in \\cite{Chung:1997rz} \\cite{Albuquerque:1998va} \\cite{Berezinsky:2001fy} \\cite{Kachelriess:2003yy}.  Some recent experimental limits from CDMS and Edelweiss data on heavy particle fluxes were reported in \\cite{Albuquerque:2003ei}.  Also, searches for neutrino showers at large zenith angles place limits on the flux of very deeply penetrating particles with cross-sections much smaller than we are considering in this work \\cite{Anchordoqui:2003jr} \\cite{Baltrusaitis:1985mt}.  Regardless of the mechanism by which the super-GZK cosmic rays evade interaction with the CMB, the properties of the air showers produced at the Earth may be quite model-dependent.  However, all current experiments analyze their air shower data with the assumption that the primary particles are ordinary hadrons.  There is no guarantee that analyzing exotic events in this way would yield sensible results.  For example, exotic events which look very different from conventional air showers could be rejected by event selection criteria designed to select precisely those events which look like ordinary air showers and not like noise.  In some cases, the data from an exotic shower could be mis-reconstructed to look somewhat like a ordinary hadronic shower.  In this case, even though the event would pass the event selection cuts, the reconstructed quantities such as the energy and arrival direction would very likely not be reliable.  In the next three sections, we will discuss how exotic events would be treated by the detectors and the data analysis techniques of the AGASA, HiRes, and Auger experiments.  We will focus specifically on the case where exotic cosmic rays with small cross-sections $\\sigma_{X-nucleon}$ produce deeply-penetrating air showers at moderate zenith angles. ",
        "conclusions": "In this note, we have described how exotic deep shower events would be treated by the experimental techniques of the AGASA, HiRes, and Auger experiments.  In performing a energy spectrum analysis, each experiment strives to develop an analysis procedure which is valid for the vast majority of cosmic ray events, which are presumed to be ordinary hadrons.  AGASA assumes that cosmic ray air showers look like those predicted by hadronic simulations.  HiRes and Auger assume that the air showers have an $X_{max}$ distribution similar to previously measured distributions when performing the fluorescence reconstruction.  Finally, Auger assumes that all cosmic ray air showers may be characterized with a single energy conversion function calibrated with hybrid data.  Because these spectrum analyses are not specifically tuned to search for a possibly small flux of exotic events, the exotic events would appear as a small amount of noise above the underlying spectrum of ordinary cosmic rays.  The effect of analyzing deep showers as ordinary showers would seem to be the same for all three experiments--in most cases, the shower energy would be overestimated.  However, the magnitude of the overestimate can be very different for each experiment.  A deep shower inducing a large factor of $\\sim$3 overestimate in AGASA would only give a $\\sim$50$\\%$ overestimate in the Auger surface detectors, due mainly to the flat dependence of the ground muon flux on the shower position.  It is also not unreasonable to expect possible factor of $\\sim$2 overestimates in the HiRes1 PCF analysis, if the mis-reconstructed deep showers still pass the event selection cuts.  In a spectrum falling rapidly as the third power of energy, a 25$\\%$ fraction of exotic cosmic rays near the GZK energy could easily produce a spectrum with no apparent GZK cutoff if their energies are overestimated on average by a factor of two.  This fraction could be even smaller if the energy systematic error is larger, or if for whatever reason some of the observed high energy events really have do have super-GZK energies.  In HiRes2, HiRes stereo and Auger hybrid, where the geometry is much better reconstructed, there should be no systematic overestimate of the energy due to geometry errors but the geometric aperture based on the average shower $X_{max}$ can easily be overestimated.  Furthermore, the event reconstruction and selection procedures may still cause the exotic events to be rejected.  We have discussed in particular the bias due to fixing the first interaction point $X_0$ in the Gaisser-Hillas fit.  There may be an additional bias from fixing the shower attenuation parameter $\\lambda\\sim$ 70 g/cm$^2$, or even assuming that the exotic profile shape can be well-described by a Gaisser-Hillas function.  In this paper, we have implicitly assumed that after the first few interactions, all showers develop in the same way and thus yield the approximately same longitudinal profile shape.  However, simulations of showers initiated by exotic heavy hadrons \\cite{Albuquerque:1998va} \\cite{Berezinsky:2001fy} indicate a tendency to broaden the longitudinal profile due to a lower energy transfer per interaction.  In future studies, the profile shape requirements may be relaxed to perform a more general search.  Specific searches for exotic events have thus far been confined to searches for high energy gammas as predicted by top-down production models.  These searches are tuned using guidance from photon-shower simulations in order to maximize their acceptance for photons, and to reduce systematic bias.  Furthermore, the null results are interpreted within the context of the shower simulations, which may have large systematic uncertainties.  The AGASA super-GZK signal may be our first indication of new physics in high energy cosmic rays, and the AGASA-HiRes discrepancy may indeed give us some hints on how to identify the new underlying physical processes.  We propose to perform model-independent searches for deep showers which may be a signal for new kinds of exotic particles.  In the case of AGASA, an analysis of the shower front curvature may reveal some outlying events which are positioned close to the ground.  In the case of HiRes1, it is probably difficult to infer any additional information from the data, given that the geometries cannot be independently measured from the pixel timing.  In the cases of HiRes2, HiRes Stereo, and Auger, where the shower geometry is measured better, it appears to be a simple matter to free the parameter $X_0$ in the Gaisser-Hillas fits in order to avoid rejecting showers whose profiles are deeper than expected.  The event selection criteria must also be carefully studied to remove any bias which may be present.  The preliminary study of the acceptance of fluorescence and surface detectors reported here can be refined with more detailed simulations.  The main difficulty is to be able to differentiate between exotic deep showers and ordinary deep showers in the exponential tail of the $X_{max}$ distribution resulting from ``punch-through'' of ordinary cosmic ray primaries.  Given enough statistics, a careful analysis of the tail events may reveal deviations from a pure exponential distribution.  The Auger observatory is producing a high rate of events with well-measured hybrid geometries and longitudinal profiles.  We are optimistic that the deep shower hypothesis can be tested fairly quickly, perhaps even with existing datasets.  \\medskip  {\\it Acknowledgments:} We would like to thank Nicolas Busca, Maximo Ave, Tokonatsu Yamamoto, and Dan Hooper, as well as the Auger groups at FNAL and the University of Chicago for interesting discussions. We would also like to thank Sergio Sciutto for generating the AIRES shower library at Fermilab.   This work has been supported by the US Department of Energy under contract No. DE-AC02-76CH03000."
    },
    "0606/astro-ph0606574_arXiv.txt": {
        "abstract": "Following the observational evidence for cosmic acceleration which may exclude a possibility for the universe to recollapse to a second singularity, we review alternative scenarios of its future evolution. Although the de Sitter asymptotic state is still an option, some other asymptotic states which allow new types of singularities such as Big-Rip (due to a phantom matter) and sudden future singularities are also admissible and are reviewed in detail. The reality of these singularities which comes from the relation to observational characteristics of the universe expansion are also revealed and widely discussed. ",
        "introduction": "It is generally agreed that we have now enough evidence for the past Hot-Big-Bang universe. Its main observational support relies on the following facts:  \\begin{itemize}  \\item The universe expands, i.e., all the galaxies move away from each other according to the Hubble law \\cite{hubble} and they experience cosmological redshift $z$ according to the relation: \\bea 1+z &=& \\frac{a(t_0)}{a(t_e)}~, \\eea where $a(t_0)$ is the scale factor at the time of observation, while $a(t_e)$ is the scale factor at the time of emission of light by a galaxy.  \\item Element abundance in the universe is: hydrogen $75 \\%$, helium $24 \\%$, and other elements $1\\%$. In particular, the amount of helium is larger than it is possible to be produced in stars, and the only solution to this problem is to assume that its abundance is primordial \\cite{gamow}.  \\item Cosmic Microwave Background (CMB) - photons once were in thermal equilibrium with charges which further decoupled and formed thermal background with blackbody radiation spectrum with temperature $T = 2.7 K$ \\cite{penzias}. The information about the density fluctuations $\\delta \\rho$ at the decoupling epoch was imprinted in the temperature fluctuations according to the formula \\bea \\label{delT} \\frac{\\delta T}{T} &\\propto& \\frac{\\delta \\rho}{\\rho} \\propto 10^{-5}~. \\eea \\end{itemize}  However, as well as about the past, it is interesting to ask about the future of the universe. The questions which naturally arise are as follows: \\begin{itemize} \\item What type of future evolution will we generally face? \\item Will the universe expand forever? Will it expand fast, faster, slower \\ldots? \\item Will we face any dramatic change of our future evolution? \\item Is it likely that we face an unexpected end of our future evolution? \\item Is there a barotropic equation of state $p(t) = w\\varrho(t)$, $w=$ const. ($p$- the pressure, $\\varrho$- the energy density) valid throughout the whole evolution, or, perhaps $w=w(t)$ so that the pressure can be expanded in series as \\bea \\label{series} p &=& p_0 + \\frac{dp}{d\\varrho}\\mid_0 (\\varrho - \\varrho_0) + \\frac{1}{2!} \\frac{d^2p}{d\\varrho^2}\\mid_0 (\\varrho - \\varrho_0)^2 + O\\left[(\\varrho - \\varrho_0 )^3\\right] \\nonumber \\\\ &=& p_0 + \\frac{\\dot{p}_0}{\\dot{\\varrho}_0} (\\varrho - \\varrho_0) + \\frac{1}{2!} \\frac{\\ddot{p}_0\\dot{\\varrho}_0 - \\dot{p}_0\\ddot{\\varrho}_0}{\\dot{\\varrho}^3} (\\varrho - \\varrho_0)^2 + O\\left[(\\varrho - \\varrho_0 )^3\\right]~, \\eea where index \"0\" refers to a quantity taken at the current moment of the evolution $t=t_0$. \\item To what extend we are able to determine $w_0 = p_0/\\varrho_0$, $w_0^{,} = (dp/d\\varrho) \\mid_0$, \\ldots? \\end{itemize}  It is generally believed that these questions may, in general, be addressed within a more fundamental framework than Einstein's general relativity theory, i.e., within the framework of fundamental theories of all physical interactions such as superstring, brane and M-theory \\cite{polchin,polch}. ",
        "conclusions": "In view of the discussion performed in this paper one has the following remarks.  Firstly, the future state of the universe may be {\\it more sudden and violent}. It means that the universe may terminate either in a sudden future singularity or in a Big-Rip singularity and this is totally different from our earlier expectations that it could terminate in an asymptotically de Sitter state or in a Big-Crunch.  Secondly, these new future singularities (Big-Rip and SFS) should not be confused - they have totally different properties with respect to geodesic completness. In particular, one can extend the evolution of the universe through a sudden future singularity, while it is not possible to do so for a Big-Rip.  Thirdly, statefinders (Hubble, deceleration, jerk, kerk etc.) may be useful to diagnose the future state of the universe. By this we mean the emergence of sudden future singularities, the emergence of generalized sudden future singularities, the evolution of the cosmic equation of state $w = w(t)$ etc.  Finally, the new energy conditions may be introduced for the sake of the proper signal for generalized sudden future singularities, or (on the contrary), for the sake of the avoidance of sudden future singularities, or generalized sudden future singularities.    \\begin{acknowledgement} This work was partially supported by the Polish Ministry of Education and Science grant No 1 P03B 043 29 (years 2005-2007). \\end{acknowledgement}"
    },
    "0606/astro-ph0606097_arXiv.txt": {
        "abstract": "A2163 was observed by the RXTE satellite for $\\sim$530 ks during a 6 month period starting in August 2004. The cluster primary emission is from very hot intracluster gas with $kT \\sim 15$ keV, but this component does not by itself provide the best fitting model. A secondary emission component is quite clearly needed, and while this could also be thermal at a temperature significantly lower than $kT \\sim 15$ keV, the best fit (to the combined PCA and HEXTE datasets) is obtained with a power law secondary spectral component. The deduced parameters of the non-thermal (NT) emission imply a significant fractional flux amounting to $\\sim 25\\%$ of the integrated 3-50 keV emission. NT emission is expected given the intense level of radio emission, most prominently from a large extended (`halo') central region of the cluster. Interpreting the deduced NT emission as Compton scattering of the radio-emitting relativistic electrons by the CMB, we estimate the volume-averaged value of the magnetic field in the extended radio region to be $B \\sim 0.4 \\pm 0.2\\, \\mu$G. ",
        "introduction": "High quality spatially resolved measurements with current X-ray satellites have clearly shown that intracluster (IC) gas is not isothermal. With increasing image detail it will likely be found that non-isothermality is a general feature of clusters. This is only to be expected, given the gas extent, origin, and the processes that have affected its evolution. The only justification for spectral fitting of the emission from a large cluster region by a single temperature plasma emission code (what has been a standard procedure) was insufficient spectral and spatial resolution.  Not only the assumption of isothermality is generally invalid when characterizing the emission from a large cluster region, the expectation that the emission is {\\it purely} thermal may also be doubtful when a wide spectral band is considered. This is particularly so in clusters with known extended regions of radio emission, where the radio emitting relativistic electrons give rise to non-thermal (NT) X-ray emission by Compton scattering off the Cosmic Microwave Background (CMB) radiation. The well known possibility that cluster X-ray spectra may have high energy power law tails (\\eg, Rephaeli 1977) has been largely ignored, also because of the low sensitivity of past X-ray satellites to the detection of the predicted low-level NT emission.  Non-isothermal gas distribution and NT emission are of considerable interest both intrinsically, for the detailed understanding of the astrophysics of clusters, as well as in the use of cluster properties and phenomena as cosmological probes. X-ray emission is currently our best probe of the properties of IC gas; the thermal structure of the gas can yield important information on energy exchange and transport processes. More globally, the detailed gas density and temperature profiles are required in analysis of measurements of the Sunyaev-Zeldovich (S-Z) effect and its use as a cosmological probe of the global parameters of the universe and its large scale structure. On the other hand, NT phenomena in clusters have the potential to contribute significantly to our understanding of the origin of relativistic particles and magnetic fields. It is obvious, therefore, that there is strong motivation for a more realistic characterization of cluster X-ray spectra.  The search for NT X-ray emission in clusters has been advanced considerably by the RXTE and BeppoSAX satellites. We have initiated and analyzed long RXTE observations of the Coma cluster (Rephaeli, Gruber \\& Blanco 1999, Rephaeli \\& Gruber 2002), A2319 (Gruber \\& Rephaeli 2002), and A2256 (Rephaeli \\& Gruber 2003). In all three clusters we found evidence for a second spectral component. While a secondary emission component would generally reflect the non-isothermality of IC gas, we found that in these three clusters the additional emission is likely to be mostly NT in origin. In each of the clusters this conclusion is based on the relative quality of the statistical fits (of thermal vs. NT models), and the high physical plausibility of NT X-ray emission in the central cluster region in which extended radio emission was measured. Similar BeppoSAX searches for NT in these and other clusters (including A119, A754, and A2199) also yielded evidence for NT emission in at least some of the clusters (Fusco-Femiano \\ea 1999, Kaastra \\ea 1999), Fusco-Femiano \\ea 2000, Fusco-Femiano \\ea 2003). Note that the significance of the BeppoSAX/PDS results has been questioned by Rossetti \\& Molendi (2003), who claim that the level of instrumental error is higher than what was previously assumed, and that the detection of NT emission in Coma is much less significant than originally reported by Fusco-Femiano \\ea (1999). However, this claim has been disputed by Fusco-Femiano \\ea (2004).  The moderately distant cluster A2163, $z=0.203$, is one of the hottest, most luminous clusters, and a prime target of radio, S-Z, and X-ray observations in recent years. Its extended central region of radio emission (Herbig \\& Birkinshaw 1994) was extensively mapped by Feretti \\ea (2001); this radio `halo' is one of the largest and most luminous. Spatial correlation between bright regions in the {\\it Chandra} and radio maps may be interpreted as evidence for intense merging activity (Markevitch \\& Vikhlinin 2001), which is thought to enhance the efficiency of particle acceleration. The gas temperature profile and a detailed temperature map were determined from {\\it XMM} (Pratt \\ea 2001) and {\\it Chandra} (Govoni \\ea 2004) measurements. Here we report the results from a very long observation of A2163, the most distant of the radio (`halo') clusters observed with the RXTE. ",
        "conclusions": "The results of the above spectral analysis clearly indicate that the emission in A2163 cannot be described by a single temperature emission model (as is apparent from positive residuals at both low and high energies). A better fit is provided by a two-temperature model, with the secondary thermal component accounting for the low energy residuals. The {\\it Chandra} temperature map of the central cluster region shows a region, roughly $2'X2'$ in area, where the temperature is $\\sim 9$ keV. The secondary thermal component we deduce has a somewhat lower value, but given the large uncertainty in the value we deduce, the difference is not significant, as is the {\\it rough} - due to the substantially larger ($58'$ FWHM) RXTE FOV, which includes the emission from the cooler outer regions - consistency of the relative flux values of the two thermal components.  However, in the best fitting model the second component is not thermal, but rather a NT power law. Indeed, NT emission is expected from the populations of \\rel electrons that emit both the extended IC radio emission, as well as the emission from the several powerful radio sources in the central $\\sim 2$ Mpc region of A2163. In Table 2 we list the measured radio fluxes and spectral (energy) indices for the dominant sources in the RXTE FOV. In addition to the remarkably luminous emission from the most extended, regularly-shaped `halo' region, we include the emission from four sources measured by Feretti et al. (2001). Fluxes and spectral indices for the `halo' and 'relic' sources are determined from VLA measurements at 6 \\& 20 cm. To estimate the expected X-ray fluxes also from the three tailed (T) sources, we need to specify their spectral indices; since these were not given by Feretti \\ea, we adopt a typical value of $0.8\\pm 0.1$.  \\begin{table*} \\caption{Values of radio parameters (and 1$\\sigma$ errors)} \\vspace{0.5cm} \\small \\vspace{0.5cm} \\begin{tabular}{ l@{\\hspace{0.3cm}} c@{\\hspace{0.4cm}} c@{\\hspace{0.4cm}} c@{\\hspace{0.4cm}} } \\hline Source & Flux Density - 20 cm (mJy) & Spectral (energy) index & Size (kpc)\\\\ \\cline{1-4} Halo & 155\\p2 & 1.6\\p0.3 & 2070\\p70 \\\\ Relic & 18.7\\p 0.3  & 2.1\\p0.3 &460\\p40 \\\\ J1615-062 (T1) & 34.5\\p0.5 & 0.8\\p 0.1 & 450 \\\\ J1615-061 (T2) & 6.0\\p0.3 & 0.8\\p 0.1 & 120 \\\\ J1616-062 (T3) & 24.9\\p0.4 & 0.8\\p 0.1 &270 \\\\ \\cline{1-4} \\end{tabular} \\tablenotetext{}{Source size is based on a Hubble constant of 70 $km\\,s^{-1}\\,Mpc^{-1}$}  \\end{table*}  On the relevant spatial scales the radiation field most tightly coupled to the electrons by Compton scattering is the CMB, coupling that is only further enhanced in A2163 ($z=0.203$) due to the $(1+z)^4$ dependence of the CMB energy density. In order to quantify the emission that effectively results from the scattering, we need to determine the spectral density distribution of the electrons. Radio emission from the central extended source dominates the overall emission, and since we expect the mean, volume-averaged field in this extended source to be much lower than field strengths in the galactic radio sources, it is clear that most of the Compton-produced X-ray emission comes from the 'halo'. Nonetheless, we have estimated the emission by \\rel electrons in the other sources listed in Table 2.  Compton fluxes of the radio sources in Table 2 can be estimated from their measured radio fluxes and estimated mean field values. The latter can be determined by assuming energy equipartition between particles and fields. This assumption may be roughly valid in the particles galactic sources, by virtue of the relatively short timescales of all the relevant processes governing the particles and fields - acceleration, couplings to the field and interstellar gas. (We note in passing that the attainment of equipartition in the IC space is questionable: IC fields and particles are likely to be of galactic origin, but the evolution during their expulsion from galaxies and in IC space is quite different; effective equilibration is not likely to occur there.) Using the deduced values of the equipartition in the radio sources we have estimated their respective Compton X-ray fluxes. These turn out to be quite small for the relic as well as the three tailed sources, whose combined flux in the 3-80 keV band is $\\leq 1\\%$ of the measured value. (Note that consideration of the dominant non-'halo' radio sources in the clusters in which we have previously found evidence for NT emission - Coma, A2256, and A2319 - similarly shows that their predicted X-ray emission contributes negligibly to the deduced NT emission in each of these clusters.)  Since the known radio sources in the central region of A2163 do not contribute appreciable X-ray emission, we attribute all the measured emission to electrons in the 'halo', and proceed to determine the mean, volume-averaged field in this region, $B_{rx}$, using the Compton-synchrotron formulae (\\eg Rephaeli 1979). For consistency with the basic premise that most of the measured Compton flux is due to electrons in the `halo', and the fact that the electron spectrum can be more precisely deduced from the radio data, we use the radio spectral index in order to infer the electron index. Doing so we compute $B_{rx}\\simeq 0.4\\pm 0.2 \\,\\mu$G. It should be emphasized that the overall uncertainty in the estimated mean field is substantially higher than this formal error due to the lack of spatial information on NT X-ray emission, which necessitates invoking the assumption that the emitting \\rel electrons and the field are co-spatial.  As a consequence of the assumption that the spatial factors in the theoretical expressions for the radio and NT X-ray fluxes are roughly equal, it follows that the mean value of the deduced magnetic field is independent of the source size and distance. The \\rel electron energy density does depend on these parameters. Scaling to the observed radius of the diffuse radio emission, and integrating the electron spectral distribution over all energies above 1 GeV, we obtain $\\rho_e \\simeq 1.4_{-0.7}^{+0.8} \\times 10^{-13} h_{70}^{-1} \\,erg/cm^3$, where $h_{70}$ is the value of the Hubble constant in units of $70$ km\\,s$^{-1}$\\,Mpc$^{-1}$). Based on the high Galactic proton to electron energy density ratio of cosmic rays, it can be conjectured that the energetic proton energy density much higher than this value. A more specific assessment cannot be made without consideration of the electron and proton origin and their respective energy losses.  IC magnetic fields can also be estimated from Faraday rotation measurements of background radio sources seen through clusters, yielding a different mean field value, $B_{fr}$. These measurements usually yield field values that are a few $\\mu$G (see, \\eg, Clarke, Kronberg, and B\\\"ohringer 2001, and the review by Carilli \\& Taylor 2002), up to an order of magnitude higher than values of $B_{rx}$. Average field strength of a few $\\mu$G over an extended cluster region would have important implications on the range of electron energies that are deduced from radio measurements, and therefore on the electron (synchrotron and Compton) energy loss times. Weaker mean (volume-averaged) fields imply higher electron energies, and therefore shorter electron energy loss times, with possibly important ramifications for \\rel electron models (\\eg, Rephaeli 1979, Sarazin 1999, Ensslin \\ea 1999, Brunetti \\ea 2001, Petrosian 2001). However, the apparent discrepancy between deduced values of $B_{rx}$ and $B_{fr}$ is usually naively interpreted. $B_{rx}$ and $B_{fr}$ actually provide very different measures of the mean field strength: $B_{rx}$ is essentially a weighted volume-average of the \\rel electron density and of $B^q$, with $q \\geq 2$; on the other hand, $B_{fr}$ is an average of the product of the line of sight component of the field and the gas density. In general, the overall spatial dependence of these averages are considerably different, and so are their deduced values. Specific examples and quantitative comparisons, including also the impact of observational uncertainties, are given by Goldshmidt \\& Rephaeli (1993), and by Newman, Newman \\& Rephaeli (2002)."
    },
    "0606/astro-ph0606212_arXiv.txt": {
        "abstract": "We analyse the environment of low redshift, $z < 0.2$, SDSS quasars using the spectral and photometric information of galaxies from the Sloan Digital Sky Survey Third Data Release (SDSS-DR3). We compare quasar neighbourhoods with field and high density environments through an analysis on samples of typical galaxies and groups.  We compute the surrounding surface number density of galaxies finding that quasar environments systematically avoid high density regions. Their mean environments correspond to galaxy density enhancements similar to those of typical galaxies.  We have also explored several galaxy properties in these environments, such as spectral types, specific star formation rates, concentration indexes, colours and active nuclei activity. We find a higher relative fraction of blue galaxies in quasar environments compared to groups and typical galaxy neighbourhoods. Consistent with this picture, the distribution of the concentration index of these galaxies also indicate a larger fraction of late-type objects. By analysing the available information of galaxy spectra we have also studied the distribution of the star formation rates of these neighbour galaxies finding that quasar environments are populated by objects with an enhanced star formation activity. An analysis of the relative flux ratios of $[OIII]\\lambda5700/H\\beta$ and $[NII]\\lambda6583/H\\alpha$ of emission-line galaxies shows no an excess of nuclei activity in quasar neighborhood with respect to the environment of a typical galaxy.  We conclude that low redshift quasar neighbourhoods ($r_p < 1\\mpc$, $\\Delta V < 500\\kms$) are populated by bluer and more intense star forming galaxies of disk-type morphology than galaxies in groups and in the field. Although star formation activity is thought to be significantly triggered by interactions, we find that quasar fueling may not require the presence of a close companion galaxy ($r_p < 100\\kpc$, $\\Delta V< 350\\kms$).  As a test of the unified AGN model, we have performed a similar analysis to the neighbours of a sample of active galaxies. The results indicate that these neighbourhoods are comparable to those of quasars giving further support to this unified scenario. ",
        "introduction": "Important clues on galaxy formation and evolution may be obtained by characterizing statistically the properties of their neighbourhoods in the local Universe. It is well known that several properties of galaxies depend on the environment where they formed and evolved where a variety of processes such as star formation, tidal stripping, merging, etc. can determine the nature of galaxies. It should also be considered the possibility of AGN feedback processes which could induce significant changes in the evolution of their companion galaxies (see for instance Croton et al. 2005).  Previous works aimed to characterize the quasar neighbourhood indicate that high redshift quasars can be used as signspots to search for rich density regions (Djorgovski 1999, Hall \\& Green 1998, Fukugita 2003). However, at lower redshifts (z $\\le $ 0.3) different studies indicate that quasars could reside in environments similar to those of normal galaxies (Smith, Boyle \\& Maddox 1995, Sorrentino \\etal 2006) or they could be located in groups (Fisher et al. 1996) or clusters of galaxies (Mclure \\& Dunlop 2001). Coldwell et al. 2002, analysing both the projected cross-correlation function and colours of galaxies around a sample of quasars and AGNs found that their typical galaxy density environment corresponds to groups of galaxies. Moreover, Coldwell \\& Lambas (2003, hereafter PaperI), using the spectral information of the 2dF Galaxy Redshift Survey (2dFGRS), found that the galaxies within $r_p < 1 \\mpc$ from quasars have a stronger star formation activity. These results are also supported by findings of Soechting \\etal (2002, 2004) who found that low redshift quasars follow the large-scale structure traced by galaxy clusters but they are not placed in the central area of galaxy clusters. More likely, they are in the cluster periphery or between two, possibly merging, galaxy clusters.  The large recent spectroscopic surveys, the Sloan Digital Sky Survey and the 2dF Galaxy Redshift Survey allow us investigate a wide range of galaxy properties such as morphology, spectral types, colours and, also, the galaxy density using volume-limited samples. In this paper we extend the work of PaperI by analysing the environment of Sloan quasars using the spectroscopic and photometric information available for the Sloan Digital Sky Survey Third Data Release (SDSS). We divided the sample in two range of redshift $0.03 < z < 0.1$ and $0.1 < z < 0.2$ in order to anaelyse different effects in these two redshift ranges populated by galaxies of different luminosities.  The layout of the paper is as follows. In section 2 we describe the data, section 3 shows the analysis of the local density estimates, section 4 and 5 describes the statistical analysis performed with the photometric and spectroscopic data respectively, and in section 6 we provide a brief discussion of the main results. ",
        "conclusions": "We have performed a statistical analysis of local quasar environments using SDSS data, our main conclusions are:  $\\bullet$ Nearby quasars systematically avoid high density regions. Local surface density estimates show that SDSS quasars reside in similar density enhancements to typical galaxies, significantly less dense than cluster environments.  $\\bullet$ Star formation activity in the surrounding of quasars is higher than in the neighbourhood of typical galaxies. This important property may provide a link between star formation and the onset of quasar activity.  $\\bullet$ Quasar environment is populated by galaxies systematically bluer, and with a disk-type morphology.  $\\bullet$ In spite of the fact that galaxy interactions efficiently trigger star formation, we find no statistical evidence that the presence of close companions is associated to quasars.  $\\bullet$The presence of quasars does not affect the fraction of AGNs in neighbour galaxies.   $\\bullet$ The results mentioned are present in the two redshift ranges explored $z<0.1$ and $0.1<z<0.2$ although they are stronger at lower redshifts, indicating that the brightest galaxies are not likely to be the major contributors to the effects. Also, there might be a signal for evolution in the sense that the effects are stronger locally.  $\\bullet$ AGN and quasar environment are similar regarding the galaxy characteristics explored and agree best in the highest redshift interval giving support to the unified scenario of AGNs and quasars.  Although quasar hosts tend to be red and bulge-dominated galaxies (Kauffmann \\etal 2003), the indication of a relative excess of quasar neighbour galaxies with a large gas fraction could suggest either an external origin for the accretion onto the central black hole, or that quasars may affect substantially the evolution of star formation up to the Mpc scale."
    },
    "0606/astro-ph0606738_arXiv.txt": {
        "abstract": "We report on non-LTE Ne abundances for a sample of B-type stellar members of the Orion Association.  The abundances were derived by means of non-LTE fully metal-blanketed model atmospheres and extensive model atoms with updated atomic data.  We find that these young stars have a very homogeneous abundance of A(Ne) = 8.27 $\\pm$ 0.05.  This abundance is higher by $\\sim$ 0.4 dex than currently adopted solar value, A(Ne)=7.84, which is derived from lines produced in the corona and active regions.  The general agreement between the abundances of C, N, and O derived for B stars with the solar abundances of these elements derived from 3-D hydrodynamical models atmospheres strongly suggests that the abundance patterns of the light elements in the Sun and B stars are broadly similar.  If this hypothesis is true, then the Ne abundance derived here is the same within the uncertainties as the value required to reconcile solar models with helioseismological observations. ",
        "introduction": "One important recent result from studies of stellar atmospheres and chemical compositions in stars is the downward revision in the abundances of carbon, nitrogen and oxygen in the solar photosphere, which was obtained from the adoption of time-dependant, hydrodynamical and 3-Dimensional model atmosphere calculations (Asplund 2005). These more realistic model atmosphere calculations indicate that the solar abundances should be lower by roughly 0.2-0.3 dex when compared to abundances derived from hydrostatic 1-D calculations. It soon became apparent that the significantly lower abundances in the Sun resulted in severe inconsistencies between the solar models and measurements from helioseismology.  Different possibilities were investigated in order to try to reconcile the solar models with the seismological observations, such as updating opacities (Bahcall et al. 2004), changing diffusion rates (Guzik et al. 2005), or significantly changing the adopted solar abundances of key elements, such as neon. Allowing for a larger neon abundance, in particular, was justified because the neon abundance in the Sun can be considered to be more uncertain given that it is not measured from lines formed in the solar photosphere. Antia \\& Basu (2005) constructed envelope models of the Sun, allowed for different abundance mixtures and focused on the density profile, which is determined from helioseismology. More complete calculations were presented in Bahcall, Basu \\& Serenelli (2005) who constructed solar models, consisting of the atmospheres plus the interior, and concluded that an adopted neon abundance A(Ne)= 8.29 $\\pm$ 0.05, would suffice in order to bring the solar models and seismological observations into an acceptable agreement. Independently, measurements of neon abundances in a sample of chromospherically active cool stars by Drake \\& Testa (2005) indirectly supported high Ne abundances in solar type stars.  It has been a long standing puzzle that the C, N and O abundances obtained for B-stars, which are young, were typically lower than the at-the-time generally accepted solar abundances from Anders \\& Grevesse (1989). These were puzzling results because from our simplest understanding of how the Galaxy chemically evolves, it is not expected that young stars in the solar vicinity would be less enriched than the Sun, which is much older. Moreover, the abundances obtained from Galactic H II regions were also lower than the accepted solar and in rough agreement with the B-star results. These inconsistencies between the abundances of young stars and H II regions, on one hand, and the Sun on the other, are reconciled nowadays with the revised solar abundances from 3-D models. In this context, it is therefore important and timely to derive accurate neon abundances in the atmospheres of early-type stars. In this study, we report on non-LTE (NLTE) neon abundance calculations for a sample of B stars members of the Orion association. The Ne abundances in young stars can independently shed light on the issue related to the reference Ne abundance in the Galaxy. ",
        "conclusions": "Measurements of neon abundances in a variety of objects that can help define the uncertain neon abundance in the Sun are potentially of great importance for solar physics. We find the Ne abundance in B stars members of the Orion association is significantly higher than the solar value by roughly 0.4 dex ($\\sim$ 2.5X). We argue that the Ne abundances measured in young OB stars should be a good representation of the solar chemical composition, as is indicated from the good agreement between the abundances in B stars and Sun for other elements such as C, N, and O. The high Ne abundances obtained here come to the rescue of the solar models that require, according to Bahcall, Basu \\& Serenelli (2005), an increase in the Ne abundance by $\\sim$2.8 $\\pm$ 0.4."
    },
    "0606/astro-ph0606354_arXiv.txt": {
        "abstract": "We present a population synthesis calculation to derive the total number of planetary nebulae (PN) in the Galaxy that descend from single stars and stars in binary systems. Using the most recent literature results on galactic and stellar formation as well as stellar evolution, we predict the total number of galactic PNe with radii $<$0.9~pc to be ($4.6\\pm1.3) \\times 10^4$.  We do not claim this to be the complete population, since there can be visible PNe with radii larger than this limit. However, by taking this limit, we make our predicted population inherently comparable to the observationally-based value of Peimbert, who determined ($7200\\pm1800$)~PNe should reside in the Galaxy today. Our prediction is discrepant with the observations at the 2.9~$\\sigma$ level, a disagreement which we argue is meaningful in view of our specific treatment of the uncertainty. We conclude that it is likely that only a subset of the stars thought to be capable of making a visible PN, actually do. In the second paper in this series, an argument will be presented that the bulk of the galactic PN population might be better explained if only binaries produce PNe.  The predicted PN formation rate density from single stars and binaries is ($1.1 \\pm 0.5$)$\\times10^{-12}$~PN~yr$^{-1}$~pc$^{-3}$ in the local neighborhood. This number is lower than the most recent PN birthrate density estimates (2.1 $\\times10^{-12}$~PN~yr$^{-1}$~pc$^{-3}$), which are based on local PN counts and the PN distance scale, but more in line with the white dwarf birthrate densities determined by Liebert et al. (($1.0 \\pm 0.25$)$\\times10^{-12}$~WD~yr$^{-1}$~pc$^{-3}$). The predicted PN birthrate density will be revised down, if we assume that only binaries make PNe. This revision will imply that the PN distance scale has to be revised to larger values. ",
        "introduction": "\\label{sec:introduction}   The origin of planetary nebulae (PNe) has been the subject of a hot debate for the last thirty years (for a list of references see \\citealt{Soker1997}). The bone of contention has been the mechanism by which the large majority of PNe achieve axi- and point-symmetric shapes. The proposed mechanisms can be divided into two classes. The first class of mechanisms involves single stars, proposing a combination of magnetic fields and/or stellar rotation to achieve an axi-symmetric mass-loss during the Asymptotic Giant Branch (AGB) phase \\citep{GarciaSegura2004, Matt2004,Blackman2004} into which the fast central star wind expands. The second class of mechanisms proposes that PNe with non spherical morphologies must have suffered an encounter with a binary companion (whether a close encounter such as a common envelope or a lesser interaction such as gravitational focussing). Depending on the binary separation, mass ratio and other parameters the PN ejecta take on different shapes.  A simple test of which of the two classes of models is more likely to be correct is to determine the PN central star binary fraction. About 10--15\\% of all PN are known to harbour close binary central stars \\citep{Bond2000}. These were discovered with photometric surveys that rely on heating of one hemisphere of the companion by the primary, or by ellipsoidal variations. Since heating effects and ellipsoidal variations quickly diminish with binary separation, only binaries with periods smaller than a few days can be detected. Surveys aimed at detecting visual binaries have been carried out \\citep{Ciardullo1999} and, once again, about 10\\% of all observed central stars appear to have a distant companion (separations $>>$ 100~AU). The companions of the latter group are unlikely to have played any significant role in shaping the PN, although it is not excluded that these wide binaries might harbor a third companion orbiting closer to the primary.  It is clear that until surveys are carried out that can systematically detect binaries with periods between a few days and a few years, we cannot know the actual PN binary fraction. Radial velocity surveys are ideal to fill the period gap, but their execution is particularly tricky, since relatively long observing runs are needed on moderate aperture telescopes with intermediate-to-high resolution spectrographs. \\citet{Mendez1989} reported on the results of a radial velocity survey of 28 central stars carried out at intermediate resolution. However, he only carried out one or two measurements per star, severely limiting the discovery potential of his survey (he had no conclusive evidence to prove the binarity of any stars in his sample) and precluding the determination of periods. \\citet{Sorensen2004} carried out a new survey of 33 central stars, at lower resolution, but where many measurements were taken for every star. They report about 40\\% radial velocity variabilty in their sample of 33 central stars, although no periods could be determined.  More recently, \\citet{DeMarco2004} and \\citet{Afsar2005} carried out radial velocity surveys with resolution comparable to the Mendez (1989) and Sorensen \\& Pollacco (2004) studies, respectively. Their samples (11 and 19 objects, respectively) provided tantalizing evidence that a much larger fraction, possibly as high as 90\\%, of central stars might have companions. Periods were not detected by these surveys, leaving the possibility that in some cases the radial velocity variability might be due to stellar wind variability. From additional observations (De Marco et al., in preparation, \\citealt{Bond2005})  at echelle resolutions  periods between 4 and 5 days were determined for both stars observed, which where drawn from the sample of De Marco et al. (2004; IC4593 and BD+33~2642). This fuels the empirically-based suggestion that  the PN central star close binary fraction might indeed be very high.  Radial velocity surveys of central stars of PN are, however,  far from being complete so that an observational answer to the question of the exact PN central star binary fraction is not yet in hand. We have therefore carried out a population synthesis calculation to revisit the issue of how many PNe are predicted in the Galaxy today.  The {\\it empirical} suggestion that most-to-all PNe are in short-period binary systems, should be supported, or at least consistent, with current theories of stellar evolution, galactic history, star formation and knowledge of related stellar populations (such as AGB stars and white dwarfs [WDs]). This suggests that  population synthesis calculations which assume that single stars {\\it do} make PNe should predict an overabundance of PNe compared to observations. In this paper (the first in a series), we address single star evolution and the total PN population. We analyze all the sources of error and state the issues involved with PN and WD counts and birth rates. In the second paper of this series (De Marco et al., in preparation, hereafter Paper~II), we will address the issue of binarity as the main evolutionary channel for the production of PNe.  In \\S~\\ref{sec:BotE} we outline the phases of our population synthesis calculation as if it were a back of the envelope calculation. In \\S~\\ref{sec:popsyn} we give a detailed explanation of our model method, as well as of our assumptions concerning the initial mass function (IMF; \\S~\\ref{ssec:IMF}), stellar lifetimes (\\S~\\ref{ssec:lifetimes}), the mass of the Galaxy (\\S~\\ref{ssec:massMW}), the star formation history (SFH; \\S~\\ref{ssec:SFH}), the age-metallicity relation (AMR; \\S~\\ref{ssec:metallicities}) and the PN visibility timescales (\\S~\\ref{ssec:PNLifetimes}). In \\S~\\ref{sec:results}, we present our results, we explain the uncertainties, and carry out a detailed comparison with observations and past results from the literature, setting the stage for the binary synthesis (Paper~II). In \\S~\\ref{sec:nrPNbulgeGCs} we compare our PN population predictions with the observed populations of the galactic bulge and galactic globular cluster (GC) system. We conclude in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have calculated the number of PNe expected to reside in the Galaxy, if PNe derive from single stars as well as from stars in binary systems. The resulting value, $(4.6 \\pm 1.3) \\times 10^4$ objects with radius smaller than 0.9~pc, is six times higher than the observationally-based estimate of 7200$\\pm$1800 objects with the same radius constraints (discrepant at the 2.9$\\sigma$ level). Although it could be argued that the actual galactic PN population is larger, we stress that our predicted value is not absolute (since it is constrained by a maximum PNe radius), but is rather designed to be compared to the determination of \\citet{Peimbert1993} and as such the comparison is much more reliable than the individual figures.  We argue that the discrepancy is due to the fact that not all stars otherwise fit to make a PN, actually do make a PN. In Paper~II we will calculate the galactic PN population deriving from binary interactions and show that, based on population synthesis arguments alone, it is more likely that most PNe could derive from such interaction rather than being produced by single stars.  We also compared our predicted PN birthrate density with that derived observationally from counts of local PNe. Our prediction (($1.1\\pm0.5)\\times$10$^{-12}$~PN~yr$^{-1}$~pc$^{-3}$) is lower than most observationally-based estimates (2.1$\\times$10$^{-12}$~PN~yr$^{-1}$~pc$^{-3}$; \\citealt{Phillips2002})\\footnote{It will strike the reader that from observational PN birthrate density values that are generally {\\it larger} than our prediction, total galactic PN populations are derived that are {\\it lower} than our estimate (see for instance \\citet{Phillips2002}). This is due to the very different methods to derive the birthrate densities in our theoretical derivation and in the observationally-based derivations found in the literature.}, although it is within the quoted uncertainty. Our estimate is similar to the more reliable observational estimate of the WD birthrate density (1.0$\\times$10$^{-12}$~WD~yr$^{-1}$~pc$^{-3}$; \\citet{Liebert2005}). The PN birthrate density from binary interaction we will derive in Paper~II will clearly be lower. A discussion of the implication of this discrepancy is left to that paper.  We predict that only 73\\% of post-AGB WDs went through the PN phase. This is in line with what we know of post-AGB evolutionary timescales (namely, that central stars with mass $<$0.56~\\msun\\ do not make a PN), the observed mass distributions of central stars (minimum central stars mass $\\sim$0.56~\\msun) and the mass distribution of WDs (which indicates that only 77\\% of post-AGB WDs have masses large enough to go through a PN phase.)."
    },
    "0606/astro-ph0606162_arXiv.txt": {
        "abstract": "We have detected FIRST J101614.3+520916 with the \\xmmn\\ X-ray Observatory.  FIRST J101614.3+520916, one of the most extreme radio-loud, broad absorption line (BAL) quasars so far discovered, is also a Fanaroff-Riley type II (FR II) radio source.  We find that, compared to its estimated intrinsic X-ray flux, the observed X-rays are likely suppressed, and that the observed hardness ratio indicates significant soft X-ray photons.  This is inconsistent with the simplest model, a normal quasar spectrum absorbed by a large neutral HI column density, which would primarily absorb the softer photons.  More complex models, involving partial covering, an ionized absorber, ionized mirror reflection, or jet contributions need to be invoked to explain this source. The suppressed but soft X-ray emission in this radio-loud BAL quasar is consistent with the behavior displayed by other BAL quasars, both radio-loud and radio-quiet. ",
        "introduction": "About 10--20\\% of quasars show broad absorption lines (BAL), especially in their ultraviolet (UV) spectra.  These absorption features usually extend to velocities as high as $\\sim 10^4$\\,\\kms\\ relative to the emission lines, indicating high-velocity outflows in the quasars.  These absorbers have been identified with winds blowing from an obscuring torus or arising from smaller scales associated with an accretion disk feeding a super massive black hole. The popular orientation model suggests that BAL quasars are normal quasars viewed along the specific line of sight, or particularly edge-on, skimming the torus or through a wind \\citep[e.g.,][]{Weym91}.  Although in this picture, quasar radio properties and BALs would seem to be independent, no radio-loud BAL quasars were found for a long time.  It remained so until deep radio surveys like the NRAO VLA Sky Survey \\citep[NVSS,][]{Cond98} and FIRST Bright Quasar Survey \\citep[FBQS,][]{Beck95,Greg96,Whit00} were conducted, surveying large areas to mJy levels, and radio-loud BAL quasars started to be identified \\citep{Beck97,Brot98,Beck00,Meno01,Brot02}. Still, BAL quasar frequency does drop significantly among the most radio-loud quasars \\citep{Beck01}. \\citet{Beck00} studied 27 BAL quasars from the FBQS sample and found that they show a wide range of radio spectral indices, from flat to steep, indicating that a range of orientations is present and therefore strongly challenging the orientation model.  So far, only a few radio-loud BAL quasars have been studied at X-ray energies (Brotherton et al. 2005).  \\balobjl\\ (hereafter \\balobj) is the first confirmed BAL quasar that has also been identified as radio-loud Fanaroff-Riley type II (FR~II) source \\citep{Greg00}. Figure~\\ref{fg:spec} shows the BALs in a rest-frame UV spectrum of \\balobj\\ and Table~\\ref{tb:para} provides its optical and radio parameters.  The radio luminosity places it among the extreme end of radio-loud BAL quasars.  Its double-lobed radio morphology and luminosity indicate a classic FR~II radio source.  \\citet{Greg00} argue that \\balobj\\ is a rejuvenated quasar, possibly through a merger or interaction.  We note that another known radio-loud, FR~II BAL quasar is LBQS~1138-0126 \\citep{Brot02}, and that the double radio lobed BAL quasar candidate PKS~1004+13 \\citep{Will99} has recently been confirmed by {\\em HST} observation \\citep{Will06} as a bone fide FR~II BAL quasar.  The other known radio-loud BAL quasars have compact structures \\citep{Beck00}.   Observations with the {\\it Chandra} X-ray Observatory and \\xmmn\\ show that BAL quasars are up to 2 orders of magnitude fainter in X-rays than non-BAL quasars of the same optical brightness \\citep{Gree01, SabHam01, Gall02,Brot05}.  Available X-ray spectral analyses of radio-quiet BAL quasars show that they appear to have normal radio-quiet X-ray photon indices ($\\Gamma \\approx 2$), partially or totally covered by absorbing columns of $N_H\\leq 10^{23}$cm$^{-2}$ \\citep[e.g.,][]{Gall02}. Although the $N_H$ derived from UV absorption lines cannot account for the absorption in the X-ray, the UV and X-ray absorbers are probably closely related \\citep*{Bran00}.  Radio-loud quasars are factors of 2--3 times brighter in the X-rays than radio-quiet quasars with the same optical magnitude \\citep{Brin00}, and tend to have harder X-ray spectra \\citep[$\\Gamma\\sim 1.6$, e.g.,][]{ReeTur00,Page05}.  These factors may make them particularly suitable for initial exploratory studies of the intrinsic X-ray properties and the properties of the line-of-sight absorbers \\citep{Brot05}.  We report the results of a short XMM-Newton observation of \\balobj\\ in this paper.   We detect the object with enough counts to compute a hardness ratio, but not enough for more detailed spectral analysis. Still, the detection can rule out some simple models and has shown us that even the most powerful radio-loud BAL quasars are weak in X-rays. ",
        "conclusions": "We have observed and detected the first confirmed radio-loud FR~II BAL quasar FIRST J1016+5209 in the X-ray with XMM-Newton for the first time. We have enough counts to derive the hardness ratio, but not enough for detailed spectral analysis.  The X-ray flux appears to be suppressed by a factor of 17 relative to the intrinsic X-rays estimated from the radio-X-ray correlation, although significant uncertainties are associated with this factor.  If the X-rays are suppressed due to absorption associated with a high column density of neutral hydrogen, the X-rays observed would be much harder, which is inconsistent with the observations.  This implies that the X-ray absorption in \\balobj\\ is more complicated, such as an ionized absorber, an ionized mirror, or neutral but partially covering the X-ray source.  Contributions from a jet are also possible.  High-quality X-ray spectra are necessary to understand the nature of the absorber."
    },
    "0606/astro-ph0606481_arXiv.txt": {
        "abstract": "Results from a large sample of hydrodynamical/N-body simulations of galaxy clusters in a \\La cosmology are used to simulate cluster X-ray observations. The physical modeling of the gas includes radiative cooling, star formation, energy feedback and metal enrichment that follow from supernova  explosions. Mock cluster samples are constructed grouping simulation data according to a number of constraints which would be satisfied by a data set of X-ray measurements of cluster \\tps as expected from {\\it Chandra} observations. The X-ray spectra from simulated clusters are fitted into different energy bands using the XSPEC {\\it mekal} model. The biasing of spectral \\tps with respect to mass-weighted \\tps is found to be influenced by two independent processes. The first scale dependency is absent in adiabatic runs and is due to cooling, whose efficiency to transform cold gas into stars is higher for cool clusters and this in turn implies a strong dependency of the spectral versus mass-weighted temperature relation on the cluster mass. The second dependency is due to photon emission because of cool gas which is accreted during merging events and biases the spectral fits. These events have been quantified according to the power ratio method and a robust correlation is found to exist between the spectral bias and the amount of cluster substructure.  The shape of the simulated temperature profiles is not universal and it is steeper at the cluster center for cool clusters than for the massive ones. This follows owing to the scale dependency introduced by cooling which implies for cool clusters higher central temperatures, in scaled units, than for massive clusters. The profiles are in good agreement with data in the radial range between $\\sim 0.1 r_{vir}$ and $\\sim 0.4 r_{vir}$; at small radii ($r\\simlt 0.1 r_{vir}$) the cooling runs fail to reproduce the shape of the observed profiles. The fit is improved if one considers a hierarchical merging scenario in which cluster cores can accrete cooler gas through merging with cluster subclumps, though the shape of the temperature profiles is modified in a significant way only in the regime where the mass of the substructure is a large fraction of the cluster mass. ",
        "introduction": "\\label{sec:introduction}   Observations show that most of the baryons present in a galaxy cluster are in the form of a hot ionized gas (intracluster medium , ICM), with temperatures in the range $\\sim 10^6 -10^7 \\gr$ \\cite{sar86}. The continuum emission is dominated by bremsstrahlung processes, and is free from the contamination effects which may arise in the optical band. For this reason, observations of galaxy clusters in the X-ray band have been extensively carried out over the past decade \\cite{hen92,ebe97,ros98} X-ray observations of the spatial distribution of cluster density and temperature  allow one to determine, under the assumption of hydrostatic equilibrium, the mass in baryons as well as the dark matter profile and the total cluster mass $M_X$. The latter quantity is connected to the global cluster temperature $T_X$, and through the $M_X-T_X$ relation the cluster mass function can then be used to find the cluster X-ray temperature or luminosity functions. Observations of these quantities are thus used to constrain theoretical models \\cite{hen91,edg90,hen97,ros98,ebe98}. Knowledge of the temperature profile is also important in order to assess the role of non-gravitational heating processes which can contribute to the budget of the gas thermal energy. This is relevant because observations show that the cluster scaling relations do not obey the self-similar behavior predicted by gravitational collapse \\cite{dav93,ala98,mar98}. Heating of the ICM by non-gravitational processes will break self-similarity and help to explain the observed relations \\cite{evr91,kai91,whi91,low96,va99,lo01}.  A powerful tool for studying the dependence of cluster observational quantities upon theoretical models is given by the use of hydro/N-body simulations. Hydrodynamic simulations have been widely used to investigate the formation and evolution of galaxy clusters \\cite{kat93,sug98,ann96,yos00,per20,lew20,bem01,mu02,lo02,va03,tor03,as03}, see also Voit (2005) and references cited therein. The advantage over analytical methods is that they can treat the gas evolution self-consistently. The validity of the numerical approach is supported by X-ray observations of surface brightness maps, which indicate the existence of merging substructure \\cite{mo95}. Moreover, high resolution observations performed with {\\it Chandra} and XMM-Newton satellites have revealed  the existence of complex cluster temperature structure, like shocks induced by mergers and 'cold front' \\cite{ma00,vi01,ma02}.   Hydrodynamic simulations in which modeling of the gas physical processes incorporates radiative cooling, star formation and energy feedback from supernovae yield a cluster luminosity-temperature relation in good agreement with data for cluster temperatures $T_X \\simgt 1 keV$ \\cite{va03,bo04}. These results support the radiative cooling model \\cite{br00}, where cluster X-ray properties are determined by the cooling efficiency to turn cold gas into stars. However, there are a number of issues for which hydrodynamic simulations that include the physical modeling of the gas described above do not produce a satisfactory agreement with data. For instance, the radial temperature profiles are not isothermal and decline with radius. At small distances from the cluster center the simulated profiles have a steep rise moving inwards. This behavior is not seen in a set of observations \\cite{ma98,al01,de02,vi05,pi05}. The data agree with the simulation results at $r \\simgt 20\\% r_{vir}$ but indicate a temperature profile which is flat or declining at very small distances. These data are also at variance  with the predictions of the standard cooling flow model \\cite{fa94}. In the very central cluster regions the amount of cool gas is smaller than that predicted by the model, with gas temperatures never below $\\simeq 1 keV$ \\cite{ka01,ta01,pe01}. This failure to reproduce temperature data has prompted many authors to consider possible heating mechanisms to explain the lack of cool gas in the cluster core. Among the proposed models, thermal conduction \\cite{voi02,jub04,do04,si06} and energy feedback from  active galactic nuclei \\cite{cu01,br01,bo02,fab02,br02,om04,re05} are the most considered.  The observed temperature profiles are then sensitive to different physical phenomena and can be used to constrain theoretical models. Moreover, as already outlined, the intracluster gas exhibits a complex thermal structure. It is then important to perform properly the comparison of X-ray observations with the results of numerical simulations, taking into account the effects of instrumental response and background contaminations on the spectral fit temperatures. As an additional effect which can bias comparisons, the spectral temperatures are measured in a specified energy bands, whereas the simulated cluster temperatures are theoretically defined according to a specified weighting scheme. This problem has been already analyzed by Mathiesen \\& Evrard (2001) and, more recently, by Gardini et al. (2004), Mazzotta et al. (2004), Rasia et al. 2005 and Vikhlinin  (2006). The aim of this paper is to investigate for systematic effects which can bias spectroscopic measurements of cluster X-ray temperatures. For this purpose, a large numerical sample of simulated clusters is used to construct spatially resolved X-ray spectra as expected from observations with the {\\it Chandra} satellite. The cluster spectral fit temperatures are found using the XSPEC library, and are compared against mean gas temperatures with theoretically defined averages. Moreover, spatially resolved X-ray spectra of the simulated clusters are also used to investigate how the measured temperature profiles differ from the projected profiles obtained directly from simulations. The comparisons are performed on a statistical basis, using the cluster numerical sample. The numerical sample is also subdivided according to the amount of substructure present in a given cluster, this allows us to investigate how spectral measurements are affected by the cluster dynamical state.  This paper can be considered a generalization of that of Mathiesen \\& Evrard (2001, hereafter ME), the main differences being the size of the numerical cluster sample used, together with a more complete physical modeling of the gas in the simulations and the X-ray spectroscopic analysis of radial temperature profiles. The paper is organized as follows. The simulations are presented in Section 2, in Section 3 the procedures with which average and spectroscopic temperatures are extracted from the simulated clusters are discussed. Section 4 is dedicated to the method used to measure substructure in the simulated cluster. The results are discussed in Section 5 and the conclusions are drawn in Section 6. ",
        "conclusions": "\\label {sec:conclusions} In this paper results from a large set of hydrodynamic/N-body simulations of cooling clusters have been used to investigate the relationships between spectral \\tps and mass-weighted or emission-weighted temperatures. X-ray spectral \\tps have been constructed according to a set of procedures in order to reproduce spectroscopic measurements as expected from {\\it Chandra} satellite. Spectral fits have been performed using the single-temperature {\\it mekal} model of XSPEC, with three free parameters: the gas temperature, the metallicity abundance and the normalization. The main difference with respect to previous works which have investigated these issues is in the physical modeling of the gas, which is allowed to cool radiatively and is subject to feedback heating by SNe. This in turn implies a scale dependency of cluster X-ray properties because of the cooling efficiency to convert cold gas into stars, which is higher for cool clusters. For instance, the relationships between the global spectral temperatures $T_S^{\\Delta}$ and the mass-weighted temperatures $T_{mw}^{\\Delta}$  are found to be significantly biased with respect to those obtained from previous runs (ME), in which the modeling of the gas was adiabatic. The scale dependency introduced by cooling has a significant impact also on the value of $T_{mw}^{\\Delta}$ as different cluster overdensities are considered, therefore the $T_S^{\\Delta}-T_{mw}^{\\Delta}$ relation is found to be strongly affected  also by the value of $\\Delta$. Moreover, the scale dependency due to cooling also enters indirectly in the determination of the shape of the  $T_S^{\\Delta}-T_{mw}^{\\Delta}$ relation because of selection effects which limit the sample construction and are dependent on the cluster mass. The biasing of spectral \\tps with respect to emission-weighted temperatures is found to be similar to that with mass-weighted temperatures, but with a wider scatter.  Another effect which governs the degree of biasing between spectral and mass-weighted \\tps is given by merging processes. Spectral temperatures, as obtained by fitting single-temperature models, can be significantly biased toward lower values than mass-weighted \\tps because of accretion of subclumps due to hierarchical clustering which contain cooler gas that modifies the photon spectrum. This source of biasing was discussed in detail by ME, here it has been analyzed using the power ratios as indicators of the cluster substructure. The main advantage of the method is the possibility to quantify the degree of substructure on a chosen scale, identified by the aperture radius $R_{ap}$. This scale dependency has been investigated by evaluating for the considered samples the cluster power ratio $\\Pi_3$  at the aperture radius $R_{ap}\\simeq R_{\\Delta}$, in correspondence to $T_S^{\\Delta}/T_{mw}^{\\Delta}$ for the cluster under consideration. The results confirm the existence of a strong correlation between the ratio $T_S^{\\Delta}/T_{mw}^{\\Delta}$ and the cluster substructure as measured by $\\Pi_3$. This correlation is found with a high significance level for all of the considered overdensities. An important result which follows from this analysis is that the biasing associated with merging events is independent from the one which follows because of the scale dependency due to cooling. Thus, the biasing of spectral \\tps in galaxy clusters is governed by two processes, independent of each other, the first being due to subclump accretion of cool gas and the second because of cooling.  This two parameter model can explain also the shape of the simulated temperature profiles. The temperature profile is however not universal even for relaxed clusters. This follows because, owing to the scale dependency introduced by cooling, the central temperature in relative units is higher  for cool clusters than for hot clusters. This in turn implies for the former steeper profiles at the cluster center  than for massive clusters. The gradients of the rescaled profiles are found to be dependent on $T_{mw}^{\\Delta}$ with a high statistical significance. This scale dependency of the shape of the profiles on the cluster mass is softened or erased entirely if the cluster has undergone a number of major merger events. The profiles of these clusters are much shallower than those of quiescent clusters. This is clearly due to the effects of remixing of the gas due to mergers and these differences have been quantified according to the cluster morphology using the power ratios. It has been found that the profiles are modified in a significant way only in the presence of major mergers, where the mass of the accreting subclump is a large fraction of the cluster mass. These results were already suggested by Allen et al. (2001b, cf. sect. 9.2), but here they are confirmed through numerical simulations.  For quiescent clusters the scaled temperature profiles are in good agreement with the considered data in the radial range $0.1 \\simlt r/r_X\\simlt 0.4$. The profiles are not consistent with those measured from a sample of massive clusters ( Allen et al. 2001a), but the disagreement can be made less severe if the simulated sample is constructed by taking into account the scale dependencies due to cooling and merging events. A more severe discrepancy is instead found at small radii ($r/r_X \\simlt 0.1$) between the measured profiles and those of the simulations. The profiles of the simulations reach their peak values at the scaled radial coordinate $r/r_X \\simeq 0.02$, which is smaller by a factor $\\sim 5$ of the peak location as derived by observations. Moreover, the peak height of the simulated temperature profiles is $\\sim 20 \\%$ higher than that measured. These discrepancies are clearly indicative that the cooling model fails to consistently reproduce  the observed temperature profiles. In a recent paper Motl et al. (2004)  argue that hierarchical merging can provide a viable scenario for the formation and evolution of cool cores in galaxy clusters. The results of Fig. 6 indicate that merging events play a key role in reshaping the cluster profiles and support the proposed model. However, it is worth noticing that a significant reshaping of the temperature peak is achieved only in the regime of strong merging ($\\Pi_3 \\simgt 75\\%$); moreover  this regime has the side effect that the central temperature drop is no longer reproduced.  The measured temperature profiles have a peak location in rescaled radial units which is located at a larger distance from the cluster center than in the simulations. This is indicative that real clusters, at some stage of their evolution, have undergone a significant diminution of their cooling rates in their cores. This is in order to provide pressure support and to prevent the inflow of gas at a higher entropy toward the cluster center. The suppression of the cooling rate is most easily achieved if the gas is heated by some energy source. The most plausible candidates are SNe whose feedback energy from their explosions can heat the gas to the desired level. However the results of the simulations performed here show that this result is not achieved even in the case of maximum theoretical efficiency of transfer of energy to the ICM. {This indicates that the complex nature of the thermal structure of the  cluster gas is not adequately described by the physical processes incorporated  in the cooling runs performed here and suggests that additional physics is clearly needed in order to prevent overcooling in cluster cores. These results agree with previous findings and have prompted many authors to consider possible physical sources of additional heating to the ICM. The most popular models are (cf. Voit 2005 and references cited therein) energy feedback from active galactic nuclei (AGN) and thermal conduction. Because of the existing temperature gradients the latter model could in principle provide the necessary heating in the cluster central regions \\cite{za03,vo04}, but is currently disfavored by more sophisticated numerical simulations which include thermal conduction \\cite{do04,po05}.  The most favored scenario is thus the AGN  model, in which a supermassive black hole at the center of the cluster is fueled by the gas of the cooling flow and can balance the energy losses of the gas through some energy transfer. A popular mechanism involves heating of the gas due to energy dissipation of buoyant bubbles, inflated by the AGN near the cluster center and rising through the ICM \\cite{cu01,br02,fab02,ru02}. The dissipation can occur through a number of processes, such as viscous dissipation or turbulence. It is clear that incorporating these physical processes  into numerical simulations of cluster formation is a necessary step in order to explain the shape of the observed profiles. This is a very demanding computational challenge which has only recently begun to be addressed \\cite{br02,be02,ba03,om04,ru04,br05,si06}. \\ In such a scenario it is difficult to assess how the main findings presented in this paper will be modified. Although a proper analysis can only be performed through numerical simulations, several qualitative arguments can nonetheless still be used to address the effect of the interaction of the AGN with the ICM on the biasing of spectral  temperatures.  In the proposed AGN heating model buoyant bubbles of plasma will rise from the cluster center toward the surrounding medium, where the transported energy will be transferred into the ICM and converted to heat. This will produce some degree of thermalization of the ICM, which in turn implies a reduction in the amount of cold gas and thus in the soft X-ray emission, thereby reducing the degree of biasing of spectral \\tps. In other words, for a given cluster mass scale, any additional source of feedback is expected to reduce the biasing. Another issue is if the scale dependency introduced by cooling will be preserved in the heating scenario. Some degree of modification will be present, since the bubble injection energy, $E_{bub}$, is scale dependent itself, according to the black hole mass. For example, in one of their models Sijacki \\& Springel (2006) propose $E_{bub} \\propto M_{200} ^{4/3}$, and therefore in  relative terms the energy feedback is higher for massive clusters. The effects on the thermal status of the ICM of energy feedback from AGN are however not expected to modify in a significant way the scale dependency due to cooling. This argument is mainly supported by recent measurements of spectral temperature profiles for a sample of cooling flow clusters (Vikhlinin et al. 2005, 2006). The radial behavior of the measured profiles for cool clusters is  significantly different and steeper than for massive clusters, which is in qualitative agreement with what found here. These conclusions are valid if the cluster possesses a regular X-ray morphology, which is indicative of a dynamically relaxed status. In the case of an ongoing merger large amount of cold gas could be supplied at the cluster center, thus replenishing the gas reservoir of the black hole and triggering AGN activity. In this scenario the thermal status of the ICM surrounding the black hole will be due to a complex interplay between the accretion of cold gas and the uplifted bubble energy. This suggests that the scale dependency depicted in Fig. \\ref{fig:dtpw} will be most likely modified in the proximity of a major merger event ($\\Pi_3 \\rightarrow 0$), with a  decrease in  the spectral bias owing to the feedback energy due to accretion and injected into the ICM. Such a complex picture can however be addressed self-consistently only through numerical simulations.   To summarize, it is worth noting, however, that any proposed heating model must be able to explain the lack of cold gas at the cluster center as well as the observed peak location of the temperature profile. Finally, the results of this paper indicate that in order to properly compare X-ray properties of simulated clusters with those observed, the mass and the morphological composition of the sample must be carefully taken into account.}"
    },
    "0606/astro-ph0606448_arXiv.txt": {
        "abstract": "{In this article we want to answer the cosmologically relevant question what, with some good semantic and physical reason, could be called the mass $M_\\mathrm{u}$ of an infinitely extended, homogeneously matter-filled and expanding universe. To answer this question we produce a space-like sum of instantaneous cosmic energy depositions surrounding equally each spacepoint in the homogeneous universe. We calculate the added-up instantaneous cosmic energy per volume around an arbitrary space point in the expanding universe. To carry out this sum we use as basic metrics an analogy to the inner Schwarzschild metric applied to stars, but this time applied to the spacepoint-related universe. It is then shown that this leads to the added-up proper energy within a sphere of  a finite outer critical radius defining the point-related infinity. As a surprise this radius turns out to be reciprocal to the square root of the prevailing average cosmic energy density. The equivalent mass of the universe can then also be calculated and, by the expression which is obtained here, shows a scaling with this critical radius of this universe, a virtue of the universe which was already often called for in earlier works by E.Mach, H.Thirring and F.Hoyle and others. This radius on the other hand can be shown to be nearly equal to the Schwarzschild radius of the so-defined mass $M_\\mathrm{u}$ of the universe. } ",
        "introduction": "The question already often has been raised by cosmologists (see e.g. Tolman, 1918, Thirring, 1918, Jordan, 1947, Mashhoon, 1988, Mashhoon et al., 1984, Rosen \\& Copperstock, 1992, Barbour, 1995) how the so-called mass $M_\\mathrm{u}$ of an expanding universe, filled with homogeneously deposited matter subject to a Hubble flow, should semantically reasonably and physically adequately be conceived and what as such it may represent in form of a quantifiable cosmic number or a cosmic reality. How should this mass $M_\\mathrm{u}$ be related to cosmological ingredients which are closer associated to the observable universe like the mean mass density, the cosmic scale variation with time or the Hubble parameter, or the so-called $\\Omega-$ value, relating cosmic mass density and the critical density $\\rho_\\mathrm{crit}$? Especially in an isotropically expanding, homogeneous universe with a vanishing or negative curvature parameter $k\\leq0$ in which the volume of the universe is infinite, the definition of such a quantity like $M_\\mathrm{u}$ at first glance does not seem to be straightforward, perhaps not making sense at all.  Starting with suggestions made by Dirac (1937) and investigations carried out by Schuecking (1954) and Einstein \\& Straus (1945) ideas were developing and concretizing that the masses of elementary particles as well as of stars might be somehow related to the large scale structure of the universe and its development in cosmic time. Much earlier already Mach (1883) had formulated as a physical requirement that the inertial mass of each particle just by its rigorous concept should somehow be related to the constellation of all other masses in the universe. Though Mach`s principle never up to the present completely entered into theoretical formulations of mass and gravity interacting with eachother, this principle already pointed to the probable fact that masses and their constellation in the universe should be related to eachother. This idea was later taken up by Thirring 1918 who was asking the question what the absolute inertial rest frame would be in the presence of rotating masses. In his view it should lead to identical geoidic deformations of the rotating earth regardless whether the earth rotates with respect to the universe with an angular frequency $\\omega,$ or whether the universe rotates with $-\\omega$ with respect to the earth at rest.  To put these two cases on a physically conceivable basis, he compared two situations: a) the earth rotating with $\\omega$ in the center of a universe at rest represented by a massive spherical shell with radius $R_\\mathrm{u}$ and mass $M_\\mathrm{u}$, and b) the earth at rest in the center of the universe represented identical to case a), however, this time rotating with $-\\omega$ with respect to the earth. The result of this comparison, when extracted from the adequate Newtonian approximation of Einstein`s field equations for rotating masses, Thirring (1918) (see also Mashhoon et al. 1984) found that at a space point close to the axis in the rotation plane the centrifugal or geodetic forces operating at the earth's equator in cases a) and b) are related by  \\begin{equation} \\textbf{K}_\\mathrm{centri}=X\\cdot\\textbf{K}_\\mathrm{geod} , \\label{1}% \\end{equation} where the factor $X$ is calculated to be: \\begin{equation} X=\\frac{4GM_\\mathrm{u}}{3c^{2}R_\\mathrm{u}}=\\frac{2}{3}\\frac{R_\\mathrm{s,u}}{R_\\mathrm{u}} , \\label{2}% \\end{equation} with the associated Schwarzschild radius $R_\\mathrm{s,u} = 2GM_\\mathrm{u} c^{-2}$ of the mass $M_\\mathrm{u}$. This expresses the fact that the equivalence principle of rotational motions can only be satisfied in the present and upcoming world represented by values $M_\\mathrm{u}$ and $R_\\mathrm{u}$, if the value for $X$ would exactly evaluate to be constant and equal to:\\textbf{ }$X=1$.\\textbf{ }To preserve $X=1$ in an expanding universe would, however, require the expression $M_\\mathrm{u}/R_\\mathrm{u}$ to be a cosmological constant\\textbf{. }This would mean that the total mass of the universe should scale with the increasing cosmic diameter $R_\\mathrm{u}$, an unusual cosmological requirement.  In the following we shall not rely on the above indicated scaling of the mass of the universe with cosmic extent, we rather shall ask the more principle question, how mass $M_\\mathrm{u}$ and extent $R_\\mathrm{u}$ of an expanding universe should be properly defined. Hereby we simply rely on the cosmological principle according to which each space point in the universe is equivalent to any other cosmic space point concerning its being surrounded by matter-filled cosmic space, and shall ask by how much instantaneously countable matter, i.e. matter distributed on a 3-dimensional spacelike hypersphere of the universe, this selected point of the universe may in fact be surrounded. This specific question in our perspectives here is not connected with the problem of how long it may take, before the presence of distant masses in the universe can be communicated by electromagnetic or gravitational waves to our local standpoint in the universe, i.e. we do not intend to sum up retarded contributions of energy depositions in space. The question rather starts from the certainty and knowledge that any spacepoint in the universe is surrounded up to large cosmic distances by cosmic matter that is subject to a cosmological expansion dynamics. What - then - may be called the instantaneous energy content of this spacepoint-related universe? ",
        "conclusions": "We have shown that with the use of an analogy to the inner stellar Schwarzschild metric applied to the cosmic matter distribution one can arrive at a reasonable definition of what could be called the mass $M_\\mathrm{u}(t)$ of the universe. Interestingly enough, this mass $M_\\mathrm{u}(t)$ scales with the critical radius $R_\\mathrm{u}$ of the universe introduced by Equ.(\\ref{10}) which was a request already since the works by Mach (1983), Thirring (1918) (see also Mashhoon et al. 1984, Barbour \\& Pfister 1995, Wesson 1999, Hoyle 1990, 1992, Jammer 2000 or Fahr \\& Heyl 2006). Furthermore the finding that the average density $\\rho_\\mathrm{0}$ of the universe turns out to be scaling with the reciprocal of the square of the above defined critical radius $R_\\mathrm{u}$ of the universe just fulfills the request for an economical universe with vanishing total energy as was discussed by Overduin \\& Fahr (2003), Fahr (2004) and Fahr \\& Heyl (2006). In the near future we have plans to also consistently include into the considerations the cosmic vacuum energy density to see what changes might result from that for the above derived concepts.  An outstanding problem perhaps may still be to better understand the natural philosophical semantics and the logical or physical implications of the above given definitions of $M_\\mathrm{u}$ and $R_\\mathrm{u}$. One should perhaps notice that in the above derived formulae no strict physical reason is presented why cosmic matter density $\\rho_\\mathrm{0}$ scales with $R_\\mathrm{u}^{-2}$. We only have derived above a critical world radius $R_\\mathrm{u}=R_\\mathrm{u}(\\rho_\\mathrm{0})$ which is given as function of the prevailing cosmic matter density $\\rho_\\mathrm{0}$ in a way such that the relation $\\rho_\\mathrm{0}\\sim R_\\mathrm{u}^{-2}$ is fulfilled. Only when by some cosmological evolution process not discussed here this density changes by its value in time, making $\\rho_\\mathrm{0}=\\rho_\\mathrm{0}(t)$, then it turns out that the accordingly changing critical radius $R_\\mathrm{u}=R_\\mathrm{u}(\\rho_\\mathrm{0}(t))$ also changes, as if the relation $\\rho_\\mathrm{0}\\sim R_\\mathrm{u}^{-2}$ would be valid. The many cosmological implications of that shall have to be discussed in more detail in forthcoming considerations."
    },
    "0606/astro-ph0606495_arXiv.txt": {
        "abstract": "A recent submillimeter line survey of Orion KL claimed detection of SiH.  This paper reports on GBT observations of the 5.7 GHz $\\Lambda$-doubling transitions of SiH in Orion.  Many recombination lines, including C164$\\delta$, are seen, but SiH is not detected.  The nondetection corresponds to an upper limit of $1.5 \\times 10^{15}$~cm$^{-2}$ ($4\\,\\sigma$) for the beam-averaged column density of SiH.  This suggests that the fractional abundance of SiH in the extended ridge is no more than twice that in the hot core. ",
        "introduction": "Interstellar hydrides are the building blocks of common molecules.  In dense regions, unsaturated hydrides represent a critical intermediate step toward the formation of saturated hydrides as well as more complex oxygenated and nitrogenated molecules.  For instance, SiO, SiN, SiS, SiC, and other molecules may be produced in large part due to gas-phase reactions with SiH \\citep{turner77,mackay95}.  But despite the important role that simple hydrides play in astrochemistry, many have never been observed \\citep{vandishoeck95}. In the case of the SiH, this may be partly due to the fact that early estimates of the $\\Lambda$-doubling frequencies were highly uncertain. The $\\Lambda$-doubling frequency of 2940 MHz reported by \\citet{douglas65} has a large (10\\%) uncertainty due to the extrapolation of results obtained experimentally at high rotational levels to low rotational levels.  Approximate calculations by \\citet{wilson75} for the ground-state triplet at 3.0 GHz are in error by more than 150 MHz.  More recent calculations by Brown, Curl, \\& Evenson (1984,1985) give the transition frequencies to an accuracy of 3 MHz or $\\sim 0.1$\\%.  SiH (also known as silylidyne and silicon hydride) is isovalent to the CH radical (methylidyne).  The ground electronic state is $^2\\Pi_{1/2}$, with a triplet of $\\Lambda$-doubling $J = 1/2$ transitions near 3000 MHz \\citep[e.g.,][]{brown85}.  This frequency is inaccessible from most present radio telescopes.  However, the first-excited state ($^2\\Pi_{1/2}, J = 3/2$) $\\Lambda$-doubling quartet is found at 5.7 GHz (Figure \\ref{fig-levels}), which is tunable at many radio telescopes.  In a recent submillimeter line survey of Orion KL, \\citet{schilke01} report the first and only tentative detection of interstellar SiH.  However, the six hyperfine $^2\\Pi_{1/2}, J = 3/2 \\rightarrow 1/2$ transitions are blended with strong emission from more common molecules (SO$_2$, CH$_3$CN, and CH$_3$OCH$_3$). \\citeauthor{schilke01} regard their detection as merely tentative and suggest that an interferometric follow-up may be required to determine whether the detected features are due to SiH.  This paper reports on observations of the 5.7 GHz $\\Lambda$-doubling lines of SiH in order to attempt to confirm the \\citeauthor{schilke01} result in the centimeter wavelength regime.  Because the density of molecular line transitions at centimeter wavelengths is much lower than in the submillimeter, detection of the 5.7 GHz lines would conclusively demonstrate that SiH is indeed present in Orion KL.  An additional motivation was the possibility of being able to determine the frequencies to much better accuracy than the 3 MHz obtained by the laboratory measurements of \\citet{brown85}.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[]{f1.eps}} \\caption{Diagram of the ground state and first rotationally excited state of SiH.  \\citet{schilke01} claim a tentative detection of the two $^2\\Pi_{1/2}, J = 3/2 \\rightarrow 1/2$ triplets.  Observations of the indicated $J = 3/2$ $\\Lambda$-doubling transitions are presented in this paper.  Frequencies are from \\citet{brown85}.\\label{fig-levels}} \\end{figure}  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[]{f2.eps}} \\caption{Spectral overview of Orion KL at 5760 MHz.  Calibrated frequency-switched data are shown with no baseline fitted.  Key recombination lines are identified.  These data were taken using a 1 MHz switch frequency to clearly resolve the H104$\\alpha$ and He104$\\alpha$ lines with minimal overlap.  Due to the method of frequency switching, approximately half-height negatives of lines appear offset by plus and minus the switch frequency.\\label{fig-overview}} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[]{f3.eps}} \\caption{Enlargement of portion of Figure \\ref{fig-overview} showing H149$\\gamma$, He149$\\gamma$, and H176$\\epsilon$ features. \\emph{Top}: Calibrated frequency-switched, Hanning-smoothed data using a switch frequency of 1 MHz.  The dashed green lines indicate the center frequencies of fitted Gaussian features.  A sinusoidal ripple has also been subtracted.  The red curve shows the sum of the Gaussian and sinusoidal fits.  \\emph{Middle}: Fit residuals. \\emph{Bottom}: Restored spectrum.  The positive Gaussian from each line fit is added back to the residuals. \\label{fig-hi}} \\end{figure} ",
        "conclusions": "A search for the 5.7 GHz $\\Lambda$-doubling lines of SiH in Orion KL did not yield a detection.  The upper limit column density of $1.5 \\times 10^{15}$~cm$^{-2}$ in the upper level suggests that the fractional abundance of SiH is not significantly higher in the extended ridge than in the hot core.  Numerous recombination lines were detected, including several from helium, carbon, and possibly sulfur.  Recombination line parameters are consistent with previous observations in the literature."
    },
    "0606/astro-ph0606340_arXiv.txt": {
        "abstract": "The direct imaging from the ground  of extrasolar planets has become today a major  astronomical and  biological  focus. This  kind  of imaging  requires simultaneously the use of a  dedicated high performance Adaptive Optics [AO] system and a differential imaging camera in order to cancel  out the flux coming from the star.  In addition, the use of  sophisticated post-processing  techniques  is mandatory  to achieve  the ultimate detection performance required. In the framework of  the SPHERE project, we  present here the development  of a new technique,  based on  Maximum A  Posteriori [MAP]  approach, able  to estimate parameters of a faint companion in  the vicinity of a bright star, using the multi-wavelength images, the  AO closed-loop data as well  as some knowledge on non-common  path and differential aberrations. Simulation  results show a $10^{-5}$   detectivity   at  $5\\sigma$   for   angular  separation   around $15\\frac{\\lambda}{D}$ with only two images. ",
        "introduction": "Today more  than 150 exoplanets have  been detected. But a  great number among them are  known by indirect gravitational  effects on their  parent star. This indirect detection and study allows one to estimate physical parameters of the companion,  like  its  orbital period  or  mass,  but  does not  indicate  its atmosphere composition or its temperature. Exoplanet direct detection from the ground represents today a great scientific gain on our knowledge of exoplanet, since  it  allows one  to  perform  spectroscopy of  the  planet.  But such  a detection needs a major improvement of technologies in use, since the star and its companion  are separated by  a fraction of  arcsecond, and the  flux ratio between   them   is   extremely   high   ($10^6$).   The   SPHERE   instrument \\cite{Beuzit-p-05}, a VLT  Planet Finder, will allow to  detect photons coming from hot Jupiter planets and will be installed on VLT in 2010. This instrument is composed  of a  high performance extreme  AO system  \\cite{Fusco-p-05a}, an optimised  coronagraphic   device  \\cite{Boccaletti-a-04}  and   a  dual  band imager\\cite{Marois-t-04}. But a  dedicated post-processing method is mandatory in order to achieve the ultimate  detection level of SPHERE. In this paper, we will  consider the case  of extreme  AO coupled  to differential  imaging. The common use of  spectral differential images is to  perform differences between images  at  different  wavelength  in  order to  calibrate  the  residuals  of aberration not corrected  by AO and the residuals  of diffraction not canceled by the  coronagraph. The  main limitation of  differential imaging  comes from differential  aberrations  between  the  two  images, or  between object images  and reference images  obtained at  different times.  The  principle of  differential imaging is  detailed in the next section.  We propose in the  third section an optimised  method  dedicated  to  our  specific  issue,  based  on  maximum  a posteriori approach  and able  to estimate the  turbulence parameters  and the object  in a  pair of  images.  We present  in the  fourth section  simulation results for the  estimation of the turbulent phase  structure function and the object. ",
        "conclusions": "We  propose  a method  which  allows  to  solve SD  limitations  (differential aberrations) and gives better results  than DD, without reference images. In a perfect  case,  detectivity  at  $5\\sigma$  reaches  less  than  $10^{-5}$  at 15$\\frac{\\lambda}{D}$.  This result  has still  to be  tested in  more complex cases  (slowly  evolving   static  aberrations  or  mis-calibration,  residual background) but  gives a rough  idea of the  potentiallity of the  method. The perspectives  for  this method  are  to perform  a  global  estimation of  the parameters  (structure   function,  object  in  the  two   images  and  static aberrations), to  process real images  obtained on a differential  imager like NACO SDI, and to generalise our approach to coronagraphy."
    },
    "0606/astro-ph0606389_arXiv.txt": {
        "abstract": "We present new \\xmm\\ observations of two luminous and high accretion-rate radio-quiet active galactic nuclei (AGNs) at $z\\sim2$. Together with archival \\xray\\ and rest-frame optical spectra of three sources with similar properties as well as 25 moderate-luminosity radio-quiet AGNs at $z<0.5$, we investigate, for the first time, the dependence of the hard (\\gtsim2\\,keV) \\xray\\ power-law photon index on the broad \\hb\\ emission-line width and on the accretion rate across $\\sim3$ orders of magnitude in AGN luminosity. Provided the accretion rates of the five luminous sources can be estimated by extrapolating the well-known broad-line region size-luminosity relation to high luminosities, we find that the photon indices of these sources, while consistent with those expected from their accretion rates, are significantly higher than expected from the widths of their \\hb\\ lines. We argue that, within the limits of our sample, the hard-\\xray\\ photon index depends primarily on the accretion rate. ",
        "introduction": "\\label{introduction}  It is widely accepted that the hard (\\gtsim2\\,keV) \\xray\\ emission from active galactic nuclei (AGNs) is primarily produced via unsaturated inverse Compton scattering of UV--soft-\\xray\\ photons from the accretion disk by a corona of hot, likely thermal, electrons (e.g., Haardt \\& Maraschi 1991; Zdziarski, Poutanen, \\& Johnson 2000; Kawaguchi, Shimura, \\& Mineshige 2001). The emitted \\xray\\ photon spectrum in the \\hbox{$\\approx$2--100\\,keV} energy range is best described by a power-law of the form $N_E\\propto E^{-\\Gamma}$, where $\\Gamma$, the photon index, takes a fairly constant value of $\\sim$2, and is predicted to be only weakly sensitive to large changes in the electron temperature and the optical depth in the corona (e.g., Haardt \\& Maraschi 1991). Throughout this work, we refer to $\\Gamma$ only as the photon index in the hard (\\gtsim2\\,keV) \\xray\\ band.  Nandra \\& Pounds (1994) found a $\\left < \\Gamma \\right >$=1.95$\\pm$0.15 for a sample of bright Seyfert\\,1 galaxies, confirming the mean predicted $\\Gamma$ and its small dispersion. Brandt, Mathur, \\& Elvis (1997) have studied the \\xray\\ spectra of a sample of nearby AGNs, including low--moderate luminosity Narrow-Line Seyfert~1 (NLS1) galaxies, defined as type~1 AGNs having \\hb\\ full width at half-maximum intensity (FWHM)\\ltsim2000\\,\\kms\\ (Osterbrock \\& Pogge 1985). The addition of NLS1s, which were previously overlooked due to observational biases, broadened the range of $\\Gamma$ values, and allowed Brandt \\et (1997) to find that $\\Gamma$ is anticorrelated with FWHM(\\hb). This extended the analogous relation between FWHM(\\hb) and the effective \\xray\\ spectral slope in the \\hbox{0.1--2.0\\,keV} band to harder \\hbox{X-rays} (e.g., Boller, Brandt, \\& Fink 1996; Laor~\\et~1997). The pseudo-thermal soft-\\xray\\ excess radiation may be responsible for cooling the corona, thus steepening the \\xray\\ spectrum (e.g., Pounds, Done, \\& Osborne 1995), and the pronounced soft excesses of NLS1s (e.g., Puchnarewicz \\et 1995) may readily explain their steep \\xray\\ spectra and the $\\Gamma$-FWHM(\\hb) anticorrelation.  The remarkable dependence of the spectral shape of the \\xray\\ emission, originating\\ltsim30\\,$R_{\\rm S}$ from the central engine, on the width of a broad-emission line region (BELR) line, emitted at $\\approx$10$^{4}\\,R_{\\rm S}$, was interpreted as a fundamental dependence of $\\Gamma$ on the accretion rate, since, as explained below, FWHM(\\hb) is considered an accretion-rate indicator in type~1 AGNs (e.g., Boroson \\& Green 1992; Brandt \\& Boller 1998). A high accretion rate would increase the disk temperature hence producing more soft-\\xray\\ radiation (the ``soft excess'') and, at the same time, increase the Compton cooling of the corona and steepen the hard-\\xray\\ power law.  Recent \\xray\\ studies of nearby ($z$\\ltsim0.5) type~1 radio-quiet quasars (RQQs) have consistently shown that $\\Gamma \\sim2.0\\pm0.5$, and some of those studies confirmed the Brandt \\et (1997) $\\Gamma$-FWHM(\\hb) relationship (e.g., Leighly 1999; Reeves \\& Turner 2000; Porquet \\et 2004; Piconcelli \\et 2005, hereafter P05; Brocksopp \\et 2006). Photon indices of $\\Gamma \\sim2.0$ are also observed for 0.5\\ltsim $z$\\ltsim6 RQQs (e.g., Reeves \\& Turner 2000; Page \\et 2005; Shemmer \\et 2005), however, the spread in their $\\Gamma$ values seems smaller than that for nearby sources (see Fig.\\,3 of Shemmer \\et 2005). On the other hand, Grupe \\et (2006) report $\\left < \\Gamma \\right > =2.21\\pm0.52$ for a sample of $z>4$ RQQs, and suggest that the relatively steep \\xray\\ spectra of such sources may be attributed to high accretion rates, in analogy with NLS1s.  \\begin{deluxetable*}{lcccccccc} \\tablecolumns{9} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablecaption{{\\sl XMM-Newton} Observation Log \\label{obs_log}} \\tablehead { \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{{\\sc Observation}} & \\multicolumn{3}{c}{{\\sc Net Exposure Time (ks) / Source Counts}} \\\\ \\colhead{{\\sc Quasar}} & \\colhead{{\\sc RA (J2000.0)}} & \\colhead{{\\sc DEC (J2000.0)}} & \\colhead{$z$\\tablenotemark{a}} & \\colhead{$N_{\\rm H}$\\tablenotemark{b}} & \\colhead{{\\sc Start Date}} & \\colhead{MOS1} & \\colhead{MOS2} & \\colhead{{\\sc pn}} } \\startdata \\qthirteenlong\\ & 13 48 44.08 & $-$03 53 25.0 & 2.370 & 2.47 & 2005 Jul 05 & 10.4 / 90 & 10.4 / 144 & 9.4 / 520 \\\\ \\hetwotwo\\ & 22 20 06.77 & $-$28 03 23.9 & 2.414 & 1.28 & 2005 Oct 31 & 25.1 / 732 & 25.1 / 790 & 21.2 / 2968 \\enddata \\tablenotetext{a}{Systemic redshift measured from the optical emission lines and obtained from S04.} \\tablenotetext{b}{Neutral Galactic absorption column density in units of $10^{20}$\\,cm$^{-2}$ obtained from Dickey \\& Lockman (1990).} \\end{deluxetable*}  In this {\\it Letter} we test the hypothesis that the accretion rate largely determines the hard-\\xray\\ spectral slope in AGNs, and argue that previous studies have not been able to break the degeneracy between the dependence of $\\Gamma$ on FWHM(\\hb) and the accretion rate, since highly luminous AGNs were left out of the analyses (e.g., Porquet \\et 2004). This is performed by investigating \\xray\\ and \\hb\\ spectral-region data for a well-defined sample of 30 moderate--high-luminosity RQQs, selected for having high-quality \\xray\\ and \\hbox{\\hb-region} spectroscopy. Five of the sources are luminous RQQs at $z\\sim2$, allowing, for the first time in this context, expansion of the AGN parameter space by $\\sim3$ orders of magnitude in luminosity and black-hole mass (\\mbh). Two of the $z\\sim2$ RQQs, \\qthirteenlong\\ and \\hetwotwo, were selected for \\xmm\\ observations from the Shemmer \\et (2004; hereafter S04) sample of luminous and high accretion-rate quasars based on their expected high \\xray\\ fluxes in that sample. In \\S\\,\\ref{observations} we present the new \\xmm\\ observations, their reduction, and the \\xray\\ spectral fitting of these two sources; their \\xray\\ and optical properties are presented in \\S\\,\\ref{individual}. In \\S\\,\\ref{indicator} we discuss our results, and in particular, the dependence of $\\Gamma$ on the accretion rate for AGNs. Throughout this work we consider only radio-quiet AGNs to avoid any contribution from jet-related emission to the \\xray\\ spectra. Luminosity distances are computed using the standard cosmological model with parameters $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$, and $H_{0}=70$\\,\\kms\\,Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0606/astro-ph0606030_arXiv.txt": {
        "abstract": "Using the \\emph{Swift} data of GRB 050315, we progress on the uniqueness of our theoretically predicted Gamma-Ray Burst (GRB) structure as composed by a proper-GRB (P-GRB), emitted at the transparency of an electron-positron plasma with suitable baryon loading, and an afterglow comprising the so called ``prompt emission'' as due to external shocks. Thanks to the \\emph{Swift} observations, the P-GRB is identified and for the first time we can theoretically fit detailed light curves for selected energy bands on a continuous time scale ranging over $10^6$ seconds. The theoretically predicted instantaneous spectral distribution over the entire afterglow is presented, confirming a clear hard-to-soft behavior encompassing, continuously, the ``prompt emission'' all the way to the latest phases of the afterglow. ",
        "introduction": "GRB 050315 \\citep{va05} has been triggered and located by the BAT instrument \\citep{b04,ba05} on board of the {\\em Swift} satellite \\citep{ga04} at 2005-March-15 20:59:42 UT \\citep{pa05}. The narrow field instrument XRT \\citep{bua04,bua05} began observations $\\sim 80$ s after the BAT trigger, one of the earliest XRT observations yet made, and continued to detect the source for $\\sim 10$ days \\citep{va05}. The spectroscopic redshift has been found to be $z = 1.949$ \\citep{kb05}.  We present here the results of the fit of the \\emph{Swift} data of this source in $5$ energy bands in the framework of our theoretical model \\citep[see][and references therein]{rlet1,rlet2,rubr,rubr2,EQTS_ApJL,EQTS_ApJL2,PowerLaws}, pointing out a new step toward the uniqueness of the explanation of the overall GRB structure. In section \\ref{model} we recall the essential features of our theoretical model; in section \\ref{fit} we fit the GRB 050315 observations by both the BAT and XRT instruments; in section \\ref{spectra} we present the instantaneous spectra for selected values of the detector arrival time ranging from $60$ s (i.e. during the so called ``prompt emission'') all the way to $3.0\\times 10^4$ s (i.e. the latest afterglow phases); in section \\ref{concl} we present the conclusions. ",
        "conclusions": "Before the \\emph{Swift} data, our model could not be directly fully tested. With GRB 050315, for the first time, we have obtained a good match between the observational data and our predicted intensities, in $5$ energy bands, with continuous light curves near the beginning of the GRB event, including the ``prompt emission'', all the way to the latest phases of the afterglow. This certainly supports our model and opens a new phase of using it to identify the astrophysical scenario underlying the GRB phenomena. In particular: \\begin{enumerate} \\item We have demonstrated that the ``prompt emission'' is not necessarily due to the prolonged activity of an ``inner engine'', but corresponds to the emission at the peak of the afterglow. \\item We have a clear theoretical prediction on the total energy emitted in the P-GRB $E_{P-GRB} = 1.98 \\times 10^{51}$ erg and on its temporal separation from the peak of the afterglow $\\Delta t^d_a = 51$ s. To understand the physics of the inner engine more observational and theoretical attention should be given to the analysis of the P-GRB. \\item We have uniquely identified the basic parameters characterizing the GRB energetics: the total energy of the black hole dyadosphere $E_{dya} = 1.46\\times 10^{53}$ erg and the baryon loading parameter $B = 4.55 \\times 10^{-3}$. \\item The ``canonical behavior'' in almost all the GRB observed by \\emph{Swift}, showing an initial very steep decay followed by a shallow decay and finally a steeper decay, as well as the time structure of the ``prompt emission'' have been related to the fluctuations of the ISM density and of the ${\\cal R}$ parameter. \\item The theoretically predicted instantaneous photon number spectrum shows a very clear hard-to-soft behavior continuously and smoothly changing from the ``prompt emission'' all the way to the latest afterglow phases. \\end{enumerate}  Only the first afterglow regime we theoretically predicted, which corresponds to a bolometric luminosity monotonically increasing with the photon detector arrival time, preceding the ``prompt emission'', still remains to be checked by direct observations. We hope in the near future to find an intense enough source, observed by the \\emph{Swift} satellite, to verify this still untested theoretical prediction.  As a byproduct of the results presented in this Letter, we can explain one of the long lasting unanswered puzzles of GRBs: the light curves in the ``prompt emission'' show very strong temporal substructures, while they are remarkably smooth in the latest afterglow phases. The explanation follows from three factors: 1) the value of the Lorentz $\\gamma$ factor, 2) the EQTS structure and 3) the coincidence of the ``prompt emission'' with the peak of the afterglow. For $\\gamma \\sim 200$, at the peak of the afterglow, the diameter of the EQTS visible area due to relativistic beaming is small compared to the typical size of an ISM cloud. Consequently, any small inhomogeneity in such a cloud produces a marked variation in the GRB light curve. On the other hand, for $\\gamma \\to 1$, in the latest afterglow phases, the diameter of the EQTS visible area is much bigger than the typical size of an ISM cloud. Therefore, the observed light curve is a superposition of the contribution of many different clouds and inhomogeneities, which produces on average a much smoother light curve \\citep[details in][]{r02,rubr}."
    },
    "0606/astro-ph0606560_arXiv.txt": {
        "abstract": "{The THz atmospheric ``windows,'' centered at roughly 1.3 and 1.5~THz, contain numerous spectral lines of astronomical importance, including three high-J CO lines, the {\\np} line at 205~$\\mu$m, and the ground transition of para-{\\hhd}. The CO lines are tracers of hot (several 100~K), dense gas; {\\np} is a cooling line of diffuse, ionized gas; the {\\hhd} line is a non-depleting tracer of cold ({\\ca}20~K), dense gas. } {As the THz lines benefit the study of diverse phenomena (from high-mass star-forming regions to the WIM to cold prestellar cores), we have built the {\\bf CO N$^+$ D}euterium {\\bf O}bservations {\\bf R}eceiver ({\\bf CONDOR}) to further explore the THz windows by ground-based observations. } {CONDOR was designed to be used at the Atacama Pathfinder EXperiment (APEX) and Stratospheric Observatory For Infrared Astronomy (SOFIA). CONDOR was installed at the APEX telescope and test observations were made to characterize the instrument. } {The combination of CONDOR on APEX successfully detected THz radiation from astronomical sources. CONDOR operated with typical T$_{rec}=1600$~K and spectral Allan variance times of {\\ca}30~s. CONDOR's ``first light'' observations of CO 13-12 emission from the hot core Orion FIR4 (= OMC1 South) revealed a narrow line with T$_{MB}\\approx 210$~K and $\\Delta V\\approx 5.4$~{\\kms}. A search for {\\np} emission from the ionization front of the Orion Bar  resulted in a non-detection. } {The successful deployment of CONDOR at APEX demonstrates the potential for making observations at THz frequencies from ground-based facilities. } ",
        "introduction": "{\\bf{CONDOR}} ({\\bf{CO N}}$^{+}$ {\\bf{D}}euterium {\\bf{O}}bservations {\\bf{R}}eceiver) is currently one of the very few instruments that can observe at Terahertz (THz) frequencies. The scarcity of astronomical data in the THz frequency regime (1-10~THz, 300-30$\\mu$m) is due to the difficulty of building receivers for these frequencies, and -- for ground-based observatories -- also due to the poor transmission of the Earth's atmosphere (e.g., Pardo et al. \\cite{pardo}).  Currently, the only astronomical, heterodyne data above 1~THz obtained from the ground are from the Heinrich Hertz Telescope (HHT) at 1.0~THz (Kawamura {\\em et al.} 2002) and the Receiver Laboratory Telescope (RLT) at 1.0 and 1.5~THz (Marrone {\\em et al.} 2004, 2006). In addition, {\\np} emission (1.5~THz) was detected with moderate spectral resolution by the South Pole Imaging Fabry-Perot Interferometer (SPIFI) from the Antarctic Submillimeter Telescope and Remote Observatory (AST/RO) (Stacey 2005). There is also a 1.3 THz and 1.5 THz heterodyne receiver for APEX under construction at Chalmers University. The Kuiper Airborne Observatory (KAO) pioneered FIR spectroscopy, initially with incoherent instruments and moderate velocity resolution (e.g. Stutzki et al. 1988, Petuchowski et al. 1994) and later  with a heterodyne receiver ({\\em e.g.}, Boreiko \\& Betz 1993). The Infrared Space Observatory (ISO) observed numerous lines in many galactic and extragalactic sources at low velocity resolution ({\\em e.g.}, van Dishoeck~2004 and references therein) and the Cosmic Background Explorer (COBE) ({\\em e.g.}, Fixsen, Bennett, \\& Mather 1999) demonstrated the large  extent of several of the important cooling lines of the ISM.  These observations have demonstrated that studies of many astronomical phenomena greatly benefit from data at THz frequencies. In order to explore the universe at THz frequencies and encouraged by both advances in mixer technology and the capabilities of 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 f\\\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}) (G\\\"usten et al. 2006) we have built CONDOR. ",
        "conclusions": "CONDOR has been successfully deployed on the APEX telescope. CONDOR operated with typical T$_{rec}\\sim1600$~K and spectral Allan variance times of 30~s. CONDOR's first light observations detected CO 13-12 emission from Orion FIR4. The line has a peak of  {\\ca}210~K and a width of {\\ca}5~{\\kms} Uncertainties in the beam shape and source  extent make the temperature scaling uncertain, but the line width is clearly smaller than that of lower-J CO lines. The core density and temperature indicated by the CO 13-12 emission are consistent with values determined from other CO observations. The narrow width of the high-J line indicates that the excitation of this warm, dense material is due to photo-heating rather than shocks. % CONDOR failed to detect {\\np} emission from the Orion Bar."
    },
    "0606/astro-ph0606756_arXiv.txt": {
        "abstract": "We present an atlas of \\spitzer\\ Space Telescope Infrared Spectrograph (\\irs) spectra of highly luminous, compact mid-infrared sources in the Large Magellanic Cloud.  Sources were selected on the basis of infrared colors and 8~\\micron\\ (\\msx) fluxes indicative of highly evolved, intermediate- to high-mass stars with current or recent mass loss at large rates.  We determine the chemistry of the circumstellar envelope from the mid-IR continuum and spectral features and classify the spectral types of the stars. In the sample of 60 sources, we find 21 Red Supergiants (RSGs), 16 C-rich Asymptotic Giant Branch (AGB) stars, 11 \\htwo\\ regions, 4 likely O-rich AGB stars, 4 Galactic O-rich AGB stars, 2 OH/IR stars, and 2 B[e] supergiants with peculiar IR spectra.  We find that the overwhelming majority of the sample AGB stars (with typical IR luminosities $\\sim 10^4$~\\lsun) have C-rich envelopes, while the O-rich objects are predominantly luminous RSGs with \\lir\\,$\\sim 10^5$~\\lsun. For both classes of evolved star (C-rich AGB stars and RSGs), we use the near- to mid-infrared spectral energy distributions to determine mean bolometric corrections to the stellar K-band flux densities. For carbon stars, the bolometric corrections depend on the infrared color, whereas for RSGs, the bolometric correction is independent of IR color.  Our results reveal that objects previously classified as PNe on the basis of IR colors are in fact compact \\protect\\htwo\\ regions with very red \\irs\\ spectra that include strong atomic recombination lines and polycyclic aromatic hydrocarbon (PAH) emission features.  We demonstrate that the \\irs\\ spectral classes in our sample separate clearly in infrared color-color diagrams that use combinations of \\tmass\\ data and synthetic \\irac/\\mips\\ fluxes derived from the \\irs\\ spectra. On this basis, we suggest diagnostics to identify and classify, with high confidence levels, IR-luminous evolved stars and compact \\protect\\htwo\\ regions in nearby galaxies using \\spitzer\\ and near-infrared photometry. ",
        "introduction": "\\label{sec:intro}   While studies in visible light, particularly spectroscopy, have been key to our understanding the nature and evolution of stars during most of their life cycles, the early and late stages of stellar evolution are best studied in the infrared.  This is because, during these stages, the stars are embedded in regions of dust and gas.  The late stages of stars of intermediate initial mass (1 -- 8~\\msun) are characterized by very high luminosities and obscuration by dusty, expanding circumstellar envelopes (see review by \\citealt{van03}).  Circumstellar dust envelopes absorb photospheric optical and ultraviolet emission and re-radiate it in the mid- to far-infrared (MFIR). Thus evolved intermediate-mass stars are luminous in the infrared (IR) but may be highly obscured at optical wavelengths.  Current theories of the later stages of the evolution of mass-losing stars therefore have been informed by \\iras, \\iso, and \\tmass\\ studies of intermediate-mass evolved stars in the Milky Way Galaxy and Magellanic Clouds.   Understanding the envelope chemistry of optically-obscured post-main sequence (MS) stars through IR studies is important for our overall view of galaxy evolution, since these stars dominate the rate of return of stellar processed material to the ISM \\citep{jur89,bar91,leb01}.  Optical surveys (e.g., \\citealt{gro99} and refs.\\ therein) can miss these objects because they are too faint to be detected. Dusty circumstellar shells in evolved stars are produced as the star loses mass via a combination of pulsations and radiation pressure on dust grains newly formed via the pulsation process. Once accelerated, these grains strip the (predominantly molecular) gas away from the stellar envelope.  While on the Asymptotic Giant Branch (AGB), intermediate mass stars undergo dredge-up and surface enrichment in products of nucleosynthesis.  The enriched material is thus returned to the ISM via the molecule-rich circumstellar dust shell.  More massive ($>$8~\\msun) stars in the red supergiant (RSG) stage may resemble high mass-loss rate AGB stars, in that they are also surrounded by optically-thick circumstellar dust shells, though such high mass-loss rate RSGs are rare.  Highly-obscured, evolved intermediate-mass (AGB) stars with large mass-loss rates likely dominate the rate of return of nuclear processed material to the interstellar medium. However, in our own Galaxy, it can be difficult to distinguish between high mass-loss AGB and RSG stars, due to their highly uncertain distances.   IR spectroscopy is particularly useful in determining the properties of circumstellar dust.  The emission and absorption properties of different dust grains have been studied in the laboratory and modeled (e.g., \\citealt{dra01,li01}).  IR-bright stars in the Galaxy have been observed with the \\iras\\ LRS and, subsequently, with the spectrometers aboard \\iso\\ (see reviews by \\citealt{wat99a, wat99b, bar99, bar97} and references therein) and classified using the dust features revealed in these spectra (e.g., \\citealt{kra02}).  The sensitivity of the \\spitzer\\ Infrared Spectrograph\\footnote{This work is based on observations made with the \\spitzer\\ Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and Ames Research Center.} (\\irs, \\citealt{hou04}) on the \\spitzer\\ Space Telescope \\citep{wer04} now makes it possible to study the IR spectra of IR-luminous, evolved stars in nearby galaxies.   In the absence of spectroscopy, photometric tools such as color-color diagrams can be used to identify and classify evolved stars.  New studies of stellar populations in the Galaxy as well as in nearby external galaxies will rely in particular on photometry obtained with the \\spitzer\\ Infrared Array Camera (\\irac, \\citealt{faz04}) and Multiband Imaging Photometer for \\spitzer\\ (\\mips, \\citealt{rie04}).  Indeed, the point sources detected in \\irac\\ and \\mips\\ imaging of nearby galaxies likely will be dominated by dust-enshrouded post-MS stars. As strong silicate and carbonaceous dust features at $\\sim$10~\\micron\\ distinguish C-rich and O-rich envelope chemistries, respectively, the use of \\irac\\ and \\mips\\ colors can, in principal, avoid the ambiguity in assigning stellar classes based on near-infrared (NIR) colors (e.g., \\citealt{ega01}, hereafter \\egant).  However, such photometric diagnostics must first be reliably associated with IR spectral properties, which are best determined from \\irs\\ spectroscopy.   As previous infrared surveys have demonstrated (e.g., \\citealt{gro95, gro98, zij96, lou97, van97, van98, tra99}), the Large Magellanic Cloud (LMC) is an ideal nearby galaxy to study the evolution of mass-losing post-MS stars. Firstly, it contains a large population of IR-luminous, mass-losing objects found at essentially the same distance, thereby alleviating the distance ambiguities that haunt \\iras\\ and \\iso\\ samples of mass-losing stars in the Galaxy. This allows the spectral properties of the LMC evolved star population to be directly related to their luminosities, providing important constraints on stellar models. Indeed, the luminosity distribution of carbon stars in the LMC has long been a key test of models of the late stages of stellar evolution in low metallicity environments (e.g., \\citealt{ibe83}).  Secondly, the low metallicities and high star formation rates of the LMC mimic those of high-redshift galaxies.  In addition, spectra of LMC objects can be used to establish ``ground truth'' for IR photometric diagnostics.  Due to the proximity of the LMC, individual objects have been detected by previous IR satellites: \\iras\\ and \\msx\\ each detected $\\sim$1800 IR sources in or towards the LMC (\\citealt{sch89}; \\egant).  The most sensitive band on \\msx, A band ($\\lambda_C$ = 8.3~\\micron), was approximately four times more sensitive than the \\iras\\ 12~\\micron\\ band, and permitted the detection of evolved stars with less extreme IR luminosities than those detected by \\iras\\ \\eganp.  \\egant\\ combined the \\msx\\ A band flux (in magnitudes) with near-IR (JHK) measurements from the \\tmass\\ survey to form various diagnostic color-color and color-magnitude diagrams for the LMC objects.  To classify their nature based upon these IR fluxes, they compared these color-color and color-magnitude diagrams to analogous diagrams constructed for various classes of stars based upon the \\citet{wai92} ``SKY'' model of the Galaxy.  On this basis, they characterized the stars in specific categories, including red supergiants, early (low mass loss) O-rich and C-rich AGB stars (\\egant\\ category 5), late (high mass loss) O-rich and C-rich AGB stars (category 6), and PNs and H~II regions (category 7).  Within these categories, the different types of objects were distinguished somewhat on the basis of colors, although only for the late C-rich and O-rich AGB stars were the objects well distinguished within a category.  Note that \\egant\\ classified $\\sim$500 of the $\\sim$1800 objects in their study as foreground Galactic objects.  To check the accuracy of this classification scheme for the LMC objects, \\egant\\ investigated the spectral types of the LMC objects using data in SIMBAD. However, for the bona fide LMC objects, only a small percentage had classifications; these were especially uncommon for categories 6 and 7. Although the limited data do not contradict their classifications, these classifications suffer from limited spectroscopic support.  This is not surprising, as the IR-luminous objects detected by \\msx\\ do not commend themselves to visible spectroscopy, and the requisite mid-IR spectroscopy has been impractical until only recently.  To date, the study by \\egant\\ stands as the most comprehensive classification of luminous 8~\\micron\\ sources in the LMC. Therefore it is important to verify and/or correct the \\egant\\ color classifications.   We have conducted a spectroscopic study with \\spitzer\\ \\irs\\ (PI: J. Kastner; PID: 3426) of a representative sample of objects selected on the basis of their IR colors and magnitudes to be the most IR-luminous post-MS stars in the LMC.  Here we present the atlas of spectra. We then use these spectra as the basis both for understanding the high-luminosity end of the LMC evolved star population and for establishing photometric diagnostics for classifying IR stars.  In \\S\\ref{sec:obs} we describe the sample selection, observations, and data reduction. In \\S\\ref{sec:res} we present the spectra and classify them according to continuum and spectral features and bolometric luminosities.  In \\S\\ref{sec:dis} we describe the classes of objects identified and compare the classifications with the types assigned by \\egant. We then use our library of \\irs\\ spectra to show how \\spitzer/\\tmass\\ colors can be used to classify IR-luminous sources in external galaxies reliably and unambiguously. On the basis of our new classifications, we discuss the implications for the late stages of stellar evolution. Detailed analysis of individual spectra and of the spectral classes established in this study will be presented in future papers. ",
        "conclusions": "\\label{sec:con}  We have obtained \\spitzer\\ low-resolution spectra for a representative sample of 60 luminous 8~\\micron\\ sources in the LMC. The sources observed were selected to include a variety of evolved, mass-losing stars, based on recently established \\tmass/\\msx\\ color criteria. The \\irs\\ spectra covered the wavelength range 5 -- 35~\\micron, allowing determination of the envelope chemistry from detailed spectral features, particularly silicate emission features at 10 and 18~\\micron\\ in O-rich objects and the SiC emission feature at 11~\\micron\\ and \\ctwohtwo\\ absorption feature at 13.7~\\micron\\ in C-rich objects. We derived IR luminosities from the combined \\tmass\\ and \\irs\\ spectral energy distributions of the objects and used these, along with the spectra, to classify the object types. We identify 16 C-rich AGB stars, 4 O-rich AGB stars, 21 RSGs, 11 \\htwo\\ regions, and 2 OH/IR stars (one supergiant and one AGB star). We also find 2 B supergiants that have peculiar IR spectra and 4 O-rich AGB stars that are in fact foreground Galactic Mira variables.  We find the vast majority of the AGB stars, with IR luminosities $\\sim 10^4$~\\lsun, have C-rich envelopes, while the O-rich objects are almost all luminous RSGs with \\lir$\\gtrsim 10^5$. The predominance of C-rich stars along the AGB in the LMC is consistent with the hypothesis that C-rich stars form more easily than O-rich stars in lower metallicity environments. However, all C-rich AGB stars in our sample have \\lir\\ $< 2 \\times 10^{4}$~\\lsun, while all four O-rich AGB stars have \\lir\\ $\\gtrsim 2 \\times 10^{4}$~\\lsun, suggesting a relatively firm upper limit ($\\sim$3.7~\\msun) for the progenitor mass of LMC Carbon stars.  The fact that all of the supergiant stars are O-rich, in turn, may indicate that massive supergiants do not have sufficiently long He-burning lifetimes to produce C-rich surfaces, and hence C-rich ejecta, despite the low metallicity of the LMC.  Comparison of the \\irs\\ classifications with the NIR color-based types \\eganp\\ revealed a number of inaccuracies in the \\tmass/\\msx\\ classification criteria. All of the sample objects with \\egant\\ type PN were reclassified, mostly as \\htwo\\ regions. Stars classified as OH/IR stars on the basis of NIR colors are mostly C-rich AGB stars, based on their \\irs\\ spectra. The O-rich AGB stars in the sample were reclassified as either RSGs or Galactic O~AGB stars, based on their \\irs\\ spectra and luminosities.  On the basis of these reclassifications, we have developed revised \\tmass/\\msx\\ color-color and color-magnitude classification criteria for LMC objects and, by extension, for luminous 8~\\micron\\ sources in other external galaxies.  For the C-rich AGB and RSG stars, bolometric corrections to the stellar K-band flux densities were derived using the observed IR SEDs. For C-rich AGB stars, the bolometric corrections were found to depend on the K-A color, allowing the bolometric luminosity to be estimated from the K and A-band fluxes with a typical uncertainty of less than 20\\%, and at most a factor of 2.  The bolometric corrections of the RSGs were found to be independent of IR color; for these objects the bolometric luminosities can be predicted from the K-band magnitude with an uncertainty of less than 15\\%.  The derived bolometric corrections could prove very useful for determining the luminosities of carbon stars and RSGs in the Galaxy and other nearby galaxies.  Infrared (\\irac/\\mips/\\tmass) colors were determined from the \\irs\\ spectra and \\tmass\\ fluxes. We found that the spectral types separate well in various mid-IR color-color diagrams, indicating that \\spitzer\\ photometry can be an effective means to infer the spectral class of dust-enshrouded objects in external galaxies. We suggest diagnostics for classifying IR-luminous sources in the Galaxy and nearby galaxies, using \\spitzer\\ photometry.  These colors may also prove useful in the analysis of \\spitzer\\ photometry of unresolved stellar populations in high-redshift galaxies."
    },
    "0606/astro-ph0606426_arXiv.txt": {
        "abstract": "\\rightskip=0pt  Detailed models of the solar cycle require information about the starting time and rise time as well as the shape and amplitude of the cycle. However, none of these models includes a discussion of the variations in the length of the cycle, which has been known to vary from $\\sim$7 to 17 years.  The focus of our study was to investigate whether this range was associated with a secular pattern in the length of the sunspot cycle.  To provide a basis for the analysis of the long-term behavior of the Sun, we analyzed archival data of sunspot numbers from 1700 - 2005 and sunspot areas from 1874 - 2005.  The independent techniques of power spectrum analysis and phase dispersion minimization were used to confirm the $\\sim$11-year Schwabe Cycle, and to illustrate the large range in the length of this cycle.  Long-term cycles were identified in archival data from 1610 -- 2000 using median trace analyses of the length of the cycle, and from power spectrum analyses of the (O-C) residuals of the dates of sunspot minima and maxima.  The median trace analysis suggested that the cycle length had a period of 183 - 243 years, while the more precise power spectrum analysis identified a period of 188 $\\pm$ 38 years.  We found that the 188-year cycle was consistent with the variation of sunspot numbers and seems to be related to the Schwabe Cycle.  We found a correlation between the times of historic minima and the length of the sunspot cycle such that the length of the cycle was usually highest when the actual number of sunspots was lowest.  The cycle length was growing during the Maunder Minimum when there were almost no sunspots visible on the Sun.  This information can now be used to improve the accuracy of the current solar cycle models, to better predict the starting time of a given cycle. ",
        "introduction": "Solar Cycle 23 was predicted to reach a maximum in early 2000 \\citep{joselynetal96}, with severe geomagnetic storms expected between 1999 and 2005.  However, this cycle had modest activity compared to the previous two cycles ({\\it e.g.}, \\citealt{detomaetal04}).  Our work was motivated by the publicity associated with this solar cycle because of the relatively high level of solar activity exhibited through flares in 2003 October and November. Balasubramaniam \\& Regan (1994) showed that the temporal behavior of solar flares was similar to a sunspot butterfly diagram, with a typical 11-year cycle.  Our interest in predicting magnetic activity cycles associated with other cool stars ({\\it e.g.}, \\citealt{richardsetal03}) led us to investigate the long-term behavior of the solar cycle.  Earlier analyses of the sunspot cycle have been summarized by \\citet{kuklin76} and \\citet{hathaway+wilson04}, and a review of the long-term solar cycle was presented by \\citet{usoskin+mursula03}.  In 1843, Heinrich Schwabe reported that he had identified a possible 10-year periodicity in the pattern of sunspots based on his 17-year study of the Sun from 1826--1843 \\citep{schwabe}.  (Note the typographical error in \\citet{wilson94} that the result was based on a 43-year study.)  In 1848, Rudolph Wolf introduced the relative sunspot number (R) and organized a program of daily observations of sunspots.  He reanalyzed all earlier data and found that the average length of a solar cycle was about 11 yrs. For more than two centuries, solar physicists have applied a variety of techniques to determine the nature of the solar cycle.  The earliest methods involved counting sunspot numbers and determining durations of cyclic activity from sunspot minimum to minimum using the ``smoothed monthly mean sunspot number'' \\citep{waldmeier61, wilson87, wilson94}.  The ``Group sunspot number'' introduced by \\citet{hoyt+schatten98} is a better documented data set than the database of relative sunspot numbers, but the two data sets since 1882 are virtually identical \\citep{hathaway+wilson04}.  Since 1874, sunspot areas have been measured in terms of the total surface area of the solar disk covered by sunspots at a given time \\citep{hath04}.  The analysis of sunspot numbers or sunspot areas is often referred to as a one-dimensional approach because there is only one independent variable, namely sunspot numbers or areas \\citep{wilson94}.  Recently, \\citet{lietal05} introduced a new parameter called the ``sunspot unit area'' in an effort to combine the information about the sunspot numbers and sunspot areas to derive the length of the cycle.  There is also a two-dimensional approach in which the latitude of an observed sunspot is introduced as a second independent variable \\citep{wilson94}.  When sunspots first appear on the solar surface they tend to originate at latitudes around 40 degrees and migrate toward the solar equator.  When such migrant activity is taken into account it can be shown that there is an overlap between successive cycles, since a new cycle begins while its predecessor is still decaying.  This overlap became obvious when \\citet{maunder04} published his butterfly diagram and demonstrated the latitude drift of sunspots throughout the cycles.  Maunder's butterfly diagram showed that although the length of time between sunspot minima is on average 11 years, successive cycles actually overlap by $\\sim$ 1 to 2 years.  In addition, \\citet{wilson87} found that there were distinct solar cycles lasting 10 years as well as 12 years.  This type of behavior suggests that there could be a cyclic pattern in the length of the sunspot cycle.  Sunspot number data collected prior to the 1700's show epochs in which there were almost no sunspots visible on the solar surface.  One such epoch, known as the Maunder Minimum, occurred between the years 1642 and 1705, during which the number of sunspots recorded were very low in comparison to later epochs \\citep{wilson94}. Other studies include the analysis of geophysical data and tree-ring radiocarbon data, which contained residual traces of solar activity \\citep{bal85}, to determine if the Maunder period truly had a lower number of sunspots or whether it was simply a period in which little data were collected or large degrees of errors existed. While the data collected before the 1700's are typically less reliable than those collected in more recent times, the timing of the Maunder Minimum is still relatively accurate because of the correlation with geophysical data.  Other epochs of minimum solar activity in the past have been noted, such as the Oort Minimum from 1010 - 1050, the Wolf Minimum from 1280 - 1340, the Sp\\\"{o}rer Minimum from 1420 - 1530 \\citep{eddy77, stuiver80, siscoe80}, and the Dalton Minimum from 1795 - 1823 \\citep{soon+yaskell03}.  These minima have been derived from historical sunspot records, auroral histories \\citep{eddy76}, and physical models linking dendrochronologically dated radiocarbon concentrations to the solar cycle \\citep{solankietal04}.  Currently, the fine details of a given solar cycle can be predicted accurately only after a cycle has begun ({\\it e.g.}, \\citealt{elling+schwentek92, hathawayetal94}; and many others) because many solar cycle models require information about the starting time and rise time as well as the shape and amplitude of the cycle.  Most of these models have focussed on the analysis of the International sunspot numbers of as many previous cycles as possible in order to predict the pattern for the current cycle.   However, none of these models includes a discussion of the variations in the length of the cycle, which has been known to vary from $\\sim$7 to 17 years.  \\begin{deluxetable}{lll} \\tablecolumns{3} \\tablecaption{Duration of the Data} \\tablewidth{0pc} \\tablehead{ \\colhead{Data Set} & \\colhead{} & \\colhead{Duration of Data}} \\startdata Spot Numbers & Daily & 1818 Jan 8 - 2005 Jan 31 \\\\ & Monthly & 1749 Jan - 2005 Jan \\\\ & Yearly & 1700 - 2004 \\\\ \\cline{1-3} Spot Areas & Daily & 1874 May 9 - 2005 Feb 28 \\\\ & Monthly & 1874 May - 2005 Feb \\\\ & Yearly & 1874 - 2004 \\enddata \\end{deluxetable}  In this paper, we have investigated  the variations in the length of the sunspot number cycle and examined whether the variability can be explained in terms of a secular pattern.  Toward this goal, we applied classical one-dimensional techniques to rederive the periodicities of solar activity using the sunspot number and area data to provide internal consistency in our analysis of the long-term behavior.  These results were used as a basis in the study of the long-term behavior of the cycle length.  In \\S2 we discussed the source of the data; in \\S3 we described the derivation of the cycle from sunspot numbers and sunspot areas using two independent techniques; in \\S4 we examined the variability in the cycle length based on the times of cycle minima and maxima using two independent techniques; and in \\S5 we discussed the results.  \\vskip20pt ",
        "conclusions": "The possible variability in the length of the sunspot cycle was examined through a study of archival sunspot data from 1610 -- 2005.  In the preliminary stage of our study, we analyzed archival data of sunspot numbers from 1700 - 2005 and sunspot areas from 1874 - 2005 using power spectrum analysis and phase dispersion minimization.  This analysis showed that the Schwabe Cycle has a duration of (10.80 $\\pm$ 0.50) years (Table 3) and that this cycle typically ranges from $\\sim$10 -- 12 years even though the entire range is from $\\sim$7 -- 17 years.  The focus of our study was to investigate whether this range was associated with a secular pattern in the length of the Schwabe cycle.  We used our derived value for the Schwabe cycle from Table 3 to examine the long-term behavior of the cycle.  This analysis was based on NGDC data from 1610--2000, a period of 386 years (using sunspot minima) or 385 years (using sunspot maxima).  The long-term cycles were identified using median trace analyses of the length of the cycle and also from power spectrum analyses of the (O-C) residuals of the dates of sunspot minima and maxima.  We used independent approaches because of the inherent uncertainties in deriving the exact times of minima and the even greater complexity in the determination of sunspot maxima.  Moreover, we derived our results from both the cycle minima and the cycle maxima.  The fact that we found similar results from the two data sets suggests that the methods used to determine these cycles (NGDC data) did not have any significant impact on our results.  \\begin{deluxetable}{lcc} \\tablecolumns{3} \\tablecaption{Derived Long Term Solar Cycles} \\tablewidth{0pc} \\tablehead{ \\colhead{} & \\colhead{Gleissberg} & \\colhead{Secular}\\\\ \\colhead{Data Set} & \\colhead{(yrs)} & \\colhead{(yrs)}} \\startdata Cycle Minima & 86.8 $\\pm$ ~8.8 & 188 $\\pm$ 40\\\\ Cycle Maxima & 86.3 $\\pm$ 18.1 & 187 $\\pm$ 37\\\\ Combined & 86.8 $\\pm$ 10.7 & 188 $\\pm$ 38\\\\ \\cline{1-3} Average & 86.6 $\\pm$ 12.5 & 188 $\\pm$ 38 \\enddata \\end{deluxetable}  The median trace analysis of the length of the sunspot number cycle provided secular periodicities of 183 -- 243 years.  This range overlaps with the long-term cycles of $\\sim$90 -- 260 years which were identified directly from the FFT and PDM analyses of the sunspot number and area data (Figures \\ref{f3} and \\ref{f4}).  The power spectrum analysis of the (O-C) residuals of the dates of minima and maxima provided much clearer evidence of dominant cycles with periods of $188 \\pm 38$ years, $87 \\pm 13$ years, and $\\sim40$ years.  These results are significant because at least two long-term cycles have transpired over the $\\sim$385-year duration of the data set.  The derived long-term cycles were compared in Figure \\ref{f8} with documented epochs of significant declines in sunspot activity, like the Oort, Wolf, Sp\\\"{o}rer, Maunder, and Dalton Minima \\citep{eddy77, stuiver80, siscoe80}.  In this figure, the modern sunspot number data were combined with earlier data from 1610-1715 \\citep{eddy76} and with reconstructed (ancient) data spanning the past 11,000 years \\citep{solankietal04}.  These reconstructed sunspot numbers were based on dendrochronologically-dated radiocarbon concentrations which were derived from models connecting the radiocarbon concentration with sunspot number \\citep{solankietal04}.  The reconstructed sunspot numbers are consistent with the occurrences of the historical minima (e.g., Maunder Minimum).  \\citet{solankietal04} found that over the past 70 years, the level of solar activity has been exceptionally strong.  Our 188-year periodicity is similar to the 205-year de Vries-Seuss cycle which has been identified from studies of the carbon-14 record derived from tree rings ({\\it e.g.}, \\citealt{wagneretal01,braunetal05}).  Figure \\ref{f8} compares the historical and modern sunspot numbers with the derived secular cycles of length (a) 183 years (\\S4.2), (b) 243-years (\\S4.2), and (c) 188 years (\\S4.4).  The first two periodicities were derived from the median trace analysis, while the third one was derived from the power spectrum analysis of the sunspot number cycle (O-C) residuals.  The fits for the 183-year periodicity all had the same amplitude, but were moderately out of phase with each other, while the fits for the 243-year periodicity were perfectly in phase for all data sets, but with different amplitudes.  This figure shows that the 183- and 188-year cycles provided a more consistent match to the sunspot number data than the 243-year cycle, especially during the Wolf, Sp\\\"{o}rer, Maunder, and Dalton Minima.  In particular, we note that the more accurately determined 188-year cycle is consistent with the behavior of sunspot numbers over time.  Moreover, the four historic minima since 1200, all occurred during the rising phase of our derived 188-year sunspot cycle when the length of the sunspot cycle was increasing.  Figure \\ref{f8} shows that, on average, the length of the sunspot cycle was highest when the actual number of sunspots was lowest.  According to our model, the length of the sunspot cycle was growing during the Maunder Minimum when there were almost no sunspots visible.  The existence of long-term solar cycles with periods between 90 and 200 years is not new to the literature.  However, the nature of these cycles is still not understood.  While \\citet{hath04} noted evidence of a cycle of $\\sim$90 years in sunspot cycle amplitude variations, our studies of the length of the sunspot cycle have shown that the dominant cycle in these amplitude variations is 188 years, with weaker periods of $\\sim$40 and 87 years.  Our results suggest that this 188-year cycle is related to the basic Schwabe Cycle. While we would like to confirm this result using a longer baseline (beyond 385 years), it should be noted that Schwabe's 10-year period was derived from only 17 years (less than two cycles) of observations \\citep{schwabe}.  In addition, it is well-known from Gauss' work on orbits that an orbit can be uniquely defined if at least three positions along that orbit are known, as long as those positions cover critical parts of the orbit.  So, our 188-year result may yet stand the test of time.  Our model predicts that the length of the sunspot number cycle should increase gradually over the next $\\sim$75 years, accompanied by a gradual decrease in the number of sunspots. This information can now be used to improve the accuracy of the current solar cycle models ({\\it e.g.}, \\citealt{dikpatietal06}) to better predict the starting time of a given cycle.  \\vskip5pt"
    },
    "0606/astro-ph0606610_arXiv.txt": {
        "abstract": "A simplified non-linear numerical model for the development of incompressible magnetohydrodynamics (MHD) in the presence of a strong magnetic field \\Bo\\ and stratification, nicknamed \\shellatm, is presented. In planes orthogonal to the mean field, the non-linear incompressible dynamics is replaced by 2D shell-models for the complex variables $u$ and $b$, allowing one to reach large Reynolds numbers while at the same time carrying out sufficiently long time integrations to obtain a good statistics at moderate computational cost.  The shell-models of different planes are coupled by Alfv\\'en waves propagating along \\Bo. The model may be applied to open or closed magnetic field configurations where the axial field dominates and the plasma pressure is low; here we apply it to the specific case of a magnetic loop of the solar corona heated via turbulence driven by photospheric motions, and we use statistics for its analysis. The Alfv\\'en waves interact non-linearly and form turbulent spectra in the directions perpendicular and, via propagation, also parallel to the mean field.  A heating function is obtained, and is shown to be intermittent; the average heating is consistent with values required for sustaining a hot corona, and is proportional to the aspect ratio of the loop to the power $-1.5$; characteristic properties of heating events are distributed as power-laws. Cross-correlations show a delay of dissipation compared to energy content. ",
        "introduction": "MHD turbulence in the presence of a mean magnetic field, with or without density and gravitational gradients, plays a role in many environments, ranging from stellar coronae and winds \\citep{kleinlw91} to the interstellar medium \\citep{desai94} and accretion disks. In such regions, energy may be transferred, accumulated and dissipated in a way which is inherently anisotropic \\citep{sheba83,oug94,kin98,muller03,oug04}.  In particular, in solar coronal physics, where one of the main problems is to understand how the corona can be sustained at more than a million of degrees, it is believed that the necessary heating could be produced at small scales generated by a non-linear cascade along a turbulent spectrum \\citep{hey92,gom92}. Furthermore, as flux tubes (\\eg in the form of coronal loops or coronal funnels) are omnipresent, the anisotropy coming from the dominant magnetic field may be a central feature of the processes governing energy dissipation, like the non-linear collisions of counter-propagating Alfv\\'en wave packets. It can thus be expected that solving the coronal heating problem, \\ie understanding how the temperature of the corona can be sustained, may require to understand the details of the turbulent dynamics of MHD in these environments.  One way to study and the dynamics of such system is to perform direct numerical simulations (DNS).  In the case of anisotropic MHD, DNS have provided insight on subjects like the anisotropy of the spectra \\citep[\\eg][]{kin98,oug04}, the parametric decay of Alfv\\'en waves \\citep[\\eg][]{delz01}, and Alfv\\'en waves filamentation \\citep[\\eg][]{passot03}.  MHD simulations are also used to study the topology of magnetic field lines and magnetic reconnection in the corona \\citep[\\eg][]{aulanier05b}. But the Reynolds numbers attained in all the DNS up to now are below $10^3$, while they are believed to be $10^{12}$ to $10^{14}$ in the corona.  DNS are very far from being able to represent all the scales of turbulence in the corona, there is a huge gap to fill. Furthermore, as statistics are of great help in the study of turbulence, attempts have been made to analyze statistically energy dissipations produced by DNS.  Distributions of events are for instance presented in \\cite{dmi98} and \\cite{gev98} from 2D DNS of reduced MHD.  But it is still difficult to get significant statistics from 3D DNS, and it is even more difficult when trying to go to higher Reynolds numbers because then the grid resolution must be higher and the computations of the model are too slow. For all these reasons, there is a need for simplified numerical models of MHD, which would run sufficiently fast to get statistics of turbulence at high Reynolds numbers, while keeping the most relevant features of MHD turbulence.  Several approaches have been used to build such simplified numerical models of MHD. For example, the Self-Organized Criticality behavior (SOC) of MHD systems can be modeled by cellular automata (CA), where the interactions of individual cells translate into a global statistic behavior of the whole system, following the first models of \\citet{luh91,luh93}. However, the need for physical realism is not entirely addressed by the cellular automata, despite efforts trying to include the constraints issued from the MHD equations \\citep{vla95, isl00, isl01, buc03}.  Another approach is to simplify the non-linear interactions by reducing the number of modes which are allowed to interact non-linearly. In the coronal loop context, a shell-model approach has been used by \\citet{nig04,nig05}. We have developed a similar numerical model independently, starting from the reduced MHD equations but allowing for the stratification of the plasma. This numerical code, nicknamed \\shellatm, allows one to reach (kinetic and magnetic) Reynolds numbers unachieved before. In this paper, we will focus on the problem of a coronal loop where energy is forced into the system by footpoint motions, describing in detail the dynamics of the heating events, turbulence spectra, statistics and scaling laws. ",
        "conclusions": "We have presented the \\shellatm\\ model, which is a generalization of shell-models \\citep{giu98} with propagation of Alfv\\'en waves along a \\Bo\\ field, with the further possibility of introducing a longitudinal stratification of physical properties (\\Bo, mass density, flux tube expansion factor).  Although the model is simple and includes only simplified physical processes, it has a very interesting complex non-linear dynamics and it is fast enough to get statistics of its fields and of the heating it produces: the simplifications we made allow to explore other properties than those accessible to classical direct numerical simulations (DNS).  While it is not meant to replace and cannot replace DNS, for example because of the lack of three-dimensional information on field-line topologies, it partially fills the huge gap between the Reynolds numbers in DNS and in the real corona: although Reynolds numbers reached in the model, of the order of $10^6$, are still lower than the ones expected in the corona, this already represents an outstanding progress compared to DNS. Furthermore this allows to explore regimes of MHD turbulence that are not accessible to DNS, for instance this allows for intermittency to appear in turbulence while having a complete description of the spectra and of the non-linear interactions in the perpendicular direction.  On the other hand, having even higher Reynolds numbers, and thus smallest scales, may require to take into account non-MHD effects, like kinetic effects, that can still be neglected in this model.   The model has been used in this paper in the case of a magnetic loop in the solar corona, in which the physical properties of the medium (namely the external longitudinal magnetic field \\Bo\\ and the mass density) are assumed to be uniform along the loop.  In this case, and thanks to the above-mentioned characteristics of the model, we could show that this model loop displays: a dynamics over a very wide range of spatial and temporal scales (4 to 5 orders of magnitude); spectra which are formed by a local cross-scale energy flux, and which have a wide inertial range in either direction, perpendicular or parallel to the external magnetic field; a clear anisotropy between the parallel and perpendicular spectra, which could be compatible with the ``critical balance'' phenomenology; a scaling of the average ratio of the magnetic energy over the kinetic energy consistent with RMHD; a heating function with multiple spatial and temporal scales; a flat longitudinal profile of the average dissipation power (although this may be dismissed by further simulations, with non-uniform physical properties of the medium along the loop); a spiky, and statistically intermittent time series of energy dissipation power; power-law distributions of the characteristics (peak energy, total energy, duration, waiting-times) of ``events'' extracted from the time series of dissipation power; finite-lifetime packets of resonant frequencies in the time series of energy, but whose frequencies are shifted from the harmonics of the linear resonant frequencies because of the non-linearities; long-range time correlations in time series; a delay of the dissipation time series compared to some energy time series; an average dissipation power that scales with the loop parameters, and that could be sufficient to sustain the high coronal temperatures.  As discussed in the text, some of these results confirm or complete the results of a similar model \\citep{nig04,nig05}.    Further directions for the study of the solar corona using this model include taking advantage of the possibility of modelling non-uniform regions to (1) allow for density gradients in a coronal loop, to seek for the preferred locations of coronal heating, and (2) study a magnetically open region such as a coronal hole.  In order to get diagnostics that can be compared to observations, this heating model can be coupled to the thermodynamics of a loop (including the cooling by conduction and radiation) upon which forward-modelling of coronal spectral lines may be carried out. We also believe that this model can be used in other heliospheric and astrophysical applications where MHD applies and where there is a strong dominant magnetic field (see Sect.~\\ref{sec:avail} in Appendix for code availability)."
    },
    "0606/astro-ph0606505_arXiv.txt": {
        "abstract": "We study the impact of setting initial conditions in numerical simulations using the standard procedure based on the Zel'dovich approximation (ZA). As it is well known from perturbation theory, ZA initial conditions have incorrect second and higher-order growth and therefore excite long-lived transients in the evolution of the statistical properties of density and velocity fields. We also study the improvement brought by using more accurate initial conditions based on second-order Lagrangian perturbation theory (2LPT). We show that 2LPT initial conditions reduce transients significantly and thus are much more appropriate for numerical simulations devoted to precision cosmology. Using controlled numerical experiments with ZA and 2LPT initial conditions we show that simulations started at redshift $z_i=49$ using the ZA underestimate the power spectrum  in the nonlinear regime by about $2,4,8\\%$ at $z=0,1,3$ respectively, whereas the mass function of dark matter halos is underestimated by $5\\%$ at $m=10^{15}M_\\odot/h$ ($z=0$) and $10\\%$ at $m=2\\times 10^{14}M_\\odot/h$ ($z=1$). The clustering of halos is also affected to the few percent level at $z=0$. These systematics effects are typically larger than statistical uncertainties in recent mass function and power spectrum fitting formulae extracted from numerical simulations. At large scales, the measured transients in higher-order correlations can be understood from first principle calculations based on perturbation theory. ",
        "introduction": "The growth of structure in the universe is a fundamental aspect of cosmology, and can be used to constrain the nature of primordial fluctuations and the energy contents of the universe (e.g. dark matter and dark energy). Numerical simulations are a very important tool in making theoretical predictions of cosmological observables  at small scales when fluctuations are nonlinear. Fluctuations are followed from their primordial seeds  through the radiation era to the matter dominated era and current accelerating phase by using linear Boltzmann equation solvers \\citep{PeYu70,BE84,MaBe95,CMBFAST,CAMB}, which are  then coupled to N-body simulations to follow their nonlinear evolution. This coupling takes the form of final conditions predicted from linear Boltzmann codes input as initial conditions into N-body simulations at some redshift $z_i$.  There are competing requirements on the best choice of $z_i$. One would like to start a simulation as early as possible, where perturbations are closest to their linear value but this can be challenging from numerical reasons since one is trying to impose a small perturbation spectrum that can be overwhelmed by sources of noise such as numerical roundoff error or shot noise of the particles. Also, when running simulations that include a baryonic component, the initial conditions simplify considerably the lower $z_i$ is, in which case the baryons had more time to fall into the potential wells of the dark matter.  Despite the rapid progress in the development of different numerical N-body techniques for solving the nonlinear growth of structure in the universe (see e.g. \\citet{Bert98} for a review), the way initial conditions  are input into these codes has essentially been the same since the first developments in cosmological simulations over twenty years ago \\citep{KlSh83,EDFW85}, by using the \\citet{ZA70} approximation (hereafter ZA).  As observations of the late universe quickly approach percent-level accuracy, it is vital to examine to what extent theoretical predictions can accurately go from primordial fluctuations to different stages of nonlinear evolution in recent epochs, e.g. redshifts $z \\la 3$. Studies of linear Boltzmann solvers \\citep{CMBcodes} have shown that independent codes agree to within 0.1\\%, whereas for N-body codes the agreement, while much more difficult to quantify, is of the order of few to ten percent depending on the statistics and scales considered \\citep{SMSS98,SBcluster,HRWH05}.  The link between Boltzmann and N-body codes, or how to set initial conditions for nonlinear evolution, has been surprisingly little studied. The accuracy of the ZA has been essentially taken for granted and instead studies of initial conditions have concentrated on other (important) issues, related to how to best implement the ZA in different situations \\citep{Bert95,Salmon96,Pen97,Bert01,YSH03,Sirko05}.  The rationale for this state of affairs seems to be that starting a simulation at  ``high enough\" redshift with linear perturbation theory (which the ZA is, in Lagrangian space) should be safe. A rule of thumb for the choice of starting redshift often used is that the {\\em rms} density fluctuations at the interparticle distance scale be of order $0.2$. As a result, low resolution simulations are typically started much later than high resolution ones.  However, one of the lessons of perturbation theory (hereafter PT, see \\cite{PTreview} for a review) is that the ZA has a very limited validity, reproducing only the correct linear growth and underestimating the skewness and higher-order moments \\citep{GW87,JBC93,B94,CM94,J95,BGE95,CBH96}, with deviations being particularly severe for velocity field perturbations because of the ZA failure to conserve momentum.  As a result of this, starting N-body simulations from ZA initial conditions leads to incorrect second and higher-order growing modes; this excites nonlinear decaying modes (called {\\em transients}) that describe how long it takes for the correct dynamics to recover the actual statistical properties of density and velocity fields  \\citep{Sco98}. This transient evolution is well understood at large scales using PT and describes  the anomalous behavior seen in measurements of higher-order moments in N-body simulations started from ZA initial conditions (as shown in \\citet{Sco98} and demonstrated in much more detail below).  The importance of transients is not just of interest for studies involving higher-order moments, as we show in this paper. Underestimation of  non-Gaussianity means that the tails of the probability distribution of densities and velocities will be artificially suppressed, which implies that rare events will be even more rare in a simulation that has transients. Since typically large volumes (and thus typically low resolution) are used to study rare events, the dark matter halo mass function at the high-mass end will be particularly affected by transients given that low-resolution simulations are often started later than high-resolution ones.   In addition, the transient behavior of large-scale higher-order moments implies that nonlinear couplings take a significant time to develop their full strength. These are the same couplings that act at small scales and give rise to the nonlinear evolution of the power spectrum, therefore transients should also affect the power spectrum at nonlinear scales.   In this paper we revisit the issue of transients from ZA initial conditions and study numerically the effects on higher-order moments, the mass function of dark matter halos and the power spectrum. As proposed in \\citet{Sco98}, transients can be reduced significantly by using second-order Lagrangian PT (hereafter 2LPT). This leads to correct second-order growth factors and greatly improves the behavior of higher-order moments. Furthermore, transients in this case decay as inverse square power of the scale factor, much faster than the inverse of the scale factor in the ZA case. We study the improvement brought by using 2LPT instead of ZA in the behavior of transients and quantify their magnitude in different statistics commonly used.  The paper is organized as follows. In section~\\ref{basics} we review how transients arise and present predictions that are contrasted with simulations later. Section~\\ref{simu} discusses the numerical simulations we run to quantify the transients in the nonlinear regime. In section~\\ref{Res} we present the main results of the paper on the magnitude of transients in higher-order cumulants, power spectrum, bispectrum, mass function and bias of dark matter halos. We summarize the results and conclude in section~\\ref{Conc}. In appendix~\\ref{sec2LPT} we discuss second-order Lagrangian perturbation theory and in appendix~\\ref{secGlass} we derive analytically the evolution of perturbations toward a glass configuration and discuss glass initial conditions from the point of view of transients. ",
        "conclusions": "\\label{Conc}   In this paper we considered the use of different approximations to set up initial conditions for cosmological simulations, concentrating on typical initial redshifts $z_i \\simeq 50$, and studied their influence on low-redshift ($z \\la 3$) statistics. The standard procedure in the literature is to use the \\citet{ZA70} approximation (ZA), a linear solution in Lagrangian space, or even linear perturbation theory in Eulerian space (typically when simulating baryons using methods which are not based on particles).  At high redshifts ($z \\sim 50$) the naive expectation is that linear perturbation theories should be accurate enough, but they are not. The reason for this is that at high redshifts nonlinear corrections are relatively much more important than at low redshift  because the relevant scales involved have a very steep spectrum (with the effective spectral index very close to $n_{\\rm eff}=-3$). An example of this is shown in Fig.~\\ref{PDFT}, where the density PDF smoothed on $10 \\Mpc$ at $z=3$ shows strong non-Gaussian features, even though naively these scales are in the linear regime (since the linear {\\em rms} density fluctuation is only $\\sigma =0.25$). Modeling these nonlinear corrections using the ZA leads to a significant underestimate of nonlinear couplings and thus non-Gaussianity. This leads to a suppression of the density PDF tail for large densities that translates into a suppression of the high-mass tail of the mass function. The weaker nonlinear couplings lead to a smaller power spectrum in the nonlinear regime.  The mismatch between the approximation used to set initial conditions and the correct dynamics leads to {\\em transients}, excitation of decaying modes that are long-lived and survive until low redshift, systematically biasing low the nonlinear couplings even after evolving with the correct dynamics of an N-body code for as much as 50-100 scale factors. Transients cannot be avoided, but they can be significantly reduced by imposing initial conditions with better approximations to the nonlinear equations of motion.   More specifically, in this paper we have shown that,  \\begin{enumerate}   \\item The measurements of statistics at low redshift for simulations started with ZA initial conditions depend significantly on the starting redshift $z_i$, signaling the presence of transients.  \\item This can be avoided to a large extent by using  2LPT initial conditions, which as we have demonstrated show very little dependence on $z_i$ even for extremely late starts such as $z_i=11.5$. Given that such runs show at most 1\\% transient effects and that transients in 2LPT decay as inverse square of the scale factor, we can estimate that our reference runs with $z_i=49$ are accurate (in terms of transients) to better than a tenth of a percent for $z\\la 3$.  \\item The simplest way to quantify transients is to follow how non-Gaussianity develops at large scales as nonlinear couplings reach their full strength during the transient evolution. Our results for higher-order cumulants at large scales confirm previous PT results \\citep{Sco98,Val02} with much better accuracy than done before.  \\item The  small scale behavior of transients is consistent with the reduction of nonlinear couplings observed at large scales, and we have checked it is not affected by the choice of softening length, force and time integration accuracy.  \\item The use of the ZA at $z_i=49$ leads to suppressions in the power spectrum of order a few ($z=0$) to almost ten percent ($z=3$) in the nonlinear regime. These results are in reasonable agreement with an estimation of transients in the power spectrum based on one-loop PT.  \\item  ZA initial conditions also lead  to a suppression in the high-mass tail of the mass function of dark matter halos. For $z_i=24$ this suppression is about 10\\% for $m=10^{15}h^{-1}M_\\odot$ at $z=0$ and $m=10^{14}h^{-1}M_\\odot$ at $z=1$. The clustering of halos is also affected at the few-percent level.  \\end{enumerate}   An obvious alternative to using 2LPT initial conditions is to start with the ZA much earlier than $z_i\\simeq 50$. However, in order to compete with the improvement brought by using 2LPT very high redshifts are required. For example, Fig.~\\ref{PkReal} shows that a $z_i=11.5$ 2LPT start has approximately a 0.5\\% transient effect on the nonlinear power spectrum at $z=0$, whereas a $z_i=49$ ZA start has a 2.5\\% effect. In order for the ZA to reduce the transients to the level of the 2LPT $z_i=11.5$ start, a $z_i=249$ ZA start would be required. Reaching the level of a 2LPT $z_i=49$ start, corresponding to our reference runs, will require about another factor of 16 in redshift (given that transients scale as $a^{-2}$ in 2LPT compared to $a^{-1}$ in the ZA). Other issues make starting simulations at very high redshifts a bit more challenging, e.g.  depending on the N-body code algorithm, numerical errors may be more difficult to avoid at very high redshifts where perturbations are so small. In addition, initial conditions for baryons at such high redshifts must be carefully included as they are still falling into dark matter potential wells.  Although we have not studied this in any detail, simulations started with linear theory instead of the ZA are expected to be much more significantly affected by transients, e.g. a simulation started with linear theory at $z_i=49$ should be roughly equivalent to one started with the ZA at $z=11.5$ (see Fig.~\\ref{S3ZALT}). Note that linear theory initial conditions are often used in codes that rely on grid methods with gas dynamics solvers. We expect such methods to show differences with e.g. SPH methods that use ZA initial conditons. When the dynamical variables are densities and velocities in a mesh (instead of information carried by particles), a way to reduce transients would be to use higher-order Eulerian perturbation theory.  While in this paper we have concentrated on statistics of the dark matter density perturbations, it would be interesting to know how much do transients affect other high-redshift probes, where transients are most likely to play an important role. In this regard, the most pressing aspect  would be to check the impact of transients in the Lyman-$\\alpha$ forest power spectrum, because of the high accuracy demanded from simulations to obtain cosmological constraints from current data sets \\citep{VHL06,SSM06}.  The use of better initial conditions should be a useful tool in reaching  future goals of simulating the nonlinear power spectrum to an accuracy of better than 1\\% percent  needed for the next-generation of weak gravitational lensing surveys \\citep{HuTa05}.  The code we used to generate ZA \\& 2LPT initial conditions in this paper is publically available at {\\tt http://cosmo.nyu.edu/roman/2LPT}."
    },
    "0606/astro-ph0606733_arXiv.txt": {
        "abstract": "The Darwin and TPF-I missions are Infrared free flying interferometer missions based on nulling interferometry. Their main objective is to detect and characterize other Earth-like planets, analyze the composition of their atmospheres and their capability to sustain life, as we know it. Darwin and TPF-I are currently in definition phase. A number of mission architectures of 3 and 4 free flying telescopes are evaluated on the basis of the interferometer's response, ability to distinguish multiple planet signatures and starlight rejection capabilities. The characteristics of the new configurations are compared also to the former, more complex Bowtie baseline architectures as well as evaluated on base of their science capability. ",
        "introduction": "The closing years of the 20th century have allowed us, for the first time, to seriously discuss interferometric instruments deployed in space achieving unprecedented spatial resolution. And thus the direct detection of Earth-like exoplanets orbiting nearby stars, and the search for key tracers of life in their atmospheres, are high-priority objectives in the long-term science plan of ESA as well as NASA. Both Agencies are working on the definition of the instruments that will meet this challenge, Darwin and Terrestrial Planet Finder (TPF) respectively with a foreseen launch in 2015+. The baseline mission duration is 5 years, extendable to 10 years in an L2 orbit. Operating in the infrared band requires that all optical components are cooled to roughly 40 K, this is achieved by passive cooling. Only the detector requires active cooling. Darwin is a major element in the Cosmic Vision 2020 program of the European Space Agency. It has the explicit purpose of detecting other Earth-like worlds, analyze their characteristics, determine the composition of their atmospheres and investigate their capability to sustain life as we know it. The Darwin mission is envisioned as four free flying spacecraft including one beam-combining spacecraft. The beam combiner and the telescope spacecraft fly in one plane with each telescope spacecraft at the same distance from the beam combiner. The resolution of the interferometer is adjusted by changing the distance between the telescope spacecrafts. A similar activity has been taking place in the United States within the context of NASA's Origins program. A science collaboration has already been established. The missions are currently in definition phase, here some characteristics of the designs are reviewed. We present the different new configuration architectures under investigation. ",
        "conclusions": "\\label{sec:concl} A number of mission architectures of 3 and 4 free-flying telescopes are evaluated on the basis of the interferometer's response, modulation efficiency, ability to distinguish multiple planet signatures and starlight rejection capabilities. Figure~\\ref{fig1} shows the characteristics of the alternative mission architectures. Even though the starlight rejection properties of the Bowtie configuration is outstanding in the comparison, the higher mean modulation efficiency of the TTN configuration and the X-array counteracts that superiority. An additional factor in mission design and complexity is the number of telescopes used. Here the Bowtie has a big disadvantage, as it needs 6 telescopes. Figure~\\ref{fig1} shows that it is essential to apply spectral channels to distinguish the signal from multiple planets due to their distinct signature as a function of distance from their star, wavelength and rotation of the array (shown here, Venus, Earth, Mars and Jupiter). Here the different architectures show their capability to disentangle multiple planetary signals. As seen in Figure~\\ref{fig1} the rectangular TTN and the X-array perform the best in this respect. Recent simulations by Velusamy et al. and Thiebaut et al (these proceedings) argue that high spatial resolution is not needed to disentangle multiple planets which leads to a similar performance of the three and four telescope architectures."
    },
    "0606/astro-ph0606219_arXiv.txt": {
        "abstract": "{} {The large majority of BL Lacertae objects belonging to the 1 Jy sample, the class prototype for radio-selected sources, are thought to emit most of their synchrotron power in the far infrared band. Ironically, this spectral region is very sparsely sampled, with only a minority of the objects having IRAS data (most of them being upper limits or low-quality detections). We aim at filling this infrared gap by presenting new, simultaneous ISOCAM and ISOPHOT observations over the $7 - 200\\mu$m range (observer's frame) for half the sample. A precise measurement of the position of the synchrotron peak frequency, $\\nu_{\\rm peak}$, can provide direct information about particle acceleration mechanisms and constrain the inverse Compton radiation that will be detected by up-coming new $\\gamma$-ray missions.} {We have observed seventeen 1 Jy BL Lacertae objects with the camera and the photometer on board the Infrared Space Observatory (ISO) satellite. Given the intrinsic variability of these sources, the data were taken by concatenating the pointings to ensure simultaneity. The ISOPHOT data reduction was done employing a novel correction, which mitigates the effect of chopping for faint sources.} {Using our new ISO data, complemented by nearly-simultaneous radio and optical observations for ten and four objects respectively, and other multi-frequency data, we have built the spectral energy distributions of our sources (plus a previously published one) and derived the rest-frame $\\nu_{\\rm peak}$. Its distribution is centered at $\\sim 10^{13}$ Hz ($\\sim 30\\mu$m) and is very narrow, with $\\sim 60\\%$ of the BL Lacs in the $1 - 3 \\times 10^{13}$ Hz range. Given our set of simultaneous infrared data, these represent the best determinations available of the synchrotron peak frequencies for low-energy peaked BL Lacs. A comparison with previous such estimates, based on non-simultaneous optical and near infrared data, may indicate strong $\\nu_{\\rm peak}$ variations in a number of sources, possibly associated with large flares as observed in the high-energy peaked BL Lac MKN 501.} {} ",
        "introduction": "BL Lacertae objects constitute one of the most extreme classes of active galactic nuclei (AGN), distinguished by their high luminosity, rapid variability, high ($> 3\\%$) optical polarization, radio core-dominance, apparent superluminal speeds, and almost complete lack of emission lines (e.g., Urry \\& Padovani \\cite{UP95}). The broad-band emission in these objects, which extends from the radio to the gamma-ray band, appears to be dominated by non-thermal processes from the heart of the AGN, undiluted by the thermal emission present in other AGN. Therefore, BL Lacs represent the ideal class to study to further our understanding of non-thermal emission in AGN.  Synchrotron emission combined with inverse Compton scattering is generally thought to be the mechanism responsible for the production of radiation over such a wide energy range (e.g., Ghisellini et al. \\cite{Ghisellini98}). The frequency at which most of the synchrotron power is emitted, $\\nu_{\\rm peak}$, ranges across several orders of magnitude, going from the far-infrared to the hard X-ray band. Sources at the extremes of this wide distribution are referred to as low-energy peaked (LBL) and high-energy peaked (HBL) BL Lacs, respectively (Giommi \\& Padovani \\cite{Giommi94}; Padovani \\& Giommi \\cite{Padovani95}). Radio-selected samples include mostly objects of the LBL type, while X-ray selected samples are mostly made up of HBL.  \\begin{table*}% \\caption{Sample properties.\\label{prop}} \\begin{center} \\begin{tabular}{lrrllrl} \\hline Name & RA(J2000) & Dec(J2000)&~~~$z$ &$\\langle R_{\\rm mag}\\rangle$& F$_{\\rm 5GHz}$ \\\\ &         &        &    &   &Jy \\\\ \\hline {PKS 0048$-$097} & 00 50 41.3 & $-$09 29 06 &$>$0.2 & 16.5 & 2.0 \\\\ {PKS 0118$-$272} & 01 20 31.7 & $-$27 01 24 &$>$0.559 & 17.0 & 1.1 \\\\ {PKS 0138$-$097} & 01 41 25.8 & $-$09 28 44 &\\magg0.733 & 17.0 & 1.2 \\\\ {PKS 0426$-$380} & 04 28 40.4 & $-$37 56 19 &\\magg1.110 & 18.0 & 1.2 \\\\ {S5 0454+844}    & 05 08 42.4 & $+$84 32 05 &$>$1.340 & 19.0 & 1.4 \\\\ {PKS 0537$-$441} & 05 38 50.4 & $-$44 05 09 &\\magg0.896& 13.5 & 4.0 \\\\ {PKS 0735+178}   & 07 38 07.1 & $+$17 42 27 &$>$0.424 & 14.5 & 2.0 \\\\ {PKS 1144$-$379} & 11 47 01.4 & $-$38 12 10 &\\magg1.048 & 16.5 & 1.6 \\\\ {OQ 530}         & 14 19 46.6 & $+$54 23 15 &\\magg0.152 & 14.5 & 1.1 \\\\ {PKS 1519$-$273} & 15 22 37.7 & $-$27 30 10 &$>$0.2 & 18.5 & 2.4 \\\\ {PKS 1749+701}   & 17 48 32.8 & $+$70 05 51 &\\magg0.770 & 16.5 & 1.1 \\\\ {PKS 1749+096}   & 17 51 32.8 & $+$09 39 01 &\\magg0.320 & 15.5 & 1.9 \\\\ {S5 1803$+$784}  & 18 00 45.7 & $+$78 28 04 &\\magg0.684 & 16.5 & 2.6 \\\\ {4C 56.27}       & 18 24 07.2 & $+$56 51 00 &\\magg0.664 & 18.5 & 1.7 \\\\ {PKS 2131$-$021} & 21 34 10.3 & $-$01 53 17 &\\magg1.285 & 20.0 & 2.1 \\\\ {PKS 2240$-$260} & 22 43 26.4 & $-$25 44 31 &\\magg0.774 & 17.5 & 1.0 \\\\ {PKS 2254+074}   & 22 57 17.3 & $+$07 43 12 &\\magg0.190 & 17.0 & 1.2 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  It is ironic that the spectrum of LBL in the infrared (IR) band, where these sources' output is thought to be largest, is very sparsely known (e.g., Sambruna et al. \\cite{Sambruna96}; Giommi et al. \\cite{Giommi02}; Massaro et al. \\cite{Massaro05}). IR data for the 1 Jy, the prototype radio-selected BL Lac sample are, for example, very sparse. The purpose of this paper is to remedy this situation by presenting ISOCAM and ISOPHOT observations of 1 Jy BL Lacs covering the $6.7 - 200\\mu$m range. Since $\\nu_{\\rm peak} \\propto \\gamma^2_{\\rm peak} \\delta B$, where $\\gamma_{\\rm peak}$ is the Lorentz factor of the electrons emitting most of the radiation, $\\delta$ is the Doppler factor, and $B$ is the strength of the magnetic field, its precise measurement can provide direct information about particle acceleration mechanisms. The position of $\\nu_{\\rm peak}$ is also important for the up-coming Astro-rivelatore Gamma a Immagini Leggero (AGILE\\footnote{\\tt http://agile.rm.iasf.cnr.it/}) and Gamma ray Large Area Space Telescope (GLAST\\footnote{\\tt http://glast.gsfc.nasa.gov/}) missions, as they will sample the $\\gamma$-ray band close to the peak of the inverse Compton emission from LBL and flat-spectrum radio quasars, whose location depends on that of the synchrotron one.  In Sect. 2 we present our sample, Sect. 3 discusses the Infrared Space Observatory (ISO) observations, data analysis, and results, while in Sect. 4 we derive the spectral energy distributions and $\\nu_{\\rm peak}$ values. Finally, Sect. 5 summarizes our conclusions. Throughout this paper spectral indices are written $S_{\\nu} \\propto \\nu^{-\\alpha}$. ",
        "conclusions": "We have presented new, simultaneous ISOCAM and ISOPHOT observations covering the $7 - 200\\mu$m range for seventeen (out of 34) BL Lacertae objects belonging to the 1 Jy sample, the prototype radio-selected (low-energy peaked) sample. These data open up a new window on the spectral energy distributions of low-energy peaked BL Lacs and fill a glaring gap in their far-infrared spectra. They also represent a sevenfold (eighteenfold) increase in the number of published ISO results for BL Lacs (LBL).  Our main results can be summarized as follows:  \\begin{enumerate}  \\item Our detection rate is good: 100\\%, 82\\%, and 94\\% at 6.7 (LW2), 9.6 (LW7), and 14.3 (LW3) $\\mu$m (ISOCAM), and 94\\%, $\\sim 65\\%$, and $\\sim 20\\%$ at 25 (P2), $65 - 100$ (C100), and $120 - 180$ (C200) $\\mu$m (ISOPHOT) respectively. To deal with the effect of chopping on our relatively faint sources, which would lead to a flux overestimate at 25 $\\mu$m, we have also employed a novel correction in the ISOPHOT data reduction;  \\item We have built the spectral energy distributions of our sources, with the addition of a previously published one, by using our new ISO data, complemented by nearly-simultaneous radio data for about half the sample (and optical data for four sources), and other non-simultaneous multi-frequency data. We have then derived the synchrotron peak frequency, $\\nu_{\\rm peak}$. This turned out to have a very narrow distribution, with $\\sim 60\\%$ of the sources having values in the range $1 - 3 \\times 10^{13}$ Hz, centered at $\\sim 10^{13}$ Hz ($\\sim 30\\mu$m). These represent the best determinations available of the synchrotron peak frequencies for low-energy peaked BL Lacs;  \\item A comparison with previous such estimates, based on non-simultaneous optical and near infrared data, may indicate strong $\\nu_{\\rm peak}$ variations in a number of sources, possibly associated with large flares as observed in the high-energy peaked BL Lac MKN 501.  \\end{enumerate}  These results, apart from characterizing synchrotron emission in blazars, can constrain particle acceleration mechanisms and are also important for the planning and data interpretation of up-coming new $\\gamma$-ray missions."
    },
    "0606/astro-ph0606443_arXiv.txt": {
        "abstract": "We study the role of massive perturbers (MPs) in deflecting stars and binaries to almost radial ({}``loss-cone'') orbits, where they pass near the central massive black hole (MBH), interact with it at periapse, and are ultimately destroyed. MPs dominate dynamical relaxation when the ratio of the 2nd moments of the MP and star mass distributions, $\\mu_{2}\\!\\equiv\\!\\left.N_{p}\\left\\langle M_{p}^{2}\\right\\rangle \\right/N_{\\star}\\left\\langle M_{\\star}^{2}\\right\\rangle $, satisfies $\\mu_{2}\\!\\gg\\!1$. We compile the MP mass function from published observations, and show that MPs in the nucleus of the Galaxy (mainly giant molecular clouds), and plausibly in late type galaxies generally, have $10^{2}\\!\\lesssim\\!\\mu_{2}\\!\\lesssim\\!10^{8}$. MPs thus shorten the relaxation timescale by $10^{2-7}$ relative to 2-body relaxation by stars alone. We show this increases by $10^{1-3}$ the rate of \\emph{large}-periapse interactions with the MBH, where loss-cone refilling by stellar 2-body relaxation is inefficient. We extend the Fokker-Planck loss-cone formalism to approximately account for relaxation by rare encounters with MPs. We show that binary stars--MBH exchanges driven by MPs can explain the origin of the young main sequence B stars that are observed very near the Galactic MBH, and increase by orders of magnitude the ejection rate of hyper-velocity stars. In contrast, the rate of \\emph{small}-periapse interactions of single stars with the MBH, such as tidal disruption, is only increased by a factor of a few. We suggest that MP-driven relaxation plays an important role in the 3-body exchange capture of single stars on very tight orbits around the MBH. These captured stars may later be disrupted by the MBH via tidal orbital decay or direct scattering into the loss cone; captured compact objects may inspiral into the MBH by the emission of gravitational waves from zero-eccentricity orbits. ",
        "introduction": "\\label{s:intro}  There is compelling evidence that massive black holes (MBHs) lie in the centers of all galaxies \\citep{fer+00,geb+03,shi+03}, including in the center of our Galaxy \\citep{eis+05,ghe+05}. The MBH affects the dynamics and evolution of the galaxy's center as a whole (e.g. \\citealt{bah+76}) and it also strongly affects individual stars or binaries that approach it. Such close encounters, which may be extremely energetic, or involve non-gravitational interactions, or post-Newtonian effects, have been the focus of many studies (see review by \\citealt{ale05}). These processes include the destruction of stars by the MBH, either by falling whole through the event horizon, or by being first tidally disrupted and then accreted (e.g. \\citealt{ree88}); tidal scattering of stars on the MBH \\citep{ale+01b}; the capture and gradual inspiral of stars into the MBH, accompanied by the emission of gravitational waves or by tidal heating (e.g. \\citealt{ale+03b,ale+03a}); or dynamical exchange interactions in which incoming stars or binaries energetically eject a star tightly bound to the MBH and are captured in its place very near the MBH (e.g. \\citealt{ale+04,gou+03}).  The interest in such processes is driven by their possible implications for the growth of MBHs, for the orbital decay of a MBH binary, for the detection of MBHs, for gravitational wave (GW) astronomy, as well as by observations of unusual stellar phenomena in our Galaxy, e.g. the puzzling young population of B-star very near the Galactic MBH \\citep{eis+05}, or the hyper-velocity B stars at the edge of the Galaxy \\citep{bro+05,fue+06,bro+06a}, possibly ejected by 3-body interactions of a binaries with the MBH \\citep{hil91}.  Here we focus on close encounters with the MBH whose ultimate outcome ({}``event'') is the elimination of the incoming object from the system, whether on the short infall (dynamical) time, if the event is prompt (e.g. tidal disruption or 3-body exchange between a binary and the MBH), or on the longer inspiral time, if the event progresses via orbital decay (e.g. GW emission or tidal capture and heating). Such processes are effective only when the incoming object follows an almost zero angular momentum ({}``loss-cone'') orbit with periapse closer to the MBH than some small distance $q$. To reach the MBH, or to decay to a short period orbit, both the infall and inspiral times must be much shorter than the system's relaxation time $t_{r}$ \\citep{ale+03b}. The fraction of stars initially on loss-cone orbits is very small and they are rapidly eliminated. Subsequently, the close encounter event rate is set by the dynamical processes that refill the loss-cone.  The loss-cone formalism used for estimating the event rate (\\citealt{fra+76,lig+77,coh+78}) usually assumes that the system is isolated and that the refilling process is 2-body relaxation. This typically leads to a low event rate, set by the long 2-body relaxation time.  Two-body relaxation, which is inherent to stellar systems, ensures a minimal loss-cone refilling rate. Other, more efficient but less general refilling mechanisms were also studied with the aim of explaining various open questions (e.g. the stalling problem of MBH binary coalescence, \\citet{ber+05,mer+05,ber+06}, or MBH feeding, \\citealt{zha+02,mir+05}) or in the hope that they may lead to significantly higher event rates for close encounter processes. These mechanisms include chaotic orbits in triaxial potentials \\citep{nor+83,ger+85,mer+04b,hol+06} (the presence of a MBH may however destroy the triaxiality near the center; \\citealt{mer+98,hol+02,sel02}); increased fraction of low angular momentum orbits in non-spherical potentials \\citep{mag+99,ber+06}; accelerated resonant relaxation of angular momentum near the MBH where the orbits are Keplerian \\citep{rau+96,rau+98,hop+06a,lev06}; perturbations by a massive accretion disk or an intermediate mass black hole (IMBH) companion \\citep{pol+94,zha+02,lev+05}. Most of these mechanisms require special circumstances to work (e.g. specific asymmetries in the potential), or are short-lived (e.g. the IMBH will eventually coalesce with the MBH).  Here we explore another possibility, which is more likely to apply generally: accelerated relaxation and enhanced rates of close encounters driven by massive perturbers (MPs). Efficient relaxation by MPs was first suggested by \\citet{spi+51,spi+53} to explain stellar velocities in the galactic disk. MPs remain an important component in modern models of galactic disk heating (see e.g. \\citealt{vil83,vil85,lac84,jen+90,han+02} and references therein). A similar mechanism was suggested to explain the spatial diffusion of stars in the inner Galactic bulge \\citep{kim+01}. In addition to dynamical heating, efficient relaxation by MPs was suggested as a mechanism for loss cone replenishment and relaxation, both in the context of scattering of Oort cloud comets to the Sun \\citep{hil81,bai83} and the scattering of stars to a MBH in a galactic nucleus \\citep{zha+02}. \\citet{zha+02} suggested MPs as a mechanism for establishing the $M_{\\bullet}/\\sigma$ relation \\citep{fer+00,geb+00} by fast accretion of stars and dark matter. They also noted the possibility of increased tidal disruption flares and accelerated MBH binary coalescence due to MPs. In this study we investigate in detail the dynamical implications of loss-cone refilling by MPs. We evaluate its effects on the different modes of close interactions with the MBH, in particular 3-body exchanges, which were not considered by \\citet{zha+02}, and apply our results to the Galactic Center (GC), where observations indicate that dynamical relaxation is very likely dominated by MPs.  This paper is organized as follows. In \\S\\ref{s:outline} we present the main concepts and procedures of our calculations. The observational data and theoretical predictions about MPs in the inner $\\sim\\!100$ pc of the GC are reviewed in \\S\\ref{s:GC_MPs}. In section \\S\\ref{s:interact} we explore the implications of relaxation by MPs for various types of interactions with the MBH. We summarize our results in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  Relaxation by MPs dominates relaxation by 2-body stellar interactions when the ratio between the 2nd moments of their respective mass functions satisfies $\\mu_{2}\\!=\\!\\left.N_{p}\\left\\langle M_{p}^{2}\\right\\rangle \\right/N_{\\star}\\left\\langle M_{\\star}^{2}\\right\\rangle \\!\\gg\\!1$. We show that Galactic MPs (stellar clusters, GMCs and smaller molecular gas clumps that exist outside the inner few pc) dominate and accelerate relaxation in the inner $\\sim\\!100$ pc of the GC. This is plausibly the case in the centers of late-type galaxies in general. There is also evidence for molecular gas in the centers of late-type galaxies (e.g. \\citealt{rup97,kna99}), which suggests that MPs may dominate relaxation there as well, and lead to the relaxation of the central regions of galactic bulges in general.  Relaxation determines the rate at which stars and binaries are deflected to near radial (loss-cone) orbits that bring them closer to the MBH than some critical periapse $q$, where they undergo a strong destructive interaction with it. The size of $q$ depends on the nature of the interaction of interest (e.g. tidal disruption, 3-body exchange). It is much larger for binaries than for single stars due to the binary's larger effective size.  We extend loss-cone theory to approximately treat rare encounters with MPs, and apply it to explore the implications of MPs on the rates of different types of close encounters. The rate reaches its maximum when loss-cone orbits are replenished by scattering within an orbital time (the full loss-cone regime). This is more easily achieved when the phase-space volume of the loss-cone is small, that is, when $q$ is small. MPs thus affect only those processes with large $q$ whose loss-cone is too large to be efficiently replenished by stellar encounters (the empty loss-cone regime).  We show that MPs will not contribute much to the disruption of single stars in the GC, whose loss-cone is efficiently replenished by stars outside the central $\\sim\\!2$ pc (MPs may accelerate the consumption of stars by more massive MBHs, where $q$ is significantly larger, or the capture of stars in accretion disks). However, MPs will enhance by factors of 10--1000 the tidal disruption rate of infalling binaries, which result in the capture of one of the stars on a tight orbit around the MBH, and the ejection at high velocity of the other star \\citep{hil91,hil92,yuq+03}. The enhancement of the event rates is dominated by the innermost MPs, and so the uncertainty in the MP distribution on the smallest scales dominates the uncertainties in the total event rate. Detailed observations of MPs in the inner GC allow us to robustly predict their effects in the Galaxy. We show that MP-induced disruptions of relatively rare massive binaries can naturally explain the puzzling presence of normal-appearing main sequence B stars in the central $0.04$ pc of the GC \\citep{eis+05}, and at the same time can account for the observed HVSs well on their way out of the Galaxy \\citep{bro+05,ede+05,hir+05,bro+06a,bro+06b,bro+06c}. Tidal disruptions of the many more faint low-mass binaries can efficiently supply single stars on very eccentric tight orbits near the MBH. Such an increase in the numbers of stars in tight orbit near the MBH may increase the rates of single star processes such as tidal disruption and heating or GW emission from tightly bound compact objects \\citep{mil+05}.  Finally, MP-induced interactions also have cosmological implications for the coalescence of binary MBHs following galactic mergers \\citep{zha+02}. We suggest that MPs can accelerate the dynamical decay of binary MBHs by efficiently supplying stars for the slingshot mechanism, and thereby help solve the {}``last parsec'' stalling problem. MP-driven loss-cone refilling will operate even in the case of a spherical potential, where other suggested mechanisms are inefficient, thus allowing MBHs to coalesce on the dynamical timescale of the galactic merger. A detailed treatment of this idea will be presented elsewhere (Perets \\& Alexander, in prep.)."
    },
    "0606/astro-ph0606396_arXiv.txt": {
        "abstract": "We report the abundance analysis of one red giant branch star in the metal-poor outer halo globular cluster NGC~5694. We obtain [Fe/H] = $-$1.93, based on the ionized lines, and our metallicity measurement is in good agreement with previous estimates. We find that [Ca+Ti/2Fe] and [Cu/Fe] of NGC~5694 are about 0.3 -- 0.4 dex lower than other globular clusters with similar metallicities, but similar to some LMC clusters and stars in some dwarf spheroidal galaxies. Differences persist, however, in the abundances of neutron capture elements. The unique chemical abundance pattern and the large Galactocentric distance (30 kpc) and radial velocity ($-138.6 \\pm 1.0$ km sec$^{-1}$) indicate that NGC~5694 had an extragalactic origin. ",
        "introduction": "The cold dark matter model for cosmology predicts a hierarchical formation mechanism for galaxies, with smaller units accreting to construct larger ones (e.g.\\ Navarro, Frenk, \\& White 1995). Signs of merger fragments have been identified kinematically, especially the Sagittarius dwarf galaxy (Ibata et al.\\ 1994), and probably the Monoceros Ring (Yanny et al.\\ 2003). More substructure is predicted by the models but confirmations have proven elusive.  With the advent of large aperture telescopes, ``chemical tagging\" becomes a powerful technique to probe past merger histories. As Freeman \\& Bland-Hawthorn (2002) discussed, stars born in galaxies whose star formation histories differ from those that have created the bulk of the Galaxy's stars may still be discernible in unusual element-to-iron ratios. For example, Cohen (2004) has found a compelling link between Palomar~12 and the Sagittarius dwarf, confirming the dynamical association found earlier by Dinescu et al.\\ (2000). Venn et al.\\ (2004) summarized the unusual abundance patterns found in the Galaxy's dwarf spheroidal (dSph) galaxy neighbors, demonstrating that the bulk of the Galactic halo did not come from such surviving systems. Unique chemical abundance patterns of globular clusters and halo field stars may become a primary method to identify common star formation origins and histories.  NGC~5694 ($\\ell = 331.1$, $b = +30.4$) is a metal-poor outer halo globular cluster lying far from the Sun (Harris 1975; Ortolani \\& Gratton 1990), as well as far from the Galactic center. Harris (1996) cites a Galactocentric distance of 29 kpc, E($B-V$) = 0.09, and a large radial velocity, $v_{\\rm rad}$ = $-$144.1 km sec$^{-1}$. The large velocity and distance led Harris (1975) and Harris \\& Hesser (1976) to suggest that NGC~5694 has a hyperbolic orbit and it is not bound to our Galaxy, or that the Galaxy contains considerable additional mass beyond the solar orbit than simple model potentials indicate. We have used the analytical model of the Galactic gravitational potential from Allen \\& Santillan (1991) to estimate a lower limit to its apogalacticon distance. We first re-estimate the cluster's distance, employing a mean horizontal branch $V$ magnitude of 18.5 (Harris 1996) and the $M_{\\rm V}$ vs.\\ [Fe/H] relation of Cacciari (2003), finding $R_{\\rm GC}$ = 30 kpc. Based only on the radial velocity, the cluster travels over 100 kpc from the Galactic center. Any significant tangential velocity will likely increase this value further. Inspired by the outer halo nature of the cluster, we have begun a high-resolution spectroscopic study of one red giant branch star in the cluster. Our results suggest that NGC~5694 has a very distinctive elemental abundance pattern, similar in some respects to those of nearby dwarf spheroidal galaxies. ",
        "conclusions": "\\subsection{Radial Velocity}  We measured the heliocentric radial velocity of the program star with respect to that of HD116713 using the IRAF FXCOR task and obtained $-$138.6 $\\pm$ 1.0 km sec$^{-1}$. Our result is in good agreement with that of Geisler et al.\\ (1995). Neglecting the three most deviant velocities, the remaining ten stars in Geisler's sample have a mean radial velocity of $-140.7 \\pm 2.4$ km sec$^{-1}$ (the error is that of the mean). Our radial velocity measurement re-confirms that star I-62 is a member of NGC~5694.  \\subsection{Elemental Abundances} Table~\\ref{tab1} summarizes the elemental abundances of NGC~5694 I-62 using the photometric surface gravity and the spectroscopic temperature. The [el/Fe] ratios for neutral elements are estimated from [el/H] and [\\ion{Fe}{1}/H] ratios. The [el/Fe] for singly ionized elements (\\ion{Ti}{2}, \\ion{Y}{2}, \\ion{Ba}{2}, \\ion{La}{2}, and \\ion{Eu}{2}) are estimated from [el/H] and [\\ion{Fe}{2}/H] ratios. This procedure follows the study of M5 giants by Ivans et al.\\ (2001), and has been employed in our prior work as well (Lee \\& Carney 2002; Lee, Carney, \\& Habgood 2005). See also Johnson et al.\\ (2006) for a discussion of the challenges presented in comparing photometric and spectroscopic temperatures and gravities. In the Table, $n$ is the number of lines used for the calculations of mean elemental abundances and $\\sigma$ is the standard deviation per line. Systematic errors, such as in adopted $gf$ values as a function of excitation potential, which could lead to systematically erroneous temperature estimates, are not included. A more detailed discussion of elemental abundances will be presented elsewhere in the future.  \\subsection{Comparisons with other Stellar Systems}  Our results are based on the analysis of only one star, and the comparisons given below must be considered suggestive rather than definitive. But NGC~5694 appears to be unusual, almost unique, in its chemical abundance pattern and warrants further study.  NGC~5694 I-62 is deficient in $\\alpha$-elements, in particular Ca and Ti, and the iron-peak element, Cu, compared with other globular clusters in our Galaxy. For [Ti/Fe], we adopt the unweighted average of [\\ion{Ti}{1}/\\ion{Fe}{1}] and [\\ion{Ti}{2}/\\ion{Fe}{2}]. Use of neutral titanium lines may suffer from NLTE effects, such as an over-ionization. However, the results from \\ion{Ti}{2} lines also yield lower titanium abundance scales in our program star, indicating that it is truly titanium deficient. In Figure~1, we show  [Ca+Ti/2Fe] and [Cu/Fe] ratios as functions of [Fe/H]. We also show those of other globular clusters in our Galaxy (Pritzl et al.\\ 2005; Simmerer et al.\\ 2003), Large Magellanic Cloud (LMC) globular clusters (Johnson et al.\\ 2006) and nearby dSph galaxies (Shetrone et al.\\ 2003). The [Ca+Ti/2Fe] ratio of NGC~5694 I-62 is very similar to those of Palomar~12 and Terzan~7, which are associated with the Sagittarius dwarf galaxy (see Dinescu et al.\\ 2000 regarding Palomar~12's association), and Ruprecht~106, which has been suggested to have been associated with the Magellanic Clouds (Lin \\& Richer 1992). On the other hand, other $\\alpha$-elemental abundances, [Mg/Fe] and [Si/Fe], appears to be normal. The LMC cluster results from Johnson et al.\\ (2005) also showed [Mg/Fe] and [Si/Fe] ratios similar to those of globular clusters in our Galaxy. Some iron-peak elemental abundances for NGC~5694 I-62, [Mn/Fe] and [Ni/Fe], appear to be normal (Sobeck et al.\\ 2006; Gratton et al.\\ 2004) relative to other clusters. However, the [Cu/Fe] ratio of NGC~5694 I-62 is $\\approx$ 0.4 dex lower than those of globular clusters studied by Simmerer et al.\\ (2003) at [Fe/H] $\\approx$ $-$2.0 dex and the nearby dSphs studied by Shetrone et al.\\ (2003). Some RGB stars in the Sculptor dSph galaxy\\footnote{ Note that Sculptor dSph is the oldest dSph galaxy studied by Shetrone et al.\\ (2003).  However, it appears to have several Gyr of active star formation and its star formation ended about 4 Gyr ago. On the other hand other dSphs have had even more extended periods of star formation, stopping only in the last 1 or 2 Gyr (Dolphin 2002; Tolstoy et al.\\ 2003). As noted by the referee, proto-galactic fragments that disrupted very early in the first few Gyr would resemble only the oldest stellar populations in dSphs, which is not the dominant population in the dSphs because of their slow, but steady, star formation rates.} and LMC appear to have similar [Cu/Fe] ratios.  The $\\alpha$-elements are predominantly synthesized during the SNe~II shell-burning at the end of the lives of massive stars. Most Cu appears to be synthesized by an $s$ process in massive stars. (The relative importance of SN~Ia for Cu abundance remains uncertain, according to Clayton 2003.) Since NGC~5694 has a very old age (De Angeli et al.\\ 2005), the contributions from SNe~Ia are unlikely to be significant, since such events are not thought to appear until $10^9$ or more years following the beginning of star formation. Further, the low [Cu/Fe] ratio of NGC~5694 cannot be understood by a metallicity-dependent yield from SNe~Ia, which appears to be more important in more metal-rich regimes (e.g.\\ McWilliam \\& Smecker-Hane 2005). One possible explanation would be that NGC~5694 formed from a proto-globular cluster cloud which was contaminated by relatively rare, massive SNe~II (e.g.\\ Tolstoy et al.\\ 2003). This suggests that NGC~5694 formed in very different environment than the bulk of globular clusters in our Galaxy.  The neutron capture elements reveal even greater complexity. Venn et al.\\ (2004) and Johnson et al.\\ (2006) have noted the low abundances of [$\\alpha$/Fe] and [Cu/Fe] for dSphs and LMC clusters compared to Galactic halo field and globular cluster stars, and Venn et al.\\ (2004) drew special attention to [Ba/Y] as a significant difference as well. The LMC clusters studied by Johnson et al.\\ (2006) do not share this trend, having solar [Y/Fe] ratios, as found in the Galactic halo, but somewhat enhanced [Ba/Fe]. NGC~5694 I-62 has a very low [Y/Fe] value, like the dwarf spheroidal galaxies, but its [Ba/Fe] ratio is lower, and lower than the LMC clusters as well. [Eu/Fe] is only slightly super-solar, resulting in a [Ba/Eu] ratio well below the LMC clusters or the dSphs.  In short, NGC~5694 is similar in some chemical abundance ratios to the LMC clusters and the dSphs, but a closer look at the neutron capture elements suggests significant differences.  Is NGC~5694 related to any of the existing dSphs? Kinematically the answer appears to be ``no\". Lynden-Bell \\& Lynden-Bell (1994) introduced the concept of alignments of orbital poles, but did not identify NGC~5694 as related to any of the dwarf spheroidal galaxies. Majewski (1994) employed radial velocities in addition to positional data and drew a similar conclusion. Finally, space velocities have been employed by Palma et al.\\ (2002) to compare the clusters' motions, but, unfortunately, NGC~5694 still does not have a measured proper motion. We await such a measurement with keen interest, given that the estimated large apogalacticon distance and unique chemical abundances suggest that NGC~5694 formed independently of the bulk of the Galaxy, and was captured subsequently."
    },
    "0606/astro-ph0606675_arXiv.txt": {
        "abstract": "We investigate how hierarchical models for the co-evolution of the massive black hole (MBH) and AGN population can reproduce the observed faint X-ray counts.  We find that the main variable influencing the theoretical predictions is the Eddington ratio of accreting sources.  We compare three different models proposed for the evolution of AGN Eddington ratio, $f_{\\rm Edd}$: constant $f_{\\rm Edd}=1$,   $f_{\\rm Edd}$ decreasing with redshift, and $f_{\\rm Edd}$ depending on the AGN luminosity, as suggested by simulations of galactic mergers including MBHs and AGN feedback.  We follow the full assembly of MBHs and host halos from early times to the present in a $\\Lambda$CDM cosmology. AGN activity is triggered by halo major mergers and MBHs accrete mass until they satisfy the observed correlation with velocity dispersion.  We find that all three models can reproduce fairly well the total faint X-ray counts. The redshift distribution is however poorly matched in the first two models. The Eddington ratios suggested by merger simulations predicts no turn-off of the faint end of the AGN optical luminosity function at redshifts $z\\ga1$, down to very low luminosity. ",
        "introduction": "Several hierarchical models for the evolution of the MBH and AGN populations \\citep[see,e.g.,][]{haehnelt1993, haiman2000, hatziminaoglou2001,  wyithe2003, Cattaneo1999, kauffmann2000, cavaliere2000, VMH, Granatoetal, lapi2006} have proved successful in reproducing the AGN optical luminosity function (OLF) in a large redshift range ($1\\la z\\la6$). Typically, these models assume that AGN activity is triggered by major mergers.  Galactic interactions  trigger gas inflows, and the cold gas may be eventually driven into the very inner regions, fueling an accretion episode and the growth of the nuclear MBH. Hydrodynamic simulations of major mergers have shown that a significant fraction of the gas in interacting galaxies falls to the center of the merged system \\citep{mihos94, mihos96}: the cold gas may be eventually driven into the very inner regions, fueling an accretion episode and the growth of the nuclear BH. This last year has been especially exciting, as the first high resolution simulations of  galactic mergers including BHs and AGN feedback, showed that the merger scenario  is generally correct \\citep{Springel2005,dimatteo2005}. In hierarchical models of galaxy formation major mergers are responsible for forming bulges  and elliptical galaxies. Support for merger driven activity therefore comes from the observed  correlation between bulge luminosity - or stellar velocity dispersion - and black hole mass, suggesting a  single mechanism for assembling black holes and forming spheroids in galaxy halos \\citep{magorrian1998, Marconietal2004, Gebhardt2000, ferrarese2000}.   Notwithstanding the success at high redshift, hierarchical models struggle to match the AGN OLF at low redshift ($z\\la1$) by overpredicting the bright end,  and underpredicting the faint end of the OLF, as the decrease of the halo merger rate with time is less dramatic than the observed fall of the AGN population. The overprediction of bright AGN can be imputed to inefficient cooling in large halos. Imposing an upper limit to the AGN host halo mass ($\\sim10^{13.5} M_{\\odot}$, \\cite{wyithe2003, Marulli2006}) in fact significantly improves the match at the bright end of the OLF.  The under-abundance of faint AGN can be instead attributed to the assumption, common to most models, of very efficient accretion, at rates close to the Eddington rate.  \\cite{merloni2003} and \\cite{merloni2004} have shown that low-redshift AGN are probably accreting inefficiently, i.e. both at an accretion rate much smaller than the Eddington rate and with a low radiative efficiency. These considerations suggest that successful predictions for the evolution of AGN luminosity should include more sophisticated models for accretion. A first step in this direction can be taken by considering the results of the recent merger simulations, which track also accretion on a central MBH. Although these simulations lack the necessary resolution for resolving the accretion process in the vicinity of the MBH, empirical models \\citep{hopkins2005a}, based on coupling results from the above simulations with the observed LF in the hard X-ray band (HXLF), have been shown to reproduce simulaneously several optical and X-ray observations.  The \\cite{hopkins2005a} empirical models, however, are not embedded in a cosmological evolutionary framework, that is they derive the MBH population properties at a given time, but not how the population of black holes at an earlier time evolves into the MBHs present at a later time. From the observed HXLF \\cite{hopkins2005a} derive the rate at which AGN of a given luminosity at the peak of activity must be created. This information is then used as a proxy for the galaxy merger rate which should provide the boundary conditions for determining the evolution of the MBH population.  We here couple the predictions from \\cite{hopkins2005a} with the merger rate expected in the currently favoured cold dark matter scenario. The main novelties of the present investigation with respect to \\cite{hopkins2005a}, are therefore that (i) the rate of mergers, which trigger AGN activity, is directly derived in the currently favoured cold dark matter (CDM) cosmology. In principle, the empirical merger rate derived by \\cite{hopkins2005a} is not granted to correspond  to the CDM one.  Also, (ii), we grow MBHs in a self-consistent way, that is we trace the whole accretion history of MBHs from early times to the present, requiring continuity in the population.  Our aim is to investigate here to which extent hierarchical models, coupled with the prescriptions based on the above simulations (\\S1), can reproduce the low-redshift evolution of AGN. We show that reproducing the HXLF is a necessary, but not sufficient, condition for matching the redshift distribution of faint X-ray counts (\\S2). We identify the redshift distribution of faint X-ray counts as the most sensitive observational result to discriminate between models (\\S3, \\S4). Finally, in \\S5 we summarize the results and discuss their implications. ",
        "conclusions": "We have attempted in this paper to place constraints on the global accretion properties of the MBH population at $z<3$.  We consider the full cosmological evolution of MBH embedded in their host halos, rather than adopting an empirical approach which takes as a starting point the observed LF in a given band in order to explain the properties of AGN in other bands. We focus here on the strength of accretion, parameterizing the accretion rate as a function of the Eddington rate, $f_{\\rm Edd}$. We show that simple models which assume $f_{\\rm Edd}=1$ (model I), although highly idealized, still are able to explain satisfactorily the growth of MBH, as traced by the OLF, in a large redshift range, as claimed previously by various investigations.  Even at low redshift, where hierarchical models start to fail in their predictions, the OLF is very well reproduced.  The HXLF probes fainter sources than the OLF, and the simplistic model I is shown to provide a poor match with the HXLF at $z<1$.  A decreasing average Eddington ratio \\citep[model II]{Shankar2004} provides a better agreement with the \\emph{shape} of the HXLF, but underestimates the \\emph{normalization}, as there is not enough time to grow high mass MBHs if the Eddington ratio is not large at $z>3$.  A model (model III) with an Eddington ratio depending on luminosity \\citep{hopkins2005a, Shankar2004} seems to be the best match at low redshift. The difference among the models is magnified when the faint X-ray counts, and in particular their redshift distribution, are calculated. Model I has a strictly hierarchical growth of MBHs, and moreover, a univocal relationship between MBH mass and AGN luminosity. Model II allows massive black holes at low redshift to shine at lower luminosity, as the accretion rate, in units of the Eddington one, decreases with time. A rigid redshift dependence, however, creates two problems: first, subEddington accretion at $z>3$ implies a much slower growth for high redshift AGN, thus causing an underestimate of the LFs \\emph{normalization} at all the considered redshifts. Second, the evolution of the Eddington rate is not fast enough, at $z<3$, to account for the faint end of the LFs, in particular the HLF.  The accretion rates, and Eddington ratios, predicted by simulations are here embedded into a cosmological framework (models IIIa and IIIb). The resulting AGN population provides a satisfactory match with  the OLF, HXLF and faint X-ray counts at low redshift, while sources at high redshift are more problematic. The low accretion rates predicted by simulations imply very long growth timescales for black holes, and therefore underestimate the occurrence of  bright quasars powered by billion solar masses black holes at high redshift ($z>2$).  It is not clear, however, if the accretion rate found in simulations, which relates to model III, i.e., strongly subEddington for low-luminosity sources (and therefore in all cases for small MBHs) indeed applies at very high-redshift.  The simulations on which \\cite{hopkins2005a} model is based assume mergers between \"normal\" galaxies. How \"normal\" are galaxies at $z>3$?  It might indeed be possible that the conditions in pre-galactic structures at high-redshift (e.g., the disturbed morphological state of galaxies) cannot be studied with simulations of mergers of evolved discs and bulges, but would require different initial conditions for the merging galaxies.  Future deep surveys can help us distinguish between various models in the high redshift Universe (see Figure \\ref{fig3}), and can probably locate the time, if any, for a transition between messy mergers with on average efficient accretion onto the MBHs, and standard galactic mergers, predicting long periods of inefficient accretion.  Finally, we computed the contribution to the XRB of faint AGN lying below the sensitivity limits of current X--ray surveys. We found that all model predict a contribution to the unresolved XRB consistent with the available limits, accounting for the whole unresolved XRB in the 2--8 keV band and for $\\sim 50$\\% in the 0.5--2 keV. The residual background intensity in the soft band may be provided by AGN shining at $z\\ga 6$ (Salvaterra et al. 2006)."
    },
    "0606/astro-ph0606439_arXiv.txt": {
        "abstract": "{We have obtained deep and wide field imaging of the Coma cluster of galaxies with the CFH12K camera at CFHT in the B, V, R and I filters. In this paper, we present the observations, data reduction, catalogs and first scientific results.  We investigated the quality of our data by internal and external literature comparisons. We also checked the realisation of the observational requirements we set.  Our observations cover two partially overlapping areas of $42 \\times 28$ arcmin$^2$, leading to a total area of 0.72 $\\times$ 0.82 deg$^2$. We have produced catalogs of objects that cover a range of more than 10 magnitudes and are complete at the 90$\\%$ level at B$\\sim$25, V$\\sim$24, R$\\sim$24 and I$\\sim$23.5 for stellar-like objects, and at B$\\sim$22, V$\\sim$21, R$\\sim$20.75 and I$\\sim$20.5 for faint low-surface-brightness galaxy-like objects. Magnitudes are in good agreement with published values from R$\\sim$16 to R$\\sim$25. The photometric uncertainties are of the order of 0.1 magnitude at R$\\sim$20 and of 0.3 magnitude at R$\\sim$25. Astrometry is accurate to 0.5~arcsec and also in good agreement with published data.  Our catalog provides a rich dataset that can be mined for years to come to gain new insights into the formation and evolution of the Coma cluster and its galaxy population. As an illustration of the data quality, we examine the bright part of the Colour Magnitude Relation (B-R versus R) derived from the catalog and find that it is in excellent agreement with that derived for galaxies with redshifts in the Coma cluster, and with previous CMRs estimated in the literature.  } ",
        "introduction": "\\label{sec:intro}   Long considered the archetype of rich and relaxed clusters, Coma is now known to have a rather complex structure, with a number of substructures detected at various wavelengths (see Biviano 1998 for a complete review on the Coma cluster up to 1995). Indeed, although a large quantity of multi-wavelength observational data are presently available for Coma, our theoretical understanding of this unique cluster is still far from complete.  Previous photometric studies of the Coma cluster were done in few colors, and were limited either to wide but shallow surveys or to small deep ones, often obtained under less-than-ideal seeing conditions. For example, very large shallow surveys (e.g. Godwin \\& Peach 1977, hereafter GP77) were the basis for important studies of the Coma cluster luminosity function and of the properties of galaxies in general. However, due to the shallowness of this catalog, very little was known about Coma's galaxy population at magnitudes fainter than R$\\simeq$20.  Recent deeper catalogs are still limited to R$\\sim$21 when covering the whole cluster (e.g. Castander et al. 2001, Terlevich et al. 2001 or Komiyama et al. 2002). Finally, very deep catalogs such as that of Bernstein et al.(1995) covered only a very small area of sky.   While many of these previous studies led to important discoveries, all suffered to some extent from insufficient depth or coverage, which made it difficult to generalize the results obtained from these studies.  Here we report a new catalog of Coma cluster galaxies, obtained with the CFH12K camera (Cuillandre et al.  2000) installed at the prime focus of the Canada-France-Hawaii Telescope (CFHT), which provides the first deep, wide-field images of Coma in the B, V, R and I filters, in seeing conditions well below 1~arcsec for a large part of the observations. As the field covered by the CFH12K is 0.32 deg$^2$ we made two overlapping pointings to cover a total field of 0.6 deg$^2$. The data obtained are significantly deeper than any of the recent CCD large field imaging of the Coma cluster; e.g. Terlevich et al. 2001 (U and V filters, depth V=20, 1~deg$^2$, 2.2~arcsec seeing) or Komiyama et al. 2002 (B and R filters, 2.25~deg$^2$, complete to R=21, seeing from 0.8 to 1.5~arcsec). A BVR trichromic image of the final result is shown in Fig. ~\\ref{pub}.  \\begin{figure*} \\caption[]{BVR color image of the Coma cluster from our CFH12K data. The size of the field is 0.72deg$\\times$0.82deg. North is up and East is to the left.} \\label{pub} \\end{figure*}   Our main scientific goals in obtaining deep wide-field multi-color images of Coma are: a)~to study environmental effects on the faint end slope of the galaxy luminosity function (as a very significant extension of, e.g., the work by Lobo et al. 1997 or by Andreon $\\&$ Cuillandre 2002); b)~to analyze Coma's morphological and dynamical structure to see how it relates to the cluster's formation history (e.g. Mellier et al. 1988 or Adami et al. 2005b), the diffuse light emission (Adami et al. 2005a) and the faint and low surface brightness galaxy distribution (Adami et al., 2006, submitted). However, because this new Coma catalog provides a rich source of data that will undoubtedly be useful for other researchers, we intend to make it publicly available.  In this paper, we describe our data acquisition, reduction, and cataloguing. Catalogs are available at http://cencosw.oamp.fr/. ",
        "conclusions": "We have presented and discussed the properties of a catalog of more than 60,000 objects in the Coma cluster based on deep CFHT observations in the B, V, R and I bands and spanning an unprecedented 10 magnitude range in one of the largest $\\sim$0.72$\\times$ $\\sim$0.82 deg$^2$ areas presently available for Coma at this depth. The catalog is complete at the 90$\\%$ level down to R$\\sim$24 for stellar-like objects and to R$\\sim$20.75 for faint low surface brightness (see Sect.~5.2) galaxy-like objects.  Astrometry is accurate to 0.5 arcsec, except for first epoch data on the edges and corners of the CCDs. Magnitude errors are smaller than $\\sim$0.3 in all bands down to R=25. Saturation effects on objects brighter than I=17.5 have been removed by using shallower published data. More generally, galaxy magnitudes are in good agreement with published data and the colors of galaxies and stars are in good agreement with synthetic models. The bright part of the CMR agrees well with that derived for galaxies with known redshifts in the Coma cluster, and with previous CMRs published in the literature.  The star-galaxy separation is robust for all objects brighter than I=21, and the star counts fit the Besan\\c con model very well.  These data have already been used to search for diffuse emission in the Coma cluster (Adami et al. 2005a) and to look for and analyze properties of faint low surface brightness galaxies (Adami et al. 2006).  We plan next to use these data to investigate the properties of the different cluster galaxy classes and to derive the luminosity functions of the Coma cluster galaxies in various bands and in various regions of the cluster. This should deepen our knowledge of environmental effects on galaxy luminosity functions, already shown to be strong in Coma (e.g. Lobo et al. 1997).  The catalog described above will be made public by the CENCOS center at http://cencosw.oamp.fr/. Individual catalogs for the second epoch R and I band data will also be available upon request."
    },
    "0606/astro-ph0606113_arXiv.txt": {
        "abstract": "{In coded mask techniques, reconstructed sky images are pseudo-images: they are maps of the correlation between the image recorded on a detector and an array derived from the coded mask pattern.} {The \\integ/\\ibis telescope provides images where the flux of each detected source is given by the height of the local peak in the correlation map. As such, it cannot provide an estimate of the flux of an extended source. What is needed is intensity sky images giving the flux per solide angle as typically done at other wavelengths.} {In this paper, we present the response of the \\integ \\ibis/\\isgri coded mask instrument to extended sources. We develop a general method based on analytical calculations in order to measure the intensity and the associated error of any celestial source and validated with Monte-Carlo simulations.}{We find that the sensitivity degrades almost linearly with the source extent. Analytical formulae are given as well as an easy-to-use recipe for the \\integ user. We check this method on \\ibis/\\isgri data but these results are general and applicable to any coded mask telescope.}{} ",
        "introduction": "\\label{s:introduction}  The usual techniques of imaging the sky by focusing the radiation become ineffective above $\\sim$ 15\\un{keV} due to technological constraints. Future observatories (\\eg Constellation-X/HXT, Hailey \\etal 2004 ; Simbol-X, Ferrando \\etal 2005) will push this limit up to 60-70\\un{keV} with the use of multi-layer grazing optics \\cite{c:ramsey02}. Collimators using the standard on/off techniques are limited by the source density at low energies (< 50\\un{keV}) and by the internal background variability at higher energies. Both problems are addressed with coded-aperture techniques, which have been employed successfully in high-energy astronomy in several previous missions such as the GRANAT/SIGMA instrument \\cite{c:paul91} or the BeppoSAX Wide Field Cameras \\cite{c:jager97}. In such telescopes, source radiation is spatially modulated by a mask consisting of an array of opaque and transparent elements and recorded by a position sensitive detector, allowing simultaneous measurement of source plus background in the detector area corresponding to the mask holes and background only in that corresponding to the opaque elements. For each point source in the field of view the two-dimensional distribution of events projected onto the detector reproduces a unique shadow of the mask pattern and the recorded shadowgram is the sum of many such distributions. The popular way to reconstruct the sky image is based on a correlation procedure between the recorded image and a decoding array derived from the spatial characteristics of the mask pattern. One can describe the usual properties of telescopes for these systems as follows: the angular resolution, defined by the angle subtended by the smallest mask element seen from the detector, and the field of view (FOV) divided in two parts: the Fully Coded Field of View (FCFOV) representing the sky region where the recorded source radiation is fully modulated by the mask and the Partially Coded Field of View (PCFOV) the region where a source projects only a part of the mask pattern on the detector.  Many mask patterns have been designed and employed in the past for optimizing the imaging quality of point-like sources in the high-energy domain. Recently, Sch\\\"afer (2004) has presented a novel method for imaging extended sources by constructing mask patterns with a pre-defined Point Spread Function (PSF). In this paper, we have developed a method to reconstruct the intensity and estimate the corresponding error of extended sources seen through a given coded mask telescope. This work is motivated by the fact that several sources of interest in \\gammaray astronomy must be considered as extended in order to properly derive their intensity such as supernova remnants (hereafter, SNRs): the Tycho SNR (with an apparent diameter $\\theta$ of $\\sim$ 8\\m), SN~1006 ($\\theta \\sim$ 30\\m), G~347.3-0.5 ($\\theta \\sim$ 1$^{\\circ}$) or RX~J0852-4622 ($\\theta \\sim$ 2$^{\\circ}$), clusters of galaxies (\\eg \\coma, $\\theta \\sim 1^{\\circ}$), and diffuse interstellar emission from various high-energy processes. We have applied this method to the \\ibis/\\isgri instrument onboard the \\integ satellite but the results presented here could be easily adapted to any coded mask telescope. We briefly describe the basic properties of \\ibis and the principles of the standard data analysis we used for simulating celestial extended sources in section \\ref{s:method}, considering the simplest case of a uniform disk. We present our results concerning the \\ibis response for this source geometry over a large range of radii (section \\ref{s:disk}) as well as in an astrophysical case: the supernova remnant SN~1006 (section \\ref{s:SN1006}). These results allow us to find a general method to reconstruct the flux and the associated error of an extended source presented in sections \\ref{s:real_flux} and \\ref{s:real_err}, respectively. We also describe the tests we performed on the Crab Nebula and compare the results of this method to those of the standard analysis (section \\ref{s:crab}). In a joint paper, we have successfully applied this method on the \\ibis/\\isgri data of the \\coma \\cite{c:renaud06}, the first extended source detected with \\ibis. Finally, we give practical tips on the use of this method for any interested observer. ",
        "conclusions": "\\label{s:discussion}  We have shown that it is always possible to obtain an estimate of the flux of any extended source with coded mask imaging. This is an intrinsic property of these imaging systems where the source flux is obtained through a linear combination of detector-pixel count rates. We have derived analytical expressions for estimating the global flux and its associated error of any extended sky source which were tested by Monte-Carlo simulations. This work was done relating to what one could expect to observe in the \\gammaray domain through a MURA coded mask instrument such as \\ibis onboard the \\integ satellite. For \\ibis, extended sources with a radius greater than $\\sim$ 4.5$^{\\circ}$ will be confused with their associated ghosts. One will need a specific method to clean the images but this is beyond the scope of this paper. The advantage of this method, as long as the integration region contains all the source, is that the observer does not have to assume any source geometry during the analysis as is the case for methods using model fitting. Of course, in principle, the optimum signal to noise ratio is attained with the smallest integration region still containing all the source. However, due to the flux dilution, the geometry of the extended source is generally not apparent in the maps and the observer may not always be provided with a precise source geometry. Nevertheless, the signal to noise ratio does not decrease rapidly with the integration region size once the source is fully contained in it. In order to check the global method on an astrophysical case, we analyzed the \\ibis/\\isgri data on the \\coma, a bright known cluster of galaxies with an \\xray size of order $30'$ diameter. We applied this method, increasing the integration region until a maximum signal to noise ratio was reached. In that way, we were able to extract for the first time the global intensity of this extended source as well as interesting constraints on its morphology. The results are presented in a joint paper \\cite{c:renaud06}. In a forthcoming paper, this method will be also used on the \\ibis/\\isgri observations on the Vela Junior SNR (Renaud \\etal 2006, in preparation).  Usually, the observer needs a lot of integration time for detecting an extended source ($\\sim$ 5 $\\times$ 10$^{5}$ secondes for the \\coma) and will have to construct mosaics of individual images. After obtaining the individual images with the standard analysis, the next step is to build mosaics as usual by weighting each image by their associated variance and transform them following the practical tips on the use of the method given in Appendix \\ref{a:practical_tips}. Any observer interested in this method may contact the first author to obtain the map $\\tilde{A}$."
    },
    "0606/astro-ph0606749_arXiv.txt": {
        "abstract": "{ We predict the survival time of initially bound star clusters in the solar neighbourhood taking into account: (1) stellar evolution, (2) tidal stripping, (3) shocking by spiral arms and (4) encounters with giant molecular clouds. We find that the predicted dissolution time is $\\tdis= 1.7 (\\Mi/10^4\\,\\Msun)^{0.67}$ Gyr for clusters in the mass range of $10^2 < \\Mi < 10^5\\,\\Msun$. The resulting predicted shape of the logarithmic age distribution agrees very well with the empirical one, derived from a complete sample of clusters in the solar neighbourhood within 600 pc. The required scaling factor implies a star formation rate of $4~10^2$ \\Msunmyr\\  within 600 pc from the Sun or a surface formation rate of $3.5~10^{-10}$ \\Msun yr$^{-1}$pc$^{-2}$ for stars in bound clusters with an initial mass in the range of $10^2$ to $3~10^4$ \\Msun. ",
        "introduction": "\\label{sec:1}  The first empirical determination of the lifetime of clusters in the solar neighbourhood is by \\citet{1958RA......5..507O}, who noticed the lack of clusters older than a few Gyr in the solar neighbourhood. Later, \\citet{1971A&A....13..309W} derived a mean dissolution time of 0.2 Gyr from the age distribution of clusters. Since most of the observed clusters within about 1 kpc from the sun have a mass in the range of $10^2$ to a few $10^3$ \\Msun\\ the value derived by Wielen is for clusters in that mass range. Theory predicts that the dissolution time of clusters depends on their initial mass in that massive clusters survive longer than low mass clusters (e.g. \\citealt{1958ApJ...127...17S}, \\citealt{1985IAUS..113..449W}, \\citealt{1990ApJ...351..121C}, \\citealt{1997ApJ...474..223G} and references therein). \\citet{2003MNRAS.340..227B} (hereafter BM03) showed from \\nbody\\ simulations that the dissolution time of clusters in the tidal field of the galaxy depends on their initial mass, $M_i$, as $\\tdis \\sim M_i^{0.62}$. Independently, this same power-law dependence was also derived empirically in a study of cluster samples in four galaxies by \\citet{2003MNRAS.338..717B}.  The dissolution time of clusters in the solar neighbourhood was recently redetermined by \\citet{2005A&A...441..117L} (hereafter L05), based on a new cluster sample of \\citet{2005A&A...438.1163K}. They found a dissolution time of $\\tdis = \\tfour (\\Mi/10^4 \\Msun)^{0.62}$ with $\\tfour = 1.3 \\pm 0.5$ Gyr for clusters with  $10^2 < M < 10^4$. This is a factor 5 shorter than the $\\tfour=6.9$ Gyr that follows from the $\\tdis(\\Mi)$ relation derived by the \\nbody\\ simulations of BM03 for clusters more massive than 4500 \\Msun\\ at a Galactocentric distance of 8.5 kpc. The simulations of BM03 include a realistic stellar mass function, stellar evolution, two-body relaxation, a detailed treatment of binary evolution and close encounters of stars. Part of this difference may be due to the fact that in low mass clusters, with lifetimes shorter than about 1 Gyr, the dynamical evolution is affected by the presence of massive stars during most of their lifetime  (see Fig. 5 of BM03). It is doubtfull that this effect alone can fully explain the difference between the results of L05 and BM03. In fact, the large discrepancy suggests that other, probably external,  disruptive effects must play an important role in destroying star clusters in the solar neighbourhood.   In this paper we explain the lifetime of clusters in the solar neighbourhood, by taking into account the combined effects of stellar evolution, tidal stripping, encounters with giant molecular clouds (GMCs) and spiral arm shocks.  We use stellar population models to describe the stellar evolution and the results of BM03 for tidal stripping. For the effects of GMCs and spiral arms we adopt the new estimates  from the recent studies by  \\citet{gieles06b} (hereafter GPZB06) and \\citet{gieles06a} (hereafter GAPZ06), which are based on \\nbody\\ simulations.   The structure of the paper is as follows. In Sect. 2 we discuss the predicted mass loss from star clusters by stellar evolution, tidal stripping, encounters with GMCs and spiral arm shocks. We calculate the mass evolution of clusters due to these four  effects. In Sect. 3 we compare our results with the observed age distribution of clusters in the solar neighbourhood. The discussion and conclusions are given in Sect. 4. ",
        "conclusions": "We studied the dissolution of star clusters in the solar neighbourhood due to four effects: stellar evolution, tidal stripping, spiral arm shocks and encounters with GMCs. For this study we adopted the descriptions of GAPZ06 and GPZB06 for the dissolution of star clusters by spiral arms and encounters with GMCs. We found that the last effect plays a  dominant role in the solar neighbourhood.    The cluster dissolution time due to spiral arms and GMCs depends on the density of the clusters, i.e. on $M/\\rh^3$. This implies that the dissolution time is $\\tdis \\sim M^{\\gamma}$ with $\\gamma={1-3\\lambda}$ if the radius of a cluster depends on its mass as $\\rh \\sim M^{\\lambda}$. We adopted $\\lambda \\simeq 0.1$ as found by \\citet{2004A&A...416..537L} for clusters in spiral galaxies and so $\\gamma \\simeq 0.7$ for dissolution by both spiral arms and GMCs. This value is very similar to $\\gamma=0.62$ predicted for dissolution by the tidal field only (BM03) and empirically derived for cluster samples in four galaxies by \\citet{2003MNRAS.338..717B}. If there was no mass radius dependence for clusters in the solar neighbourhood, e.g. $\\lambda=0$, then the predicted age distribution would have a shallower slope than shown in Fig. \\ref{fig:3}, since more old (massive) clusters would have survived.   Our calculated dissolution times of clusters in the solar neighbourhood are about a factor five smaller than predicted by BM03 for clusters in the tidal field of our Galaxy, with stellar evolution, binaries and two-body relaxation taken into account. This is reminiscent of the short dissolution time of clusters in the central region of the interacting galaxy M51, where the empirical dissolution time is even ten times shorter than can be explained by stellar evolution and tidal fields \\citep{2005A&A...441..949G}. GMCs severely limit the lifetime of clusters in that galaxy also (see the discussion in GPZB06).  The steep drop in the observed age distribution at $t \\simeq 2$ Gyr can be explained by an upper mass limit for the initial cluster mass in the solar neighbourhood of about $3~10^4$ \\Msun. However, this value is uncertain because it depends crucially on the completeness of the used sample at ages above 1 Gyr. (The sample contains only six clusters older than 1 Gyr.)  The mass versus age distribution of our adopted sample, shown in Fig. 8 of L05, suggests that the lower mass limit of the observed clusters increases steeply for clusters older than 1 Gyr. Since the predicted value of $dN/dt$ at any age depends on $M_{\\rm lim}$ and $M_{\\rm max}$ as given in Eq. \\ref{eq:Nt}, and $M_{\\rm lim}$ increases when the lower mass limit increases, an increase in this limit implies an increase in $M_{\\rm max}$ derived from the observed age distribution. Based on this argument and the small number of clusters older than about 1 Gyr in the observed sample, we  conclude that the derived value of $M_{\\rm max}= 3~10^4\\,\\Msun$ should be considered as a lower limit of the real maximum initial mass.   The vertical shift applied to the  predicted age distributions to match the observed one indicates a star formation rate of $4~10^2$ \\Msunmyr\\ in bound clusters of $\\Mi>100\\,\\Msun$ within a distance of 600 pc, corresponding to a surface formation rate of $3.5~10^{-10}$ \\Msun yr$^{-1}$pc$^{-2}$. This is a factor 2 to 3 smaller than the SFR derived from the study of embedded stars by \\citet{2003ARA&A..41...57L} because many of the stars are born in unbound clusters that dissolve within 10 Myr.  The very good agreement between the predicted and observed age distribution of clusters shows that dissolution of clusters in the solar neighbourhood is dominated by encounters with GMCs, as was already suggested by \\citet{1958RA......5..507O}. In fact, the good agreement may be slightly fortuitous because we have underestimated the dissolution by two-body relaxation (see \\S 1) and slightly overestimated the dissolution by encounters with GMCs, because we adopted the midplane density of GMCs whereas clusters may spend a faction of their lifetime above or under the galactic disk. Both effects are expected to be smaller than a factor two and may partially cancel out."
    },
    "0606/astro-ph0606277_arXiv.txt": {
        "abstract": "It has been known for more than 30 years that star formation in giant molecular clouds (GMCs) is slow, in the sense that only $\\sim 1\\%$ of the gas forms stars every free-fall time. This result is entirely independent of any particular model of molecular cloud lifetime or evolution. Here we survey observational data on higher density objects in the interstellar medium, including infrared dark clouds and dense molecular clumps, to determine if these objects form stars slowly like GMCs, or rapidly, converting a significant fraction of their mass into stars in one free-fall time. We find no evidence for a transition from slow to rapid star formation in structures covering three orders of magnitude in density. This has important implications for models of star formation, since competing models make differing predictions for the characteristic density at which star formation should transition from slow to rapid. The data are inconsistent with models that predict that star clusters form rapidly and in free-fall collapse. Magnetic- and turbulence-regulated star formation models can reproduce the observations qualitatively, and the turbulence-regulated star formation model of Krumholz \\& McKee quantitatively reproduces the infrared-HCN luminosity correlation recently reported by Gao \\& Solomon. Slow star formation also implies that the process of star cluster formation cannot be one of global collapse, but must instead proceed over many free-fall times. This suggests that turbulence in star-forming clumps must be driven, and that the competitive accretion mechanism does not operate in typical cluster-forming molecular clumps. ",
        "introduction": "\\citet{zuckerman74} first pointed out that star formation in giant molecular clouds (GMCs) happens surprisingly slowly. Comparing the mass of GMCs in the Galaxy with the total Galactic star formation rate implies that no more than $\\sim 1\\%$ of the gas can form stars for each cloud free-fall time. This result is sufficiently surprising that numerous theories have been proposed to explain it, ranging from the idea that strong magnetic fields \\citep[e.g.][]{allen00} or turbulence \\citep[e.g.][]{krumholz05c} within clouds inhibit star formation to the idea that galactic-scale gravitational instability regulates star formation \\citep[e.g.][]{li05a} to the idea that GMCs are, contrary to most observational estimates to date \\citep{blitz06a}, actually gravitationally unbound \\citep[e.g.][]{clark04}.  Following \\citet{krumholz05c}, we define the dimensionless star formation rate per free-fall time $\\sfrff$ as the fraction of an object's mass that it converts into stars per free-fall time at the mean density of that object. The \\citet{zuckerman74} argument shows that $\\sfrff\\approx 0.01$ for GMCs. This provides a powerful constraint on models of GMCs. For example, it rules out the early GMC model of \\citet{goldreich74}, a simplified version of which is that GMCs are spheres of gas of density $\\rho$ that are in free-fall collapse. The clouds reach a singularity in a time $t_{\\rm ff}(\\rho)=[3\\pi/(32 G \\rho)]^{1/2}$, at which point their mass is converted into stars. For a population of such clouds $\\sfrff=1$, which is inconsistent with the \\citet{zuckerman74} result.  An important observational question, and a crucial test for theories of how star formation is regulated, is to what densities and length scales $\\sfrff$ remains much smaller than unity. In other words, is there a density at which something like the \\citet{goldreich74} free-fall collapse model becomes reasonable? As an example, consider observing a star-forming region using a molecular tracer sensitive to gas at densities of $n_H\\gtsim 10^{12}$ cm$^{-3}$, where $n_H$ is the number density of hydrogen nuclei, to estimate the total mass of gas at such high densities. This is larger than the mean density of ``cores'' seen both observationally \\citep[e.g.][]{barranco98} and in simulations of star formation regulated by turbulence \\citep[e.g.][]{jappsen05}, and is roughly the density at which models of magnetically-regulated star formation predict that gas will completely decouple from the magnetic field and enter free-fall collapse \\citep[e.g.][]{desch01}. At such high densities protostellar outflows can probably stop at most half the gas from reaching a star \\citep{matzner00}, and the thermal pressure in gas at $10^{12}$ cm$^{-3}$ is considerably higher than the typical ram pressure of the turbulence in GMCs, so gas at such high densities is largely impervious to external perturbations.  Thus, regardless of the model of star formation one adopts, one would expect that almost all of the gas at densities $\\gtsim 10^{12}$ cm$^{-3}$ is part of gravitationally-bound, collapsing objects that have largely decoupled from the background turbulent flow. In the absence of effective internal support or external disturbance, order unity of the gas at such high densities is likely to be incorporated into a star within one free-fall time. For this reason, essentially all models of star formation predict that the total mass of gas at densities $\\gtsim 10^{12}$ cm$^{-3}$, divided by the free-fall time of this gas, should yield a value comparable to the total star formation rate in the region over which the mass is measured. Instead of $\\sfrff\\sim 0.01$ as for GMCs, i.e. slow star formation, one would obtain $\\sfrff\\sim 1$, i.e. rapid star formation.  However, different models make different predictions about the shape of the curve of $\\sfrff$ versus density in between $\\sim 1\\%$ at the characteristic density of GMCs, $\\sim 100$ cm$^{-3}$, and $\\sim 1$ at a density $\\gtsim 10^{12}$ cm$^{-3}$. These different predictions correspond to different models of the physical scale at which gas both decouples from the background flow and ceases to be supported by internal feedback mechanisms, and thus transitions from slow to rapid star formation. At one extreme, magnetic regulation models such as those of \\citet{desch01}, neglecting for the moment turbulent enhancement of the ambipolar diffusion rate \\citep[e.g.][]{heitsch04}, predict that star formation only becomes rapid once gas decouples from the magnetic field, a process that does not even begin until densities of $\\sim 3\\times 10^{10}$ cm$^{-3}$. At the other extreme, \\citet{bonnell03} argue that star clusters form from gas clumps at densities of $\\sim 5\\times 10^{4}$ cm$^{-3}$ that undergo a free-fall collapse in which at least $\\sim 30\\%$ of their mass is converted into stars \\citep{kroupa01b}, so star formation should be rapid at this density or higher. In this model, the decoupling scale corresponds to the transition from globally unbound structures (GMCs) to globally bound structures (protoclusters). Thus, extending the \\citet{zuckerman74} calculation of star formation rate divided by free-fall time to higher densities, in hopes of identifying a scale at which there is a transition from slow to rapid star formation, provides a means of distinguishing between models of how star formation is regulated.  It is important at this point to differentiate the concepts of the \\textit{rate} and \\textit{efficiency} of star formation, and the \\textit{lifetime} of star-forming clouds. Unfortunately these terms are often confused in the literature, and there are no standard definitions, so we describe here the definitions we use in this paper. The star formation rate $\\dot{M}_*$ is the easiest to define, since it is simply the instantaneous conversion rate of gas into stars within some volume. If we pick a density threshold $\\rho$, we can then define the dimensionless star formation rate per free-fall time for the gas above that density threshold, $\\sfrff = \\dot{M}_*/[M(>\\rho)/t_{\\rm ff}(\\rho)]$, where $M(>\\rho)$ is the mass of material inside the volume of density $\\rho$ or greater. In contrast, the lifetime $t_{\\rm cl}(\\rho)$ of star-forming cloud is somewhat more ambiguous. We take it to mean the total duration during which a cloud is visible in a tracer sensitive to densities of $\\rho$ or more. Note that our definition neglects the complication that something that starts as a single cloud may in the course of its evolution break up into multiple pieces, so the visible lifetime and the dynamical lifetime may be different. Finally, by the efficiency $\\epsilon(\\rho)$ we mean the fraction of gas mass $M(>\\rho)$ that is converted into stars by a cloud over its lifetime $t_{\\rm cl}(\\rho)$. Again, we neglect the complication that clouds are not closed boxes, so $M(>\\rho)$ is likely to be time-dependent due to continuing accretion or mass loss. Ideally $\\epsilon(\\rho)$ should be computed using a Lagrangian definition of the cloud mass, i.e.\\ all the mass that reached density $\\rho$ or more at some point. However, this is generally not a direct observable.  Roughly speaking, the rate, efficiency, and lifetime are related by $\\dot{M}_* \\approx \\epsilon(\\rho) M(>\\rho)/t_{\\rm cloud}(\\rho)$. Consequently, one can define a rough time-averaged value for $\\sfrff$ in a cloud $\\left<\\sfrff(\\rho)\\right> \\approx \\epsilon(\\rho)/[t_{\\rm cl}(\\rho)/t_{\\rm ff}(\\rho)]$, i.e.\\ the mass fraction converted into stars divided by the cloud lifetime in free-fall times. One could also describe $\\left<\\sfrff\\right>$ as being the efficiency per free-fall time, since it measures a fraction of mass converted into stars. However, we refer to it as a rate because it is measured in amount per unit time, whereas efficiency in the literature most commonly refers to the total fraction of mass converted into stars, not the amount per unit time. Moreover, since the time-averaged definition is ambiguous anyway due to uncertainties in exactly what is meant by the efficiency and the lifetime, we will generally use the instantaneous value of $\\sfrff$, which is well defined and, as we show, directly observable.  The reason for making all these definitions explicit it to point out that observational constraints on $\\sfrff$ by themselves do not directly constraint $t_{\\rm cl}$ or $\\epsilon$, and vice versa. Numerous authors have used various observational techniques to estimate $t_{\\rm cl}$ in clusters and GMCs \\citep[e.g.][]{elmegreen00, hartmann01, tassis04, mouschovias05, tan06a, ballesterosparedes06, blitz06a}.  However, with the exception of \\S~\\ref{clustertime}, we will not discuss the question of cloud lifetime at all in this paper. When we refer to rapid versus slow star formation, what we mean is $\\sfrff\\sim 1$ or $\\ll 1$, not is $t_{\\rm cl}\\sim t_{\\rm ff}$ or $\\gg t_{\\rm ff}$. These questions are conceptually distinct. As an example, note that in the \\citet{goldreich74} picture of GMCs as collapsing spheres, the evolution time scale is always the cloud free-fall time, $t_{\\rm cl} \\sim \\tff$. However, the model still gives $\\sfrff=1$, and we would therefore describe it as a rapid star formation. In contrast, \\citet{ballesterosparedes06} argue that star formation lasts only one crossing time of a molecular cloud, so again $t_{\\rm cl}\\sim \\tff$, but that during this time only $\\sim 1\\%$ of the mass turns into stars. We would describe this as slow star formation, since $\\sfrff\\sim 0.01$, even though the cloud evolution time is similar to that in the \\citeauthor{goldreich74} model. The \\citeauthor{zuckerman74} argument rules out the \\citeauthor{goldreich74} free-fall collapse model for GMCs, but is consistent with the \\citeauthor{ballesterosparedes06} model. Thus, we emphasize that \\textit{in this paper we remain completely agnostic on the question of molecular cloud lifetime}.  In this paper we consider star formation in several classes of object. Infrared dark clouds (IRDCs) are regions of high extinction seen in absorption against the Galactic infrared background \\citep{egan98, carey00, simon06}. IRDCs are clearly associated with star formation, and in at least some cases IRDCs have massive protostars embedded within them \\citep{rathborne05}. Several authors have suggested \\citep[e.g.][]{menten05,tan05b,rathborne06a} that IRDCs are the progenitors of star clusters. Within IRDCs, at still higher densities, are dense molecular clumps. These objects may be observed in a variety of molecular transitions with high critical densities, and we consider two here: HCN(1-0) \\citep{gao04b,gao04a,wu05} and CS(5-4) \\citep{plume97,shirley03}. Molecular clumps seen in these two transitions are often associated with water masers and other signs of massive, clustered star formation. Our goal is to determine $\\sfrff$ for each of these increasingly dense gas tracers. We also determine this quantity for the Orion Nebula Cluster using a completely different method, which provides an independent check on our estimates.  The remainder of this paper proceeds as follows: in \\S~\\ref{sfrest} we use a variety of observations to derive $\\sfrff$ for our objects and construct a plot of $\\sfrff$ versus characteristic density to search for signs of a transition from slow to fast star formation. In \\S~\\ref{implications}, we compare our results to theoretical models for the regulation of star formation, and also point out some implications for the process of star cluster formation. Finally, in \\S~\\ref{conclusions} we summarize our conclusions, and suggest directions for future work. ",
        "conclusions": "\\label{conclusions}  We present observational evidence for two surprising conclusions, one shown quite strongly and the other more tentatively. First, and very clearly, \\textit{star formation in dense gas is slow.} The time required to convert all the gas into stars, the depletion time, is longer than the free-fall time by at least and order of magnitude, and probably closer to two orders of magnitude. There is no evidence that gas at densities up to $n_H \\sim 10^5$ cm$^{-3}$ resides primarily in objects that are undergoing collapse and rapid star formation. Second, and more tentatively, \\textit{the ratio of free-fall time to the depletion time is independent of the characteristic density of the star-forming object in question}. The data at present are uncertain, and are insensitive to changes of less than roughly an order of magnitude. Nonetheless, it is interesting that this ratio varies by no more than an order of magnitude from the typical density of GMCs to the typical density of HCN(1-0)-emitting gas, a range of nearly 3 decades in density.  These observations are a strong constraint for theories of star formation, and seem difficult to explain in the context of models in which there is a transition from slow, unbound star formation to rapid, bound star formation somewhere between the GMC scale and the protocluster scale. Slow star formation in cluster-forming gas also implies that clusters require many free-fall times to assemble, as recently argued by \\citet{tan06a} on other grounds. Models in which star formation takes place in virialized objects and is inhibited by strong magnetic fields are qualitatively consistent with the data, and models in which star formation is inhibited by turbulence are both qualitatively and quantitatively consistent with observations.  In the future it would be extremely useful to improve the data on which Figure \\ref{sfrffn} is based. One way to do this would be to perform detailed studies of other young clusters and obtain data comparable to that for the ONC. This would provide a method of estimating $\\sfrff$ that is independent of luminosity conversions and does not suffer from concerns about the completeness of galactic surveys. Another improvement in the data could come from an unbiased survey of CS(5-4) emission in the Milky Way or in another galaxy, which would allow us to replace the upper limit on $\\sfrff$ we derive here with an actual estimate. This would be particularly valuable because the upper limit on $\\sfrff$ from CS(5-4) is well above the estimate from HCN(1-0), even though the densities are not very different. Unbiased CS observations could likely bring down this point.  Another improvement would be extend the data to higher densities. To accomplish this will require observations either of external galaxies or relatively complete surveys of the Milky Way in molecular transitions that trace densities $\\gtsim 10^6$ cm$^{-3}$. Determining masses from the luminosities in these transitions will probably require high resolution follow-up observations of Galactic sources in optically thin isotopomers so that the luminosity may be correlated against a virial mass. While this is a significant observational challenge, such surveys might make it possible to identify a scale at which star formation transitions from slow to fast, a crucial datum in understanding the star formation process."
    },
    "0606/astro-ph0606088_arXiv.txt": {
        "abstract": "We analyze the three-year \\emph{WMAP} temperature anisotropy data seeking to confirm the power spectrum and likelihoods published by the \\emph{WMAP} team.  We apply five independent implementations of four algorithms to the power spectrum estimation and two implementations to the parameter estimation.  Our single most important result is that we broadly confirm the \\emph{WMAP} power spectrum and analysis. Still, we do find two small but potentially important discrepancies: On large angular scales there is a small power excess in the \\emph{WMAP} spectrum (5--10\\% at $\\ell \\lesssim 30$) primarily due to likelihood approximation issues between $13 \\le \\ell \\lesssim 30$. On small angular scales there is a systematic difference between the V- and W-band spectra (few percent at $\\ell \\gtrsim 300$). Recently, the latter discrepancy was explained by \\citet{huffenberger:2006} in terms of over-subtraction of unresolved point sources.  As far as the low-$\\ell$ bias is concerned, most parameters are affected by a few tenths of a sigma. The most important effect is seen in $n_{\\textrm{s}}$. For the combination of \\emph{WMAP}, Acbar and BOOMERanG, the significance of $n_{\\textrm{s}} \\ne 1$ drops from $\\sim2.7\\sigma$ to $\\sim2.3\\sigma$ when correcting for this bias.  We propose a few simple improvements to the low-$\\ell$ \\emph{WMAP} likelihood code, and introduce two important extensions to the Gibbs sampling method that allows for proper sampling of the low signal-to-noise regime. Finally, we make the products from the Gibbs sampling analysis publically available, thereby providing a fast and simple route to the exact likelihood without the need of expensive matrix inversions. ",
        "introduction": "Very recently, the \\emph{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}) team released the results from three years of observations of the cosmic microwave background \\citep{hinshaw:2006, page:2006, spergel:2006}. In addition to improving upon the successful first-year release \\citep{bennett:2003a}, this new data set includes the first full-sky polarization maps at millimeter wavelengths \\citep{page:2006}, and is destined to be of great importance for the CMB community for many years to come.  Given the dominant position of the \\emph{WMAP} experiment in our current understanding of cosmology, it is imperative that all of the results are verified by several external groups using independent techniques.  In this paper we begin these efforts by re-computing the angular power spectrum and likelihoods\\footnote{We use version v2p1 of the \\emph{WMAP} likelihood in this paper; http://lambda.gsfc.nasa.gov/}, arguably the two most important products of any CMB experiment. Similar work was carried out after the first-year \\emph{WMAP} release in 2003 by several groups (e.g., Slosar, Seljak and Makarov, 2004 -- the importance of foreground marginalization at large scales; O'Dwyer et al.\\ 2004 -- the importance of exact likelihood/posterior evaluation at large scales; Fosalba and Szapudi 2004 -- pixel space versus harmonic space computations).  These analyses significantly contributed to the improvements that were implemented for the three-year \\emph{WMAP} data release \\citep{hinshaw:2006}.  Our philosophy in the present paper is that of multiple redundancy and cross-checking: For both the angular power spectrum and the likelihoods (or posteriors) we apply several different algorithms and in some cases different implementations of the same algorithm. Additionally, the data were downloaded, prepared and analyzed by five independent groups, further reducing the risk of both programming errors, simple misunderstandings and algorithmic peculiarities. This increases our confidence in the final results.   One question that will be given particular attention is the issue of statistical and numerical procedure. Examples are, how should a high signal-to-noise but low-$\\ell$ likelihood be regularized at its small scale cut-off in order to avoid numerical artifacts? And in what $\\ell$-range does a pseudo-spectrum estimator provide an adequate approximation to the full posterior, and where should an exact approach (such as Gibbs sampling) be preferred?   A separate question is the problem of foreground modeling and subtraction. This will be considered in greater detail in a future publication, and for now we mainly adopt the approaches used by the \\emph{WMAP} team, with one or two notable exceptions: A power spectrum that does not rely on external information \\citep{saha:2006} is established and used as a cross-check on the results that are obtained from template-corrected maps, and we compare spectra computed for different frequencies in order to search for frequency-dependent signatures.  Another major part of the three-year \\emph{WMAP} release are measurements of the CMB polarization \\citep{page:2006}. This is a very complex data set, and requires a high degree of algorithmic sophistication for proper interpretation. The necessary extensions to the methods described in the present paper are currently being developed, and the results from the corresponding analyses will be reported upon as soon as they are completed. In the present paper, we focus on temperature information only.  The rest of the paper is organized as follows. In Section \\ref{sec:methods} we briefly list the methods we use, and in Section \\ref{sec:data} we describe the data. Then, we focus on the low-$\\ell$ part of the power spectrum in Section \\ref{sec:large_scale} and the high-$\\ell$ part in Section \\ref{sec:small_scale}. Cosmological parameters are considered in Section \\ref{sec:parameters}, before drawing conclusions in Section \\ref{sec:conclusions}. Algorithmic details are deferred to two Appendices. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have made an extensive re-analysis of the \\emph{WMAP} temperature anisotropy data. We have re-computed the power spectrum using six different codes, and the cosmological parameters with two different codes. Five different groups have contributed with separate computations. In total, the amount of cross-checks provided by this effort leads us to conclude that most important analysis issues concerning the three-year temperature power spectrum are well understood.  Our main conclusion from this work is that we confirm the \\emph{WMAP} temperature power spectrum on most scales. However, subtle discrepancies are found both at large and small scales, and these can be summarized as follows: \\begin{enumerate} \\item There is a small but significant bias at large angular scales (5--10\\% at $\\ell \\le 30$) in the \\emph{WMAP} power spectrum and published likelihood code. This is primarily due to statistical issues in the form of an inadequate likelihood approximation, and secondarily due to residual foregrounds. \\item There is a systematic bias between the V- and W-bands. This has recently been identified by \\citet{huffenberger:2006} to be at least partially caused by over-correction of unresolved point sources. However, even after correcting for this, a small discrepancy is present, and this could possibly indicate further systematics. A possible candidate worth considering could be uncorrected beam asymmetries. \\item The low-$\\ell$ likelihood bias has a noticable effect of some parameters. Most interestingly, it increases $n_{\\textrm{s}}$ by $0.4\\sigma$, reducing the nominal detection of $n_{\\textrm{s}} \\ne 1$ from $\\sim2.7$ to $\\sim2.3\\sigma$. Further, as reported by \\citet{huffenberger:2006}, correcting for the point source over-correction increases $n_{\\textrm{s}}$ by an additional $\\sim0.3\\sigma$, to a final significance of $n_{\\textrm{s}} \\ne 1$ of $\\sim2.0$. \\end{enumerate}  In addition to these cosmologically important results, we make a few algorithmic recommendations based on the work presented here regarding the \\emph{WMAP} analysis. First and foremost, an exact likelihood evaluation should be used at least up to $\\ell \\le 30$. Either direct evaluation or Gibbs sampling can be used for these scales, but of course, the negligible computational cost of the latter, when incorporated into an MCMC sampler, makes it the preferred choice. Second, we advocate using the more conservative Kp2e cut at very low l as a hedge against possible foreground contamination just outside the directly downgraded Kp2 mask. However, this has a negligible effect in terms of cosmological parameters.   The products from the Gibbs sampling analysis (with a uniform prior) are made publically available\\footnote{Available from http://www.astro.uio.no/$\\sim$hke/wmap3} in the form of three data sets. First, the most relevant for the applications presented in this paper is the collection of $4\\times 1000$ sky signal spectrum samples. These form the basis of the Blackwell-Rao estimator, and may be used for parameter estimation as demonstrated in this paper. A corresponding Fortran 90 module is also provided. Second, a set of 4000 ensemble averaged spectrum samples are presented. These may be used for visualization purposes of the power spectrum. Third, a set of sky map samples are provided for $\\ell \\le 50$. These may be used for analyses that require phase information, for instance studies of non-Gaussianity.     Thus, we conclude that although fast methods such as MASTER are very useful in the early stages of analyzing a new experiment, when fast turn around times are essential, the extensions to Gibbs sampling (as described in the Appendices) have now rendered an exact method quite tractable with currently available computing resources. Based upon our experience with these methods, we believe that Gibbs sampling can and will play a significant role in the future analysis of Planck data."
    },
    "0606/astro-ph0606331_arXiv.txt": {
        "abstract": "We have used the NICMOS NIC1 camera on the Hubble Space Telescope to obtain high angular resolution images of 52 ultracool dwarfs in the immediate Solar Neighbourhood. Nine systems are resolved as binary, with component separations from 1.5 and 15 AU. Based on current theoretical models and empirical bolometric corrections, all systems have components with similar luminosities, and, consequently, high mass ratios, $q > 0.8$. Limiting analysis to L dwarfs within 20 parsecs, the observed binary fraction is $12^{+7}_{-3}\\%$. Applying Bayesian analysis to our dataset, we derive a mass-ratio distribution that peaks strongly at unity. Modelling the semi-major axis distribution as a logarithmic Gaussian, the best fit is centered at log$a_0 = 0.8$ AU ($\\sim6.3$ AU), with a (logarithmic) width of $\\pm0.3$. The current data are consistent with an overall binary frequency of $\\sim24\\%$. ",
        "introduction": "Over the last four years we have been undertaking a census of the lower-mass constituents of the immediate Solar Neighborhood (Reid \\& Cruz, 2002 - Paper I), concentrating, in particular, on ultracool dwarfs (spectral types M7 and later) within 20 parsecs of the Sun (Cruz et al, 2003; Cruz et al, 2006; Reid et al, in prep.). As part of that survey, we have compiled an all-sky catalogue of 87 L dwarfs in 80 systems with formal distance estimates less than 20 parsecs. This sample offers an opportunity to investigate the statistical characteristics of the local L dwarf population.  Binarity is a key property of low-mass stars and brown dwarfs. Both the overall frequency of binary systems and the distribution of their properties (particularly mass ratios, $q$, and separations, $\\Delta$) have emerged as potential tests of various formation theories. Ultracool dwarfs have been the targets of numerous high-resolution imaging surveys, both using adaptive optics on ground-based telescopes and with the Hubble Space Telescope (HST). As summarised most recently by Burgasser et al (2006), the results of those surveys indicate an observed frequency of $\\sim15\\%$; this compares with an overall binary frequency of 30 to 40\\% for M dwarfs (Fischer \\& Marcy, 1992; Reid \\& Gizis, 1997) and 60 to 70\\% for G dwarfs (Duquennoy \\& Mayor, 1991). The overwhelming majority of ultracool binaries have small separations, $\\Delta < 15$ AU. The nearest ultracool dwarfs are therefore the prime targets for multiplicity surveys, since they provide the optimal resolution in linear units. Those systems also provide the best sensitivities for detecting very low luminosity companions, although there is growing evidence for a tendency towards mass ratios close to unity among ultracool binaries.  We have been using the NICMOS camera on HST to search for binary systems among the nearest L dwarfs. To date, we have acquired observations of 52 ultracool systems, of which 49 are classed as spectral type L and three as late-type M dwarfs; nine are resolved as binaries, including the L/T system 2MASSW J22521073-1730134 (Reid et al, 2006 - hereinafter RLBC06). High spatial-resolution images are available from the literature for a further 4 L-dwarf systems within 20 parsecs. We present our observations in the following section; the characteristics of the candidate binaries are discussed in \\S3; we consider the statistical properties of the full dataset, and the implications in \\S4; the final section summarises our results and conclusions.  \\section {Observations}  The nearby late-M and L dwarfs imaged in this program were observed as part of a Cycle 13 HST SNAPSHOT program. Most of these systems have spectroscopic distance estimates, and while all were placed within 20 parsecs at the outset of our observational program, a handful have revised distances that lie beyond that limit. All targets were observed with the NIC1 camera and the F110W and F170M filters using the same exposure sequences. The observations in both filters consist of a pair MULTIACCUM exposures, nodding 2.0 arcseconds between the two exposures. The total exposure times are 284 seconds at F110W and 896 seconds in the F170M filter. As discussed in \\S2.2, the combined data give limiting magnitudes of m$_{110} \\sim 21.9$ and m$_{170} \\sim 20.0$ magnitudes (on a Vegamag system) for isolated point sources. Late-type T dwarfs have (m$_{110}$ - m$_{170}$) colours of $\\sim1.3$ magnitudes; thus, the F110W data offer the highest sensitivity for the detection of faint companions to the targeted L dwarfs.  \\subsection {Identifying binary systems}  The NICMOS data were processed through the standard HST pipeline, and we have analysed the final mosaiced image using standard {\\sl IRAF} routines\\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.}. The NIC1 data have a plate-scale of 0.043 arcseconds pixel$^{-1}$, while the formal resolution of the HST data is 0.09 arcseconds with the F110W filter and 0.14 arcseconds with the F170M filter. Since the F110W images have both higher angular resolution and better sensitivity, we have concentrated on those data in our search for faint companions.  We have searched for potential binary companions using a variety of techniques. First, visual inspection of the images reveals a number of systems with obvious close companions. Second, we have used the {\\sl imexam} and {\\sl daophot} routines in {\\sl IRAF} to measure the point spread function (psf), searching for sources with broad full-width half-maxima or unusual profiles. Finally, as is evident from the images of the candidate binaries, unresolved point sources possess a strong Airy disk, with the flux level rising to $\\sim10\\%$ of the peak flux at radial separations of 0.205 arcseconds in the F110W data and 0.24 arcseconds at F170M. This obviously affects the potential detection of companions at those radii.  We have analysed these data using the same techniques outlined in our discussion of the L/T binary, 2M2252-1730 (RLBC06). One of the sources observed in our program is 2M0825+2115. This L7.5 dwarf has previous HST observations with WFPC2 (Reid et al, 2001), which demonstrate that the image profile matches a single point source at optical wavelengths. We have therefore taken this object as the psf template for the NIC1 observations\\footnote{In RLBC06 we used our data for 2M0045+1634 as the template; this source is unresolved by NICMOS, and subtracting the 2M0825+2115 data shows no evidence for any significant residuals.}. We have used the {\\sl imshift} IRAF routine to align the 2M0825 image with each of the other L dwarf targets; scaled the reference data to match the peak flux; and subtracted the 2M0825 data, leaving a `cleaned' image of the environs of the target. There are imperfections in most subtractions, since the NICMOS psf profile changes on relatively short timescales, but none of the low-level residuals have profiles resembling a very faint companion.  Based on our analysis, forty-three ultracool dwarfs show no significant evidence for binarity. Most have psf profiles with full-width half-maxima (FWHM) of 2.3 to 2.4 pixels (0.099 to 0.103 arcseconds). Three L dwarfs have slightly broader profiles: 2M1507-1627 and 2M1936-5502, with FWHM = 2.47 pixels (0.106 arcseconds); and 2M0036+1820, with FWHM = 2.56 (0.110 arcseconds). Subtracting the 2M0825 template psf show no evidence for the presence a secondary component, and the broader profiles are probably an instrumental effect. Pertinent data for the unresolved ultracool dwarfs are given in Table 1. In most cases, the distance estimates rest on the Cruz et al (2003) spectral-type/M$_J$ relation, and therefore have uncertainties of $\\sim15\\%$.  The remaining nine dwarfs observed in this program are identified as probable binaries. Figure 1 presents NICMOS F110W images of seven sources: higher resolution ACS images of the LHS 102BC (GJ 1001BC) system are discussed by Golimowski et al (2004b), while images of the 2M2252 system are presented by RLBC06. Data for all the candidate binaries are given in Table 2. We discuss these systems in more detail in \\S3.  \\subsection {Photometry}  We have utilised two techniques to determine instrumental magnitudes from the NICMOS images. First, we used the {\\sl phot} routine in {\\sl daophot} to determine aperture photometry for well-isolated sources (the `single' objects and candidate binaries with separations exceeding 0.5 arcseconds). In these cases, we adopted an aperture size of radius 9 pixels (0.36 arcseconds). For the close binaries, we use a smaller aperture size, correcting to 9-pixel photometry using aperture corrections derived from measurements of 2M0825+2115. In the latter cases, we have also estimated the relative magnitudes of the two sources by measuring the peak flux of each component using the IRAF {\\sl imexam} profile-fitting routine; combining these data with aperture photometry of both components gives the magnitudes of the two components.  Our aperture photometry is tied to the Vega magnitude scale using the standard HST flux calibration and flux zeropoints of 1786 Janskys and 946 Janskys at F110W and F170M, respectively. We have also used the results given by Schultz et al (2005) to apply appropriate corrections to adjust our photometry to infinite aperture. The resultant magnitudes, m$_{110}$ and m$_{170}$, are listed in Tables 1 and 2. All of these dwarfs have JHK$_S$ photometry from the 2MASS database; indeed, we have used this photometry to derive color terms between the F110W/F1170M and J/H magnitude systems (see RLBC06). The 2MASS data are also given in Tables 1 and 2, and we have used the relative magnitudes in the HST systems to estimate J and H magnitudes for the individual components of the candidate binary systems.  \\section {Ultracool binaries}  \\subsection {The present sample}  We have identified nine ultracool dwarfs as probable binaries. Table 2 gives the observed properties for these systems, and Table 3 lists the intrinsic properties inferred for the individual components. We have computed luminosities for each component by applying J-band bolometric corrections as a function of spectral type. Figure 2 shows the basis for our calibration, plotting data for M, L and T dwarfs from the analysis of Golimowski et al (2004a). As discussed further in the following section, we estimate masses from the M$_{bol}$ estimated using the models computed by Burrows et al (1997).  The probable (or confirmed) binary systems are as follows: \\begin{description}  \\item[LHS 102BC/GJ 1001BC:] The companion to the nearby M3.5 dwarf LHS 102 was discovered originally by Goldman et al (1999), and the L dwarf was itself revealed as double in NICMOS observations obtained as part of a snapshot survey of stars within 10 parsecs of the Sun (Golimowski et al, 2004b). The system is barely resolved with NICMOS, but the binary status was confirmed through optical observations with the Advanced Camera for Surveys. This system is a classical ultracool binary, with near equal-magnitude components. Golimowski et al point out that the L dwarf properties appear inconsistent with the trigonometric parallax of $104\\pm11$ milliarcseconds cited by van Altena et al (1995) for the primary, and they suggest that a distance closer to 15 parsecs is more plausible. New trigonometric parallax measurements are currently being undertaken by the CTIOPI consortium, and we refer the interested reader to Golimowski et al (2004b) for further discussion. It is clear that the system lies within 20 parsecs and, for present purposes, we adopt the larger distance in computing the intrinsic parameters listed in Table 3.  \\item[2M0025+4759:] Originally classed as L5 based on near-infrared data, optical spectra indicate a type of L4, consistent with a distance of 23 parsecs for a single dwarf. The HST observations resolve the system into two near-equal luminosity components, implying a distance of 31$\\pm$7 parsecs. The system has also been resolved in ground-based observations with the Keck Laser Guide Star AO system (Liu, 2006, priv. comm.). 2M0025+4759 lies only $\\sim3.5$ arcminutes from HD 2057 (G171-58/G217-47), a solar metallicity F8 dwarf (Carney et al, 1994) with an Hipparcos parallax placing it at a distance of 42$\\pm$2 parsecs. HD 2057 itself is likely a close binary (Latham et al, 2002; Balega et al, 2004). Based on matching Str{\\\"o}mgren photometry against isochrones, Nordstr{\\\"o}m et al (2004) estimate an age of $\\sim1.1$ Gyrs for HD 2057, with an upper limit of 3.6 Gyrs and no specified lower limit. \\\\ S. Schmidt (2006, priv. comm.) has combined 2MASS near-infrared and POSS II I-band images (5.9 year baseline) to derive proper motions of ($\\mu_\\alpha, \\mu_\\delta) = (+0.312\\pm0.034, -0.009\\pm0.044)$ arcsec yr$^{-1}$ for the L dwarf. Those data are in reasonable agreement with the Hipparcos astrometry of HD 2057, ($\\mu_\\alpha, \\mu_\\delta) = (+0.274\\pm0.001, +0.011\\pm0.001$) arcsec yr$^{-1}$. If 2M0025+4759 {\\sl is} a wide companion of HD 2057, it lies at a separation of $\\sim8,800$ AU. This would make it the widest known binary with an ultracool component, but the separation lies within the span of other binaries of comparable total mass (see Figure 8 in Reid \\& Walkowicz, 2006). Finally, strong lithium absorption is evident in the combined spectrum, indicating that both components are brown dwarfs with $M < 0.065 M_\\odot$ (see Schmidt \\& Cruz, 2006, for further discussion). Matched against either the Burrows et al (1997) or Baraffe et al (1998) models, the absence of significant lithium depletion implies an age less than 1 Gyr, broadly consistent with the age estimated for HD 2057.  \\item[2M0147-4954:] This dwarf was targeted for observation based on a preliminary spectral type of L0. We have since revised the classification to M8, pushing the system beyond the 20-parsec limit even as a single dwarf; as a binary, with likely spectral types of M8 and L2, we estimate the distance as $\\sim33$ parsecs. The flux ratio of the components is similar to that in the 2M0429-3123 system, and the Burrows et al (1997) models indicate a similar mass ratio.  \\item[2M0429-3123:] This M7.5 dwarf was resolved originally by Siegler et al (2005) using adaptive optics on the ESO VLT. We estimate a distance of 11.5 parsecs based on the J magnitude of the primary (J$_1$) and the spectral-type/M$_J$ calibration from Cruz et al (2003). There are no indications that the system is particularly young. Siegler et al estimate a spectral type of L1 for the secondary.  \\item[2M0700+3157:] This is the only L dwarf in the present sample with a direct trigonometric parallax measurement (Thorstensen \\& Kirkpatrick, 2003). With M$_J \\sim 14$, the secondary is probably spectral type $\\sim$L6.  \\item[2M0915+0422:] This ultracool dwarf has been identified independently as a binary system through ground-based AO observations (M. Liu, 2006, priv. comm.). Like LHS 102BC, the two components are almost equal in magnitude and therefore have identical masses. Both components are likely to be brown dwarfs.  \\item[2M1707-0558:] This system was first resolved {\\sl via} ground-based observations with Spex on the IRTF by Burgasser et al (2004).  Originally classed as spectral type L1 based on the combined optical spectrum, McElwain \\& Burgasser (2006) have obtained resolved NIR spectroscopy of this system and derive spectral typee of M9/L0 and $\\sim$L3. Their observations also confirm that the components have common proper motion.  \\item[2M2152+0937:] This is another equal-luminosity/equal-mass ultracool binary. As with 2M0147-4954, the identification of this dwarf as a binary system removes it from the 20-parsec sample.  \\item[2M2252-1730:] This is one of the handful of L/T binary systems currently known. As discussed in RLBC06, the secondary is noticeably fainter with respect to the primary in the F170M filter than in the F110W passband. Infrared spectroscopy confirms that this is due to the presence of significant methane absorption. Both components are likely to be of substellar mass.  \\end{description}  In most of the observations, the targeted L dwarf is the only object visible in the NICMOS image. There are seven cases, however, where other point sources are visible in the $\\sim10 \\times 10$ arcsecond NIC1 field of view. With one exception, these sources lie more than 2 arcseconds from the L dwarf and are either bright (J$<16$) and detected in ground-based observations, or extremely faint ($J>19$). These candidate wide companions all have colours and magnitudes that are inconsistent with very low-mass ultracool dwarfs.  The exception is 2M1705-0516. As Figure 1 shows, this L0.5 dwarf has a faint candidate companion at a separation of 1.36 arcseconds and PA=-5$^o$. The object is unresolved (FWHM$\\sim0.10$ arcseconds at F110W), and has m$_{110}=18.00$ and m$_{170} = 16.76$, corresponding to J$\\sim17.4$, (J-H)$\\sim0.7$. With a galactic latitude of $b \\sim +20^o$, this is unlikely to be a reddened background source. The (J-H) colours are consistent with either a mid-type M dwarf at a distance of 1-2 kpc or an early-type T dwarf which, at a distance of 19.5 parsecs, would have M$_J \\sim 16.0$. At present, we cannot distinguish between these two possibilities. Follow-up imaging at a later epoch will confirm whether the candidate companion shares the proper motion of the putative primary. For current purposes, we treat 2M1705 as a single ultracool dwarf.  \\subsection {Observations of additional systems}  Four L dwarfs from the 20-parsec sample have been observed at high spatial resolution in the course of other binary search programs. These dwarfs are identified in Table 4, where we list relevant data.  \\begin{description} \\item[Denis-P J0205:] One of the three L dwarfs discovered by the DENIS brown dwarf mini-survey (Delfosse et al, 1997), Denis-P J0205 was identified as a binary by Koerner et al (1999) based on K-band imaging with the Keck telescope. Bouy et al (2005) have recently suggested that the brighter component is itself double, and the system consists of two late L dwarfs and a T dwarf.  \\item[SDSS 0423-0414:] Originally identified from SDSS observations, this dwarf was classed as type T0 based on its near-infrared spectrum (Geballe et al, 2002). Hawley et al (2002) and Cruz et al (2003) obtained optical spectra in the course of their surveys, and both class the dwarf as type L7.5. Burgasser et al (2005) have resolved the discrepancy; NICMOS observations show that SDSS 0423 is an L/T binary, with properties similar to 2M2252-1730.  \\item[2M0746+2000:] Lying at a distance of $\\sim12$ parsecs, this is the brightest L dwarf currently known. Reid et al (2000) originally noted that the dwarf appeared to be overluminous, and high-resolution optical imaging with Wide-Field Camera 2 on HST confirmed that the system is a near equal-magnitude binary (Reid et al, 2001).  \\item[Kelu 1:] The first isolated L dwarf to be discovered (Ruiz, Leggett \\& Allard, 1997), trigonometric parallax measurements indicated that Kelu 1 was overluminous, but high-resolution follow-up observations with HST showed no evidence for binarity (Mart\\'{\\i}n et al, 1999). Those initial observations, however, suffered from bad timing: Liu \\& Leggett (2005) have recently resolved the system using ground-based adaptive optics observations, as orbital motion has separated the components. Those observations also resolve a long-standing conundrum: Kelu 1 exhibits weak lithium absorption, implying that, as a single star, it had been caught just as lithium was being depleted; as a binary, these observations are explained due to dilution of the (full-strength) lithium line in component B by continuum from the higher-mass component A. The inferred age for the system is $<800$ Myrs. \\end{description}  All four of these L dwarfs are multiple systems, but this high proportion is not entirely surprising: three are among the brightest L dwarfs known (in apparent magnitude), while the fourth, SDSS 0423, has an unusual spectrum; those properties are correlated directly with binarity, so it is not surprising that these systems were targeted through HST observations. We have therefore not included these systems in our analysis of binary frequency among L dwarfs. A further 31 L dwarfs with formal distances less than 20 parsecs currently lack high resolution imaging data.  \\section {Discussion}  \\subsection {Companion detection limits}  Our main goal is the detection of low luminosity companions to these ultracool dwarfs. At small separations, the detection limit is set by the psf of the primary; at larger separations, the limit primarily reflects the signal-to-noise of the observations. Figure 3 presents azimuthally-averaged radial profiles in both the F110W and F170M filters for a representative unresolved ultracool dwarf, 2M0523-1403 (m$_{110}=13.66$, m$_{170}=12.32$; J=13.12, H=12.22); the upper panels show the linear profiles (in counts sec$^{-1}$ pix$^{-1}$), and the lower panels plot more extended profiles in a logarithmic scale.  Figure 3 illustrates both the broader psf in the F170M passband and the higher sensitivity of the F110W imaging. The faintest isolated sources detected in our data have peak count-rates of $\\sim0.12$ counts sec$^{-1}$ pix$^{-1}$ in F110W and $\\sim0.15$ counts sec$^{-1}$ pix$^{-1}$ in F170M, corresponding to the dashed lines plotted in the lower panels of Figure 3. These peak count-rates correspond to magnitudes of m$_{110}=21.9$ (J$\\sim$21.5) and m$_{170}=20.0$ (H$\\sim$20.0) for point sources. Extremely cool brown dwarfs are expected to have neutral colours in (J-H), and are therefore easier to detect in the F110W passband.  The effective sky background increases, and the ability to detect a companion decreases correspondingly, within $\\sim$1.25 arcsecond of the central star. As noted above, the typical full-width half-maximum of the F110W psf is $\\sim$0.10 arcseconds. Only equal-magnitude binaries, such as LHS 102BC, are detectable at such small separation. LHS 102BC is clearly elongated in our images, but effectively marks the small-separation limit of our survey.  \\subsection {The L dwarf binary frequency}  Previous analyses of binarity in ultracool dwarfs (Reid et al, 2001; Gizis et al, 2003; Bouy et al, 2003; Burgasser et al, 2003; Siegler et al, 2005) are based on magnitude-limited samples. We can cast the present analysis in terms of a volume-limited sample, even though we have only observed a subset of that sample. The formal distance limit of the parent ultracool dwarf survey is 20 parsecs; five L dwarf binaries and 41 unresolved objects (including 2M1705-0516) have formal distance estimates within this limit, a binary fraction of $10.9_{-3}^{+6}$\\%, where the uncertainties are derived using the formalism outlined by Burgasser et al (2003).  In most cases, however, the distances listed in Tables 1 and 2 are based on spectroscopic parallaxes; those estimates should be corrected for Malmquist bias in statistical analyses. The M$_J$/spectral-type relation from Paper V has a dispersion of $\\sim0.35$ magnitudes; this corresponds to an absolute magnitude correction of $\\Delta$M$_J = -0.12$ magnitudes, effectively reducing the limit in apparent distance to 19 parsecs. Five binaries and 38 single stars fall within this limit, giving an observed binary fraction of $11.6_{-3}^{+7}$\\% for the present sample. This result is formally lower than previous estimates (see Burgasser et al, 2006), although consistent within the (substantial) uncertainties.  \\subsection {Masses and mass ratios for the L dwarf binaries}  Our observations measure luminosity ratios for ultracool binaries. Evolutionary effects complicate calculation of the corresponding mass ratios, since brown dwarfs cool and fade at rates that increase with decreasing mass. This is illustrated in Figure 4, where we plot (M$_{bol}$, mass) isochrones for low-mass star/brown dwarf models by Burrows et al (1997) models and by Chabrier et al (2000). Even though the latter \"dusty\" models (which do not extend below $\\sim900$K or $\\sim$T6) are a poor match to the colors of late-type L and T dwarfs, the predicted bolometric magnitudes are in reasonable agreement with the Burrows et al dataset. The labeled horizontal lines mark the locations of the binary components listed in Tables 2 and 3, components with spectral types ranging from M7.5 for 2M0429A to $\\sim$T2 for 2M2252B.  Figure 4 clearly shows the strong age dependence of brown dwarf masses. Two further points can be made regarding mass ratios of L dwarf binaries. First, old systems must have high mass ratios; thus, a 5-Gyr-old system with an M8 primary (comparable to 2M0429A) and a T2 secondary (like 2M2252B) has components of mass of 0.07 and 0.085$M_\\odot$, and a mass ratio of $q \\sim 0.82$. Second, at younger ages ($\\tau \\le 1$ Gyrs), the isochrones have similar slopes through the L-dwarf r\\'egime in the (M$_{bol}$, mass) plane; this implies that $q$ decreases with decreasing age. Thus, the same M8/T2 system has $q \\sim 0.5$ at $\\tau = 1$ Gyr (0.04 \\& 0.08 $M_\\odot$), and $q \\sim 0.38$ at $\\tau = 0.3$ Gyrs (0.06 and 0.023 $M_\\odot$).  All binaries in the current sample are field dwarfs. Consequently, the only direct means of constraining mass/age is the presence of lithium absorption.  2M0700+3157 is the only L dwarf with detected lithium absorption, indicating that both components have a mass below 0.065$M_\\odot$ and an age less than $\\sim1$ Gyr. It is likely that the absence of other lithium detections reflects the relatively low signal-to-noise of the optical spectra; nonetheless, the net result is that we have no direct age estimates for almost all the sample.  Under these circumstances, we must use models to estimate the likely age distribution. Allen et al (2005) have undertaken this type of calculation, basing their analysis on the Burrows et al (1997) evolutionary models. The results depend on the star formation history adopted for the Galactic disk and, to a lesser extent, the form adopted for the underlying mass function. Figure 5 shows the predicted cumulative age distributions for L0, L5 and L6-8 dwarfs for a constant star formation rate and a power-law mass function, $\\Psi(M) \\propto M^{-1}$. The three distributions are very similar at young ages, with $\\sim30\\%$ of each sample younger than $\\sim1$ Gyrs. The curves diverge at larger ages, with an increasing fraction of older dwarfs at earlier spectral types. Thus, half of the local L0 dwarfs are expected to be younger than $\\sim3$ Gyrs, while the 50$^{th}$ percentile mark is reached at age $\\sim1.7$ Gyrs for L6=L8 dwarfs.  As a qualitative guide to the mass ratios of the binaries consider here, we have used the Burrows et al (1997) mass-luminosity relations to estimate component masses for ages of $\\tau = 1$ and 3 Gyrs. The exceptions are 2M0700 and Kelu 1, where the presence of strong lithium absorption indicates ages less than 1 Gyr. Those data, and the corresponding mass ratios, are listed in Tables 3 and 4. Under these assumptions, all systems have high mass ratios, $q > 0.6$ for $\\tau = 1$ Gyr and $q > 0.85$ for $\\tau = 3$ Gyrs.  \\subsection {The distribution of mass ratios and separations}  All of the L dwarf binaries listed in Tables 2 and 4 have components with relatively high flux ratios. 2M1707AB exhibits the largest magnitude difference, with $\\Delta$J = 1.75 magnitudes, and most systems have $\\Delta$J$<0.4$ magnitudes. How do these flux ratios compare with the detection limits of the NICMOS observations?  Figure 6 plots the F110W psf in magnitudes, scaling the measurements relative to the peak brightness. We mark the location of the ultracool companions listed in Tables 2 and 4. The dotted lines mark the effective detection limits spanned by the present set of NICMOS observations. It is clear that all the detected companions are well above those limits. Moreover, as found in previous binary surveys, all of the detected companions lie at relatively small separations.  To set a rough mass scale for these comparisons, we have used the Burrows et al (1997) models to predict flux ratios for a 0.07$M_\\odot$ primary and companions with $0.2 < q < 0.9$ and $\\tau = 0.5, 1$ and 5 Gyrs. We choose this value for the primary mass since Allen et al (2005) estimate average masses of $\\langle M \\rangle = 0.074M_\\odot$ for field L0 dwarfs, $0.067M_\\odot$ at L5 and $0.063M_\\odot$ for L6-L8. High mass (long-lived) brown dwarfs (0.06 to 0.075 $M_\\odot$) contribute disproportionately to the local L dwarf population. The resulting flux ratios, shown in Figure 6, suggest that we ought to be able to detect systems with mass ratios as low as $q \\sim 0.2$ with the present set of NICMOS observations. As a guide, a $q=0.2$ system comprises an M8 primary (T$_{eff}\\sim2440$K, M$_{bol}=13.45$) and $\\sim$T7 secondary (T$_{eff}=740$K, M$_{bol}=18.4$) at age 0.5 Gyrs, and  a $\\sim$T1 primary (1230K, 16.8) and room-temperature Y-type secondary (350K, 21.8) at age 5 Gyrs\\footnote{We assume BC$_J$=2.0 for dwarfs later than T8.}.  We can quantify our estimates of the underlying mass-ratio and separation distributions through the Bayesian analysis techniques described by Allen et al (2005). Given a particular model for the companion distribution, we can use a disparate set of observations to derive the {\\sl posterior distribution}, $P(\\theta | D)$, the probability of the model given the data. By Bayes rule, the posterior distribution is the convolution of the likelihood distribution (the likelihood of the data given the model) and the prior distribution (the initial probability of the model).  Allen et al's (2005) analysis centres on the substellar mass function, but the same techniques can be used to probe the mass function of binary companions. Initial results are included in Burgasser et al's (2006) review, analysing data from previous binary surveys (Koerner et al, 1999; Reid et al, 2001; Gizis et al, 2003; Close et al, 2003; Bouy et al, 2003 and Siegler et al, 2005). In those calculations, the semi-major axis distribution is characterised as a Gaussian in $\\log{a}$, with central value $a_0$ and half-width $\\sigma_a$, while the mass-ratio distribution is defined as a power-law, index $\\gamma$, for a binary fraction $N$. We follow the same approach here, adopting the posterior distribution from the analysis cited in Burgasser et al. (2006) as the prior distribution for our analysis.  Figure 7 shows the probability distributions derived for each parameter. Expressing $a$ in astronomical units, the best-fit values are $\\log{a_0} = 0.8^{+0.06}_{-0.12}$, $\\sigma_a = 0.28\\pm0.4$, $\\gamma = 3.6\\pm1$ and $N = 24^{+6}_{-2}\\%$. This analysis reinforces the results outlined in Burgasser et al. The best-fit power-law index, $\\gamma$, indicates a steep mass function for L dwarf companions, implying a mass-ratio distribution with a strong preference for equal-mass systems. The semi-major axis distribution peaks at $\\sim6$ AU, with the model predicting very few systems at separations either beyond $\\sim20$ AU or within $\\sim1$ AU.  Figure 8 compares the mass-ratio distribution and semi-major axis distribution derived from the present analysis (both the data and the best-fit model) against the results derived by Fischer \\& Marcy (1992) for M dwarfs and Duquennoy \\& Mayor (1991) for G dwarfs. Clearly, in both cases, the G-dwarf distributions are radically different, while the M dwarf results are closer to our L dwarf analysis. At small separations, imaging data, even with HST, set weaker constraints, leading to the extended tail in the best-fit probability distribution of $a_0$. Maxted \\& Jefferies (2005) have argued that significant numbers of spectroscopic binaries remain hidden in L dwarf samples, although their hypothesis currently lacks substantial observational support. With that caveat, our analysis indicates an overall binary fraction of $\\sim24\\%$, continuing the trend of decreasng binary frequency with decreasing mass. ",
        "conclusions": "We have presented analysis of high spatial-resolution NICMOS images of 52 ultracool dwarfs in the immediate Solar Neighbourhood. Nine systems are resolved as binary, with component separations from 0.1 to 1.0 arcseconds corresponding to linear separations between 1.5 and 15 AU. Based on current theoretical models and empirical bolometric corrections, all systems have high mass ratios; none includes components with magnitude differences greater than 1.5 magnitudes at J. This is consistent with previous surveys for ultracool binaries in the general field. The observed binary frequency, limiting analysis to stars with Malmquist-corrected distances within 20 parsecs, is of $12_{-3}^{+7}$\\%.  Following Allen et al (2005) and Burgasser et al (2006), we have used Bayesian analysis to quantify these results. We derive a mass-ratio distribution that peaks strongly at unity, and matching the semi-major axis distribution with a logarithmic Gaussian gives a best-fit value of $\\log{a_0} = 0.8$, or $\\sim6.3$ AU. Our analysis indicates that the current data are consistent with an overall L-dwarf binary frequency of $\\sim24\\%$.  Acknowledgements: The observations described in this paper are associated with HST program \\#10143, and those data were obtained {\\sl via} the Hubble Space Telescope data archive facilities maintained at the Space Telescope Science Institute. Support for this research was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. KLC is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0401418. \\\\ This publication makes use of data from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center, and funded by the National Aeronautics and Space Administration and the National Science Foundation. 2MASS data were obtained from the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
    },
    "0606/astro-ph0606107_arXiv.txt": {
        "abstract": "Wide-field, very long baseline interferometry (VLBI) observations of the nearby starburst galaxy NGC 253, obtained with the Australian Long Baseline Array (LBA), have produced a 2.3 GHz image with a maximum angular resolution of 15 mas (0.3 pc). Six sources were detected, all corresponding to sources identified in higher frequency ($>5$ GHz) VLA images. One of the sources, supernova remnant 5.48$-$43.3, is resolved into a shell-like structure approximately 90 mas (1.7 pc) in diameter. From these data and data from the literature, the spectra of 20 compact radio sources in NGC 253 were modelled and found to be consistent with free-free absorbed power laws. Broadly, the free-free opacity is highest toward the nucleus but varies significantly throughout the nuclear region ($\\tau_0\\sim 1->20$), implying that the overall structure of the ionised medium is clumpy. Of the 20 sources, nine have flat intrinsic spectra associated with thermal radio emission and the remaining 11 have steep intrinsic spectra, associated with synchrotron emission from supernova remnants. A supernova rate upper limit of 2.4 yr$^{-1}$ is determined for the inner 320 pc region of the galaxy at the 95\\% confidence level, based on the lack of detection of new sources in observations spanning almost 17 years and a simple model for the evolution of supernova remnants. A supernova rate of $>0.14 (v/ 10^{4})$ yr$^{-1}$ is implied from estimates of supernova remnant source counts, sizes and expansion rates, where $v$ is the radial expansion velocity of the supernova remnant in km s$^{-1}$. A star formation rate of $3.4 (v/10^{4}) < SFR(M\\geq5M_{\\Sun})<59$ M$_{\\Sun}$ yr$^{-1}$ has been estimated directly from the supernova rate limits and is of the same order of magnitude as rates determined from integrated FIR and radio luminosities. ",
        "introduction": "\\label{sec:introduction}  Starburst galaxies are those galaxies containing regions (usually associated with the nucleus) currently undergoing high rates of star formation, much higher than seen in galaxies like our own and such that a significant fraction of the available gas for star formation will be exhausted within a single dynamical timescale \\citep{leh96}.  A starburst is triggered by the rapid accumulation of gas in a small volume, causing large numbers of massive stars to form.  The accumulation of gas can be caused by interactions between galaxies, for example the merger of gas rich spiral galaxies can result in the accumulation of 60\\% of the gas within the inner 100 pc of the merger product \\citep{bar96}.  Gas can be funnelled into the nuclear region of a galaxy by dynamical processes associated with bar instabilities \\citep{com00}.  Once formed, the starburst hosts large numbers of massive stars that ionise their environment and evolve quickly to form supernovae \\citep{eng98}.  The supernova remnants can drive activity in the nuclear region and can produce strong winds out of the disk of the galaxy \\citep{doa93}.  Starbursts are therefore the result of processes acting on a wide range of scales and themselves drive energetic activity on a wide range of scales.  By direct observation of the supernova remnants in starburst regions at radio wavelengths (to mitigate against severe extinction at optical wavelengths and provide very high angular resolution) it is possible to investigate the individual remnants and use them to reconstruct the supernova and star formation history of the starburst.  Furthermore, it is possible to use radio observations to investigate the ionised gaseous environment of the starburst.  Such observations provide a link between large-scale dynamical effects in the galaxies, activity in the star forming region itself, and the energetic phenomena in turn driven by the starburst.  NGC 253 is a prominent nearly edge-on starburst galaxy in the Southern Hemisphere ($\\delta\\sim-25^{\\circ}$), its proximity making it an ideal target for high spatial resolution observations. Previous studies of NGC 253 have assumed a distance of 2.5 Mpc \\citep{tur85}, based on the estimate of \\citet{dev78}. Both this and the revised estimate of 2.58$\\pm$0.7 Mpc \\citep{puc88} were derived from the Tully-Fisher relation. More recently \\citet{kar03} determined the magnitude of the tip of the red giant branch from Hubble Space Telescope/WFPC2 images and arrived at a more reliable distance estimate of 3.94$\\pm$0.5 Mpc. In this paper this new distance estimate is adopted.  In many respects NGC 253 is a twin of it's Northern Hemisphere counterpart M82. Both are of similar size, are nearly edge-on, barred, and have similar infrared spectra and luminosities. \\cite{eng98} confirmed the existence of a large-scale bar $\\sim$7 kpc in diameter crossing a circumnuclear ring of similar size in a de-reddened K-band image of NGC 253. Using a combination of optical, infrared and millimetre observations, \\citet{for00} concluded that a ring of young star clusters $\\sim$80 pc in diameter defined the inner edge of a cold gas torus centred on the assumed nucleus. The torus is fed by gas channelled in by the large-scale bar and deposited at the Inner Lindblad Resonance \\citep{arn95}.  Australia Telescope Compact Array (ATCA) images of para-NH$_{3}$ (1,1) and (2,2) and ortho-NH$_{3}$ (3,3) and (6,6) were used to probe the temperature of the nuclear regions and highlight some of the inner structure of the torus where the temperatures are greatest \\citep{ott05}. The full extent of the torus is seen more clearly in CO and [\\ion{C}{1}] emission line maps which trace cool gas and dust \\citep{isr95}, these observations showing that the torus is clumpy and appears to have a diameter of $\\sim900$ pc.  Associated with the nuclear region of the galaxy is gas producing strong optical emission lines such as H$\\alpha$, [\\ion{S}{2}], and [\\ion{O}{3}] \\citep{for00,car94}.  The environment is heavily ionised and the emission line region has a complex morphology \\citep{for00}.  Also associated with the nuclear region, and having its origin in the inner starburst region, is a powerful wind that has formed a bipolar nuclear outflow cone perpendicular to the disk of the galaxy. The outflow has been investigated by \\citet{hec90} and \\citet{sch92} using H$\\alpha$ spectral line measurements but is seen more clearly in X-ray observations by Chandra \\citep{wea02, str02}.  \\ion{H}{1} plumes are also seen bordering the H$\\alpha$ and X-ray halo emission and are also believed to be associated with the inner starburst region and active star formation in the disk \\citep{boo05}. It is thought that the wind is matter blown from the central starburst region by the ensemble of massive stars and supernova remnants within the inner 80 pc.  Figure \\ref{fig:figmultiwav} presents images of the galaxy at various wavelengths, on different spatial scales, to illustrate of the relationships between the galaxy bar (IR), the molecular torus (CO), the inner starburst region (radio), and the wind eminating from that region (H$\\alpha$ and x-ray).  The supernova remnants that are a primary subject of this paper are concentrated within the inner 80 pc of the galaxy but are also distributed widely outside the inner starburst region.  At radio wavelengths, \\citet[hereafter U\\&A97]{ulv97} imaged NGC 253 with the VLA at frequencies between 1.5 GHz and 23 GHz over a period of eight years and identified 64 individual compact sources within 320 pc of the nucleus. The identification of sources at lower frequencies ($<10$ GHz) was hindered by confusion and the diffuse emission of the galaxy. Spectral indices, with an error less than 0.4, could therefore only be determined for 17 compact sources, mainly those isolated sources outside the nuclear region. In the absence of any new sources or any significant fading of existing sources over the course of their observations, U\\&A97 derived a supernova rate of no more than 0.3 yr$^{-1}$. At 23 GHz the assumed nucleus (5.78-39.4) is the strongest source and is unresolved. Approximately half of the 17 detected compact sources with measured spectral indices were shown by U\\&A97 to be associated with \\ion{H}{2} regions. The remainder are most likely to be supernova remnants.  In the first attempt to resolve the inner starburst region of NGC 253 at low radio frequencies, \\citet{tin04} used the Australian Long Baseline Array (LBA) in the first multi-baseline VLBI observation of NGC 253. At 1.4 GHz the observation approximately matched the angular resolution of the VLA at 23 GHz, thus overcoming confusion problems. Two sources were detected, neither of which corresponded to the assumed nucleus. \\citet{tin04} showed that the radio emission from compact sources in NGC 253 is partially absorbed by ionised gas with free-free optical depths ranging between $\\tau_{0}=2.5$ and $\\tau_{0}>8$ ($\\tau_{0}$ denotes the free-free optical depth at 1 GHz).  Further evidence for free-free absorption in NGC 253 comes from measurements of the nuclear starburst region between 300 MHz and 1.5 GHz by \\citet{car96}, showing a dramatic spectral flattening at low frequencies. The modelling of radio-recombination line emission from the nuclear region of NGC 253 suggests that the lines originate from a gas with an electron density of between $\\sim7\\times10^{3}$ and $\\sim1.7\\times10^{4}$ cm$^{-3}$ and with a temperature of $7.5\\times10^{3}$ K, distributed in a spherical structure of uniform density and a diameter of 3$-$4 pc \\citep{moh02}. \\citet{tin04} found that these results were consistent with a high degree ($\\tau_{0} \\sim 36$) of free-free absorption towards the assumed nucleus, 5.79$-$39.0.  Similar results for the supernova remnants in M82 have been obtained by \\citet{ped99} and \\citet{mcd02}.  In M82 30 supernova remnants have been detected from a population of 46 compact radio sources \\citep{mcd02} and several resolved \\citep{ped99, mux05}.  \\citet{mcd02} found a free-free optical depth toward supernova remnants of $\\tau_{0}\\sim3$.  Motivated by the strong indicators of free-free absorption in NGC 253, we have embarked on a program of VLBI observations of this galaxy and other prominant Southern Hemisphere starburst galaxies. This paper reports results from a new observation of the nuclear region of NGC 253 at 2.3 GHz. The resulting data have better resolution and sensitivity than the 1.4 GHz observations  of \\citet{tin04} and we have used them to investigate in more detail the free-free absorbed spectra of supernova remnants and \\ion{H}{2} regions in the NGC 253 starburst (\\S~\\ref{sec:ffmodel}). In particular we are interested in probing the structure of the ionised environment of the supernova remnants and the relationship between the ionised medium and the other components of the galaxy (\\S~\\ref{sec:spectra}), further constraining the supernova rate in NGC 253 using previously published images and our new VLBI images (\\S~\\ref{sec:snrate}) and re-evaluating estimates of the star formation rate based on our results (\\S~\\ref{sec:sfrate}).  A subsidiary focus of this project is to achieve high sensitivity, high fidelity, high resolution, and computationally efficient imaging over wide fields of view (at least compared to traditional VLBI imaging techniques).  As such, the associated exploration of the relevant techniques is a small first step toward the much larger task of imaging with the next generation of large radio telescopes, for example the Square Kilometre Array (SKA: \\citet{hal05, car04}).  The SKA will use baselines ranging up to approximately 3000 km and cover fields of view greater than one square degree instantaneously at frequencies of around 1 GHz.  The computational load required for this type of imaging task is vast and will drive the development of novel imaging algorithms, using techniques such as w-projection\\footnote{EVLA Memo 67:  \\url{http://www.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf}} and the use of super-computing facilities.  Investigating the performance of these techniques now, under the most challenging current observing conditions, is therefore a useful activity. ",
        "conclusions": "\\subsection{Radio spectra and free-free absorption modelling of the compact sources} \\label{sec:ffmodel} U\\&A97 measured the spectral indices of 23 compact radio sources in NGC 253 using VLA observations between 5 GHz and 23 GHz (17 with spectral index errors of less than 0.4). Assuming that the spectral indices apply down to 2.3 GHz and that the sources are not resolved out by the higher resolution available with the LBA, 18 of these sources should be visible above the 5 $\\sigma$ detection limit of our 2.3 GHz LBA observations. These sources are listed in Table \\ref{tab:tabflux} with flux density measurements from VLA observations at 5 GHz, 8.3 GHz, 15 GHz and 23 GHz (U\\&A97), together with LBA observations of two sources (5.48$-$43.3 and 5.62$-$41.3) at 1.4 GHz and flux density upper limits of 1.8 mJy at 1.4 GHz for the remaining sources \\citep{tin04}.  Comparing this list with the sources in Table \\ref{tab:tabsources} reveals that only 4 sources (5.48$-$43.3, 5.62$-$41.3, 5.72$-$40.1 and 5.79$-$39.0) were detected with the LBA at 2.3 GHz. The 2.3 GHz flux densities for all 18 sources are listed in Table \\ref{tab:tabflux} with upper limits of 1.2 mJy for non-detections. A further two sources, 5.805$-$38.92 and 5.97$-$39.7, were detected with the LBA at 2.3 GHz and also by U\\&A97 using the 15 GHz VLA A and 23 GHz VLA (A+B) configurations, and are also listed in Table \\ref{tab:tabflux}. However, due to mismatches in resolution these source fluxes are only considered as approximate measurements by U\\&A97. As a check to ensure that our measured fluxes at 2.3 GHz were not an underestimate as a result of higher resolution, total flux density measurements were compared between the 6 telescope LBA (Parkes, Narrabri, Mopra, Tidbinbilla, Hobart and Ceduna) observations and the 4 telescope array (excluding the longer baselines associated with Ceduna and Hobart) observations and were found to be the same within the measurement errors. Furthermore, U\\&A97 found no significant variability in individual radio sources in NGC 253 over a period of 8 years, so it can be assumed that these measured fluxes, at multiple wavelengths, over multiple epochs, can be used to construct the spectra.  \\citet{tin04} argued that free-free absorption was the cause for the sharp downturn observed in the spectra of the compact radio sources at 1.4 GHz. At 2.3 GHz the effect of free-free absorption is not as great as at 1.4 GHz, so the detection of two sources at 1.4 GHz and a further four sources at 2.3 GHz appears to support this argument. The non-detection of the remaining 14 sources thus suggests that the effect of free-free absorption is still significant at 2.3 GHz. To test this quantitatively, an analysis similar to that undertaken by \\citet{mcd02} for M82 was performed.  \\citet{mcd02} constructed the spectra of 20 compact radio sources in M82 from existing and new observations between 1.42 GHz and 15 GHz using the VLA and MERLIN. The spectra were investigated using three models as described by the equations  \\begin{equation} S(\\nu)=S_{0}\\nu^{\\alpha}, \\end{equation} \\begin{equation} S(\\nu)=S_{0}\\nu^{\\alpha}e^{-\\tau(\\nu)}, \\end{equation} \\begin{equation} S(\\nu)=S_{0}\\nu^{2}(1-e^{-\\tau(\\nu)}), \\end{equation} where \\begin{equation} \\tau(\\nu)=\\tau_{0}\\nu^{-2.1}. \\end{equation}  Equation 1 represents a simple power-law spectrum, equation 2 a free-free absorbed power-law spectrum, and equation 3 a self-absorbed bremsstrahlung (free-free absorbed) spectrum. In these expressions $\\alpha$ is the optically thin intrinsic spectral index, $\\tau_{0}$ is the free-free optical depth at 1 GHz and $S_{0}$ is the intrinsic flux density of the source at 1 GHz. The combination of high frequency VLA data and 1.4 GHz LBA data, together with the detection of six sources at 2.3 GHz, and the more stringent upper limit for non-detections at 2.3 GHz, allow us to use these models to constrain the free-free parameters of the 20 compact radio sources listed in Table \\ref{tab:tabflux}.  The spectrum of each source was tested against each of the three models described above using a reduced$-\\chi^{2}$ criterion to determine the best fit. For all sources, the free-free absorbed power-law model produced a significantly better fit than the self-absorbed and power-law models, for example giving a mean reduced$-\\chi^{2}$ fit of 0.8, 23, and 40 for models 2, 3 and 1 respectively. Upper limits of 1.8 mJy and 1.2 mJy were used where there were non-detections from the LBA observations at 1.4 GHz and 2.3 GHz, respectively.  The resulting free parameters $S_{0}$, $\\tau_{0}$ and $\\alpha$ from the model fits are listed in Table \\ref{tab:tabff}. The free-free absorbed power law model for each source is shown against the corresponding measured spectrum in Figure \\ref{fig:figff}. The data shown in Table \\ref{tab:tabff} and Figure \\ref{fig:figff} confirm the free-free absorption interpretation of \\citet{tin04}. In particular, the new 2.3 GHz data for 5.48$-$43.3 and 5.62$-$41.3 complete spectra that are very well fit by free-free absorbed power laws and are close to those adopted by \\citet{tin04}. Furthermore, the results for the six sources listed in Table 2 of \\citet{tin04} are consistent with the more stringently constrained free-free absorbed power law models listed in Table \\ref{tab:tabff}.  Using similar criteria to that used by U\\&A97, of the 20 sources listed in Table \\ref{tab:tabff}, 9 have flat intrinsic power law spectra ($\\alpha > -0.4$) indicative of \\ion{H}{2} regions dominated by thermal radio emission (T), the 11 remaining sources have steep intrinsic spectra (S), as normally associated with supernova remnants \\citep{mcd02}. The approximately equal number of thermal versus non-thermal sources is consistent with the findings of U\\&A97 in NGC 253, however a slightly larger proportion of non-thermal sources is present in M82 \\citep{mcd02}.  \\subsection{Comparison with multi-wavelength datasets} \\label{sec:spectra}  When the modelled free-free opacity is illustrated against the source location (Figure \\ref{fig:figngc253}) it is clear that it does not vary smoothly across the central region of the galaxy, but is rather clumpy. Since free-free absorption is associated with ionised gas, it is worthwhile investigating tracers of ionised gas at other wavelengths in an attempt to understand the distribution of gas in these regions.  HST images of NGC 253 at optical wavelengths provide the closest match in resolution to those obtained with the LBA and VLA. Of particular interest is the H$\\alpha$ emission line as it traces diffuse ionised gas and is prominent in pre and post-main sequence stars. The \\ion{S}{3} line has been used together with the H$\\alpha$ line to highlight regions with photo-ionization from \\ion{H}{2} regions and OB stars \\citep{for00}. HST H$\\alpha$ and [\\ion{S}{3}] emission line images were made from archival HST data and the images shifted so that the source co-ordinates were aligned with those listed by \\citet{for00}. The two images were then correlated against the free-free opacities modelled at the compact source locations. As many sources only had lower limits for the free-free optical depth, simple statistical techniques to test the relationship between measured intensity and the free-free opacity would not fully utilise all of the data available. Instead, 'survival analysis' methods were employed to take into consideration the lower limits when performing linear regression tests on the data.  The survival analysis was performed on the data using tasks in the ASURV package \\citep{lav92}. The resulting correlation tests indicated that no trend was apparent between the H$\\alpha$ or [\\ion{S}{3}] emission line intensity and the modelled free-free optical depth for any of the sources. The lack of correlation is most likely due to the high levels of extinction that are apparent at optical wavelengths and are characteristic of starburst galaxies \\citep{eng98}. This extinction makes properties deduced from optical measurements unrepresentative \\citep[e.g.][]{rie80, pux91}.  To overcome the problems associated with extinction it is necessary to observe at wavelengths, such as X-ray and radio, that can penetrate gas and dust. Radio observations in particular also need to be toward the high frequency end of the spectrum to avoid the effects of free-free absorption. Australia Telescope Compact Array (ATCA) images of the Ammonia \\citep{ott05} and HCN emission lines \\citep{ott04} and \\emph{Chandra} hard x-ray images \\citep{wea02} were obtained and compared against the modelled free-free optical depth at each of the source locations. Ammonia (NH$_{3}$) is a good temperature indicator whereas HCN is a good tracer for high density molecular gas ($>10^{4}$ cm$^{-3}$). The hard x-ray observations show regions of photo-ionization by massive young stars or SNR-driven shocks. All three images were analysed using a survival analysis against the modelled free-free optical depth with no clear trend evident. Here the lack of correlation is most likely due to the mismatches in resolution. In comparison to the 2.3 GHz data, the ATCA and \\emph{Chandra} images were of substantially lower resolution and so were not able to resolve sufficient detail to provide a useful comparison with the VLBI data.  High resolution radio recombination line (RRL) observations with the VLA provide a reasonable match in resolution to the VLBI observations. Furthermore, the high frequency H92$\\alpha$ (8.3094 GHz) and H75$\\alpha$ (15.2815 GHz) lines are not adversely affected by the effects of free-free absorption and can be used to estimate temperatures, densities, and kinematics \\citep{lob04}. The hydrogen radio recombination line (RRL), which occurs in thermal gas, is insensitive to dust extinction and can be used to probe highly obscured star formation.  \\citet{moh05} have imaged RRL emission in NGC 253 at arcsecond resolution and detected 8 sources. By modelling the emission, \\citet{moh05} provide size and electron density estimates for these sources.  When the RRL source associated with the core is overlaid with the core position in the 2.3 GHz LBA continuum map, 6 of the sources modelled for free-free absorption in Table \\ref{tab:tabff} are found to be located near 5 of the RRL sources. By assuming that the sources are centrally located within a uniform density sphere of ionised gas of temperature 7500 K, the expected free-free opacities have been estimated, the results listed in Table \\ref{tab:tabrrlff}.  Four of the VLA sources (5.49$-$42.3, 5.59$-$41.6, 5.72$-$40.1 and 5.78$-$39.4) have free-free opacities that are within the limits imposed by the RRL models (labelled as 'low' and 'high' in Table \\ref{tab:tabrrlff}). The first three sources are identified as thermal regions. However 5.78$-$39.4 is non-thermal ($\\alpha\\sim-0.8$). The close proximity of 5.78$-$39.4 to the assumed core 5.79$-$39.0, where dense ionised gas is expected, together with the observed RRL emission, suggests that it may be a supernova remnant embedded in or behind a \\ion{H}{2} region.  The two VLA sources associated with the 5.795$-$39.05 RRL source, the assumed core (5.79$-$39.0) and a nearby \\ion{H}{2} region (5.805$-$38.92) have free-free optical depths that fall well below the limits imposed by the RRL models. This suggests that the sources may not be centrally located within the gas distribution. If the sources are in front of or to the side of the high density region then they may exhibit the lower free-free optical depth that is inferred.  Many of the other VLA sources exhibit a significant degree of free-free absorption but do not have any associated RRL emission. Most of these sources are non-thermal and so may not necessarily have a RRL component. However, there are five thermal sources that do not have RRL sources associated with them. Four of these are weak and may fall below the sensitivity limit of the RRL observation. The remaining source has a combination of thermal and non-thermal emission (5.62$-$41.3) and it is possible that the thermal component is not strong enough to be detected within the sensitivity limits of the RRL observation. Interestingly, 5.62$-$41.3 does lie within a beam width of the RRL emission associated with 5.59$-$41.6 and were it associated with that emission it would have a free-free optical depth that is consistent with that expected from the source parameters modelled by \\citet{moh05} from the RRL emission.  \\subsection{Comments on individual compact radio sources}  \\subsubsection{Resolved SNR 5.48$-$43.3}  Source 5.48$-$43.3 dominates observations at 1.4 GHz and 2.3 GHz. In our high resolution image (Figure \\ref{fig:figsnr}) 5.48$-$43.3 appears to be resolved into a shell of approximately 90 mas in diameter corresponding to an extent of approximately 1.7 pc. If we assume an average expansion velocity of $v=10,000$ km s$^{-1}$ relative to the shell centre, similar to that measured for the supernova remnant 43.31+592 in M82 \\citep{ped99}, this implies an age of $\\sim 80 (10^{4}/v)$ years for this remnant. The eastern side of the remnant appears 2 to 3 times brighter than the western side. This may be the result of interaction with denser interstellar material in that direction. The circular symmetry of the remnant may indicate that the interaction has only recently begun or that the density of the ISM is only marginally greater.  With a total flux of 32 mJy at 2.3 GHz, 5.48$-$43.3 is approximately half the size and 25 times brighter than Cas A would be at this frequency if it were at the same distance. Furthermore, with the assumed expansion velocity, 5.48$-$43.3 is one quarter the estimated age of Cas A.  \\subsubsection{Assumed Core 5.79$-$39.0}  The measured size of 5.79$-$39.0 in the 2.3 GHz LBA image is approximately 80$\\times$60 mas, consistent with the lower limit of 30 mas measured with the Parkes-Tidbinbilla Interferometer \\citep{sad95} and the upper limit of $\\sim100$ mas measured by U\\&A97 with 8 GHz VLA observations. 5.79$-$39.0 is not detected in the full resolution LBA image at 2.3 GHz as would be expected given its size and flux density.  5.79$-$39.0 is the strongest source in 5$-$23 GHz VLA images but is not visible in 1.4 GHz LBA observations \\citep{tin04}. In our 2.3 GHz LBA image the source is detected but is weaker than expected based on its high frequency spectral index, suggesting that there is significant absorption at lower frequencies. Our source modelling (\\S\\ref{sec:ffmodel}) shows that the source is best fit with a free-free absorbed power law model with optical depth $\\tau_{0}\\sim17$. Evidence of free-free absorption in the nuclear region has also been found by \\citet{rey83}, \\citet{car96} and \\citet{tin04}.  \\subsubsection{Resolved SNRs 5.62$-$41.3 and 5.79$-$39.7}  The source 5.62$-$41.3 lies in a region that has the highest HCN/CO ratio and thus the highest average density ($\\sim10^{5}$ cm$^{-3}$) of any region studied by \\citet{pag95}. This is one of only two sources detected at 1.4 GHz with the LBA \\citep{tin04}. U\\&A97 suggest that the source is a mix of supernova remnants and \\ion{H}{2} regions. In our 2.3 GHz LBA observations the source is extended ($130\\times85$ mas) in the low resolution image (Figure \\ref{fig:figngc253}) and is not detected in the full resolution image. It does not appear to be greatly affected by free-free absorption and so may lie in front of the dense gas detected by \\citet{pag95}. Its modelled spectral index of $-0.31$ is at the low end of that measured by U\\&A97.  Source 5.79$-$39.7 is the smallest supernova remnant in our 2.3 GHz LBA image but is not detected in the higher resolution image. It also appears to be in a region with significant levels of ionised gas as it is heavily affected by free-free absorption.  \\subsubsection{The \\ion{H}{2} regions 5.72$-$40.1 and 5.805$-$38.92} Source 5.72$-$40.1 is a flat spectrum source which has been observed to vary in its flux significantly at 5 GHz and 8.3 GHz from VLA observations (U\\&A97). Our measured size of $60\\times50$ mas is significantly smaller than the size of $230\\times120$ mas measured at 23 GHz and $200\\times90$ mas measured at 15 GHz using the VLA (U\\&A97). It also has a position offset that is not as consistent as the other sources when compared against VLA positions. These two facts suggest that it may have structure that varies significantly with frequency. This source is associated with RRL emission that is consistent with a free-free absorbed \\ion{H}{2} region.  Source 5.805$-$38.92 is extended in size ($130\\times50$ mas) and not detected in higher resolution LBA images. It is associated with RRL emission but exhibits less free-free absorption than expected given the electron densities and source sizes predicted from RRL emission models.  \\subsection{The supernova rate in NGC 253} \\label{sec:snrate} With new observations made since those of U\\&A97 and with an improved distance determination to NGC 253 \\citep{kar03} it is worth reinvestigating the limits that may be placed on the supernova rate, discussed in detail by U\\&A97.  The increased distance estimate to NGC 253 suggests that earlier spatial measures were underestimated by a factor of $\\sim1.6$. Similarly, luminosity and energy estimates would have been underestimated by a factor of $\\sim2.5$. The overall effect on the population of observed supernova remnants in NGC 253 is that they are not only older but also more luminous than previously estimated.  An upper limit to the supernova rate can be made by assuming that all of the FIR luminosity is reprocessed supernova energy, that is, all energy released by supernovae is eventually thermalized. Using this assumption \\citet{ant88} determined an upper limit of 3 yr$^{-1}$ for the supernova rate. Adjusting the original FIR luminosity measured by \\citet{tel80} for a distance of 3.94 Mpc gives a total FIR luminosity of $3.8\\times10^{10} L_{\\Sun}$ or $\\sim1.4\\times10^{44}$ ergs s$^{-1}$ for NGC 253. Dividing this by the canonical $10^{51}$ ergs of total energy output per supernova results in a maximum rate of $\\sim4.5$ supernovae per year. This method has been shown to grossly overestimate supernova rates since FIR emission also results from stellar heating of dust and can be contaminated by the presence of an AGN \\citep{con92}.  A lower limit on the supernova rate can be made based on the number of detected remnants, their size and an assumed expansion rate. \\citet{ulv94} assumed that of order 100 compact sources were supernova remnants with upper size limits of 2$-$5 parsecs based on a distance of 2.5 Mpc to NGC 253. Furthermore, assuming an expansion at 5,000 km s$^{-1}$ (rate of expansion of the FWHM of a Gaussian model), \\citet{ulv94} arrived at a supernova rate of 0.1$-$0.25 yr$^{-1}$. As only $\\sim50\\%$ of the compact radio sources in NGC 253 are associated with supernova remnants (U\\&A97) and the revised source sizes are now of order 3$-$8 parsecs, a supernova rate of $>0.14-0.36$ is expected given an assumed expansion rate of 10,000 km s$^{-1}$. Since the mean expansion rate of the supernova remnants in NGC 253 is unclear, we can rewrite the limit as $\\nu_{SN}>0.14 (v/10^{4})$ yr$^{-1}$, where $v$ is the shell radial expansion velocity in km s$^{-1}$. This is of the same order of magnitude as the supernova rate determined for M82, using similar methods, $\\sim0.07$ yr$^{-1}$ \\citep{ped03}.  Having observed NGC 253 at a second epoch, \\citet{ulv91} found that no new sources stronger than 3 mJy at 5 GHz were detected after a period of 18 months. They statistically modelled the effect of such non-detections on the supernova rate. Their model assumed a hypothetical population of supernovae that peaked at 5 GHz 100 days after their optical maxima, had a 5 GHz peak flux that was with equal probability either 5, 10, 15, or  20 times that of Cas. A and decayed as the $-$0.7 power of time. \\citet{ulv91} determined that approximately $2/3$ of supernovae that occur during the 18 month period between epochs should be detectable. Further assuming that the occurrence of supernovae in NGC 253 are Poisson distributed, \\citet{ulv91} determined with 95\\% confidence that the upper limit of the supernova rate was 3.0 yr$^{-1}$. The same model was used to determine an upper limit of $\\sim1.4$ yr$^{-1}$ at the end of a third epoch \\citep{ulv94} and $\\sim0.3$ yr$^{-1}$ at the end of the fourth epoch (U\\&A97).  Since the fourth epoch two further high resolution radio observations of NGC 253 have been made. The first was four years later by \\citet{moh05} at 5.0 GHz and the second was the observation made 4.8 years later again at 2.3 GHz as described in this paper. Although the \\citet{moh05} image can be directly compared to the historical data at the same frequency, the observations at 2.3 GHz require extra consideration of the nature of supernova remnants at lower frequencies, the differing sensitivity limits,  and the effects of free-free absorption. Regardless of these differences, what is clear is that no new sources are evident in either the \\citet{moh05} images or our 2.3 GHz image.  To place limits on the supernova rate based on these observations, we have developed a new model, based in principle on that of \\citet{ulv91}. At 5 GHz the hypothetical supernova for this model is identical to the one described in \\citet{ulv91}, however at 2.3 GHz it is assumed that the supernova remnant will peak 200 days after the optical peak, albeit at the same flux density as at 5 GHz - this approximately follows the SN 1980k light curve determined by \\citet{wei86}.  A Monte-Carlo simulation was used to drive the model by randomly choosing a uniformly distributed peak luminosity between 5 and 20 times that of Cas A., which was then scaled to flux density for a given distance to NGC 253 and also adjusted for free-free absorption using Eqn 2. A Poisson distributed random number was then used to define how many supernovae, given a specified supernova rate, would occur between two epochs and at which uniformly distributed times these supernovae would occur. Given the peak flux density and timing of the supernova, a test was made to determine if each supernova remnant could be detected when observed at the subsequent epoch given the sensitivity limit for the observation. By executing the Monte-Carlo simulation over 10,000 iterations, the proportion of supernova remnants detected at the end of each epoch was determined. The simulation was seeded with a supernova rate of 0.1 yr$^{-1}$ across all epochs, the resulting output giving the confidence level for a non-detection at each epoch at that rate. Linear interpolation was then used to determine a new rate that would drive the simulation towards a 95\\% confidence level and the simulation was repeated until the 95\\% level was reached.  Table \\ref{tab:tabsnrate} shows the results of three separate tests with the Monte-Carlo simulation. In each test, a run was made for each epoch to place a limit on the supernova rate at the time of that observation. Also, given that a supernova remnants can peak and fade beneath the sensitivity limit between two epochs, the proportion of supernova remnants ($\\beta_{SN}$) that can actually be detected at the end of each epoch was also noted.  In test 1 the distance to NGC 253 was set to 2.5 Mpc and no free-free absorption was assumed in order to provide a direct comparison to the \\citet{ulv91, ulv94, ulv97} results.  At epoch 2, the first test resulted in a supernova rate of $<2.4$ yr$^{-1}$ compared to $<3.0$ yr$^{-1}$ obtained by \\citet{ulv91}. This difference can be attributed to the proportion detected ($\\sim83$\\%) being greater than initially anticipated with a simpler model ($\\sim67$\\%). The proportion detected at epoch 3 ($\\sim68$\\%) is similar to that predicted by \\citet{ulv91} however the cumulative effect of a greater rate of detection at the second epoch results in a supernova rate of only $<1.0$ yr$^{-1}$ compared to $<1.4$ yr$^{-1}$ obtained by \\citet{ulv94}. At the fourth epoch a lower proportion of detections are predicted ($\\sim48$\\%) as a result of an increased number of supernova remnants fading below the sensitivity limits over the longer period between observations, this results in a supernova rate of $<0.62$ yr$^{-1}$ which is almost twice the maximum predicted by U\\&A97. Given further observations by \\citet{moh05} and those presented in this paper it is possible to add two further epochs to the test which in the absence of further detections results in a further lowering of the maximum allowable supernova rate to $<0.44$ yr$^{-1}$ and $<0.26$ yr$^{-1}$ for epochs 5 and 6 respectively. The results of our simulations show that the model originally used by \\citet{ulv91} does not accurately determine the supernova rate when the time between epochs exceeds approximately 2 years, under the assumptions used.  In test 2 the Monte-Carlo simulation was repeated for NGC 253 at the revised distance of 3.94 Mpc. The increased distance results in more supernova remnants falling below the sensitivity limits of the observations. For observations at 5 GHz this results in the possibility of only 8$-$22\\% of all supernovae occurring between epochs being detected. The overall effect is to increase the upper limit of the supernova rate to 3.0 yr$^{-1}$ at the end of epoch 4 and to 0.6 yr$^{-1}$ at the end of epoch 6.  It has been shown in \\S~\\ref{sec:spectra} that the supernova remnants in NGC 253 appear to be situated in or behind a screen of ionised gas. In test 3 the effects of free-free absorption on the upper limit of the supernova rate are taken into consideration with the Monte-Carlo simulation. For this test a median value of $\\tau_{0}=6$ was taken based on free-free modelling in \\S~\\ref{sec:spectra}. The effect of free-free absorption at 5 GHz is not greatly significant but is sufficient to drop the proportion of detections to between 5\\% and 13\\%. At 2.3 GHz the effect of free-free absorption is more significant with the proportion of supernova remnants that can be detected dropping from 70\\% to 10\\%. Overall, this results in upper limits on the supernova rate of 5.2 yr$^{-1}$ at epoch 4 and 2.4 yr$^{-1}$ at epoch 6.  The overall supernova rate upper limit of 2.4 yr$^{-1}$ is below the upper limit determined from FIR emission, however it is more than an order of magnitude greater than that determined from source counts and sizes alone. U\\&A97 observed no significant source fading over a period of 8 years and on this basis they could further limit the supernova rate to no greater than 0.3$-$0.6 yr$^{-1}$. The apparent lack of significant variability on time-scales of many years has also been observed for a majority of compact radio sources in M82 by \\citet{kro00}. It is possible that the dense environment around the nuclear regions of NGC 253 and M82 inhibit the expansion and fading of the remnants. If this is the case then the remnants may be older than they appear and this could reduce the supernova rate further. It is also possible that we are only seeing the tip of the iceberg, there may be a large background supernova rate but we only observe the rare bright events.  To improve the supernova rate estimates further it would be useful to refine the Monte-Carlo simulator to include a more realistic probability distribution for the supernova remnant peaks and make better use of the \\citet{wei86} radio supernova light curves. A reduction in the upper limit of the supernova rate can be achieved by increasing the probability of detecting new supernovae as they occur. This may be achieved with frequent high sensitivity observations of NGC 253. Furthermore, through multi-epoch observations, it should be possible to measure the rate of expansion of the brighter more resolved remnants and so allow for age determination at a greater level of accuracy and subsequently improved estimates of the supernova rate.  \\subsection{The star formation rate in NGC 253} \\label{sec:sfrate}  The star-formation rate (SFR) of a star-forming galaxy is directly proportional to its radio luminosity $L_{\\nu}$ at wavelength $\\nu$ \\citep{con92, haa00}: \\begin{equation} \\left( \\frac{SFR(M\\geq5M_{\\Sun})}{M_{\\Sun}\\mathrm{yr}^{-1}} \\right)=Q \\left \\{ \\frac{ \\frac{L_{\\nu}}{\\mathrm{W Hz^{-1}}}}{\\left[ 5.3\\times10^{21} \\left( \\frac{\\nu}{\\mathrm{GHz}}\\right) ^{-0.8} + 5.5\\times10^{20} \\left( \\frac{\\nu}{\\mathrm{GHz}} \\right) ^ {-0.1} \\right]} \\right \\}. \\end{equation} \\citet{con92} derives this relation purely from radio considerations by calculating the contribution of synchrotron radio emission from SN remnants and of thermal emission from \\ion{H}{2} regions to the observed radio luminosity. The factor $Q$ accounts for the mass of all stars in the interval $0.1-100 M_{\\Sun}$ and has a value of 8.8 if a Saltpeter IMF ($\\gamma=2.5$) is assumed - this IMF is assumed for the remainder of this section. \\citet{ott05} used a measure of the 24 GHz continuum emission in NGC 253 to derive a SFR which, when adjusted for a distance of 3.94 Mpc, gives a SFR of 10.4$\\pm$1 M$_{\\Sun}$ yr$^{-1}$.  A large proportion of the bolometric luminosity of a galaxy is absorbed by interstellar dust and re-emitted in thermal IR. As a result FIR emission is a good tracer for young stellar populations. By modelling the total energy emission of a massive star and assuming that the contribution of dust heating by old stars is negligable, \\citet{con92} determined the following relation between the star-formation rate of a galaxy and the FIR luminosity $L_{IR}$:  \\begin{equation} \\left( \\frac{SFR(M\\geq5M_{\\Sun})}{M_{\\Sun}\\mathrm{yr}^{-1}} \\right)=9.1\\times10^{-11}L_{FIR}/L_{\\Sun}. \\end{equation} Using \\citet{rad01} FIR measures for NGC 253, adjusted for a distance of 3.94 Mpc, gives a SFR of 1.8$-$2.8 M$_{\\Sun}$ yr$^{-1}$ for the inner $\\sim300$ pc nuclear region and 3.5$-$4.3 M$_{\\Sun}$ yr$^{-1}$ for the entire galaxy.  Finally, the star formation rate can be determined from the supernova rate directly. For most galaxies the total radio non-thermal and thermal luminosities and the FIR/radio ratio is proportional to the average star formation rate for $M\\geq5M_{\\Sun}$ \\citep{con92}. Furthermore, as all stars more massive than $8M_{\\Sun}$ become radio supernova, the star formation rate can be determined using the following relationship:  \\begin{equation} \\left( \\frac{SFR(M\\geq5M_{\\Sun})}{M_{\\Sun}\\mathrm{yr}^{-1}}\\right) \\sim 24.4 \\times \\left[ \\frac{\\nu_{SN}}{\\mathrm{yr}^{-1}}\\right]. \\end{equation} Using the supernova rate limits of $0.14 (v/10^{4})<\\nu_{SN}<2.4$ yr$^{-1}$ (\\S~\\ref{sec:snrate}) the limits on the star formation rate are $3.4 (v/10^{4}) < SFR(M\\geq5M_{\\Sun})<59$ M$_{\\Sun}$ yr$^{-1}$. This is of the same order of magnitude as the nuclear SFR determined from FIR emissions and in agreement with that determined from 24 GHz continuum emissions alone. Using equations 6 and 7, the supernova rate can be estimated from the FIR-based SFR estimate giving a value of 0.07$-$0.11 yr$^{-1}$. This figure is in line with measurements based on source sizes and suggests that the simple model used to determine the supernova rate from FIR emission in \\S~\\ref{sec:snrate} was grossly overestimating the supernova rate. It is encouraging that the three separate estimates of the SFR in NGC 253 all produce results that are only a factor of a few different. These estimates are similar to star-formation rates derived from the supernova rate and FIR emission in M82 which give values of $\\sim1.8$ and $\\sim2.0$ M$_{\\Sun}$ yr$^{-1}$ respectively \\citep{ped03}."
    },
    "0606/astro-ph0606661_arXiv.txt": {
        "abstract": "{ {\\it Context.} Using Particle-In-Cell simulations i.e. in the kinetic plasma description, the discovery of a new mechanism of parallel electric field generation was recently reported. \\\\ {\\it Aims.} We show that the electric field generation parallel to the uniform unperturbed magnetic field can be obtained in a much simpler framework using the ideal magnetohydrodynamics (MHD) description. \\\\ {\\it Methods.} We solve numerically ideal, 2.5D, MHD equations in Cartesian coordinates, with a plasma beta of 0.0001 starting from the equilibrium that mimics a footpoint of a large curvature radius solar coronal loop or a polar region plume. On top of such an equilibrium, a purely Alfv\\'enic, linearly polarised, plane wave is launched. \\\\ {\\it Results.} In ideal MHD the electric field parallel to the uniform unperturbed magnetic field appears due to fast magnetosonic waves which are generated by the interaction of weakly non-linear Alfv\\'en waves with the transverse density inhomogeneity. In the context of the coronal heating problem  a new  two stage mechanism of plasma heating is presented by putting emphasis, first, on the generation of parallel electric fields within an  ideal MHD description directly, rather than focusing on the enhanced dissipation mechanisms of the Alfv\\'en waves and, second, dissipation of these parallel electric fields via {\\it kinetic} effects. It is shown that a single Alfv\\'en wave harmonic  with frequency $\\nu = 7$ Hz and longitudinal wavelength  $\\lambda_A = 0.63$ Mm, for a putative Alfv\\'en speed of 4328 km s$^{-1}$, the generated parallel electric field could account for 10\\% of the necessary coronal heating requirement. It is also shown that the amplitude of the generated parallel electric field exceeds the Dreicer electric field by about four orders of magnitude, which implies realisation of the run-away regime with associated electron acceleration.\\\\ {\\it Conclusions.} We conjecture that  wide spectrum (10$^{-4}-10^3$ Hz) Alfv\\'en waves, based on the observationally constrained spectrum, could provide the necessary coronal heating requirement. The exact amount of energy that could be deposited by such waves through our mechanism of parallel electric field generation can only be calculated once a more complete parametric study is done. Thus, the \"theoretical spectrum\" of the energy stored in parallel electric fields versus frequency needs to be obtained. ",
        "introduction": "The coronal heating problem, the puzzle of what maintains the solar corona 200 times hotter than the photosphere, is one of the main outstanding questions in solar physics (see e.g.  \\citet{t05} for a recent brief review on the subject). A significant amount of work has been done in the context of heating of open magnetic structures in the solar corona \\citep{hp83,nph86,p91,nrm97,dmha00,bank00,tan01,hbw02,tn02,tna02,tnr03}. Historically, all phase mixing studies have been performed in the Magnetohydrodynamic (MHD) approximation; however, since the transverse scales in the Alfv\\'en wave collapse progressively to zero, the MHD approximation is inevitably violated. Thus, \\citet{tss05a,tss05b} studied the phase mixing effect in the kinetic regime, using Particle-In-Cell simulations, i.e. beyond an MHD approximation, where a new mechanism for the acceleration of electrons due to the generation of a parallel electric field in the solar coronal context was discovered. This mechanism has important implications for various space and laboratory plasmas, e.g. the coronal heating problem and acceleration of the solar wind. It turns out that in the magnetosphere a similar parallel electric field generation mechanism in transversely inhomogeneous plasmas has been reported \\citep{glm04,glq99}. See also \\citet{mgl06} and references therein.  This new mechanism occurs when an Alfv\\'en wave moves along the field in a plasma with a transverse density inhomogeneity. The progressive distortion of the Alfv\\'en wave front due to differences of local Alfv\\'en speed then generates the parallel electric field. In this work we show that the electric field generation parallel to the uniform unperturbed magnetic field can be obtained in a much simpler framework using an ideal MHD description, i.e. without resorting to complicated wave particle interaction effects  such as ion polarisation drift and resulting space charge separation, which seems to be the fundamental cause of electron acceleration. See also the Discussion section for a more detailed account of what led us to consider the MHD approximation for parallel electric field generation.  In this paper, we also explore the implications of this effect for the coronal heating problem. Preliminary results have been reported elsewhere \\citep{t06}. The present study explores the importance of MHD versus kinetic effects in solar plasmas, which has recently attracted considerable attention due to apparent difficulties in the coronal heating problem, as well as the natural tendency of wave heating models to proceed to progressively small spatial scales. Also, recent observations have indicated further sub-structuring of the coronal morphology \\citep{ineke2006}, which is relevant to our earlier work \\citep{t05}. ",
        "conclusions": "After the comment paper by \\citet{mgl06} we came to the realisation that the electron acceleration seen in both series of works \\citep{tss05a,tss05b,glm04,glq99} is a non-resonant wave-particle interaction effect. In works by \\citet{tss05a,tss05b} the electron thermal speed was $v_{th,e}=0.1c$ while the Alfv\\'en speed in the strongest density gradient regions was $v_A=0.16c$; this unfortunate coincidence led us to the conclusion that the electron acceleration by parallel electric fields was affected by the Landau resonance with the phase-mixed Alfv\\'en wave. In works by \\citet{glm04,glq99} the electron thermal speed was $v_{th,e}=0.1c$ while the Alfv\\'en speed was $v_A=0.4c$ because they considered a more strongly magnetised plasma applicable to Earth magnetospheric conditions. However, the interaction of the Alfv\\'en wave with a transverse density plasma inhomogeneity when the Landau resonance condition $\\omega=k_\\parallel c_A(x)$ is met can be quite important for the electron acceleration \\citep{chaston2000,hs76}. \\citet{chaston2000} assert that the electron acceleration observed in density cavities in aurorae can be explained by the Landau resonance of the cold ionospheric electrons with the Alfv\\'en wave. \\citet{hs76} also established that at resonance the Alfv\\'en wave fully converts into the kinetic  Alfv\\'en wave with a perpendicular wavelength comparable to the ion gyro-radius. We can then conjecture that {\\it because of  kinetic Alfv\\'en wave front stretching} (due to phase mixing, i.e. due to the differences in local Alfv\\'en speed), this perpendicular component {\\it gradually realigns} with the ambient magnetic field and hence creates the time varying parallel electric field component. This points to the importance of the Landau resonance for electron acceleration when the resonance condition is met. But as seen in works of  \\citet{glm04,glq99}, even when the resonance condition is not met, electron acceleration is still possible.  There were three main stages that lead to the formulation of the present model:  (i) The realisation that the parallel electric field generation (and particle acceleration) is a {\\it non-resonant} wave-particle interaction effect lead us to the question: could such  parallel electric fields be generated in a MHD approximation?  (ii) If one considers {\\it non-linear} generation of the fast magnetosonic waves in the transversely inhomogeneous plasma, then $\\vec E = - (\\vec V \\times \\vec B)/c$ contains a non-zero component parallel to the ambient magnetic field $E_z = - (V_x B_y - V_y B_x)/c$.  (iii) From previous studies \\citep{bank00,tan01} we knew that the fast magnetosonic waves ($V_x$ and $B_x$) did not grow to a substantial fraction of the Alfv\\'en wave amplitude. However, after reproducing the old parameter regime (k=1, i.e a frequency of 0.7 Hz), the case of k=10, i.e a frequency of 7 Hz was considered, which showed that  fast magnetosonic waves and in turn a parallel electric field were more efficiently generated.    There are two main issues that need to be discussed to address the plausibility of the proposed model: (i) the plausibility of parameters used in the model and (ii) the relation to the observations.  First, the parameter space of the problem is quite large. The level at which the fast magnetosonic wave ($V_x$ and $B_x$) and hence the parallel electric field ($E_z$) amplitudes saturate depends on several parameters. As indicated previously \\citep{bank00,tan01}, this level depends on the strength of the transverse density gradient and the wavelength (width in the case of a Gaussian Alfv\\'enic pulse) of the Alfv\\'en wave. In particular it was found that the stronger density gradients yield lower saturation levels of the fast magnetosonic wave amplitude (due to the fact that destructive wave interference starts earlier when the density gradients are strong), while shorter wavelengths (width in the case of the Gaussian Alfv\\'enic pulse) of the Alfv\\'en wave generate higher levels (e.g. Figs. 11 and 12 in \\citet{tan01}). We have not done a full parameter space investigation to demonstrate the effect, but instead we fixed the transverse density gradient guided by observations. In particular, the observed length scale of the density inhomogeneities  in loops vary between 0.15 and 0.5 loop radii \\citep{g2002}. In our model the length scale of the density inhomogeneity (half-width) is 3 Mm and the loop radius is 10 Mm (see Fig.~1) which makes the ratio 0.3. This is the median value in the observed range (0.15 -- 0.5). This eliminates one parameter of variability in the parameter space. For the Alfv\\'en speed, we used a putative value of  4328 km s$^{-1}$ and performed two numerical runs for two frequencies, 0.7 and 7 Hz. A full investigation should map a range of frequencies. Such an analysis is pending, but see below for preliminary estimates in the context of coronal heating.   \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics{4816f7.eps}} \\caption{Plot of the flux carried by  Alfv\\'en waves from $10^{-4}$ Hz  up to a frequency $\\nu$ versus that frequency as inferred from the empirically constrained spectrum of \\citet{cfk99}. The solid line corresponds to Eq.(18), while the dashed line shows te coronal heating requirement for a temperature of 1 MK. } \\end{figure}  Second, Alfv\\'en waves as observed in situ in the solar wind always appear to be propagating away from the Sun and it is therefore natural to assume a solar origin for these fluctuations. However, the precise origin in the solar atmosphere of the hypothetical source spectrum for Alfv\\'en waves (turbulence) is unknown, given the impossibility of remote magnetic field observations above the chromosphere-corona transition region \\citep{vp97}. Studies of ion cyclotron resonance heating of the solar corona and high speed winds exist which provide important spectroscopic constraints on the Alfv\\'en wave spectrum \\citep{cfk99}. Although the spectrum can and is observed at distances of 0.3 AU, it can be projected back to the base of corona using empirical constraints, see e.g. the top line in Fig. 5 from \\citet{cfk99} (see also the more elaborate model of \\citet{cvb05}). Using the latter figure we can make the following estimates. Let us look at single harmonic, first. At a frequency of 7 Hz (used in our simulations), the magnetic energy of Alfv\\'enic fluctuations is $E_\\nu^{(\\rm 7 \\, Hz)} \\approx 10^7$ nT$^2$ Hz$^{-1}$. For a single harmonic with $\\nu=7$ Hz this gives for the energy density $E^{(\\rm 7 \\, Hz)} \\equiv \\nu E_\\nu^{(\\rm 7 \\, Hz)}/ (8 \\pi) \\approx 7 \\times 10^{-3}$ G$^2/(8 \\pi)\\approx 2.8 \\times 10^{-4}$ erg cm$^{-3}$. Surprisingly this {\\it semi-observational} value  is quite close to the {\\it theoretical} value given by Eq.(13). As we saw above, such a single harmonic can provide approximately 10 \\% of the coronal heating requirement. Next let us look at how much energy density is stored in the Alfv\\'en wave spectrum based on the empirically guided top line in Fig. 5 from \\citet{cfk99}. Their spectral energy density (which they call \"power\") is approximated by the so-called $1/f$ spectrum, i.e. $E_\\nu=0.6 \\times 10^8 / \\nu$ nT$^2$ Hz$^{-1}$. In proper energy density units this is $E_\\nu=2.4 \\times 10 ^{-4} / \\nu$ erg cm$^{-3}$ Hz $^{-1}$. Thus, the flux carried by  Alfv\\'en waves from say $10^{-4}$ Hz  up to a frequency $\\nu$ would be \\begin{equation} F^{\\rm AW}= \\int_{10^{-4} \\, {\\rm Hz}}^{\\nu}E_\\nu c_A^0 d \\nu= 1.04 \\times 10^5 \\ln(\\nu/10 ^{-4}) \\end{equation} $\\left[{\\rm erg \\, cm^{-2} s^{-1}}\\right]$. It is instructive to look at this equation graphically in Fig.(7): the Alfv\\'en wave spectrum from $\\nu=10^{-4}$ Hz up to about a few $\\times 10^3$ Hz carries a flux that is nearly as much as the coronal heating requirement (Eq.(6)) with $T=1$ MK. We do not consider higher frequencies because at about $10^4$ Hz, ions become resonant with circularly polarised  Alfv\\'en waves and dissipation proceeds though Landau resonance -- a well-studied mechanism, but quite different from our non-resonant mechanism of parallel electric field generation.  There are several possibilities of how this flux carried by Alfv\\'en waves (fluctuations) is dissipated. If we consider the regime of frequencies up to $10^3$ Hz, the ion cyclotron resonance condition is not met and hence dissipation is dominated by the mechanism of parallel electric field dissipation formulated in this paper. However, at this stage it is unclear how much energy could actually be dissipated. This is due to the fact that we only have two points, 0.7 Hz and 7 Hz, in our \"theoretical spectrum\". As we saw, a single Alfv\\'en wave harmonic with frequency 7 Hz can dissipate enough heat to account for 10\\% of the coronal heating requirement. Eq.(18) shows how much flux is carried by the  Alfv\\'en waves, while in order to calculate how much of it is actually dissipated depends on {\\it what level $F_E$ from Eq.(14) will attain for each harmonic}. Hence, the flux dissipation through our mechanism  would be \\begin{equation} F^{\\rm D}= \\int_{10^{-4} \\, {\\rm Hz}}^{\\nu}E_\\nu c_A^0 D(\\nu)d \\nu \\end{equation} where $D(\\nu)$ is the \"theoretical spectrum\" of the energy stored in parallel electric fields. For a 7 Hz single harmonic it can be obtained from \\begin{equation} {\\frac{2.4 \\times 10 ^{-4}c_A^0 D(\\rm 7 \\, Hz)}{(\\rm 7 \\, Hz)}}({\\rm 7 \\, Hz})=E_E c_A^0 \\end{equation} with $E_E=3.7471 \\times 10^{-4}$ $\\left[{\\rm erg \\, cm^{-3} }\\right]$ from Eq.(13) rendering $D(\\rm 7 \\, Hz)=1.56$. For a 0.7 Hz single harmonic $E_E$ would be different because as can be seen from Fig.(3) $E_z\\approx 3 \\times 10^{-5}$ (as opposed to $E_z=0.001$ for $\\nu=7$ Hz). Since $E_E$ scales as $E_z^2$, then $D(\\rm 0.7 \\, Hz)=1.56 \\times (3 \\times 10^{-5}/0.001)^2 \\approx 1.4 \\times 10^{-3}$. More numerical runs are needed to fully map $D(\\nu)$. This is deferred until further work is done.  A clear distinction should be made between our model and the ones that use ion cyclotron resonance damping of Alfv\\'en waves. Low frequency (0.001-0.1 Hz) Alfv\\'en waves dominate the power spectrum of fluctuations in the solar wind at and beyond 0.3 AU. These waves are able to transport a significant amount of energy \\citep{cfk99}. On the contrary, high frequency (10-10000 Hz) Alfv\\'en waves are known to damp more easily than their low frequency counterparts, but they are not expected to contain much power. This is problematic for ion cyclotron resonance damping models because Landau resonance of ions occurs at high frequencies (few $10^4$ Hz). Our model on the contrary is of a non-resonant nature. Thus, it is hoped that it can provide enough heating once the frequency range 10$^{-4}-10^3$ Hz at the base of corona is mapped numerically (i.e. once the level attained by the parallel electric field amplitudes for each frequency is determined). The proposed wide spectrum idea for Alfv\\'en waves is not as \"theoretical-rather-than-demonstrable\" as the one proposed by \\citet{tn01} for the slow magnetosonic waves. One could question the validity of the wide spectrum idea for the slow magnetosonic waves in coronal loops as predominantly single harmonics (with periods of 3 and 5 minutes, etc.) are observed. At the same time it is not possible to observe the high frequency waves that  are dissipated already. However, in the case of Alfv\\'en waves, the  presence of a wide spectrum with a frequency range of 10$^{-4}-10^3$ Hz at the base of corona, which is actually observed at 0.3 AU,  seems to be more likely. The possibility of meeting the full coronal heating requirement with the wide spectrum Alfv\\'en waves via the proposed two stage mechanism of parallel electric field generation needs further investigation."
    },
    "0606/astro-ph0606382_arXiv.txt": {
        "abstract": "We report the detection of six discrete, low-luminosity ($L_X < 10^{33}~\\ergs$) X-ray sources, located within 12\\arcsec\\ of the center of the collapsed-core globular cluster M30 (NGC~7099), and a total of 13 sources within the half-mass radius, from a 50~ksec \\chandra\\ ACIS-S exposure.  Three sources lie within the very small upper limit of 1\\farcs9 on the core radius.  The brightest of the three core sources has a luminosity of $L_X\\,\\mbox{(0.5--6~keV)} \\approx 6\\times10^{32}~\\ergs$ and a blackbody-like soft X-ray spectrum, which are both consistent with it being a quiescent low-mass X-ray binary (qLMXB)\\@.  We have identified optical counterparts to four of the six central sources and a number of the outlying sources, using deep \\emph{Hubble Space Telescope} and ground-based imaging.  While the two proposed counterparts that lie within the core may represent chance superpositions, the two identified central sources that lie outside of the core have X-ray and optical properties consistent with being CVs. Two additional sources outside of the core have possible active binary counterparts.  We discuss the X-ray source population of M30 in light of its collapsed-core status. ",
        "introduction": "M30 (NGC~7099) is one of 21 Galactic globular clusters that show strong evidence of having undergone core collapse \\citep{Djorgovski86,Lugger95}.  The collapsed state of the core has been confirmed by \\emph{Hubble Space Telescope} (\\hst) imaging \\citep{Yanny94,Sosin97,Guhathakurta98}.  \\citet{Sosin97} has placed an upper limit of 1\\farcs9 (0.08~pc) on the core radius of M30 from high-resolution \\hst\\ Faint Object Camera imaging.  The extraordinarily high central density of M30, which may exceed $\\sim 10^6~\\msun\\,{\\rm pc}^{-3}$, makes its core and the surrounding power-law cusp one of the highest density environments in the Galaxy. A high rate of stellar interactions is expected in the core and cusp regions, resulting in blue straggler formation via stellar mergers, binary formation via tidal capture and/or 3-body interactions, and binary interactions such as exchange encounters \\citep{Hut92}.  M30 shows some of the strongest evidence of stellar interactions seen in any globular cluster.  It has a very high blue straggler frequency, a very large bluer-inward color gradient, and a large deficit of bright giants in the inner region \\citep{Guhathakurta98,Howell00}.  Thus, the central region of M30 is expected to be a conducive environment for the formation of X-ray binaries.  Like the nearby collapsed collapsed-core globular cluster NGC~6397, M30 does not contain any bright ($L_X \\gtrsim 10^{35}~\\ergs$) low-mass X-ray binaries (LMXBs).  This opens the possibility of studying the distribution of low-luminosity sources near the center of M30, which would be hindered by the presence of bright sources.  Previous \\rosat\\ observations of M30 detected low-luminosity X-ray emission from the vicinity of the core, with about 10$''$ positional accuracy \\citep{Johnston94,Verbunt01}.  Since this is comparable to the size of the cusp region, the \\rosat\\ detection is consistent with either a single source or a centrally concentrated source population.  In order to further investigate the X-ray source population of M30, we have obtained and analyzed a medium-depth \\chandra-ACIS exposure.  As expected, we detected a central population of X-ray sources.  We describe the X-ray data, analysis method, and results in the following sections. ",
        "conclusions": "It is interesting to ask whether the observed size of the X-ray source population in M30 accords with expectation.  \\citet{Pooley03} have shown that for 12 globular clusters observed by \\chandra, the X-ray source population size scales approximately linearly with the predicted stellar encounter rate.  They interpret this as strong empirical evidence that the long-observed overabundance of exotic binaries in clusters is primarily the result of dynamical formation in the dense core regions.  Following \\citet{Verbunt87}, they define a normalized encounter rate $\\Gamma$ as the volume integral to the half-mass radius of $\\rho^2/v$, where $\\rho$ is the total mass density, $v$ is the velocity dispersion, and single-mass \\citet{King66} model structure is adopted for the radial profiles of $\\rho$ and $v$.  Most of the contribution to $\\Gamma$ comes from the core and the immediately surrounding region; a simple computation using an analytic King model shows that about 90\\% of the interactions occur within $2 r_c$.  As noted by \\citet{Pooley03}, this formalism leads to the scaling $\\Gamma \\propto \\rho_0^2 \\, r_c^3/v_0$.  M30 is one of the clusters included in the \\citet{Pooley03} sample and their estimate of its background-corrected population size of $\\sim$5--6 \\chandra\\ sources within the half-mass radius agrees well with their prediction of $\\sim$5 sources from their empirical $N$--$\\Gamma$ relation.  Our estimate of the background-corrected source population size is twice as large, at $10\\pm2$, where the error represents the uncertainty in the background correction only, as do the error bars in Fig.~2 of \\citet{Pooley03}. Our background-corrected source count suggests that M30 has a source excess relative to its computed interaction rate $\\Gamma$, although small number statistics are clearly an issue here.  Somewhat surprisingly, M30 has the fourth smallest interaction rate in the \\citet{Pooley03} sample, notwithstanding its very high central density.  This is a consequence of its small core size. \\citet{Pooley03} adopted values of core radius, central density, and central velocity dispersion from the \\citet{Harris96} compilation (2003 February online version); the values for M30 are $r_c = 3\\farcs6$ and $\\rho_0 = 1.1\\times10^5~\\msun~{\\rm pc}^{-3}$.  We note that given the collapsed-core status of M30, these values are rather uncertain.  \\citet{Sosin97} has placed an upper limit of $1\\farcs9$ on the core radius, suggesting a preferred value of $1\\arcsec$, and has estimated a central density of $8.5\\times10^5~\\msun~{\\rm pc}^{-3}$ from dynamical modeling.  These large changes in $r_c$ and $\\rho_0$ produce partially compensating changes in the value of $\\Gamma$, increasing it by a factor of 1.3 for $v_0$ held fixed and by a factor of 1.7 for $v_0 \\propto \\rho_0^{0.5} r_c$ (as for a King model). Another source of uncertainty is that a single-mass King model does not give a good description of the collapsed-core structure of M30\\@. It is interesting to note that another of the three core-collapsed clusters in the \\citet{Pooley03} sample, NGC~6397, significantly deviates from the empirical $N$--$\\Gamma$ relation in the sense of having an overabundance of X-ray sources relative to the encounter rate computed from the \\citet{Harris96} values of core parameters, including a core radius of $3\\arcsec$.  We note that even with an HST-based analysis of the star count profile of NGC~6397, the core radius value remains uncertain.  \\citet{Taylor05} find a best fit value of $4\\farcs4 \\pm 3\\farcs2$ (1-$\\sigma$ errors).  This uncertainty range leads to a corresponding uncertainty in $\\Gamma$, since $\\Gamma \\propto r_c$ if the central surface brightness and central velocity dispersion are fixed.  Given the uncertainties arising from the core-collapsed status of M30 and its consequently small core size, we suggest that its relation to the other clusters in the \\citet{Pooley03} sample is not entirely clear.  Simple models for core collapse produce the scalings $\\rho_0 \\propto r_c^{-2.2}$ and $v_0 \\propto r_c^{-0.05}$ \\citet{Cohn80}. These result in $\\Gamma \\propto r_c^{-1.4}$, indicating that clusters undergoing core radius oscillations should experience episodic bursts of strongly enhanced binary formation and ejection as the core passes through its densest phases.  Thus, the size of the X-ray binary population in a collapsed-core cluster may undergo substantial fluctuations in time.  In light of these issues, detailed dynamical simulations of the evolution of binary populations in clusters undergoing core collapse would be very helpful."
    },
    "0606/astro-ph0606457_arXiv.txt": {
        "abstract": "Recent ultraviolet spectra from the {\\em Hubble Space Telescope} (HST) and {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) satellites demonstrate that UV line and continuum fluxes observed from Mira~B are increasing back towards the levels that the {\\em International Ultraviolet Explorer} observed in 1979--1980 and 1990--1995, after having been found in a very low state by HST and FUSE in 1999--2001.  The UV emission is associated with accretion of material onto Mira~B from Mira~A's massive wind, so the variability is presumably due to variations in accretion rate.  From wind absorption features, we estimate a Mira~B mass loss rate of $2.5\\times 10^{-12}$ M$_{\\odot}$~yr$^{-1}$, indicating that Mira~B's wind has increased in strength along with the accretion rate.  The UV variability may be associated with a previously reported 14-year periodicity in Mira~B's optical emission. ",
        "introduction": "Mira~AB is the nearest symbiotic system at a distance of only $\\sim130$~pc, \\citep{macp97}, and the only symbiotic binary in which the components of the system have been resolved at wavelengths ranging from X-rays to radio \\citep{mk97,mk05,ldm06}.  The primary (Mira, o~Ceti, HD~14386) is a very large pulsating AGB star (photospheric diameter of $\\sim500$~R$_{\\odot}$), with a massive, cool wind of $\\sim 10^{-7}$ M$_{\\odot}$~yr$^{-1}$ \\citep{pfb88}. The companion (Mira~B, VZ~Ceti) is accreting material from Mira~A, but the exact nature of the accretor is uncertain \\citep{mj84}; it has usually been assumed to be a white dwarf \\citep{ahj54,mk05}. The current angular separation between the components is $\\sim$ $0.6^{\\prime\\prime}$ \\citep{mk97}.  The orbit of the binary is poorly known, but \\citet{jlp02} estimate a long orbital period of $\\sim 500$ years, and this orbit suggests that the distance from Mira~A to Mira~B is close to the apparent separation in the plane of the sky, about 70~AU.  This interacting binary, being one of the few accretion systems in which the components of the system are resolvable, provides a unique opportunity to study the spectral energy distributions of the components of an interacting binary individually, as well as being a laboratory for testing accretion models.  In 1995, the Faint Object Camera (FOC) instrument on the {\\em Hubble Space Telescope} (HST) resolved the Mira system for the first time spatially and spectrally at UV and optical wavelengths \\citep{mk97}. Since then, the system has been monitored from UV to radio wavelengths.  In the past few years, we have been witnessing dramatic changes in Mira~B, especially at UV wavelengths \\citep{mk97,bew01}. Spectra from the Space Telescope Imaging Spectrograph (STIS) instrument on HST imply a substantial decrease in accretion luminosity in the late 1990s, compared with fluxes observed at earlier times by the {\\em International Ultraviolet Explorer} (IUE) and HST/FOC.  Furthermore, a forest of Ly$\\alpha$-fluoresced H$_{2}$ emission lines dominated the HST spectra of Mira~B in 1999, despite not being seen at all in any previous IUE or HST spectra \\citep{bew01,bew02}. Observations from the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) demonstrated that the low accretion luminosity and prominent H$_{2}$ emission persisted into 2001 \\citep{bew04}. On 2003 December 6, {\\em Chandra} observations detected an  unprecedented soft X-ray outburst from Mira~A \\citep{mk05}, further evidence of remarkable dynamicism in the Mira~AB system.  The outburst was also detected in contemporaneous H$\\alpha$ observations and is likely associated with a magnetic flare followed by a mass ejection or jet-like activity \\citep[][Karovska 2006, in prep.]{mk05}.  We here present recent UV observations of Mira~B from HST and FUSE, obtained following the X-ray outburst. We compare these data with past observations by these satellites and by IUE.  We discuss the long term variations of Mira~B and a possible correlation with long term optical variability. ",
        "conclusions": ""
    },
    "0606/astro-ph0606727_arXiv.txt": {
        "abstract": "{ESO spectra of HD\\,65949 show it to be unlike any of the well-known types within its temperature range $\\approx$ 13600K.  It is neither a silicon, nor a mercury-manganese star, though it has a huge Hg II line at $\\lambda$3984.  We estimate $\\log({\\rm Hg/H}) + 12.0 \\approx 7.4$.  This is higher than any published stellar mercury abundance. HD\\,65949 is a member of a nearby open cluster, NGC\\,2516, which is only slightly older than the Pleiades, and has been of recent interest because of its numerous X-ray emission stars, including HD\\,65949 itself, or a close companion. A longitudinal magnetic field of the order of $-290$ Gauss at the 4.7~$\\sigma$ level was very recently diagnosed from accurate circular spectropolarimetric observations with FORS 1 at the VLT.  The spectral lines are sharp, allowing a thorough identification study. Second spectra of Ti, Cr, and Fe are rich.  Mn II is well identified but not unusually strong.  Numerous lines of S II  and P II are found, but not Ga II.  The resonance lines of Sr II are strong.  While many Y II lines are identified, and Nb II is very likely present, {\\it no} Zr II lines were found.  Xe II is well identified.  Strong absorptions from the third spectra of the lanthanides  Pr, Nd, and Ho are present, but lines from the second spectra of lanthanides are extremely weak or absent. Among lines from the heavier elements, those of Pt II are clearly present, and the heaviest isotope, $^{198}$Pt, is indicated.  The uncommon spectrum of Re II is certain, while Os II and Te II are highly probable.  Several of the noted anomalies are unusual for a star as hot as HD 65949. ",
        "introduction": "\\label{se:intro}  HD 65949 may have the highest mercury abundance of any known star.  The unusual strength of Hg II $\\lambda$3984 was already noted by Abt and Morgan (1969) in their study of the open cluster NGC 2516.  They did not call it a manganese star.  The age of the cluster was estimated as log\\,$t$ = 8.2 $\\pm$ 0.1 (Sung, et al. 2002). Several studies indicate an X-ray source closely coincident with the position of HD 65949 (cf. Damiani, et al 2003, Table 4, Source No. 44).   Chemically peculiar (CP) stars in the temperature range corresponding to mid-B to A0 are usually either silicon or mercury-manganese stars.  The former belong to the magnetic sequence, designated CP2 by Preston (1974).    Line identifications for HD 65949 were made with the help of predicted line strengths from a model atmosphere with $T_{\\rm eff} = 13600K$ and $\\log(g) = 4.0$, derived from Str\\\"{o}mgren colors and H$\\beta$ photometry (cf. Moon, 1984). Artificially high abundances of ``metals'' were assumed in order to avoid missing lines. Central intensities greater than 0.001 were derived for all lines ($\\lambda\\lambda$3500-9400) in VALD (Kupka, et al. 2000). The technique was presented by Cowley (1995) and is similar to one available to users of VALD.  A few abundances are derived below.  These are based on a provisional or temporary model with the same $T_{\\rm eff}$ and log($g$) from Str\\\"{o}mgren photometry, but with abundances enhanced by a factor of 3 for atomic numbers (Z) from 3 to 35, a factor of 10 for Z from 36 to 72, and 100 for Z above 73.  Following Dworetsky (2005), we assumed a depleted helium abundance, He/H = 0.0085.  We have placed a detailed line identification list with additional references online: http://www.astro.lsa.umich.edu/~cowley/hd65949. We also provide links to two tables, highlighting the Pt II and Re II spectra. ",
        "conclusions": "HD 65949 is one of a number of CP stars in the cluster NGC 2516, which is slightly older than the Plieades.  It is unclear how this fact might be related to the peculiarities discussed here.   Although HD 65949 probably has the highest mercury abundance of any known CP star, it is not a mercury-{\\it manganese} star. Typically, HgMn stars with similar temperature have the greatest overabundances of manganese.  Whatever mechanisms operate to produce the overabundances of the heaviest elements, have failed to produce the usual anomalies near the iron peak.  Clues to the unusual nature of HD 65949 may lie hidden among the abundances of trace and marginally identified species for which superior observational material is required."
    },
    "0606/astro-ph0606511_arXiv.txt": {
        "abstract": "The study of cosmic microwave background (CMB) polarization is still in a pioneering stage, but promises to bring a huge advancement in cosmology in the near future, just as high-accuracy observations of the anisotropies in the total intensity of the CMB revolutionized our understanding of the universe in the past few years. In this contribution, we outline the scientific case for observing CMB polarization, and review the current observational status and future experimental prospects in the field. ",
        "introduction": "Strong theoretical arguments suggest the presence of fluctuations in the polarized component of the cosmic microwave background (CMB) at a level of 5-10\\% of the temperature anisotropy. A wealth of scientific information is also expected to be encoded in this polarized signal. However, while the existence of anisotropies in the temperature of the Cosmic Microwave Background (CMB) has now been firmly established and accurately measured by several experiments \\cite{boom, max, dasi, wmap}, bringing us high-precision constraints on the parameters of the standard cosmological model, the investigation of the polarized component of CMB anisotropy is still in its infancy.  It is then interesting to review, although briefly, the status of CMB polarization research, both theoretically and observationally. CMB polarization theory is well developed, while observations are complicated and hard to perform. In this contribution, we attempt to give a coverage of both issues: the basics of CMB polarization theory are outlined in Section~\\ref{theory}, while the current and future observational situation is described in Section~\\ref{observations}. ",
        "conclusions": "We have presented a brief overview of the status and prospects of CMB polarization research. The scientific potential of investigating polarization cannot be underestimated: we can obtain crucial information on such topics as the physics of reionization, the energy scale of inflation, we can get much better constraints on cosmological parameters that are currently affected by severe degeneracies and even measure accurately the mass of the neutrino. The field is already in a phase of hectic activity, data of increasing quality are being released and many new missions devised. It is likely that if B modes exist and their level is not surprisingly low, they will be detected from a sub orbital effort within the next few years. However, in order to fully exploit the information in CMB polarization we will have to wait for a new generation satellite, that could hardly be in operation before the 2020s. The long road to this mission will be paved by the experiments under development."
    },
    "0606/astro-ph0606583_arXiv.txt": {
        "abstract": "We review work on the evolution of planetary nebulae and proto-planetaries via magneto-rotational mechanisms showing that a dynamo generated magnetic field can produce the energy and momentum needed to drive pPN and PNe outflows. Angular momentum considerations lead to the conclusion that single stars may not be capable of supporting strong fields for long times. Thus we take the working hypothesis that most PN may form via binary stars. We propose that the grand challenge for PN studies is fully understanding the diverse physical processes at work in binary late stage evolution including the development of disks, fields and outflows. ",
        "introduction": "In recent years the field of Planetary Nebula research has approached something of a cross-roads.  The impressive success of radiation transfer, micro-physical modeling of nebular plasma conditions and 1-D radiation hydro-dynamic models applicable to spherical nebula have provided confidence that an fairly complete understanding of parameters within the nebula had been achieved. This success has also raised the question of the direction for the field. In the absence of an overarching science question, the field of PNe studies is in danger of becoming eclipsed by other domains of interest in astronomy which appear to have more visible and compelling frontiers.  The impression that PNe is a study of \"questions mostly answered\" is wrong however. The last decade has also shown that, in spite of great progress, our understanding of both meso-physical (Huggins, these proceedings) and macro-physical (Balick -these proceedings) nebular characteristics remain far from certain. High resolution images from platforms like the HST have shown that we still lack a well tested model which can explain the shapes of PNe and link them, in a consistent way, to the evolution of the central star.  In this contribution I review some outstanding issues in pPNe and PNe shaping and suggest how these provide us with a new way of looking at PNe studies. The perspective I advocate is that the frontiers of PNe studies are broad and, most importantly, provide links to some of the most pressing questions in a variety of other astronomical environments.  \\subsection{Arguments for MHD Launching {\\it and} Collimation} The need for magnetic fields to aid in shaping the nebula has been discussed before (Balick \\& Frank 2002).  The work of Garcia-Segura, Lopez and others (\\cite{ga99}) has demonstrated that including magnetic fields on nebular scales $R>10^{16} cm$ leads to a wide variety of morphologies including well collimated jets and point symmetry. What these models do not address however is the launching of the wind.  The work of Bujarrabal, Alcolea, Sanchez-Contreras and others makes \"weak field\" models untenable and provides one of the most important challenges to PNe studies.  It also offers a cogent direction forward for the future. In \\cite{buj01} the momentum and energy budgets for both AGB and pPN outflows where derived. Their conclusion was that while AGB outflows could be satisfactorily explained via radiation driving, pPNe outflows showed momentum \"excesses\" that where $10^3$ or higher. Such large values of the ratio of outflow momentum to what could be supplied by photons, $P = L/c$, could not be accounted for via multiple scattering. Thus another source of energy is needed to produce the wind. \\cite{buj01} also claimed that the acceleration timescales for the outflows should be short ($t < 1000$ y).  Thus pPN, and subsequent, PNe may initially form via an outburst or explosion rather a continuous wind.  It also noteworthy that the statistics of pPN and young PN morphologies appear to differ dramatically from fully developed PNe. The samples of pPNe of Bujarrabal and collaborators and of young PNe (Sahai \\& Trauger 1998) are both dominated by strongly bipolar outflows. It remains to be seen if this is a observational bias or not.  If such a distinction between older and younger population maintains after further study, it will indicate that the early shaping of PNe occurs via forces which differ as the nebula matures. ",
        "conclusions": ""
    },
    "0606/astro-ph0606256_arXiv.txt": {
        "abstract": "We describe our program to search for and study the kilo-parsec scale radio jets in a sample of high-redshift (greater than 3.4), flat spectrum quasars using new and archival VLA data.  Two of these radio jets have been imaged with Chandra, and have X-ray counterparts and are briefly discussed. These high-redshift sources are important targets for testing current X-ray jet emission models for kpc-scale jets and follow-up multi-wavelength observations will shed light on this problem. ",
        "introduction": "Chandra X-ray Observatory and Hubble Space Telescope observations have established that X-ray and optical emission from kpc-scale radio jets in active galactic nuclei (AGN) are common (e.g. Stawarz 2003).  In quasars, the X-rays are widely interpreted as inverse Compton (IC) scattered emission off the CMB by electrons in the jet emitting synchrotron radiation at very low radio frequencies (Tavecchio et al. 2000; Celotti et al. 2001). If this is the dominant X-ray production mechanism, it implies quasar jets have large bulk velocities (Lorentz factors, $\\Gamma$$\\sim$3--15) on the observed 10's -- 100's kpc-scale, so that the electrons in the jet frame see a sufficiently boosted source of seed photons to explain the X-rays.  A natural consequence of the IC/CMB model is that high-z quasars should have prominent X-ray jets (Schwartz 2002a) because of the strong dependence of the CMB energy density on redshift: $f_{X}/f_{r}\\propto$ U$_{\\rm CMB}$$\\propto$(1+z)$^{4}$.  The recent detection of a bright X-ray jet in the z=4.3 quasar GB~1508+5714 (Yuan et al. 2003, Siemiginowska et al. 2003), with only a faint radio counterpart (\\S~2), lends support to this scenario.  This jet sees 1--2 orders of magnitude times greater energy density from the CMB than jets at lower redshift (z$\\simlt$2), so its extreme redshift may account, to first order, for its large X-ray to radio luminosity ratio (Cheung 2004). In the framework of the IC/CMB model, variations in $B$ and jet Doppler factor, $\\delta$, can cause significant spread in the observed $f_{X}/f_{r}$ (both along a given jet, and source-to-source), and would smear the possible relation with redshift if only a limited redshift-range is studied.  It is important to identify more jets in high-redshift quasars to observe with both Chandra and the VLA in order to compare with lower-redshift examples. Observations of these systems could help to distinguish between the different models proposed to account for the X-ray emission -- one would not, for instance, expect a z-dependence in the synchrotron X-ray jets.  Unfortunately, very little is known about the extended radio structure of high-redshift quasars. Recent Chandra imaging of a number of z$>$4 radio loud quasars do not reveal significant extended X-ray emission (Bassett et al. 2004; Lopez et al. 2004), although there is no pre-existing information on possible radio structures in these samples.  We recently began a program to search for and study kpc-scale radio jets in high-redshift quasars.  The overall goal of our work is to compile a comparison sample to the X-ray jet detections at lower-redshift to test for a bulk redshift dependence of $f_{X}/f_{r}$.  We do not aim for our sample to be a complete one since the lower-z detections are inhomogenous -- we require only that our high-z sample be comparable in radio luminosity and that the jets are beamed (i.e. are flat-spectrum core-dominated quasars). Ongoing work from this program is discussed.  \\begin{figure*}[t] \\centering \\includegraphics[width=72mm,angle=-90]{fig1.eps} \\caption{[LEFT] Chandra X-ray image of GB~1508+5714 (color) with VLA 1.4 GHz contours overlaid (taken from Cheung 2004). [RIGHT] New VLA 4.9 GHz map confirming the faint radio/X-ray feature $\\sim$2.5'' west of the core (see left panel). The lowest contour level is 67 $\\mu$Jy/bm (2$\\sigma$), and subsequent ones increase by factors of $\\sqrt{2}$ up to a peak of 0.323 Jy per beam (1.50''$\\times$1.05'' at PA=--57.8 deg). } \\label{f1} \\end{figure*} ",
        "conclusions": "Jets are common features of radio-loud AGN, at radio and X-ray wavelengths, at least locally (z$\\simlt$2). Our work suggests that this may be true also at high-redshift (out to z=4.7) at radio wavelengths. Proposed Chandra observations will tell us if this is true also at X-rays. Extending studies to these high-redshifts should help us to distinguish between the competing synchrotron and IC models to determine if there is a redshift dependence in the X-ray jet emission.  \\begin{table*}[t] \\begin{center} \\caption{High-redshift (z$>$3.4) Quasars with (publicly) Known Radio Structures} \\begin{tabular}{|c|c|l|} \\hline\\hline Name       & z   & Description [References]\\\\ \\hline 0201+113 &   3.61  & known 4'' radio jet \\cite{sta90}\\\\ 1239+376 &   3.818 & known $\\sim$6'' radio jet \\cite{tay96} \\\\ 1351--018 &   3.707 & VLBI-scale jet \\cite{fre97} \\\\ 1422+231 &   3.62  & gravitational lens \\cite{pat92} \\\\ 1508+572 &   4.301 & known $\\sim$2.5'' X-ray/radio jet \\cite[][This work]{sie03,yua03,che04} \\\\ 1630--003 &   3.424 & gravitational lens \\cite{win02} \\\\ 1745+624 &   3.889 & VLBI jet \\cite{tay96}; $\\sim$2.5'' X-ray/radio jet \\cite[][This work]{bec92} \\\\ 2000--330 &   3.773 & VLBI jet \\cite{fom00}; new $\\sim$4'' radio jet (This work)\\\\ 2215+020 &   3.572 & VLBI jet; diffuse 7'' radio extension \\cite{lob01} with possible X-ray emission \\cite{sch02b} \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   \\bigskip"
    },
    "0606/astro-ph0606126_arXiv.txt": {
        "abstract": "Many astrophysical models predict a diffuse flux of high-energy neutrinos from active galactic nuclei and other extra-galactic sources. At muon energies above 1 TeV, the upward-going muon flux induced by neutrinos from active galactic nuclei is expected to exceed the flux due to atmospheric neutrinos. We have performed a search for this astrophysical neutrino flux by looking for upward-going muons in the highest energy data sample from the Super-Kamiokande detector using 1679.6 live days of data. We found one extremely high energy upward-going muon event, compared with an expected atmospheric neutrino background of $0.46 \\pm 0.23$ events. Using this result, we set an upper limit on the diffuse flux of upward-going muons due to neutrinos from astrophysical sources in the muon energy range 3.16--100~TeV. ",
        "introduction": "\\label{sec:Introduction}  The GeV-PeV energy range is unexplored territory for neutrino astronomy --- observations of neutrinos at these energies will open a new window on the high energy universe. A wide variety of astrophysical phenomena are expected to produce extremely high energy neutrinos, ranging from active galactic nuclei (AGNs) and gamma-ray bursts (GRBs; \\citealt{2002RPPh...65.1025H,1995PhR...258..173G}) to more exotic sources such as dark matter annihilation or decays of topological defects \\citep{2004IJMPA..19..317S}.  The flux of neutrinos at such high energies is quite small; therefore, large-scale detectors are required. One effective technique for observing high-energy neutrinos with an underground detector is to look for muons produced by $\\nu_{\\mu}$ or $\\bar{\\nu}_{\\mu}$ interacting in the surrounding rock. (Throughout this paper, ``muons'' will refer to both $\\mu^{+}$ and ${\\mu^{-}}$.) The muon range in rock increases with muon energy, which expands the effective interaction volume for high-energy events. Downward-going neutrino-induced muons cannot be distinguished from the much larger flux of downward cosmic-ray muons, but since cosmic ray muons cannot travel through the entire Earth, upward-going muons are almost always neutrino-induced. Thus, upward-going muons provide a suitable high-energy neutrino signal.  At muon energies above $1-10{\\rm {\\; TeV}}$, the upward-going muon flux due to neutrinos from AGNs is expected to exceed the flux due to atmospheric neutrinos~\\citep{Stecker:1995th,Mannheim:1998wp}. This cosmic neutrino flux could be detected either by searching for point sources of high-energy neutrinos or by detecting a diffuse, isotropic flux of neutrinos from unresolved astrophysical sources. A diffuse cosmic neutrino flux would be observed as an excess to the expected atmospheric neutrino flux at high energies. In this analysis, we focus on searching for a diffuse flux of upward-going muons due to neutrinos from astrophysical sources using the highest energy data sample in Super-Kamiokande (Super-K).  This study complements other Super-K searches for astrophysical point sources of high energy neutrinos that use data over a larger energy range \\citep{2006ApJ...652..198A}. In this paper we describe a search for evidence of a high energy astrophysical neutrino flux in Super-K's highest energy upward-going muon sample. In \\S~\\ref{sec:Detector} we describe the Super-Kamiokande detector, and in \\S~\\ref{sec:Event-Selection} we give the details of how we selected candidate events from Super-K's ultra--high-energy sample. We evaluate our selection process with Monte Carlo in \\S~\\ref{sec:High-Energy-Isotropic-MC} and calculate the observed upward-going muon flux in \\S~\\ref{sec:Flux-Calculation}. Sections \\ref{sec:Expected-Atmospheric-Background} and \\ref{sec:Analytical-Estimate} discuss the background due to the atmospheric neutrino flux. Based on the results, we set an upper limit in \\S~\\ref{sec:Upper-Limit} and conclude in \\S~\\ref{sec:Conclusions}. Any necessary estimates and approximations have been made so that they lead to a conservative result for this upper limit.   %auto-ignore ",
        "conclusions": ""
    },
    "0606/astro-ph0606310_arXiv.txt": {
        "abstract": "We study acoustic oscillations (eigenfrequencies, velocity distributions, damping times) of normal crusts of strange stars. These oscillations are very specific because of huge density jump at the interface between the normal crust and the strange matter core. The oscillation problem is shown to be self-similar. For a low (but non-zero) multipolarity $l$ the fundamental mode (without radial nodes) has a frequency $\\sim 300$ Hz and mostly horizontal oscillation velocity; other pressure modes have frequencies $\\gtrsim 20$~kHz and almost radial oscillation velocities. The latter modes are similar to radial oscillations (have approximately the same frequencies and radial velocity profiles). The oscillation spectrum of strange stars with crust differs from the spectrum of neutron stars. If detected, acoustic oscillations would allow one to discriminate between strange stars with crust and neutron stars and constrain the mass and radius of the star. ",
        "introduction": "Strange stars are hypothetical compact stars which are built entirely or almost entirely of strange quark matter (containing light $u$, $d$, and $s$ quarks and possibly electrons). The hypothesis of strange stars is based on the idea of \\citet{Witten} that the strange quark matter is the absolutely stable form of matter even at zero pressure. The hypothesis cannot be definitely confirmed or refuted by available theories and experimental data. Strange stars are attracting permanent attention \\citep[see, e.g., a recent review of][]{Weber2005}. One cannot exclude that at least some stars, which are currently thought to be neutron stars, are in fact strange stars. The theories predict the existence of bare strange stars (with the strange matter extended to the very surface) and strange stars with a thin crust of normal matter extended not deeper than the neutron drip point \\citep*[e.g.,][]{Alcock1986}. The normal crust may occur, for instance, due to accretion of normal matter onto a bare strange star.  Strange stars of masses $M$ much lower than the ``canonical'' neutron-star mass, $M \\sim 1.4\\, M_\\odot$, are predicted to have radii $R \\ll 10$ km, essentially smaller than the radii of neutron stars. However, strange stars with $M \\sim 1.4\\, M_\\odot$ have nearly the same radii, $R \\sim 10$ km, as typical neutron stars.  We will show that strange stars with crusts could be potentionally distinguished from neutron stars by their oscillation spectra. Oscillations of strange stars have been analyzed by a number of authors. \\citet*{Yip1999} studied three types of quadrupole oscillations of strange stars and hybrid neutron stars. They found some differences in oscillation frequencies of such stars. Also, they demonstrated that the damping times of oscillations are sensitive to the model of strange matter. Recently \\citet{Benhar} showed that the combined knowledge of the frequency of emitted gravitational waves and the mass or radius of a compact object would allow one to discriminate between a strange star and a neutron star and set stringent bounds on the parameters of quark matter. Both groups of authors analyzed global oscillations of strange stars.  In this paper we focus on pressure oscillations (p and f modes) of strange star crusts. In our previous papers \\citep{Chug2005,Chug06} we have studied oscillations localized in a neutron star crust due to large multipolarity $l \\gtrsim 100$. Here we show that acoustic oscillations in a strange star crust are very specific even for low $l$ because of a huge density jump at the interface between the normal crust and the quark core. ",
        "conclusions": "We have studied pressure oscillations of strange star crusts.  Our main conclusions are as follows.  (1) The oscillations are almost insensitive to the various modifications of the EOS in the normal crust. The polytropic EOS provides approximately the same eigenfrequencies as the EOS of the accreted crust, except for density discontinuous g modes, which are absent for the polytropic EOS (see Section \\ref{SecNumRes} and Figure \\ref{Fig_Freq}).  (2) The oscillation problem for acoustic modes is self-similar (in the plane-parallel approximation). Once the problem is solved for one stellar model, it can easily be rescaled to strange star models with any mass and radius (but the same EOS and the maximum crust density; see Section \\ref{SecFormal}).  (3) For a thin polytropic crust, the oscillation problem is solved analytically (Sections \\ref{SubSecNonRadOsc} and \\ref{SubSecRadOsc}).  (4) The oscillation spectrum of a strange star crust is {\\it specific}. The frequencies of fundamental modes depend linearly on $l$; the frequencies of p modes are almost independent of $l$. These features are {\\it unmistakable seismic signatures of strange stars with crust} (Section~\\ref{SecNumRes}).  (5) A detection and identification of one fundamental mode and one p mode would enable one, in principle, to infer the mass and radius of a strange star (if the maximum crust density is known) or at least to obtain corresponding upper limits (Section \\ref{SubSecStarParams}).   Therefore, oscillation modes of strange stars with crust are potentially good tools to distinguish these strange stars from neutron stars and to determine their masses and radii. The oscillation frequencies could be detected by radio-astronomical methods very precisely.  A search for these oscillation modes could be useful. Some of them do not damp quickly and can survive for a long time (Section \\ref{SecNumRes}). Pressure modes are robust because they are relatively independent of the thermal state of the crust, and they should not be strongly affected by the crustal magnetic field."
    },
    "0606/astro-ph0606476_arXiv.txt": {
        "abstract": "We present the results of the $UBVR_{\\rm \\scriptscriptstyle C}I_{\\rm \\scriptscriptstyle C}$ surface brightness profiles decomposition of the Seyfert galaxies III~Zw~2, Mrk~506 and Mrk~509. The profiles were modelled as a sum of a Gaussian, a S\\'ersic law and an exponent. A Ferrers bar and a Gaussian ring were added to the model profiles of III~Zw~2 and Mrk~506, respectively. The parameters and the total magnitudes of the structural components were derived. ",
        "introduction": "We present here the first results of an outgoing study aimed to do a detailed decomposition of Seyfert galaxies $UBVR_{\\rm \\scriptscriptstyle C}I_{\\rm \\scriptscriptstyle C}$ surface brightness profiles (hereafter SBPs). Our advantages over Seyfert galaxies SBP decomposition found in the literature are the following: (1) we model explicitly the active nucleus in opposition to some authors who avoid nucleus modelling (e.g. \\cite{chatzichristou01}), (2) we use S\\'ersic rather than de~Vaucouleurs law for bulge, (3) we use a truncated exponential law defined in \\cite{kormendy77} along with a pure exponent in disk modelling, and (4) we model bar/oval/lens/ring components that have been generally skipped by Seyfert galaxies SBPs decomposers (e.g. \\cite{boris02}).  All galaxies to be decomposed were observed at Rozhen NAO of Bulgaria with the 2-m telescope and Photometrics AT200 CCD camera ($0.309\\,\\rm arcsec\\,\\rm px^{-1}$) through a standard Johnson-Cousins $UBVR_{\\rm \\scriptscriptstyle C}I_{\\rm \\scriptscriptstyle C}$ set of filters. The SBPs of the galaxies were extracted fitting ellipses to the galaxian isphotes by means of ${\\rm \\scriptstyle FIT/ELL3}$ command of the ${\\rm \\scriptstyle SURFPHOT}$ context of ${\\rm \\scriptstyle ESO-MIDAS}$ package (see \\cite{slavcheva05} for details).  Our first decomposition results concern the Seyfert galaxies III~Zw~2 (Sy1.0), Mrk~506 (Sy1.5) and Mrk~509 (Sy1.2); the Seyfert types are taken from NED. ",
        "conclusions": ""
    },
    "0606/astro-ph0606195_arXiv.txt": {
        "abstract": "Spherical (nonrotating) accretion flows with small-scale magnetic fields have been investigated using three-dimensional, time-dependent MHD simulations. These simulations have been designed to model high-resolution (quasi) steady accretion flows in a wedge computational domain that represents a small fraction of the full spherical domain. Subsonic and supersonic (super-fast-magnetosonic) accretion flows have been considered. Two accretion regimes have been studied: conservative, or radiatively inefficient, and nonconservative, in which the heat released in magnetic reconnections is completely lost. The flows in both regimes are turbulent. They show the flattened radial density profiles and reduction of the accretion velocities and mass accretion rates in comparison with hydrodynamic Bondi flows. In the conservative regime, the turbulence is more intensive and supported mostly by thermal convection. In the nonconservative regime, the turbulence is less intensive and supported by magnetic buoyancy and various magnetic interactions. We have concluded that steady, supersonic spherical accretion cannot be developed in the presence of small-scale magnetic fields. ",
        "introduction": "Mass accretion onto a central gravitational body plays an important role in the formation and evolution of a large variety of astrophysical objects such as planets, stars, galactic nuclei, galaxies, and clusters of galaxies (see Frank, King \\& Raine 1992). Depending on the amount of the angular momentum $\\ell_0$, which the accreting mass carries at the outer boundary radius $R_{\\rm out}$, accretion flows can take either a disk-like or spherical-like form. Flows with relatively high angular momentum, $\\ell_0\\simeq\\sqrt{GMR_{\\rm out}}$, where $G$ is the gravitational constant and $M$ is the central mass, form centrifugally supported accretion disks (e.g., Shakura 1972). Flows with low angular momentum, $\\ell_0\\ll\\sqrt{GMR_{\\rm out}}$, can form spherical or quasi-spherical accretion flows at radii $R\\ga \\ell_0^2/GM$ in which the centrifugal force is weak and not sufficient to balance gravity (Illarionov \\& Sunyaev 1975). The theoretical study of spherical accretion flows is based on an analytical solution discovered several decades ago by Bondi (1952). This solution describes the idealized case of isentropic nonmagnetic accretion flows. Since that time, the theory of spherical accretion has been significantly developed so that it now includes an understanding the role of different physical mechanisms (such as radiative cooling and heating, magnetic field dissipation, thermal and radiation transports, etc.) and more realistic inner and outer boundary conditions (for reviews, see Frank {\\it et al.} 1992; Kato, Fukue, \\& Mineshige 1998). More recent studies of spherical accretion considered convection (Markovi\\'c 1995), accretion onto a magnetic dipole (Toropin {\\it et al.} 1999), magnetic diffusivity due to turbulence (Shadmehri 2004), and vorticity (Krumholz, McKee, \\& Klein 2005).  This paper continues the numerical study of radiatively inefficient spherical (nonrotating) accretion flows with magnetic fields conducted by Igumenshchev \\& Narayan (2002, hereafter IN; also see Pen, Matzner, \\& Wong 2003 for a similar study). IN have demonstrated, with help of three-dimensional MHD simulations, that the effects of a magnetic field can significantly modify the structure of Bondi-type flows. They argued that even initially weak magnetic fields can produce dramatic changes. The main reason for these changes is the local, nonuniform release of the thermal energy during the dissipation of tangled magnetic fields in reconnections. This release suppresses the inward motion of mass and results in the development of turbulence that is mainly supported by thermal convection. The dissipation of a magnetic field in reconnections is compensated by the efficient amplification of the radial field component in spherical convergent flows (Shvartsman 1971). IN have developed a simple analytical theory of spherical convection-dominated accretion flows (spherical CDAFs\\footnote{Note that IN used instead the name ``convection-dominated Bondi flows\" which we do not adopt here.}), which is similar, with regards to the involved basic physics, to the theory of rotating CDAFs (Narayan, Igumenshchev, \\& Abramowicz 2000; Quataert \\& Gruzinov 2000). This theory predicts the flattened radial density profile $\\rho\\propto R^{-1/2}$ in contrast to the steeper density profile in asymptotic Bondi flows, $\\rho\\propto R^{-3/2}$ (see the appendix, eq.~[A6]). Because of this flattened profile, the mass accretion rate in spherical CDAFs is expected to be significantly lower than the Bondi mass accretion rate in flows with the equivalent outer boundary conditions; \\begin{equation} \\dot{M}\\sim\\dot{M}_{\\rm Bondi}\\left({R_{\\rm in}\\over R_{\\rm out}}\\right), \\end{equation} where $R_{\\rm in}$ is the inner radius of the flows. These properties make the spherical CDAF solution to be a prominent candidate for explaining the phenomenon of dim galactic nuclei contained supermassive black holes, including our Galactic center Sgr A$^*$ (e.g., Melia 1992; Baganoff {\\it et al.} 2001, 2003; Quataert 2002; Ghez {\\it et al.} 2003; Ho, Terashima, \\& Okajima 2003; Soria {\\it et al.} 2006), in which accretion disks are typically not observed. Also, spherical CDAFs can be employed to explain the problem of missing isolated neutron stars in our Galaxy (Treves \\& Colpi 1991; Blaes \\& Madau 1993; Turolla {\\it et al.} 1994; Belloni, Zampieri, \\& Campana 1997; Toropina {\\it et al.} 2003; Perna {\\it et al.} 2003) and the observed properties of isolated stellar mass black holes (e.g., Fujita {\\it et al.} 1998; Agol \\& Kamionkowski 2002).  The generality of the numerical results of IN and Pen {\\it et al.} (2003) were limited in the important aspect that no steady or quasi-steady accretion flows were obtained. These authors studied transient states of the flows, which originated because of specific initial conditions and involved moving outward shocks. These results also suffered from insufficient numerical resolution, especially in the innermost region of the flows. In addition, the assumption of an initial bipolar magnetic field in IN resulted in the domination of a large-scale poloidal field and doughnut-like density distribution near the black hole at the later evolution time (note that similar field topology was proposed by Bisnovatyi-Kogan \\& Ruzmaikin 1974). Unless such large-scale poloidal field can be naturally present in some objects, other possible field configurations, such as a small-scale field with zero net magnetic flux, can likely be natural too in spherical accretion flows. The main goal of the present study is the investigation of the role of small-scale magnetic fields in such flows.  In the numerical aspect, we overcome some of the limitations of the IN's approach employing a new simulation design. In this design, we assume a permanent injection of mass and magnetic field into the computational domain that allows us to obtain steady or quasi-steady accretion flows after performing long-time evolution simulations. The numerical resolution has been improved by adopting spherical coordinates and conducting simulations in the wedge computational domain, which represents a small fraction of the full spherical domain (see Fig.~1). The injected field is assumed to have a small-scale component in the form of radially extended magnetic loops with a zero net magnetic radial flux (see Fig.~2). Modifications of the simulation technique allow us to investigate two limiting energetic regimes: conservative (or radiatively inefficient) and non-conservative (in which the heat from magnetic reconnections is completely lost).  The paper is organized as follows. In Section 2 we describe the simulation technique, initial and boundary conditions, and algorithm of the mass and magnetic field injection. Section 3 presents numerical results, and in Section 4 we discuss these results and make final conclusions. We reproduce some analytic solutions of hydrodynamic accretion flows, including Bondi solution, in the appendix. ",
        "conclusions": "We have numerically investigated quasi-steady MHD spherical accretion flows with imposed small-scale magnetic fields. We have confirmed the previous theoretical expectation and numerical results such that the flows are turbulent and have radial structures different from the Bondi-type accretion flows (e.g., Shvartsman 1971; IN). Turbulence in the MHD flows is developed and supported by the interchange instability, thermal convection, and various magnetic interactions, including magnetic reconnections and buoyancy. The highly nonuniform release of energy in reconnections, which is, in fact, the release of the gravitational energy converted and stored in the form of magnetic field, makes these flows so special and different from the laminar and stable Bondi-type flows (however, see Kovalenko \\& Eremin 1998). We have found that magnetic buoyancy, in addition to thermal convection, can play an important role in the modification of the flow structure, especially in the case of nonconservative (or radiatively efficient) flows in which the convection could not be developed.  The most important consequence of turbulence in our MHD models, both conservative and nonconservative, is the modification of the radial flow structure. Figure~3 demonstrates this modification. Turbulent subsonic MHD Models~B1 and B2 have the flattened time-averaged density profiles, higher temperatures, and lower accretion velocities in comparison with their hydrodynamic counterpart Model~A1. These properties make Models~B1 and B2 more like spherical CDAFs (see IN) than Bondi accretion flows. The theory of spherical CDAFs predicts the flattened power-law density profile $\\rho\\propto R^{-\\sigma}$ in which $\\sigma=0.5$. Models~B1 and B2 have profiles close to this, but more steep, $\\sigma\\approx 0.7$. These steeper profiles can be explained by the influence of the inner boundary condition in our numerical models, whereas the analytic CDAF solution is ``boundary free.\" Abramowicz {\\it et al.} (2002) argued that the proximity of the black hole absorbing boundary makes the inner regions of CDAFs to be advection dominated. In particular, they demonstrated that viscous rotating CDAFs become advection dominated inside $\\sim 50$ to 100 gravitational radii. Our numerical models have a rather small radial range, $R_{\\rm out}/R_{\\rm in}=10$, and, therefore, the effects of the inner boundary can be significant.  Turbulent Models~B1 and B2 have shown the reduction of the accretion rates in comparison with laminar Model~A1 (see right lower panel in Fig.~3). This result qualitatively confirms the predicted reduction of the accretion rate in spherical CDAFs (see eq.~[1]). However, the actual reduction of the accretion rates, which is about 10\\%, shows a poor quantitative agreement with estimate (1). We explain this poor agreement by the limited radial range in our models, whereas estimate (1) was obtained in the limit $R_{\\rm out}\\gg R_{\\rm in}$.  The conservative and nonconservative subsonic MHD models (Models~B1 and B2, respectively) are turbulent, however, the turbulence properties in these models are different. In the conservative model, the turbulent motions are more intensive and mainly supported by thermal convection, which makes this model similar to spherical CDAFs. The nonconservative model has less intensive turbulence, which is supported through various magnetic interactions in which magnetic buoyancy seems to be dominated. In some respect, the magnetic buoyancy acts similar to the thermal convection; it also transports (magnetic) energy outward. This transport can explain the flattened density profile in the nonconservative model, like the convection in spherical CDAFs (see IN). However, the energy transport provided by magnetic buoyancy is less efficient and much weaker than the convection transport.  Our supersonic MHD accretion flows (Models~C1 and C2) have been initiated by injecting a low entropy matter into the computational domain. These flows form quasi-steady Alfvenic shocks that separate the laminar super-fast-magnetosonic outer inflows and the post-shock turbulent nearly Alfvenic inner inflows. We have found that the averaged flow pattern in these models is not steady on the large time scale $\\gg t_{\\rm ff}$. The post-shock regions gradually expand, forcing the Alfvenic shocks to move outward. These regions are quite similar to the subsonic MHD flows in Models~B1 and B2, and, therefore, the gradual expansion of these regions can be explained by the outward energy flux provided by convection and/or magnetic buoyancy. We have found that the Alfvenic shock moves significantly faster in the case of Model~C1, which develops a convection in the post-shock region. The latter is consistent with our observation that the more intensive turbulence and, respectively, the larger outward energy flux happens in convective Model~B1, rather than in nonconvective Model~B2.  Based on our study of the supersonic MHD models, we conclude that {\\it stationary} supersonic (or super-fast-magnetosonic) accretion flows cannot be realized in the presence of small-scale magnetic fields. This conclusion is equally applied to radiatively efficient and inefficient flows. Such supersonic flows will unavoidably create shocks at the equipartition radius and these shocks will be moved outward because of the action of convection and magnetic buoyancy, which are developed in turbulent post-shock regions of these flows. These post-shock regions will continue to expand and, on the large time scale, fill the entire accretion domain, causing the flows to be subsonic everywhere.  The wedge computational domain of our models has the limited opening angle (see Fig.~1), and we use the specific boundary conditions in the angular directions (see Section~2.1) to minimize the influence of this limited-size domain. Definitely, the used boundary conditions could have some effects on the properties of the simulated MHD turbulence, for example, limiting the spatial scales and effecting the motions in the vicinity of the polar sliding boundaries. Other consequence of the employing this domain is the inability to simulate large scale magnetic structures, which can be developed during the reverse cascade of energy from small to large spatial scales in a MHD turbulence. This issue of the limited computational domain should be addressed in future works.  The accretion of mass in our MHD models is accompanied by the reconnection dissipation of the magnetic field in turbulent flows. The rate of dissipation is consistently regulated through feedback mechanisms that also regulate the intensity of the turbulence. To illustrate the dissipation process quantitatively, we consider the nonconservative MHD models in which the change of the time-averaged radial flux of the total energy \\begin{equation} F_R=\\int_{\\Omega_0} R^2 q_R d\\Omega \\end{equation} directly corresponds to the amount of dissipated magnetic energy. The integral in equation (13) is taken over the solid angle of the computational domain $\\Omega_0$ and the radial flux per unit square (see eq.~[3]) is \\begin{equation} q_R=\\rho v_R\\left({v^2\\over 2}+\\epsilon +{P\\over\\rho} +{B^2\\over 4\\pi\\rho}-{GM\\over R}\\right)- {B_R\\over 4\\pi}({\\bf v}\\cdot {\\bf B}). \\end{equation} The solid lines in Figure~12 represent the radial dependence of the total energy fluxes obtained in Models~B2 and C2. These fluxes have been normalized to the flux $GM\\dot{M}_{\\rm in}/R_{\\rm in}$. The dashed lines in Figure~10 correspond to the flux conservation $F_R={\\rm const}$. The difference between the dashed and solid lines at a given $R$ represents the amount of the energy dissipated in the radial range from $R_{\\rm out}$ to $R$. Note that the energy dissipates in the whole volume in Model~B2, whereas it dissipates only in the post-shock region and no energy dissipates in the pre-shock region in Model~C2. The total amount of reconnection losses given in units of gravitational energy at $R_{\\rm in}$ is about 6.5\\% in Model~B2 and about 5\\% in Model~C2."
    },
    "0606/astro-ph0606640_arXiv.txt": {
        "abstract": "We analyze the linear stability of a stalled accretion shock in a perfect gas with a parametrized cooling function ${\\cal L}\\propto \\rho^{\\beta-\\alpha} P^\\alpha$. The instability is dominated by the $l=1$ mode if the shock radius exceeds $2-3$ times the accretor radius, depending on the parameters of the cooling function. The growth rate and oscillation period are comparable to those observed in the numerical simulations of Blondin \\& Mezzacappa (2006). \\\\ The instability mechanism is analyzed by separately measuring the efficiencies of the purely acoustic cycle and the advective-acoustic cycle. These efficiencies are estimated directly from the eigenspectrum, and also through a WKB analysis in the high frequency limit. Both methods prove that the advective-acoustic cycle is unstable, and that the purely  acoustic cycle is stable. Extrapolating these results to low frequency leads us to interpret the dominant mode as an advective-acoustic instability, different from the purely acoustic interpretation of Blondin \\& Mezzacappa (2006).\\\\ A simplified characterization of the instability is proposed, based on an advective-acoustic cycle between the shock and the radius $r_\\nabla$ where the velocity gradients of the stationary flow are strongest. The importance of the coupling region in this mechanism calls for a better understanding of the conditions for an efficient advective-acoustic coupling in a decelerated, nonadiabatic flow, in order to extend these results to core-collapse supernovae. ",
        "introduction": "The recent discovery of a strong $l=1$ instability of stalled accretion shocks in the context of core collapse supernovae (Blondin \\etal 2003, Scheck \\etal 2004, Ohnishi \\etal 2006, Burrows \\etal 2006) has revived the interest in the fundamental stability properties of accretion shocks. This instability could be a major ingredient in the mechanism of acceleration of neutron stars (Scheck \\etal 2004, Janka \\etal 2004, Scheck \\etal 2006a). It was also considered as a means to instigate g-mode dipole oscillations of the accreting neutron star (Burrows et al. 2006). While most of these authors recognized the presence of an advective-acoustic cycle similar to the one found by Foglizzo (2002, hereafter F02) in a different context, Blondin \\& Mezzacappa (2006, hereafter BM06) challenged this interpretation and advocated a purely acoustic mechanism. Understanding the mechanism at work in this instability is a crucial step towards correctly extrapolating its consequences in a more realistic astrophysical situation.  The physics of the advective-acoustic cycle is based on the linear coupling between acoustic and advected perturbations through the flow gradients: both entropy and vorticity perturbations act as source terms for the acoustic wave equation (Foglizzo 2001, hereafter F01, and Foglizzo \\& Galletti 2003). This has been known for decades in the field of jet engines, since the pioneering works of Candel (1972), Howe (1975), Marble \\& Candel (1977) and Abouseif et al. (1984). In the subsonic flow below the stalled shock, this linear coupling is due to the gradients associated with the convergence of the flow, its deceleration, gravity and cooling. The interaction between the shock and the flow gradients gives birth to an advective-acoustic cycle, in which an advected perturbation generates a pressure feedback which triggers, at the shock, a new advected perturbation. Although Foglizzo (2002) proved the instability of the advective-acoustic cycle in an accelerated isothermal flow, the fate of such cycles in a cooled decelerated flow is still an open question: can they account for the instability observed in the simulations of BM06 ?  The first aim of the present study is to clarify the instability mechanism at work, using perturbative techniques. The accretion flow is idealized as a perfect gas passing through a stationary shock, and subject to cooling processes schematically described by a cooling function ${\\cal L}\\propto \\rho^{\\beta-\\alpha} P^\\alpha$, mimicking in the simplest manner the neutrino cooling in the core-collapse context. It is the first time that a linear approach has been used to understand the mechanism of this nonradial instability in the core-collapse context.  A first step is to confirm that the dominating $l=1$ mode identified by BM06 in the linear phase of their numerical simulation indeed corresponds to the most unstable eigenmode of the linear problem.  Beyond the determination of the eigenspectrum and the validation of numerical simulations, we wish to address the question of the instability mechanism, using techniques similar to F02. These techniques allow for a direct interpretation of the full eigenspectrum in terms of the efficiencies of the acoustic cycle and the advective-acoustic cycle. Alternatively, these efficiencies can also be computed in the WKB approximation. We wish to use both methods in order to check whether the instability is of acoustic or advective-acoustic nature. Understanding the nature of the instability leads us to construct, in a companion paper (Foglizzo \\etal 2006), a simple toy model which can be solved analytically and allows us to reach a fundamental understanding of some of the properties of the instability.  The present  paper is organized as follows: the boundary value problem associated with this stalled accretion shock is described in Sect.~\\ref{sect_perturb}, where we establish the boundary conditions at the shock and compare them to those used by Houck \\& Chevalier (1992, hereafter HC92) in the different context of supernova fallback. We determine in Sect.~\\ref{sect_spectrum} the eigenfrequencies of the flows studied by BM06, and compare them with the linear phase of their numerical simulations. Then we investigate in Sect.~\\ref{sect_mechan} the mechanism responsible for this instability using the techniques of F02. The purely acoustic cycle is shown to be stable, and the advective-acoustic cycle is shown to be unstable with respect to $l=1$ perturbations, in the range of validity of our approximations. These results are extrapolated to very low frequency perturbations in Sect.~\\ref{sect_discuss}. The arguments of BM06 are reconciled with the advective-acoustic interpretation of the instability in Sect.~\\ref{sect_BM06}. Results are summarized in Sect.~\\ref{sect_conc}. ",
        "conclusions": "}  This work is the first characterization, through a linear study, of the advective-acoustic instability in a decelerated accretion flow involving cooling processes. Our formulation of the boundary conditions at the shock corrects an error in HC92 concerning non radial perturbations. The numerical solution of this problem confirms the existence of an $l=1$ instability, as found in the numerical simulations of Blondin \\etal (2003), for the two types of cooling functions studied by BM06. A detailed comparison of the growth time and the oscillation period of the dominant mode of the instability revealed discrepancies which can reach $\\sim30\\%$, possibly due to the numerical difficulty of advecting vorticity waves towards the region of deceleration without artificially damping them by numerical viscosity. The optimal grid size near the accretor surface is estimated as a fraction $0.1-0.2\\%$ of the shock distance from the accretor.  \\\\ The main purpose of this study was to clarify the instability mechanism at work for $l=1$ perturbations, with the following results for both types of considered cooling functions $\\alpha=3/2$, $\\beta=5/2$ and $\\alpha=6$, $\\beta=1$: \\par (1) We have proven, for the first time, that an advective-acoustic instability of the $l=1$ mode takes place in a decelerated accretion flow involving cooling processes. The WKB approximation used in this proof required that the shock radius exceeds $\\sim 10$ times the accretor radius.  \\par (2) The low frequency $l=1$ instability occurring when the shock distance is moderate ($r_\\sh/r_\\nabla\\ge2$) has also been interpreted as an advective-acoustic instability, using our conclusion (1) together with a continuity argument: the instability of the low frequency modes can be interpreted in continuity with the instability at higher frequency in a series of flows with larger shock radii. This continuity argument is based on a comparison of both the eigenspectra (Fig.~\\ref{figr5_30}) and the eigenfunctions (Fig.~\\ref{figpression}) of the unstable modes.  \\par (3) The purely acoustic cycle is very stable ($|{\\cal R}|\\le 0.5$) in the range of shock radii and frequencies allowed by our approximations. This result disfavours of the acoustic interpretation of BM06.  \\par (4) The efficiency of the advective-acoustic cycle is an increasing function of the shock distance, which reaches an efficiency $|{\\cal Q}|\\sim 4-5$ for a shock radius $r_\\sh/r_\\nabla\\sim 10$.  \\par (5) We have proposed a simple approximation of the advective-acoustic cycle timescale $\\tau_{\\cal Q}$ based on $\\tau_\\nabla$, defined as the advection time from the shock to the radius $r_\\nabla$ where the velocity gradient is strongest, plus a radial acoustic feedback for the sake of simplicity. This estimate is quite accurate at high frequency, but overestimates the cycle timescale at low frequency by $\\sim 20\\%$.  \\par (6) We have shown that the oscillation period of the fundamental mode is a measure of the timescale $\\tau_{\\cal Q}$. The oscillation period of the most unstable mode is comparable to $\\tau_{\\cal Q}$,  $\\tau_{\\cal Q}/2$ or $\\tau_{\\cal Q}/3$ depending on the shock radius (Fig.~\\ref{figa6b1wcut_wopt}). \\\\  Our efforts to understand the instability mechanism at work in a simplified flow aim at guiding our intuition when interpreting more complex numerical simulations of astrophysical flows (\\ie Scheck \\etal 2006b). The general description of the instability is globally satisfactory, but some fundamental questions are still unanswered. The following three questions may be answered by further studies: \\par (i) What is the maximum efficiency $|{\\cal Q}|_{\\rm max}$ of an advective-acoustic cycle with cooling ? We observed in Figs.~\\ref{figQ0} that $|{\\cal Q}|_{\\rm max}$ increases slowly with the shock distance for both types of cooling function: is this due to the influence of an extended quasi adiabatic region where enthalpy gradients contribute to the advective-acoustic coupling (F01) ? Or is this related to the cooling mechanism in the deceleration region ? How does $|{\\cal Q}|_{\\rm max}$ depend on the adiabatic index $\\gamma$ ? Answering these questions should help us estimate the efficiency $|{\\cal Q}|$ in astrophysical flows with a more realistic equation of state and more elaborate cooling processes. \\par (ii) What are the conditions for the dominance of a $l=2$ mode, as observed in Fig.~\\ref{fig43a32b52l06} for $\\alpha=3/2$, $\\beta=5/2$ for $1.5<r_{\\rm sh}/r_*<1.9$ ? How does the maximum efficiency $|{\\cal Q}|_{\\rm max}$ depend on the degree $l$ of the perturbation ? This question could be important with respect to the asymmetry of the explosion, and the subsequent kick of the neutron star (Scheck et al. 2006a). \\par (iii) Can we better understand the conditions under which the instability is dominated by the fundamental mode, its first or second harmonic ? This question may be important with respect to the explosion mechanism proposed by Burrows \\etal (2006). This mechanism requires the nonlinear transfer of energy from an unstable flow above the neutron star to a gravity mode of the neutron star, whose frequency can be a factor 10 higher. We may expect that the higher the frequency of the advective-acoustic cycle, the easier the excitation of the gravity mode of the neutron star through nonlinear processes.\\\\ Some of these questions can be addressed by further simplifying the accretion flow in order to allow for analytical calculations of the advective-acoustic cycle, free from the low frequency limitations inherent to the WKB approximation. This is the purpose of a companion paper (Foglizzo \\etal 2006)."
    },
    "0606/astro-ph0606530_arXiv.txt": {
        "abstract": "To explain the properties of the most massive low-redshift galaxies and the shape of their mass function, recent models of galaxy evolution include strong AGN feedback to complement starburst-driven feedback in massive galaxies. Using the near-infrared integral-field spectrograph SPIFFI on the VLT, we searched for direct evidence for such a feedback in the optical emission line gas around the z$=$2.16 powerful radio galaxy MRC1138-262, likely a massive galaxy in formation.  The kpc-scale kinematics, with FWHMs and relative velocities $\\lesssim 2400$ \\kms\\ and nearly spherical spatial distribution, do not resemble large-scale gravitational motion or starburst-driven winds. Order-of-magnitude timescale and energy arguments favor the AGN as the only plausible candidate to accelerate the gas, with a total energy injection of $\\sim few \\times 10^{60}$ ergs or more, necessary to power the outflow, and relatively efficient coupling between radio jet and ISM. Observed outflow properties are in gross agreement with the models, and suggest that AGN winds might have a similar, or perhaps larger, cosmological significance than starburst-driven winds, if MRC1138-262 is indeed archetypal. Moreover, the outflow has the potential to remove significant gas fractions ($\\lesssim 50$\\%) from a ${>\\cal L}^{*}$ galaxy within a few 10 to 100 Myrs, fast enough to preserve the observed [$\\alpha$/Fe] overabundance in massive galaxies at low redshift. Using simple arguments, it appears that feedback like that observed in MRC1138-262 may have sufficient energy to inhibit material from infalling into the dark matter halo and thus regulate galaxy growth as required in some recent models of hierarchical structure formation. ",
        "introduction": "Models of structure formation and galaxy evolution have reached a state where the impact of large-scale baryonic feedback on galaxy evolution can no longer be neglected. However, the physics of such processes are complex, and as yet not very well understood, especially at high redshift, where the most massive galaxies formed most of their stars. Therefore, simulations of the ensemble of galaxies typically incorporate these processes, which originate on scales beyond the resolution of typical large scale simulations, as simple relationships and parameterizations, which are in turn tuned to approximately match observed galaxy properties. In spite of these limitations, baryonic feedback plays a crucial role for our understanding of galaxy evolution, perhaps most dramatically at high redshift \\citep[e.g.,][]{croton06}  The impact of feedback related to intense star-formation is now observationally well-established at low and high redshift \\citep[e.g.,][]{heckman03, lehnert96a}, and simulations improved considerably when including it. But significant and likely fundamental discrepancies remain, e.g., when comparing the {\\it observed} mass function of galaxies with the {\\it predicted} mass function of dark matter halos at the upper end of the mass function \\citep{benson03}. From low-redshift studies, starburst-driven ``superwinds'' are known to be efficient in driving out gas only from the comparably shallow gravitational potentials of low mass galaxies \\citep[e.g.,][]{martin99, lehnert99, heckman00}. However, these winds are mainly energy driven, therefore susceptible to radiative losses, and they probably lack the power, efficiency, and velocities necessary to remove significant gas fractions from the deep potential wells of the most massive galaxies and their dark-matter halos. As a consequence, they cannot explain the overabundance particularly of massive dark-matter halos compared to the co-moving galaxy density.  Therefore another powerful feedback mechanism is recently gaining in popularity among researchers modeling and simulating the ensemble properties of galaxies, which are vigorous outflows driven by powerful active galactic nuclei at high redshift \\citep[AGN; e.g.,][]{silk98,springel05}.  Energy- and momentum-driven AGN winds may reach very high energy output through radiation pressure and particle ejection.  The most powerful AGN may eject enough energy and momentum to unbind significant fractions of the total interstellar medium (ISM) of the host galaxy. As they slow and expand, such winds are able to efficiently heat gas in the inter-cluster medium (ICM) and even the intergalactic medium \\citep[IGM; e.g.,][]{nath02}. However, due to the clumpy nature of the ISM in galaxies, rather small working surface of the jet, and AGN producing ``light jets'', they are not generally thought to couple strongly with the ambient ISM in their hosts \\citep{begelman89}.  The observed characteristics of massive galaxies support the hypothesis that strong AGN winds at high redshift may have a significant impact on galaxy evolution and their ensemble characteristics.  Spheroidal and black hole masses in nearby massive systems follow the tight M$_{BH}$-$\\sigma$ relation \\citep{ferrarese02, tremaine02}, which perhaps signals co-eval growth of SMBH and the bulge self-regulated through negative AGN feedback \\citep{silk98}. The most massive galaxies, virtually all spheroidals, are ``old, red, and dead''. They are metal-rich and have luminosity weighted ages consistent with massive star-formation at high-redshift and subsequent passive evolution, perhaps indicating that the gas was effectively removed during the most massive burst \\citep[e.g.,][]{thomas99,thomas05,romano02}.  Maybe the tightest constraint comes from the metal abundance ratios, especially $[\\alpha/Fe]$, which indicate intense star-formation truncated by strong feedback within less than few $\\times 10^8$ yrs \\citep[e.g.,][]{thomas99}. In addition, the exponential cut-off in the bright end of the galaxy luminosity function may be a direct result of the efficiency of AGN feedback \\citep{croton06}. Furthermore, \\citet{best05} find evidence that more moderate feedback of radio-loud AGN might balance gas cooling in the halos of the most massive galaxies at low redshift.  A promising way to identify the fingerprints of a highly efficient mode of AGN feedback at high redshift is by observing the gas kinematics in rapidly growing, massive galaxies at high redshift which host AGN. Powerful high-z Radio Galaxies (HzRGs) fit these requirements: They reside in dense environments \\citep{venemans02, kurk04, best03, stevens03}. Their magnitudes, colors, and continuum morphologies are consistent with large stellar masses and passive evolution at $z\\lesssim 2-3$ \\citep{vanbreugel98, best98}. HzRGs at $z > 2-3$ also contain large masses of molecular gas \\citep[$10^{11}$ \\msun; e.g.,][]{rottgering97}, dust \\citep[$10^{8-9}$ \\msun; e.g.,][]{archibald01, reuland04}, and UV/optical line-emitting gas \\citep[$10^{8-9}$ \\msun; e.g.,][]{vanoijk97}. At the highest redshifts, they are also likely to be forming stars at prodigious rates \\citep[$100-1000$ \\msun\\ yr$^{-1}$; e.g.,][]{dey97}. Curiously, HzRGs are known from longslit-spectroscopy to often have spatially extended broad emission lines with FWHM$ > 1000$ \\kms, indicative of extreme kinematics \\citep[e.g.,][]{mccarthy96,tadhunter91,villar99}. However, with the slits typically aligned along the radio axis, the global role of these kinematics could not be examined more closely, highlighting the need for integral-field observations of HzRGs, preferably at rest-frame optical wavelengths (i.e., in the observed near-infrared) to investigate the overall gas dynamics.  An excellent candidate for such a study is the $z=2.16$ MRC1138-262, with $H=18.04$ and stellar mass $M\\sim 5\\times 10^{11}$ \\msun. VLA interferometry reveals 2 radio jets with complex morphologies \\citep{carilli02}. Its nebulosity extends over $\\sim 20-25$ kpc radius\\footnote{Using a flat $H_0 = 70$\\kms, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmololgy, size scales at $z=2.16$ are 8.3 kpc arcsec$^{-1}$, the luminosity distance is 17070 Mpc, the angular size distance is 1710 Mpc. Cosmic age at that redshift is 3.0 Gyrs.}, and has a significant surrounding over-density of \\ha\\ emitters which indicates it might be in the center of a forming galaxy cluster \\citep{kurk04}. Rest-frame UV spectroscopy implies star-formation rates of $40-70$ \\msun\\ yr$^{-1}$.  By examining the spatially-resolved kinematics and spectral properties of the rest-frame optical emission line gas in MRC1138-262, we wish to investigate the influence of the AGN on the excitation and kinematics of the emission line gas using data obtained with the near-infrared integral-field spectrograph SPIFFI. ",
        "conclusions": "We presented a study of the spatially-resolved dynamics in the optical emission line gas around the z$=$2.16 powerful radio galaxy MRC1138-262, using the near-infrared integral-field spectrograph SPIFFI on UT1 of the VLT. The large-scale kinematics in the emission line nebula of this galaxy are not consistent with the signatures of large-scale gravitational motion or starburst-driven winds. Velocities and FWHMs of $\\sim 800- 2400$ \\kms\\ indicate a vigorous outflow with an estimated total energy which is orders of magnitude higher than suggested for starburst-driven winds. Based on timescale and energy arguments, we conclude that the observed outflow, which is a strict lower limit to the intrinsic amount of outflowing gas, is most plausibly driven by the AGN. This implies a relatively efficient interaction between AGN and interstellar medium, with efficiency $\\epsilon = {\\cal O}(0.1)$, mostly mediated by mechanical energy injection through the radio jet. Radiation-driven winds do not appear powerful enough to explain the observed velocities over the size of the optical line emitting gas.  For MRC1138-262 we estimate from simple energetic arguments that the total energy needed to drive the observed outflow of optical emission line gas is of-order few $\\times$ $10^{60}$ ergs with a mass outflow rate, $\\dot{M}\\gtrsim 400$ \\msun\\ yr$^{-1}$.  If we assume that every luminous QSO near the peak of their co-moving density at z$\\sim$2 has such an outflow phase with similar characteristics to MRC1138-262 we find that powerful AGN eject $\\gtrsim 5 \\times 10^{56}$ ergs Mpc$^{-3}$ of energy into the IGM and $\\gtrsim 0.6-14 \\times 10^{6}$ \\msun\\ Mpc$^{-3}$ of gas, including $\\sim 1-30\\times 10^{4}$ \\msun\\ Mpc$^{-3}$ in metals. These energy and masses are significant compared to the the total binding energy and mass budget of ${\\cal L}^{\\star}$ early-type galaxies in the local Universe. If MRC1138-262 is indeed archetypal, then AGN winds have the potential to be of similar, perhaps even larger, cosmological signficance than starburst-driven winds, especially for the most massive galaxies. They accelerate the outflowing gas to much higher velocities than starburst-driven winds, and much of this material and energy are likely to escape the potential of even the most massive individual galaxy halo.  Outflows like in MRC1138-262 have the potential to explain the observed properties of low-redshift massive galaxies, namely their old, and narrow range of ages and low gas fractions. Moreover, the large mass loss rates and short timescales of their outflows can naturally explain the observed enhancement of [$\\alpha$/Fe] relative metal abundances observed in the most massive galaxies at low redshift. AGN-driven outflows appear to be a plausible mechanism that efficiently suppresses star-formation within a few $\\times 10^8$ yrs.  They also have high enough energies to perhaps stop the accretion flow of matter from the IGM into their dark matter halos.  Thus, if accretion is followed by AGN activity, then AGN feedback may be an effective mechanism for regulating galaxy growth \\citep{croton06}.  Our study comprises only one source, but we have indirect evidence for MRC1138-262 being far from unique. Velocities of $\\sim 1000$ \\kms\\ are not uncommon in the extended emission line gas of powerful radio galaxies, but are generally extracted from longslit spectra aligned with the radio jet axis only. Simulating a longslit spectrum from the SPIFFI data cube, we find good agreement between MRC1138-262 and these samples. The bolometric and radio jet luminosities of MRC1138-262 are not unusual, and the source follows the $K-z$ relationship of high-redshift radio galaxies, if correcting for the luminous broad \\ha\\ line emission. We are currently acquiring rest-frame optical, integral-field data for a larger sample of HzRGs, to make statistically more robust predictions. Preliminary results strongly support the present analysis, giving further evidence that AGN feedback plays a major role in the evolution of the most massive galaxies in the universe."
    },
    "0606/astro-ph0606706_arXiv.txt": {
        "abstract": "TeV gamma rays emitted by GRBs are converted into electron-positron pairs via interactions with the extragalactic infrared radiation fields. In turn the pairs produced, whose trajectories are randomized by magnetic fields, will inverse Compton scatter off the cosmic microwave background photons. The beamed TeV gamma ray flux from GRBs is thus transformed into a GeV isotropic gamma ray flux, which contributes to the total extragalactic gamma-ray background emission. Assuming a model for the extragalactic radiation fields, for the GRB redshift distribution and for the GRB luminosity function, we evaluate the contribution of the GRB prompt and scattered emissions to the measured extragalactic gamma-ray flux. To estimate this contribution we optimistically require that the energy flux at TeV energies is about 10 times stronger than the energy flux at MeV energies. The resulting gamma-ray diffuse background is only a small fraction of what is observed, allowing blazars and other sources to give the dominant contribution. ",
        "introduction": "The nature of the extragalactic gamma ray background emission has been a topic of great interest since EGRET collaboration evaluated its spectrum in the range from 30 MeV to 100 GeV \\citep{Sreekumar:1997un}. The diffuse emission coming from beyond the galaxy was determined by subtracting the contributions of resolved point sources, the diffuse galactic emission, and the instrumental background from the gamma-ray intensities observed by EGRET. The emission is found to be a power law in energy \\begin{equation} \\frac{dN_{\\gamma}}{ dA \\, dt  \\, d\\Omega \\, dE} = (7.32\\pm0.34) \\times {10}^{-6} \\quad {(\\frac{E}{0.451 {\\rm GeV}})} ^{-2.10 \\pm 0.03}  \\, {\\rm {cm}^{-2} {s}^{-1} {sr}^{-1} {GeV}^{-1}} \\end{equation} and is highly isotropic on the sky. \\citet{Mukherjee:1999it} suggested that blazars can explain up to 25 per cent of the extragalactic emission. The hypothesis that the flux is predominantly due to blazars seems to be reinforced by a new evaluation of the extragalactic emission \\citep{Strong:2004ry} which is slightly lower and steeper than that found by \\citet{Sreekumar:1997un}. The result is not consistent with a power law and shows some positive curvature which the authors relate to an origin in blazars emission. \\citet{Kneiske:2004rf} recently suggested that up to 85 per cent of the extragalactic emission could arise from blazars. According to \\citet{Stecker:1996, Stecker:2001} blazars can account for the entire extragalactic $\\gamma$-ray background observed by EGRET. Above 100 MeV normal galaxies contribute from 3 to 10 per cent of the observed diffuse extragalactic flux. However from the analysis of \\citep{Erlykin:1995} the spectrum of normal galaxies seems to differ from the extragalactic diffuse emission spectrum. \\citet{Dar:1995} have suggested that the extragalactic diffuse arises from cosmic rays interacting with intergalactic gas. This contribution seems to disagree with the diffuse gamma ray spectrum according to \\citet{Stecker:1996}. \\cite{Loeb:2000} suggested fossil radiation from shock accelerated cosmic rays during structure formation as possible contribution. \\citet{Stawarz:2006} have estimated that inverse Compton scattering of starlight photon fields by the ultrarelativistic electrons in kiloparsec-scale jets in FR I radio galaxies contributes about one percent to the EGRET extragalactic flux. \\cite{Chi:1989} and \\cite{Wdowczyk} have shown that upscattering of CMB to $\\gamma$-ray energies by cosmic ray electrons and protons does not provide a sinificant contribution to the diffuse extragalactic emission.  However, the question about the origin of the extragalactic emission is still open and can possibly be solved by admitting that different sources contribute to it. In fact all unresolved discrete sources outside the Galaxy contribute to the extragalactic background emission; the problem is clearly to understand how big the different contributions are. As already pointed out by \\cite{Hartmann:2003}, who considered GRBs as a source for the diffuse gamma-ray emission at MeV energies, prompt and delayed emissions from GRBs should also contribute to the diffuse extragalactic emission, especially if some GRBs emit photons in the GeV-TeV energy. In fact, outside the GRB source, due to interactions with cosmic infra-red background photons, most of the high energy GRB photons produce high-energy electron-positron pairs. The pairs inverse Compton scatter off CMB photons and produce secondary photons, which in turn interact with IR photons and generate other pairs. Multiple inverse Compton scatterings occur until the energy of the secondary photons is no longer large enough to trigger pair production with the IR photons. When the energy of the scattered photons is not sufficient to produce a subsequent pair, the photons travel to the Earth without undergoing further absorption. After multiple pair-production and inverse Compton processes, the initial energy of the TeV photons escaping the GRBs will have been shifted to MeV-GeV energies \\citep{Coppi:1997,Plaga:1995,Dai:2002gx,Wang:2004,Razzaque:2004cx}. The data on extragalactic diffuse emission in the MeV-GeV energy range are thus useful to put constraints on high energy emission from gamma ray bursts at GeV-TeV energy.  In the following we assume a flat cosmological model with normalized Hubble constant $h=0.71$, with $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$.  Cosmological distances and volumes are calculated following \\cite{Hogg:1999ad}. ",
        "conclusions": "Current limits for the contribution of blazars and other sources to the extragalactic diffuse emission indicate that between 25 and 50 per cent of the extragalactic diffuse emission is not explained yet and, in principle, can be due to any other unresolved source outside the Galaxy \\citep{Sreekumar:1997un, Mukherjee:1999it, Hartmann:2003, Strong:2004ry, Kneiske:2004rf, Stecker:1996, Stecker:2001, Dar:1995, Loeb:2000, Stawarz:2006}. The prompt and scattered emissions from GRBs should contribute to the diffuse extragalactic emission, too. In fact, outside the GRB source, due to interactions with cosmic infra-red background photons, most of the high energy GRB photons produce produce high-energy electron-positron pairs. The pairs inverse Compton scatter off CMB photons and produce electromagnetic cascades. In this way the beamed GeV-TeV GRB emission is re-processed and converted into an isotropic MeV energy emission, which contributes to the extragalactic diffuse emission.  We have modeled the emission from GRBs as due a low energy synchrotron spectrum and a higher energy IC spectrum. The IC spectrum extends up to Klein-Nishima limit. Also we have assumed the higher energy IC emission from GRBs to be 10 times stronger than the lower energy standard synchrotron emission. This assumption is equivalent to requiring that the GRB inverse Compton Y-paramater is equal to 10. Using the BATSE peak flux distribution \\citet{Guetta:2004fc} derived the GRB luminosity function assumed for our estimates. The luminosity function can be approximated by a broken power law with a break peak luminosity of $4.4 \\times {10}^{51} {\\mathrm {ergs}} {\\mathrm {s^{-1}}}$, a typical jet angle of 0.12 rad, and a local GRB rate of $0.44 h^3 {\\mathrm{{Gpc}^{-3}}} {\\mathrm {{yr}^{-1}}}$. Using the redshift and luminosity distributions derived by \\citet{Guetta:2004fc} we have summed up the prompt and scattered emissions from all GRBs in the universe. The sum of the prompt and scattered emission from all GRBs is shown in Fig.~\\ref{fig5} for a particular choice of GRB parameters. From our optimistic model the $\\gamma$-ray emission from prompt and scattered GRB emissions can provide only a small fraction of the extragalactic diffuse emission at EGRET energies and cannot therefore explain the missing flux. From Fig.~\\ref{fig5} the sum of prompt and scattered emission from synchrotron and 10 times stronger IC emissions from all GRBs is less than the part of the extragalactic diffuse emission not explained by blazars, admitting that the blazars explain between 25 and 50 percent of the isotropic extragalactic emission. Our optimistic GRB model does not overproduce the extragalactic $\\gamma$-ray background. We can thus conclude that TeV energy prompt fluxes from GRBs can be at least 10 times bigger than the synchrotron flux without violating the limit imposed by the extragalactic diffuse emission. In allowing TeV energy fluxes from GRBs of comparable or even stronger intensity to the MeV fluxes, our result is consistent with those models \\citep{Peer} which predict fluences at TeV energies similar to those at MeV energies, the observations of high energy emission from GRB 970417a with Milagrito \\citep{Atkins:2002ws} and the analysis of GRB 941017 done by \\citet{Gonzales}, which pointed out the existence of a second flux from GRB 941017 that cannot be explained as synchrotron emission.  Experiments like Milagro and future missions like GLAST or mini-HAWC should therefore devote part of their efforts to investigate VHE emissions from GRBs.  We note that the recent Swift detection of GRB 060218 \\citep{Campana:2006} suggests that there might be a low-luminosity population of GRBs with a much higher event rate \\citep{Liang:2006}. On the other hand, these bursts have low luminosities and tend to be X-ray flashes \\citep{Campana:2006}, which would compensate their high event rate. Future analyses are needed to reveal the contribution from this category of GRBs.  \\clearpage \\begin{figure}[ht] \\begin{center} \\includegraphics[angle=0,width=12cm]{f5.eps} \\caption {GRB Total Emission from bursts having average bulk Lorentz factors 316, time variability 1 second and duration 20 seconds.  The slopes $\\alpha=-1$ and $\\beta=-2$ for both synchrotron and IC spectra. The ratio between the synchrotron peak flux and the TeV emission peak flux is assumed to be 10. The flux is attenuated by $\\gamma \\gamma$ absorption inside and outside the sources. \\label{fig5}} \\end{center} \\end{figure} \\clearpage"
    },
    "0606/astro-ph0606238_arXiv.txt": {
        "abstract": "We present \\chandra\\ observations of two relatively high redshift \\frii\\ radio galaxies, \\ecaez\\ and \\eclpl\\ ($z=1.785$ and $z=1.956$ respectively), both of which show extended \\xr\\ emission along the axis of the radio jet or lobe.  This \\xr\\ emission is most likely to be due to \\ic\\ scattering of Cosmic Microwave Background (CMB) photons. Under this assumption we estimate the minimum energy contained in the particles responsible.  This can be extrapolated to determine a rough estimate of the total energy. We also present new, deep radio observations of \\eczpa, which confirm some association between radio and \\xr\\ emission along the NE-SW radio axis and also that radio emission is not detected over the rest of the extent of the diffuse \\xr\\ emission. This, together with the offset between the peaks of the \\xr\\ and radio emissions may indicate that the jet axis in this source is precessing. ",
        "introduction": "Because to its unprecedented sensitivity and spatial resolution, \\chandra\\ has imaged a wide range of low- and high-redshift radio sources, revealing their detailed \\xr\\ emission. Several processes can produce \\xr s over large scales from these objects: thermal emission from shocks, as well as \\sync\\ radiation and \\ic\\ scattering of a seed photon field (e.g.  \\citet{HK02}).  The latter can originate from a variety of sources.  Synchrotron photons produced in the jet are up-scattered in the process known as Synchrotron Self-Compton (\\ssc). Thermal photons, reprocessed from the nuclear region, can be scattered into the jet \\citep{scharf4c4117}. A further source of photons is the Cosmic Microwave Background (CMB), which becomes more important with redshift \\citep{feltenrees}. All such photon fields, once up-scattered by the relativistic particles in the jet, produce spectra with the same spectral shape as the synchrotron radiation, because they scatter off the same population of electrons. If \\ic\\ is responsible for the X-ray emission, the radiative lifetimes of (highly energetic) radio \\sync -emitting electrons are typically shorter than for the (less energetic) electrons which give \\ic\\ emission in the \\xr\\ band, so \\ic\\ emission traces an older population of particles and may therefore be offset from, and less spatially compact than, the radio emission. For \\ic\\ scattering to dominate over \\sync\\ emission for a particular Lorentz factor $\\gamma$, the magnetic field energy density must be less than the energy density of the seed photon field (i.e. $\\frac{B^2}{8\\pi}\\le\\mathcal{U}_{\\rm phot}$).  Assuming that this criterion is met, the photon field with the highest energy density will be the most important source of \\ic\\ \\xr s.  At redshift $z<1$, the energy density of the CMB is low and \\ic\\ of the CMB (\\icb) emission is likely to be most important in sources only if the jets are relativistic (\\citealt{0637-752}; \\citealt{celottietal2001}). In sources that are not beamed, \\icb\\ may be dominated by other \\xr\\ emission processes \\citep{sambrunasurvey}. \\xr\\ emission from an increasing number of jets and lobes at these redshifts has been detected and is often explained by \\icb\\ (e.g. \\citealt{sambrunasurvey}, \\citealt{HK02}, \\citealt{elena04}, \\citealt{overzier} and \\citealt{kands}). Some \\xr -bright jets appear to be closely aligned to the line-of-sight, relativistic and extending to distances of the order of 100\\kpc\\ (\\citealt{tavecchiosub2kpc} and \\citealt{marshallsurvey}) and appearing, in some cases (e.g. MRC 1136-135, \\citet{sambruna2jets}), to be systematically decelerating.  Since the energy density of the CMB increases as ${(1+z)}^{4}$, this counterbalances surface brightness dimming with redshift. Therefore, \\icb\\ in powerful high-redshift objects should be detectable across the universe, even if beaming is not common \\citep{beacons}. So far, the search for \\icb\\ in very high-redshift jets has proved elusive: current snapshot surveys of high-redshift ($z > 3.5$) radio-loud quasars have generally failed to detect extended \\xr\\ jet emission (\\citealt{noicz4}, \\citealt{laura} and \\citealt{beacons}).  This paper presents \\chandra\\ observations of two radio galaxies in the redshift range $1.75-2$ (\\ecaez\\ and \\eclpl), with the aim of detecting and characterising their \\xr\\ emission.  These sources, and \\eczpa\\ at $z=1.786$ which is re-examined in this paper, are the highest redshift \\frii\\ radio galaxies (i.e. \\citet{FR} class II objects) to be detected in \\xr s to date\\footnote{see http://hea-www.harvard.edu/XJET/index.cgi as of 2006 Jan $31$}, with the exception of 6C\\thinspace 0905$+$3955 \\citep{6c0905} and 4C\\thinspace 41.17 \\citep{scharf4c4117}.  New, deep radio observations of \\eczpa\\ are also discussed.  The galaxies were chosen because of their power and spatial extent, meaning that they have a large reservoir of relativistic particles resolvable by \\chandra\\ and so are ideal candidates in the search for \\icb. Preliminary reports of these results appear in \\citet{carolinAN} and \\citet{xrayuniverse}.  Throughout this paper, all errors are quoted at $1\\sigma$ unless otherwise stated and the cosmology is $H_{0} = 71$\\kmpspmpc, $\\Omega_{0}=1$ and $\\Lambda_{0} = 0.73$. ",
        "conclusions": "\\label{sec:conclusions}  The \\chandra\\ observations of two relatively high-redshift \\frii\\ radio galaxies, \\ecaez\\ and \\eclpl, show extended \\xr\\ emission which is most convincingly explained by inverse-Compton scattering of CMB photons.  \\eclpl\\ appears to produce \\xr\\ emission from beyond the end of the radio jet, which appears to be harder than the rest of the extended emission.  The possibility of a coincidental background or foreground source cannot be excluded.  New radio data for \\eczpa\\ is presented, confirming the results from \\citet{3c294second} and illustrating that the \\xr\\ emission is much more extended than the radio emission. This and the offset detected between the radio jet axis, which represents current jet activity and the \\xr\\ emission, which traces the jet's previous position, provides some evidence that this source is precessing.  The \\xr\\ luminosity of the extended regions of these sources is used to calculate the energy contained in the electrons responsible for the \\xr\\ emission. This gives a lower limit to the amount of energy in the lobe of the order of a few times $10^{59}$\\erg\\ if \\xr\\ producing electrons are assumed to have a Lorentz factor $\\gamma \\sim 10^3$. The energy estimate increases by a factor of $\\sim 1.5$ when a power-law distribution of electrons is considered and possibly by up to $10^3$ when the proton component is included. \\\\"
    },
    "0606/astro-ph0606712_arXiv.txt": {
        "abstract": "The {\\it Microvariability and Oscillations of Stars (MOST)} satellite observed the B supergiant HD 163899 (B2 Ib/II) for 37 days as a guide star and detected 48 frequencies  $\\la$ 2.8 c~d$^{-1}$ with amplitudes of a few milli-magnitudes (mmag) and less.  The frequency range embraces g- and p-mode pulsations.  It was generally thought that no g-modes are excited in  less luminous B supergiants because strong radiative damping is expected in the core. Our theoretical models, however, show that such g-modes are excited in massive post-main-sequence stars, in accordance with these observations. The nonradial pulsations excited in models between $20M_\\odot$ at $\\log T_{\\rm eff} \\approx 4.41$ and $15M_\\odot$ at $\\log T_{\\rm eff} \\approx 4.36$ are roughly consistent with the observed frequency range. Excitation by the Fe-bump in opacity is possible because g-modes can be partially reflected at a convective zone associated with the hydrogen-burning shell, which significantly reduces radiative damping in the core.  The {\\it MOST} light curve of HD 163899 shows that such a reflection of g-modes actually occurs, and reveals the existence of a previously unrecognized type of variable, slowly pulsating B supergiants (SPBsg) distinct from $\\alpha$ Cyg variables.  Such g-modes have great potential for asteroseismology. ",
        "introduction": "The B star region in the HR diagram contains a wide range of variables.  The most luminous  (luminosity classes of 0-Ia; $\\ga 10^5 L_\\odot$) are known as $\\alpha$ Cyg variables and show semi-regular light and radial-velocity variations on timescales of 1 -- 6 weeks \\citep[e.g.,][]{van98}.  Those variations are thought to be a combination of pulsation and rotational modulation of winds \\citep[e.g.,][1997]{van91,Kau96}.  Pulsations in such very luminous stars can be excited mainly by the Fe-bump of opacity \\citep{Rog92} at a temperature of about $2\\times10^5$K.  The excitation mechanism is, however, more affected by the strange-mode character in more luminous stars \\citep{Kir93,Gla96,Dor00}. The target star discussed in this paper, HD163899 (B2Ib/II), is a supergiant that is less luminous and less affected by winds than the $\\alpha$ Cyg variables.   There are also many less luminous B-type variable stars whose light and radial-velocity variations are caused by radial and nonradial pulsations. These pulsations are excited by the classical $\\kappa$-mechanism at the Fe-bump of opacity. The type of mode excited depends mainly on the effective temperature of the stars; radial and nonradial p-mode pulsations are excited in hotter stars and nonradial g-mode pulsations are excited in relatively cooler stars \\citep[see e.g.,][]{Pam99}. Most of the stars are main-sequence stars. Hotter p-mode pulsators ($\\beta$ Cephei stars) are more luminous or more massive ($\\ga 8M_\\odot$) than the relatively cooler (B3 -- B8) g-mode pulsators, slowly pulsating B (SPB) stars \\citep{Wae91}. Mainly based on line-profile variations, g-modes are known to be excited also in early-type Be stars \\citep[e.g.,][]{Riv03}.   The $\\beta$ Cephei stars have been known for more than a century, but the excitation mechanism was identified only in the early 1990s \\citep{Kir92,Mos92} after the emergence of OPAL opacity tables \\citep{Rog92}. The observational properties of the $\\beta$ Cephei stars are thoroughly discussed by \\citet{Sta05}.  The SPB stars are nearly perfectly confined to the main-sequence band \\citep[e.g.,][]{deC04}, in accordance with the theoretical prediction that the radiative core in a post-main-sequence star strongly damps g-modes where they have very short wavelengths  \\citep{Gau93,Dzi93}. In this paper, however, we show that reflection of pulsation modes at the convective zone associated with the hydrogen-burning shell quenches the radiative damping in the core, and hence the Fe-bump can excite some g-modes in a post-main-sequence massive star with a radiative core.  The {\\it Microvariability \\& Oscillations of STars (MOST)} photometric satellite was launched on June 2003 and is fully described by \\citet{MOST}.  The first scientific results were published by \\citet{Mat04}.  {\\it MOST} observations have significantly clarified the properties of pulsating B stars.  {\\it MOST} detected clear light variations in the rapidly rotating O9Ve star $\\zeta$ Oph \\citep{Wal05a}, which revealed the presence of $\\beta$ Cephei-type pulsations in this line-profile variable.  {\\it MOST} found that the mono-periodic $\\beta$ Cephei star $\\delta$ Ceti is actually multiperiodic with three additional frequencies \\citep{Aer06a}.  {\\it MOST} discovered a new SPB star, HD 163830 (B5II/III), near the edge of the main-sequence band, detecting 20 g-mode frequencies \\citep{Aer06b}.  {\\it MOST} detected g-modes and possibly r-modes in the rapidly rotating Be star HD 163868 \\citep{Wal05b}. The latter observations showed, for the first time, that many g-modes can be excited in a rapidly rotating Be star, and provided a natural interpretation of $\\lambda$ Eri-type variables as SPBe stars. This finding suggests that nonradial pulsations might play a crucial role in Be phenomena.  In this paper, we report another important discovery by the {\\it MOST} satellite~--~g-modes in the B supergiant HD 163899. ",
        "conclusions": "The {\\it MOST} satellite discovered that the B type supergiant HD 163899 pulsates in g-modes as well as in p-modes, showing, for the first time, SPB-type g-mode pulsations in a less luminous B supergiant. We have performed pulsation stability analyses for evolutionary models in the mass ranges of $7\\le M/M_\\odot \\le 20$, and found that g-modes are excited by the Fe-bump of opacity in post-main-sequence stars in a wide range of effective temperatures if the stellar mass is larger than $\\sim 12M_\\odot$.  Our results indicate the existence of a previously unrecognized group of variable stars; slowly pulsating B supergiants (SPBsg) distinct from $\\alpha$ Cyg variables.  The excitation of g-modes in B supergiants becomes possible by a reflection of g-modes at the fully-developed convection zone associated with the hydrogen-burning shell.  This reduces significantly the radiative damping in the radiative core of the star. Since the frequencies of excited g-modes depend on the structure around the convective shell, SPBsg have a great potential for asteroseismology.  HD 163899 is the first member of the SPBsg family.  Strictly speaking, HD 163899, showing both g- and p-modes, is a hybrid of SPBsg and $\\beta$ Cephei-type.  The B supergiant $\\iota$ CMa (B3 Ib/II) could be another SPBsg member.  \\citet{Bal85} obtained a 0.717 c~d$^{-1}$ periodicity for $\\iota$ CMa.  Since the light curve looks irregular, it is likely multi-periodic.  Another possible candidate could be HD 98410 (B2/B3 Ib/II) which is listed in \\citet{Sta05} as one of the ``poor and rejected $\\beta$ Cephei candidates'' with a Hipparcos-deduced period of 1.453 days.  We also note that \\citet{Wae98}'s list for $\\alpha$ Cyg-type supergiants includes some less luminous B supergiants or bright giants such as HD54764 (B1Ib/II), HD141318 (B2II), and HD168183 (B1Ib/II).  These stars could also be SPBsg candidates.  Further observations of these stars are needed to confirm their membership.   Also, for a detailed asteroseismological study of HD 163899, spectroscopic observations are important to properly identify the pulsation modes that are responsible for the photometrically observed frequencies."
    },
    "0606/astro-ph0606654_arXiv.txt": {
        "abstract": "Simulated dark matter profiles are often modelled as a `NFW' density profile rather than a single power law. Recently, attention has turned to the rather rigorous power-law behaviour exhibited by the `pseudo phase-space density' of the dark matter halo, which is defined dimensionally in terms of the local density and velocity dispersion of the dark matter particles. The non-power-law behaviour of the density profile is generally taken to exclude simple scale-free, in-fall models; however the power-law behaviour of the `pseudo-density' is a counter indication. We argue in this paper that both behaviours may be at least qualitatively understood in terms of a dynamically evolving self-similarity, rather than the form for self-similar infall that is fixed by cosmological initial conditions. The evolution is likely due to collective relaxation such as that provided by the radial-orbit instability on large scales. We deduce, from a distribution function given by first order coarse-graining, both the NFW-type density profile and the power-law pseudo-density profile. The results are not greatly sensitive to variation about $3$ in the power of the velocity dispersion used in the definition of the phase space pseudo-density. We suggest that the power $2$ may create the more physical quantity, whose deviations from a power-law are a diagnostic of incomplete relaxation. ",
        "introduction": "Frequently the self-similar in-fall model of structure formation is associated solely with the spherically-symmetric, power-law, purely radial, dynamics that was conclusively defined in the seminal papers by \\citet{fg84} and \\citet{bert85}. In such a restricted formulation, despite the non-linear exactness of the results, this model is not considered to have much application to the hierarchical-merging theory of dark matter halo formation. Moreover the model is generally thought to fail to account for the universality of form found in the simulated dark matter halos, since it predicts instead a memory of local cosmological conditions in the ultimate density power-law.  However this memory is not in fact very far from universality. If the density of the final equilibrium halo, $\\rho$, is written as a power-law in the radius $r$ in terms of the `self-similar index' (see additional discussion below) $\\ad$ as $ \\rho\\propto r^{-2\\ad}$, then the index is in fact only weakly dependent on `reasonable' cosmological conditions \\citep{fg84, hw99} so that there is effective universality.  This is reflected in the relation for the cosmologically fixed index \\citep{fg84, hw99} $\\ad=3\\epsilon/(2(\\epsilon+1))$, where $-\\epsilon$ is the effective power of the initial cosmological peturbation and a `reasonable' value is close to $2$.  Our quantity $\\epsilon$ is defined in terms of the density profile while that of \\citet{fg84} is given in terms of the mass profile. This means that $\\epsilon =3\\epsilon_{fg}$, where $\\epsilon_{fg}$ is the parameter used in \\citet{fg84}. Moreover all strictly radial-orbit simulations show that the index $\\ad$ tends universally to $1$ if it is initially smaller than $1$. This means that the halo density profile can never be flatter than $r^{-2}$, and it is always close to this value for $\\epsilon\\approx 2$.  As an example of how the self-similar index may be coupled to standard cosmological initial conditions, we refer for example to \\citet{peacock99} (p494) where the linear density profile around a maximum in the density field with power spectrum $P(k)\\propto k^n$ gives $2\\ad=3(3+n)/(4+n)$ or $\\epsilon =3+n$. This assumes that the self-similarity is already established in this early collapse phase. The estimate works well on the galactic scale where $n\\approx -1$, so that $\\epsilon \\approx 2$ as does $2\\ad$. In \\citet{hw99} the authors suggest instead $\\epsilon =(3+n)/2$, which is the expected linear profile around an $n\\sigma$ peak rather than about a strict maximum \\citep{peebles93} (p546). This gives better results for cluster scale halos where $n\\approx 1$ and $\\epsilon$ is again nearly $2$. It seems that primordial clusters need not dominate their surroundings in the same way as do primordial galaxies.  The strict power-law behaviour is not expected to be maintained at the edge of the virialized system \\citep{hw99, hld02,henrik04}. This can be due either to mass exhaustion or to tidal truncation \\citep{henrik04, henrik06}. For example in a mass exhaustion situation where $\\epsilon\\rightarrow \\infty$, the $\\ad\\rightarrow 3/2$ and the density profile power goes to $-3$ \\citep{hw99}. The tidally truncated edge goes as $r^{-4}$ \\citep{henrik04}.  It remains evident despite the preceeding caveats that the cosmologically determined self-similarity class can not be extended  all the way to the central regions of the halo , since the density there has been established as flattened relative to $-2\\ad$ \\citep{nfw} ( hereafter NFW), \\citep{taylor01,moore98}. It seems to be the strictly radial nature of the orginal self-similar infall models that prevents them from experiencing certain kinds of relaxation \\citep{mwh06,hjs,barnes05} and hence from evolving a central curved density profile.  It is useful to recall in what follows that the self-similarity index or `class' is defined more generally in \\citet{CH91} and \\citet{henrik97} (see also \\citet{hw95} for the steady case). In these papers it is shown to give the ratio of the length to time dimensions  (that is $L^\\delta/T^\\alpha$) of whatever quantity is currently controlling the self-similarity. The mass dimension is no longer independent because of the need to  maintain $G$ constant during the scaling\\citep{hw95}. This formulation has the advantage of allowing the self-similar index, and thus the controlling quantity, to be a parameter in the relevant equations; which parameter  may change dynamically.    In  \\citet{hld02} and \\citet{henrik04}  an asymmetric version of the self-similar infall model is treated. In general one can treat anisotropic self-similarity in both velocity and real space. Rather than solve the resulting equations by simulation, a series expansion is used in terms of the reciprocal resolution in phase space. The smallness of this parameter (essentially $1/\\alpha$ with $\\ad$ fixed: see details below) is related to the degree of coarse graining both in phase-space and in time. As $\\alpha\\rightarrow \\infty$ the coarse-graining is maximal and the solution is steady. These papers have shown that it is  the maximally  coarse-grained, steady, limit of self-similar infall wherein the density behaves in simple power-law fashion. Higher order descriptions of the dynamics reveal a transitory central `flattening' ( always relative to the pure power-law, purely radial profile above) that is presumably due to collective relaxation. These terms are not always active however, being notably ineffective for the purely radial infall solution. The ultimate central density profile must arise from the dynamical evolution of the self-similar index $\\ad$ (since the steady state density power-law is just the negative of twice this value, just as for radial orbits discussed above) , which phenomenon we refer to as a `running' index  in this paper.   One might protest that the running index is really a way of saying that the self-similar symmetry is lost `en route' to the relaxed state of the central regions. However simulations show an indication of an underlying self-similarity in other ways. One indication is the proportional growth of the virial radius and the NFW scale radius (i.e. constant NFW `concentration') \\citep{mwh06}). Another indication is the power-law behaviour of the dimensional or pseudo phase-space density (the real density being the distribution function or DF) \\citep{taylor01,dehnen05,barnes05, barnes06,austin05}. If then the scaling symmetry is broken during the dynamical relaxation of the central halo, it should only be weakly broken in the sense that we can expect a solution in the form of a perturbed self-similarity. If the evolution is slow relative to a particle crossing time, then at each epoch the system may be regarded as self-similar with the current similarity index. The running index thus corresponds to a kind of `adiabatic' self-similarity.It is precisely in this form that we seek a solution for the DF to first order.  So holding the similarity index strictly constant, as is done in the radial self-similar infall models, does not  allow for effective internal relaxation. There are formal higher order terms available for the radial self-similar infall \\citep{hld02} which in principle describe the approach to a relaxed state, but  they do not arise in the absence of a suitable mechanism. This has been explicitly confirmed in the simulations by \\citet{mwh06}. Analytically, as is emphasized in section 4.1 of this paper, a radially biased DF (correct to first order) does not allow a curved density profile to coexist with a power-law pseudo-density. Such coexistence  is however characteristic of the different self-similarity in the halo centre. The purely radial-shell instability detected in \\citet{hw97} for the original models is thus not the effective mechanism relaxing the halo, which is consistent with the moderate redistribution of energy found in these cases \\citep{hw99}.  In \\citet{hw95} a class of spherically symmetric, stationary, isotropic power-law solutions permitting flatter profiles than the cosmologically fixed radial infall models was found. These solutions include the \\citet{EC97} profile and the NFW profile \\citep{henrik06}, which the cosmologically fixed radial self-similarity can not do. Simulations show \\citep{hjs,mwh06} that such solutions are closer to the state of  the halo centres, although the envelopes are closer to the cosmologically fixed radial self-similar infall \\citep{hw99,lu} until edge effects are encountered. Once again a `running' self-similar index is implied, but such steady  solutions do not describe the temporal approach to equilibrium during which the relaxed state is established. This requires the higher order terms in our series expansion.  As a result of these preliminary considerations, two questions are to be addressed in this paper:  (i) Does the `similarity class' (the value of the parameter $\\ad$ ) evolve dynamically during the halo growth due to collective relaxation in anisotropic systems, and  (ii) can this evolution be such as to destroy the power-law density profile while maintaining a power-law pseudo-density and constant concentration?  The answer to the first of these questions was implicitly assumed to be yes in the discussion of \\citet{henrik06}, since there it was argued that the central similarity class of an isolated core could be set by conserving particles and phase-space volume rather than by the cosmological outer conditions.  A `running' self-similar index is in fact compatible with the formulation of self-similarity pioneered in \\citet{CH91} (see also \\citet{henrik97} and comment below).    In this paper therefore we explore the predictions of the first order term in the coarse-graining series of the self-similar system \\citep{hld02}. This term is time dependent and includes in a transitory fashion (until it becomes too large) the approach to the equilibrium state due to collective relaxation.  In addition the dynamical system is regarded as being adiabatically self-similar so that the index $\\ad$ is free to vary from a value slightly greater than $1$  (where it is set by reasonable scale free cosmological conditions \\citet{fg84} and above) to a value approaching $0.5$ or less  as the density profile power passes through $-1$. It is important to understand however that before this ultimate steady state is reached, the first order terms can curve the density profile in such a fashion as to fit the simulations over a limited radial range. Our conclusion wll be that these same terms do not perturb the pseudo phase-space density so strongly in the same radial range.   In the next section after a brief review  of the method employed, we calculate the first order coarse-grained terms for spherically symmetric self-similarity with velocity anisotropy \\citep{hw95, hld02,henrik04}. The spherical symmetry should not be thought of as an essential limitation since similar equilibrium solutions exist with general symmetry \\citep{henrik04} (appendix A).  Spherical symmetry does however allow additional integrals of the motion which increase the choices available for the DF \\citep{kl92,hw95,henrik04}. Since there are no explicit non-radial forces in spherical symmetry, we are forced to manually reproduce the evolution of the DF from radial orbits to isotropy. The increased choice of DF in spherical symmetry permits this. Thus we are able to choose a core DF in section (4.2) that corresponds to the spherically averaged equilibrium DF, which form persists even in general symmetry \\citep{henrik04}.  Section 3 discusses the calculation of the pseudo phase-space density (referred to frequently as `pseudo-density' for brevity)  in terms of the higher order DF, and subsequently detailed results for several DFs are presented in a section 4. Section five draws our conclusions and gives a brief speculative discussion. ",
        "conclusions": "In this paper we have used two analytic ideas. One was the introduction of the running or adiabatic self-similarity, while the other was  the series coarse-graining of the self-similarity which allows a description of the approach to the final steady state due to internal relaxation. It would in principle predict the relaxation in non-spherical symmetry, particularly if a renormalization improvement were found.  We have used the second technique to deduce a DF to zeroth and first order that can apply in the flat density, isotropic, limit. Such is expected to be the case inside the NFW scaling radius of a halo. This DF appears quite naturally in many places \\citep{Tretal, hw95} in zeroth order. We used this DF to calculate the first order corrections to the density and to the phase-space pseudo-density (generalized slightly by the parameter $e$). We have found a range of the running index $a(e)$ over which the first order correction to the density may be substantial while that of the pseudo-density is negligible.  More generally, the transitory first order term together with the running index exert compensating effects that can maintain the slope of the pseudo-density sensibly constant over a wider range of $a$. This is not expected to change until, and if, a rigorously flat core is detected in the simulations. The special case where $e=2/3$, so that we consider in fact the function $\\rho/\\sigma^2$, should show a very good power-law with slope $-2$ between the virial radius and the scale radius, after  which the slope will show an oscillation (ultimately flattening ) relative to this power-law.  The region of predominatly radial orbits between the NFW scale radius and the virial radius (ignoring edge effects) should be described accurately by the original radial in-fall models \\citep{fg84,bert85}, which predict in our terms a slope of $-((2-3e)a+3e)$ for the pseudo phase-space density. This is  $-2$ for $a=1$, and $-15/8$ for $a=9/8$ and $e=1$. It is also $-2$ for $e=2/3$.  It is important to note that not just any DF allows the pseudo-density to be a power-law while the density flattens. In particular neither a steep density, isotropic DF, nor a radially biased, steep density DF does so. Neither the density nor the pseudo-density flattens markedly in the envelope where the latter DF is relevant, if edge effects are ignored (e.g. \\citep{mwh06}).  Thus we conclude that the remarkable power-law behaviour in the pseudo-density \\citep{taylor01} and its related forms is a consequence of the isotropic DF produced in the core by  processes of relaxation, such as the radial orbit instability \\citep{pol81,me85}  . Because the behaviour is found to be similar for a reasonable range of $e$ about $e=1$ (especially on the low side) we do not think that the form $\\rho/\\sigma^3$ is especially significant (see also \\citet{hansen04}, \\citet{dehnen05}, \\citet{barnes06}). It should be admitted however that currently most simulations \\citep{dehnen05, barnes06} do reveal $e=1$ as the superior power law although other values of $e$ produce power laws superior to that of the density (e.g. \\citep{barnes06}) as we would expect. Should there prove ultimately to be a  preference for $e=1$, it would remain a puzzle according to the present analysis. It might be attributed to higher order behaviour since after all this is a first order calculation. We give an alternate suggestion in our speculation below.  The approximate constancy of the NFW concentration is guaranteed in our treatment only by the fact that the evolution is taken to be adiabatically self-similar throughout. Together with the previous paragraph, we have thus answered in the affirmative, at least tentatively, the questions posed in the introduction.  It seems that we have not penetrated very far into the details of the underlying relaxation however, beyond suggesting that the equations do know about it! But we note that our explanation of the power-law behaviour in the pseudo phase-space density invokes a running or adiabatic self-similarity. This is thought to arise as collective effects change the dominant conserved quantities from those set by initial conditions to those compatible with an isotropic distribution function and an isolated core \\citep{henrik06}. The radial orbit instability is one collective effect that achieves this \\citep{hjs,barnes05, mwh06}, but in spherical symmetry we can only imitate such effects by using appropriate DF's at the beginning and end of the relaxation process.   Finally on a speculative note, we observe that $\\rho/\\sigma^2$ is predicted here to have a slope rigorously $-2$ until significantly inside the NFW scale radius. Such a relation implies that the local Jeans' length (ignoring the difference between the total dispersion and the radial dispersion for the moment) scales linearly with radius. Moreover, as discussed in \\citep{henrik06}, we might expect the collective relaxation to proceed by Landau-damping on waves created below or at this scale. Hence this relation states that the relaxation is effective up to a fixed fraction, probably close to unity, of the local scale. This would be perceived as violent `virialization' on the large scale and asymptotic `thermalization' on the small scale.  The flattening deviation from the rigorous inverse square behaviour of this quantity would  indicate the absence of such relaxation (the Jeans' length becomes greater than the local scale). The apparent superiority of $e=1$ in the current simulations may be due to the expected radial variation of $\\sigma\\propto r^{(1-a)}$. Thus while $\\rho/\\sigma^2$ flattens in the absence of complete relaxation, dividing by an additional factor $\\sigma$ at small $r$ for $a<1$ acts against this flattening since $\\sigma$ is small there. We would predict on these grounds that in simulations with higher resolution $\\rho/\\sigma^2$ should become a better power law than $\\rho/\\sigma^3$.  On this view the essential physics is in the behaviour of $\\rho/\\sigma^2$ and the physical running value of $a$ would be set by setting $C=0$ for $e=2/3$. This suggests that the physical core value of $a$ would be $\\approx 0.8$ in a perfect simulation.  Subsequent flattening of the density would be due to first and higher order terms. Fig.(\\ref{fig:curvefit}) is not changed essentially by using $a=0.8$, so that such flattening is at least possible in a transitory fashion. Ultimately the renormalization group approach used by \\citet{McD06} may allow a longer lived description."
    },
    "0606/astro-ph0606148_arXiv.txt": {
        "abstract": "N-body simulations sample their initial conditions on an initial particle distribution, which for cosmological simulations is usually a glass or grid, whilst a Poisson distribution is used for galaxy models, spherical collapse etc. These pre-initial conditions have inherent correlations, noise due to discreteness and preferential alignments, whilst the glass distribution is poorly defined and computationally expensive to construct.  We present a novel particle distribution which can be useful as a pre-initial condition for N-body simulations, using a simple construction based on a ``quaquaversal'' tiling of space. This distribution has little preferred orientation (i.e. is statistically isotropic), has a rapidly vanishing large scale power-spectrum ($P(k) \\sim k^{4}$), and is trivial to create. It should be particularly useful for warm dark matter and cold collapse simulations. ",
        "introduction": "Numerical simulations have become a very powerful tool for investigating non-linear gravitational phenomenon, such as understanding the evolution and properties of cosmological structures. Since the simulations contain a rapidly increasing number of particles, and probe structures on ever smaller scales, it is timely to address some of the fundamental aspects of the initial conditions used in these simulations.  One such aspect is that the cosmological standard model assumes that the early Universe is statistically isotropic. This is in contrast with the usual choice of {\\em pre-initial} conditions for N-body simulations, most often given by placing particles on a uniform grid, which is intrinsically anisotropic.  While it has been shown that this anisotropy can produce non-physical effects in simulations, such effects are difficult to quantify. For cold dark matter simulations it is thought that the physically relevant correlations quickly grow and dominate over fluctuations due to discreteness. The same is not true for warm dark matter simulations for example.  To address this and related questions numerically it is useful to have alternative pre-initial conditions (PreICs). We construct a novel pre-initial condition by making use of a tiling of three dimensional space called the {\\em quaquaversal tiling} \\citep{conrad}. We will show that this new particle distribution, which is statistically isotropic (i.e. has little intrinsic directionality), has mass fluctuations which decay as rapidly as in a grid. At the same time it is a deterministic structure, and can be trivially generated. We present a C-code for the generation of these structures, which can be downloaded from {\\tt http://krone.physik.unizh.ch/{}\\~{}hansen/qua/}. ",
        "conclusions": "We have studied an alternative to the standard pre-initial conditions for N-body simulations of cosmological structures.  The standard {\\em grid} has strong orientations, and the {\\em glass} is poorly defined and computationally expensive to create. We have therefore considered a particle distribution created starting from an equal volume tiling of space, called the {\\em quaquaversal tiling}. The particle distribution is trivial to create, has virtually no orientation (is statistically isotropic), and has rapidly vanishing large scale power-spectrum. We provide a C-code for the generation of these structures on {\\tt http://krone.physik.unizh.ch/{}\\~{}hansen/qua/}."
    },
    "0606/astro-ph0606181_arXiv.txt": {
        "abstract": "One of the biggest challenges associated with a nulling interferometer-based approach to detecting extra-solar Earth-like planets comes from the extremely stringent requirements of pathlength, polarization and amplitude matching in the interferometer. To the extent that the light from multiple apertures are not matched in these properties, light will leak through the nuller and confuse the search for a planetary signal. Here we explore the possibility of using the coherence properties of the starlight to separate contributions from the planet and nuller leakage. We find that straightforward modifications to the optical layout of a nulling interferometer will allow one to measure and correct for the leakage to a high degree of precision. This nulling calibration relaxes the field matching requirements substantially, and should consequently simplify the instrument design. ",
        "introduction": "One suggested method for finding Earth-like planets orbiting other stars involves a space-based nulling interferometer operating in the thermal IR \\citep{brace78,woolf98,noeck99}. The National Aeronautics and Space Agency has been considering such a mission: ``Terrestrial Planet Finder'' (TPF; \\nocite{coulter} Coulter, 2004), while the European Space Agency is planning an equivalent mission: Darwin \\citep{kalt03}.  The primary challenge in detecting such planets is the extreme contrast ratio between the planet and its primary star ($\\sim 10^{6-7}$ at a wavelength of $10 \\mu$m) and close angular separation ($\\sim 0.1$ arcsec). To prevent the planet light from being overwhelmed by photon noise from the primary, some way of selectively nulling the starlight is required. However, the extreme nulling ratio desired sets very stringent limits on the allowable instrument performance, specifically in terms of pathlength error ($\\delta {\\phi}^2$), polarization ($\\delta {\\alpha}^2$) and amplitude mismatch ($\\delta a^2$), all of which cause ``leakage'' in the null. \\cite{lay04} has shown that for certain nulling configurations, in addition to simple leakage terms of the form $\\delta \\phi^2$ or $\\delta a^2$, second-order coupling effects of the form $\\delta \\phi \\delta a$ lead to stability requirements on the order of $\\sim 1.5$ nm in path and $\\sim 0.1$\\% in amplitude over hour-long timescales.  These performance levels are almost an order of magnitude more stringent than previous estimates, and pose severe challenges to the TPF mission design.  In this paper we explore a simple way of determining the nulling leakage based on exploiting the fact that the leakage light is coherent with the starlight, but not with the planet light.  By making a simple modification to the interferometer back-end it becomes possible to measure the leakage terms and hence calibrate the null. This idea is an adaptation of the ``Synchronous Interferometric Speckle Subtraction'' concept proposed by \\cite{guy04} for coronographic instruments. In the Section 2 we describe the original Bracewell nuller concept. We then introduce the nulling calibration technique and derive its expected performance, as well as show simple simulations of the application. In Section 3 we introduce the more complicated dual-Bracewell architecture actually being considered for TPF, as well as show how the calibration concept could be applied to it. We discuss the effects this might have on the TPF system design in Section 4. ",
        "conclusions": "We discuss the concept of coherent calibration of an interferometric null and its application to the Terrestrial Planet Finder instrument. We find that such an approach to calibration greatly relaxes the required levels of electric field matching, and associated stability requirements. This should improve instrument performance, particularly at shorter wavelengths where the effects of path-length control limitations are most severe."
    },
    "0606/astro-ph0606462_arXiv.txt": {
        "abstract": "Excited Si$^+$ and Fe$^+$ species are routinely observed in the host environment of $\\gamma$-ray burst (GRB) afterglows, but are not commonly seen in other extragalactic locations.  Their presence signals unusual properties in the gaseous environment of these GRB hosts that arise either as a result of the intense ionizing radiation of the afterglow or through collision excitation in a dense cloud.  In particular, the photon pumping scenario has explicit expectations for temporal variation in the strength of the excited lines, owing to the decline in the ionizing flux of the GRB afterglow.  We analyze afterglow spectra of GRB\\,020813 obtained in two epochs of $\\approx$\\,16 hours apart, and examine transitions from the first excited state of Fe$^+$ at $J=7/2$ in these two sets of data. We report a significant decline by at least a factor of five in the equivalent width of the \\ion{Fe}{ii}\\,$\\lambda$2396 transition, the strongest from the $J=7/2$ state. We perform a Monte-Carlo analysis and determine that this temporal variation is present at more than 3\\,$\\sigma$ level of significance. This observation represents the first detection in the temporal variation of the excited Fe$^+$ states in the GRB host ISM, a direct influence of the burst itself on its environment.  We further estimate that the Fe$^+$ gas resides in 50-100 pc from the afterglow, based on the afterglow lightcurve and the presence and absence of the excited \\ion{Fe}{ii}\\,$\\lambda$2396 in the two-epoch observations. ",
        "introduction": "Long-duration $\\gamma$-Ray Bursts (GRBs) are believed to trace the death of massive, short lived stars \\citep[see][ for review]{woosley06}. The brightness of GRB afterglows for a brief period makes them visible throughout most of the observable Universe \\citep[e.g.][]{kawai06} and therefore a sensitive probe of both the interstellar medium (ISM) of their host galaxies and the intervening gas \\citep[e.g][]{vreeswijk04,chen05,prochter06}. The study of metal absorption lines detected in the afterglow spectra illuminates the physical conditions in the ISM hosting the GRB, such as the \\ion{H}{i} column density, the gas metallicity, the dust-to-gas ratio, and the kinematics.  Moreover, the access to features belonging to the circumstellar medium of the progenitor star of the GRB may yield strong constraints on the progenitor and may help to understand the impact of the GRB on the local environment (e.g.\\ Ramirez-Ruiz \\etal\\ 2005; van Marle, Langer, \\& Grac\\'ia-Segura 2005; Prochaska, Chen, \\& Bloom 2006, hereafter PCB06).  Interestingly, the transitions from excited states of C$^+$, Si$^+$, and O$^0$ appear to be a generic feature in GRB host environments, while they are observed through only rare sightlines, such as broad absorption-line (BAL) quasars \\citep{hall02}, and in gas near stars like $\\eta$~Carinae and $\\beta$~Pictoris \\citep{gull05,lagrange88}. In addition, a large fraction of the GRB hosts also show excited Fe$^+$ lines \\citep{prochaska06}. Does the detection of these excited state absorption lines reveal extreme (pre) circumburst environments or are they the result of work done by the afterglow in early times after the burst?  To answer this we need to consider competing scenarios for the production of excited-state transitions, and particularly for Fe$^+$: (1)~collisional excitation and (2)~photon pumping. PCB06 have recently shown that in the presence of an intensified UV radiation field from the afterglow --~known to be stronger than $10^5$ to $10^6$ times the ambient far-UV radiation field in the ISM of the Milky Way~-- indirect UV pumping alone can produce the excited states observed in GRB hosts without imposing extreme gas density and temperature\\footnote{A third scenario is for the bright progenitor star, such as a WR star, to UV pumping these excited Fe$^+$. However, as we will show below, at the distance allowed for Fe$^+$ to survive the intense afterglow UV radiation ($\\sim 50$ pc) the excitation rate from the WR star is insufficient to produce the observed excited states. In addition, the excitation rate due to direct IR pumping by the afterglow is orders of magnitude smaller than the excitation rate due to afterglow UV pumping.}. Under the photon pumping conditions, temporal variation in the absorption-line strength of excited state transitions is expected, because of the decline in the ionizing flux of the afterglow. \\citet{perna98} and \\citet{bottcher99} predicted time-dependent resonance absorption features in the GRB afterglow spectra, such as the \\ion{Mg}{ii} doublet, over a timescale of a few days to a few weeks. Their analysis also indicated that important constraints on the extent and gas density of the circumburst medium can be derived. However, searches for absorption-line variability have so far yielded null results \\citep{vreeswijk01,mirabal02}.  Under the photon pumping scenario, temporal variation in excited transitions on scales of a few hours is expected, because the decay timescale for the excited states ranges from 8 to 15 minutes. Given the new insights into the origin of excited lines, in this {\\em Letter} we analyze afterglow spectra of GRB\\,020813 obtained in two different epochs of $\\approx\\,16$ hours apart.  These data sets have a suitable time coverage to serve the purpose to search for variation in the excited Fe$^+$ lines.  We report a significant decline in the equivalent width of the strongest transition from the excited state Fe$^+$ $J=7/2$. This observation represents the first detection in temporal variation of the excited Fe\\,II states in GRB host ISM, lending strong support for the indirect UV pumping scenario as the primary excitation mechanism of this gas. ",
        "conclusions": "We have compared afterglow spectra of GRB\\,020813 obtained roughly 16 hours apart.  Significant variation in the strength of the excited Fe\\,II\\,$\\lambda$2396 feature identified in the GRB host is detected at $>3$\\,$\\sigma$ level of significance.  A natural consequence of the excited lines being produced by UV pumping is that these excited levels are populated according to the flux of the radiation field. As the GRB afterglow fades, one must witness the decay of the excited states (PCB06).  Our analysis of the afterglow spectra of GRB\\,020813 presents the first supporting evidence for the UV pumping model.  For the Fe$^+$ excited levels, the half-life for spontaneous emission is $\\approx 10$ minutes, significantly smaller than the characteristic decay time of the GRB afterglow (Fig.~\\ref{lightcurve}).  Therefore, the population of the excited states of Fe$^+$ is expected to track the UV flux.  At $t_{\\rm obs} \\approx 4.7$ and $t_{\\rm obs} \\approx 20.8$ hours, the excitation rates of the \\ion{Fe}{ii} $J=7/2$ level due to indirect UV pumping are $\\approx 25$ s$^{-1}$ and 5 s$^{-1}$, respectively.  The excitation rates are estimated adopting the temporal declining index $\\alpha=-1.09$ from \\citet{li03} and a spectral index $\\beta=-1$ from \\citet{covino03}.  They are roughly consistent with the measured equivalent width and limit from the two-epoch observations.  Following PCB06, we find that the Fe$^+$ gas must lie beyond 40 pc from the GRB afterglow in order to survive the intense ionizing radiation field, but within 100 pc of the GRB afterglow such that it will significantly populate the Fe$^+$ levels at $t_{\\rm obs} \\approx 4.7$ hours. In the case of GRB\\,020813, the absence of \\ion{Fe}{ii}\\,$\\lambda$2396 at $t_{\\rm obs} \\approx 20.80$ hours further constrains the Fe$^+$ gas at $> 50$ pc from the afterglow. This is consistent with the presence of Mg$^0$ identified in the host environment \\citep{barth03,savaglio04} which implies a distance $>30$ pc (PCB06). The results presented here stress the value of obtaining both early and late-time observations of GRB afterglows.  While acquiring the observations using the same telescope and instrument is preferred, we demonstrate that precise equivalent width measurements using data from different facilities can also yield strong constraints for the temporal evolution of the excited levels.  Finally, a time-dependent calculation of radiative transfer and ion excitation may help further reveal the geometry and physical conditions of the gas."
    },
    "0606/hep-ph0606157_arXiv.txt": {
        "abstract": "At ultra-high density, matter is expected to form a degenerate Fermi gas of quarks in which there is a condensate of Cooper pairs of quarks near the Fermi surface: color superconductivity. In this chapter we review some of the underlying physics, and discuss outstanding questions about the phase structure of ultra-dense quark matter.  We then focus on describing recent results on the crystalline color superconducting phase that may be the preferred form of cold, dense but not asymptotically dense, three-flavor quark matter.  The gap parameter and free energy for this phase have recently been evaluated within a Ginzburg-Landau approximation for many candidate crystal structures.  We describe the two that are most favorable. The robustness of these phases results in their being favored over wide ranges of density.  However, it also implies that the Ginzburg-Landau approximation is not quantitatively reliable.  We describe qualitative insights into what makes a crystal structure favorable which can be used to winnow the possibilities. We close with a look ahead at the calculations that remain to be done in order to make quantitative contact with observations of compact stars. ",
        "introduction": "\\label{sec:intro}  The exploration of the phase diagram of matter at ultra-high temperature or density is an area of great interest and activity, both on the experimental and theoretical fronts. Heavy-ion colliders such as the SPS at CERN and RHIC at Brookhaven have probed the high-temperature region, creating and studying the properties of quark matter with very high energy density and very low baryon number density similar to the fluid which filled the universe for the first microseconds after the big bang. In this paper we discuss a different part of the phase diagram, the low-temperature high-density region. Here there are as yet no experimental constraints, and our goal is to understand the properties of matter predicted by QCD well enough to be able to use astronomical observations of neutron stars to learn whether these densest objects in the current universe contain quark matter in their core. We expect cold, dense, matter to exist in phases characterized by Cooper pairing of quarks, i.e. color superconductivity, driven by the Bardeen-Cooper-Schrieffer (BCS)\\cite{BCS} mechanism. The BCS mechanism operates when there is an attractive interaction between fermions at a Fermi surface. The QCD quark-quark interaction is strong, and is attractive in many channels, so we expect cold dense quark matter to {\\em generically} exhibit color superconductivity. Moreover, quarks, unlike electrons, have color and flavor as well as spin degrees of freedom, so many different patterns of pairing are possible. This leads us to expect a rich phase structure in matter beyond nuclear density.  Calculations using a variety of methods agree that at sufficiently high density, the favored phase is color-flavor-locked (CFL) color-superconducting quark matter\\cite{Alford:1998mk} (for reviews, see Ref.~\\refcite{Reviews}).  However, there is still uncertainty over the nature of the next phase down in density. Previous work\\cite{Alford:2003fq,Alford:2004hz} had suggested that when the density drops low enough so that the mass of the strange quark can no longer be neglected, there is a continuous phase transition from the CFL phase to a new gapless CFL (gCFL) phase, which could lead to observable consequences if it occurred in the cores of neutron stars.\\cite{Alford:2004zr} However, it now appears that some of the gluons in the gCFL phase have imaginary Meissner masses, indicating an instability towards a lower-energy phase.\\cite{Huang:2004bg,Casalbuoni:2004tb,Giannakis:2004pf,Alford:2005qw,Huang:2005pv,Fukushima:2005cm,Gorbar:2005rx,Kryjevski:2005qq,Gorbar:2006up,Iida:2006df,Fukushima:2006su,Gerhold:2006dt} Analysis in the vicinity of the unstable gCFL phase cannot determine the nature of the lower-energy phase that resolves the instability. However, the instability is telling us that the system can lower its energy by turning on currents, suggesting that the crystalline color superconducting phase,\\cite{Alford:2000ze} in which the condensate is modulated in space in a way that can be thought of as a sum of counterpropagating currents, is a strong candidate. ",
        "conclusions": ""
    },
    "0606/astro-ph0606409.txt": {
        "abstract": "{An analysis is presented of VLT-FLAMES spectroscopy for three Galactic clusters, NGC\\,3293, NGC\\,4755 and NGC\\,6611. Non-LTE model atmosphere calculations have been used to estimate effective temperatures (from either the helium spectrum or the silicon ionization equilibrium) and gravities (from the hydrogen spectrum). Projected rotational velocities have been deduced from the helium spectrum (for fast and moderate rotators) or the metal line spectrum (for slow rotators). The origin of the low gravity estimates for apparently near main sequence objects is discussed and is related to the stellar rotational velocity. The atmospheric parameters have been used to estimate cluster distances (which are generally in good agreement with previous determinations) and these have been used to estimate stellar luminosities and evolutionary masses. The observed Hertzsprung-Russell diagrams are compared with theoretical predictions and some discrepancies including differences in the main sequence luminosities are discussed. Cluster ages have been deduced and evidence for non-coeval star formation is found for all three of the clusters.  Projected rotational velocities for targets in the older clusters, NGC\\,3293 and NGC\\,4755, have been found to be systematically larger than those for the field, confirming recent results in other similar age clusters. The distribution of projected rotational velocities are consistent with a Gaussian distribution of intrinsic rotational velocities. For the relatively unevolved targets in the older clusters, NGC\\,3293 and NGC\\,4755, the peak of the velocity distribution would be 250 \\kms with a full-width-half-maximum of approximately 180 \\kms. For NGC\\,6611, the sample size is relatively small but implies a lower mean rotational velocity. This may be evidence for the spin-down effect due to angular momentum loss through stellar winds, although our results are consistent with those found for very young high mass stars. For all three clusters we deduce present day mass functions with $\\Gamma$-values in the range of -1.5 to -1.8, which are similar to other young stellar clusters in the Milky Way. ",
        "introduction": "As part of a European Southern Observatory Large Programme, we are using the Fibre Large Array Multi-Element Spectrograph (FLAMES) at the Very Large Telescope to survey the hot stellar populations in our Galaxy and in the Magellanic Clouds. We have observed in excess of 50 O-type stars and 500 B-type stars, in a total of seven clusters. These data have been supplemented with observations of our brighter targets using the Fibre-Fed Extended Range Optical Spectrograph (FEROS) at La Silla. Evans et al. (\\cite{Eva05}; hereafter Paper I) have given an overview of the scientific goals of the project, discussed the target selection and data reduction techniques, and presented some preliminary analysis for three Galactic clusters. The scientific aims of the survey as discussed in Paper I, included understanding how rotation and metallicity influence stellar evolution and stellar winds, calibration of the wind momentum-luminosity relationship and the nature of supernova progenitors. In the case of stellar evolution, theoretical studies (see, for example, Heger \\& Langer \\cite{Heg00}; Maeder \\& Meynet \\cite{Mae00}) provide predictions of the surface enhancements expected for helium and nitrogen. Of particular relevance to our survey are the larger logarithmic enhancements predicted for lower metallicity regimes (Maeder \\& Meynet \\cite{Mae01}). Although there are many previous detailed quantitative studies of O- and B-type stars (Crowther at al. \\cite{Cro02, Cro05}; Bouret et al. \\cite{Bou03}; Hillier at al. \\cite{Hil03}; Evans et al. \\cite{Eva04}; Trundle et al. \\cite{Tru04}; Walborn et al. \\cite{Wal04}; Dufton et al. \\cite{Duf05}; Trundle \\& Lennon \\cite{Tru05}), these have tended to concentrate on narrow lined (and hence possibly slowly rotating) stars, whilst the sample sizes have been relatively small. The current survey will sample both the main sequence and supergiant regimes and include rapidly rotating stars, thereby extending the scope of such studies.  In order to have a sample of stars in the Milky Way to compare with the FLAMES datasets for the SMC and LMC we have obtained spectra of likely cluster members in the three Galactic Clusters, NGC\\,3293, NGC\\,4755 and NGC\\,6611. These clusters were chosen so as to have a similar age range to their lower metallicity Magellanic Cloud counterparts. Although the Milky Way has an abundance gradient (see, for example, Afflerbach et al. \\cite{Aff97}, Rolleston et al. 2000, Gummersbach et al. 1998, Daflon \\& Cunha 2004) these clusters are relatively close to the solar position and appear to have a similar metallicity to that of the Sun from a preliminary model atmosphere analysis of selected narrow line targets (see Hunter et al. 2006 for details). Previous studies have discussed the HR-diagrams, star formation timescales, initial mass funtions and rotational velocity distributions in young Galactic clusters and have highlighted the value of spectroscopic analysis of cluster members rather than spectral type assignments from photometric colours (e.g. see Massey et al. \\cite{1995ApJ...454..151M}, Hillenbrand et al. \\cite{Hil93}). Additionally there has been an revival of interest in the rotational velocity distribution of large samples of massive stars for comparison with the predictions of theoretical evolutionary models and also to probe initial conditions during star-formation. Strom et al. (\\cite{Str05}) have measured projected rotational velocities for approximately 200 B-type stars in h\\&$\\chi$\\,Per and suggest that differences between the mean values for the cluster and field stars are due to the star-formation process rather than being a consequence of angular momentum evolution as they have moved from the ZAMS.  Huang \\& Gies (\\cite{Hua05a}) have presented rotational velocities of stars in nine Galactic open clusters with age ranges of approximately 6 to 73 Myr, and have provided an interpretation on the basis of evolutionary arguments. Additionally the extensive field star sample of  Abt et al. (\\cite{2002ApJ...573..359A}) remains a benchmark study with which to compare cluster rotational velocities. All of these studies illustrate the importance of large stellar samples in order to ameliorate statistical significance problems arising from the random inclination angle of the rotational axes. The availability of the FLAMES spectrograph on an 8m telescope has allowed studies of large samples to be extended to the Magellanic Clouds. For example, Martayan et al. (\\cite{2006A&A...445..931M}) have studied the effects of metallicity on the rotational velocities of B-type and Be stars in the LMC and find that the LMC stars rotate faster than Be stars at solar metallicity.  In this paper we analyse the early type stars in the three Galactic clusters observed with the FLAMES spectrograph. Stellar atmospheric parameters are deduced using non-LTE model atmosphere calculations and rotational velocities are estimated by comparing observed and theoretical rotationally broadened spectra. The former are discussed with particular reference to the surface gravities found for near main sequence objects; evidence is presented that relates the lower gravity estimates to large stellar rotational velocities. Our projected rotational velocities are discussed both in the context of possible systematic differences between field and cluster populations and to infer the underlying rotational velocity distributions. The estimates of the atmospheric parameters are used to deduce cluster distances leading to estimates of stellar luminosities and evolutionary masses.  The former are then used to generate Hertzsprung-Russell diagrams and the agreement with theoretical predictions is discussed, while present day mass functions are estimated from the latter.  %----------------------------------------------------------------- sec: obs ---- ",
        "conclusions": "\\label{sec:concl} We have analysed high resolution spectroscopy of targets in three Galactic clusters to estimate atmospheric parameters and projected rotational velocities. These have been used to derive stellar luminosities and evolutionary masses. Our principle conclusions are as follows: \\begin{enumerate} \\item Atmospheric parameters have been deduced from non-LTE model atmosphere calculations and should be more reliable than those inferred from photometry or spectral types. Hence the observed HR-diagrams should be most the reliable yet determined for these clusters. \\item A significant number of relatively cool low gravity objects have been identified and are found on average to have large projected rotational velocity. The low gravities may then reflect at least in part the significant centrifugal accelerations experienced by these stars. \\item Observed Hertzsprung-Russell diagrams have been compared with theoretical predictions for both non-rotating and rotating models. The observed main sequence luminosities are in general larger than predicted. For NGC\\,6611, this may reflect non co-eval star formation. Other possible explanation include stellar rotation and the presence of a significant number of binary systems. \\item Cluster ages have been deduced and some evidence for non-coeval star formation is considered especially for the young cluster NGC\\,6611. Indeed for none of the clusters does a single isochrone provide a good fit to the observed HR-diagram. \\item Projected rotational velocities for targets in the older clusters, NGC\\,3293 and NGC\\,4755, have been found to be systematically larger than those for the field, confirming the conclusion of Strom et al.\\ (\\cite{Str05}) for h and $\\chi$\\,Persei.  Contrary to the finding of Strom et al., this phenomenon appears to be present at all the B-type spectral types present in our sample. \\item The distribution of projected rotational velocities are consistent with a Gaussian distribution of current rotational velocities. For the relatively unevolved targets in the older clusters, NGC\\,3293 and NGC\\,4755, the peak of the velocity distribution would be 250\\kms with a full-width-half-maximum of approximately 180 \\kms. For NGC\\,6611, the corresponding values are 175\\kms and 165 \\kms. \\item We find that the mean projected rotational velocity of stars which have strong winds (15-60\\msun) is lower than that of the lower mass stars. This is consistent with predictions of the rotating models which show the spin-down effect due angular momentum loss through stellar winds. \\item For all three clusters we deduce present day mass functions with $\\Gamma$-values in the range -1.5 to -1.8. \\end{enumerate}"
    },
    "0606/astro-ph0606304_arXiv.txt": {
        "abstract": "{\\large} We use deep optical, infrared and radio observations to explore the symbiosis between the nuclear activity and galaxy evolution in the southern compact radio source PKS1549-79 ($z=0.1523$). The optical imaging observations reveal the presence of tidal tail features which provide strong evidence that the host galaxy has undergone a major merger in the recent past. The merger hypothesis is further supported by the detection of a young stellar population, which, on the basis of spectral synthesis modelling of our deep VLT optical spectra, was formed 50 -- 250~Myr ago and makes up a significant fraction of the total stellar mass (1 -- 30\\%). Despite the core-jet structure of the radio source, which is consistent with the idea that the jet is pointing close to our line of sight, our HI 21cm observations reveal significant HI absorption associated with both the core and the jet. Moreover the luminous, quasar-like AGN ($M_V < -23.5$) is highly extinguished ($A_v > 4.9$) at optical wavelengths  and shows many properties in common with narrow line Seyfert 1 galaxies (NLS1), including relatively narrow permitted lines ($FWHM \\sim$1940~km s$^{-1}$), highly blueshifted [OIII]$\\lambda\\lambda$5007,4959 lines ($\\Delta V \\sim 680$~km s$^{-1}$), and evidence that the putative supermassive black hole is accreting at a high Eddington ratio ($0.3 < L_{bol}/L_{edd} < 11$). The results suggest that accretion at high Eddington ratio does not prevent the formation of powerful relativistic jets.  Together, the observations lend strong support to the  predictions of some recent numerical simulations of galaxy mergers in which the black hole grows rapidly through merger-induced accretion following the coalescence of the nuclei of two merging galaxies, and the major growth phase is largely hidden at optical wavelengths by the natal gas and dust. Although the models also predict that AGN-driven outflows will eventually remove the gas from the bulge of the host galaxy, the {\\it visible} warm outflow in PKS1549-79 is not currently capable of doing so. However, much of the outflow may be hidden by the material obscuring the quasar and/or tied up in hotter or cooler phases of the interstellar medium.  By combining our estimates of the reddening of the quasar with the HI column derived from the 21cm radio observations we have also made the first direct estimate of the HI spin temperature in the vicinity of a luminous AGN: $T_{spin} > 3000$K. ",
        "introduction": "Powerful radio galaxies have long been considered as useful signposts to massive structures in the distant universe. However, there is now increasing speculation that AGN and associated jet activity are not merely signposts but, in fact, intimately linked to the galaxy evolution process. The main evidence for such links is twofold. On the one hand the shape of the redshift evolution in the number density of quasars and radio galaxies shows marked similarities to that of the star formation history of the universe (Dunlop \\& Peacock 1990, Madau et al. 1996), while on the other close correlations have been found between the masses of the supermassive black holes that produce the prodigious nuclear activity and the properties of the bulges of the host galaxies (e.g. Tremaine et al. 2002).  The correlations between the redshift evolution of the radio source population and the star formation history of the field galaxy population can be explained in terms of hierarchical galaxy evolution models (e.g. Kauffmann \\& Haenelt 2000). It is well known from studies of nearby merging galaxies that stars are formed as part of the merger process. Numerical simulations of mergers of gas rich galaxies also show that the tidal torques associated with the mergers concentrate gas in the circumnuclear regions of the galaxies (Barnes \\& Hearnquist 1996, Mihos \\& Hearnquist 1996);  some of this gas may be available to fuel circumnuclear starburst, quasar and jet activity. Therefore it is likely that AGN activity and near-nuclear starbursts in massive galaxies are a direct consequence of hierarchical galaxy evolution; AGN may be a symptom of the overall galaxy evolution process.  In this context, radio galaxies are key objects for understanding the link between star formation and AGN activity in individual galaxies. Not only is there clear morphological evidence that a significant subset of powerful radio galaxies --- $\\sim$50\\% of those with strong emission lines --- have undergone recent mergers (Heckman et al. 1996, Smith \\& Heckman 1989), but many of the radio galaxies with morphological evidence for mergers also show signs of recent star formation (Tadhunter et al. 1996, Aretxaga et al. 2001, Tadhunter et al. 2002, 2005, Wills et al. 2002, 2004, Johnston et al. 2005). Despite these results, considerable uncertainties remain about the relationship between star formation and AGN/jet activity in radio galaxies. For example, it is not certain at what stage of the merger the AGN activity is triggered, nor whether the star formation is exactly coeval with the AGN activity. Indeed, some recent studies of the young stellar populations in radio galaxies suggest that the activity may be triggered relatively late in the merger sequence, a significant time after the major episode of merger-induced star formation (Tadhunter et al. 2005, Emonts et al. 2005, Johnston et al. 2005). However, the sample sizes are currently small and further detailed studies of the young stellar populations in powerful radio galaxies are required to put our understanding of the link between star formation and AGN activity on a firmer footing.  While the links between star formation and AGN/jet activity provide evidence that AGN are a bi-product of the overall hierarchical galaxy evolution process, it has been suggested that the correlations between the supermassive black holes masses and the properties of the bulges of the host galaxies are a consequence of the feedback effects of AGN-induced outflows (Silk \\& Rees 1998, Fabian 1999, di Matteo et al. 2005). To date, most of the evidence for such outflows has been provided by observations of narrow and broad absorption lines observed in the spectra of AGN at UV and X-ray wavelengths (see Crenshaw, Kraemer \\& George 2003 for a review). The outflows associated with these absorption features clearly have an impact on the near-nuclear environments on a scale $<$10pc,  but the impact they have on the 1 -- 10 kpc scale of galaxy bulges is less clear. Although outflows have also been detected in both ionized and neutral gas on a kpc-scale in powerful radio galaxies, uncertainties remain about the general importance of the outflows and their driving mechanism (quasars, jets or starbursts).  The subject of this paper --- PKS1549-79 --- is a key object for studies of the links between galaxy evolution and AGN activity, because it shows clear evidence of both the merging process that may trigger the activity, and the feedback effect of the AGN activity on the surrounding ISM.  PKS1549-79 is a compact flat spectrum radio source. Previously unpublished VLBI radio images of the source show that it has the compact core-jet structure expected for such sources, consistent with the idea that its jet is pointing close to the line of sight. Remarkably for such a source, low-resolution optical spectra fail to reveal the broad permitted lines and strong non-stellar continuum characteristic of quasars and BL Lac objects (Tadhunter et al. 2001). Therefore this object does not readily fit in with the simplest versions of the unified schemes which propose that lines of sight close to the radio axis have relatively unobscured views of the AGN (Barthel 1989).  Despite the absence of strong quasar continuum and broad emission line features, optical spectra of the source do reveal  high ionization forbidden emission lines which suggest that the source contains a powerful quasar-like AGN, even if this AGN is not directly visible at optical wavelengths.  However, the high ionization emission lines are unusually broad, and blueshifted by $\\sim$650 km/s relative to the low ionization lines. In addition, this object has HI 21cm absorption detected against its radio core, at the same redshift as the low ionization emission line system. Further support for the idea that the quasar nucleus is obscured by the circum-nuclear material in this source is provided by K-band spectroscopic observations that detect a broad Pa$\\alpha$ emission line and strong non-stellar continuum characteristic of quasars (Bellamy et al. 2003).  Aside from its unusual emission line properties, PKS1549-79 belongs to the subset of $\\sim$30 -- 50\\% of powerful radio galaxies for which young stellar populations make a substantial contribution to the optical continuum emission (Tadhunter et al. 2002). Further evidence for star formation in this object is provided by its unusually strong far-IR emission, which leads to its classification as an ultra luminous infrared galaxy (ULIRG: see Sanders \\& Mirabel 1996 for a definition) with $L_{ir} =1.6\\times10^{12}$~L$_{\\odot}$ . This evidence for recent star formation is consistent with the idea that the host galaxy of this source has been involved in a merger which has also triggered a major episode of star formation.  In this paper we present new deep spectroscopic and imaging observations of PKS1549-79 at optical, infrared and radio wavelengths that allow us to investigate the triggering of its AGN activity, characterise its circumnuclear outflows, and deduce the nature of its AGN. In this way we can place this key object in the context of evolutionary scenarios for the origin and evolution of the activity.  We assume cosmological parameters of $H_0 =75$~km s$^{-1}$ and $q_0 = 0.0$ throughout the paper. For these parameters 1 arcsec corresponds to 2.39~kpc at the redshift of the host galaxy of PKS1549-79. ",
        "conclusions": "The new observations presented in this paper confirm that PKS1549-79 represents a luminous quasar observed at an early stage of its evolution. The quasar and jet activity have been triggered in a major merger which has also triggered substantial star formation. Accreting at close to the Eddington rate, the black hole in this system is growing rapidly and simultaneously driving relativistic jets. The growth phase of the black hole is substantially obscured at optical wavelengths in this object, consistent with recent galaxy evolution models (Hopkins et al. 2005). However, although the optical emission lines provide evidence for a high ionization, warm outflow associated with the  AGN, the warm outflow currently observed at optical wavelengths is not in itself sufficient to remove all the gas from the central regions of the galaxy.  PKS1549-79 is a key object because it links together several classes of active galaxies, including radio galaxies, quasars, ULIRGs, NLS1 and AGN with high ionization absorption line systems. Moreover it shows many properties in common with the recently discovered population of high redshift AGN detected at sub-mm wavelengths (Alexander et al. 2005). The common thread is that all these classes of object are associated with the triggering and feedback effects of major galaxy mergers. Although much attention is paid to the links between star formation and activity in objects at high redshift, this study of PKS1549-79 serves to emphasise that some galaxies are evolving rapidly in the local universe, and that it is possible to use such objects to study the coevolution of black holes and galaxy bulges in depth.  Several further observational studies of PKS1549-79 would enhance our understanding of the relationship(s) between AGN activity and galaxy evolution. They include the following. \\begin{itemize} \\item [-] Near-IR measurements of the stellar bulge luminosity and velocity dispersion to improve estimates of the black hole mass. \\item [-] High quality optical spectroscopy observations of a wider range of emission line diagnostics to measure the electron density accurately, and thereby improve our determination  of the properties of the warm gas outflow. \\item [-]Deep X-ray spectroscopy observations to refine estimates of the quasar bolometric luminosity and absorbing HI column. Such observations would also allow  the links between PKS1549-79, NLS1 and sub-mm galaxies to be further investigated. \\item [-]Multi-epoch VLBI observations to search for superluminal motions in the jet, and accurately determine the inclination of the jet to the line of sight. \\item [-]Deeper HI 21cm measurements and high resolution optical observations of the NaID line to investigate the importance of neutral gas outflows. \\end{itemize}  \\subsection*"
    },
    "0606/astro-ph0606418_arXiv.txt": {
        "abstract": "We present constraints on the evolution of the virial to stellar mass ratio of galaxies with high stellar masses in the redshift range $0.2<z<0.8$, by comparing weak lensing measurements of virial mass $M_{\\rm vir}$ to estimates of stellar mass $M_{\\rm star}$ from COMBO-17. For a complete sample of galaxies with $\\log ( M_{\\rm star} / M_\\odot) > 10.5$, where the majority show an early-type morphology, we find that the virial mass to stellar mass ratio is given by $ M_{\\rm vir}/M_{\\rm star} = 53^{+13}_{-16} $. Assuming a baryon fraction from the concordance cosmology, this corresponds to a stellar fraction of baryons in massive galaxies of $\\Omega_b^*/\\Omega_b = 0.10\\pm 0.03$. Analysing the galaxy sample in different redshift slices, we find little or no evolution in the virial to stellar mass ratio, and place an upper limit of $\\sim 2.5$ on the growth of massive galaxies through the conversion of gas into stars from $z=0.8$ to the present day. ",
        "introduction": "In the standard cold dark matter paradigm, luminous galaxies reside in massive dark matter halos.  The large-scale clustering and formation of these luminous galaxies is driven by the clustering and formation histories of their dark matter halos. In contrast, the visible properties of such galaxies are a strong function of the complex physics of baryonic matter, in particular the laws governing stellar and AGN feedback, gas cooling and star formation.  The luminosities and colors of galaxies are a strong function of the efficiencies of these processes, and our inability to robustly model them is a critical limitation of current models of galaxy formation and evolution \\citep[as discussed, for example, by][]{Benson03}. A key way to constrain the efficiency of such processes is to measure empirically the ratio of stellar mass to the total virial mass of the halo and galaxy system, for certain subsets of the galaxy population.  Assuming a baryon to dark matter fraction from the concordant cosmology, it is then possible to determine empirically the fraction of baryons which manage to cool and condense into stars in halos of a given mass.  Such measurements of the mass out to $\\ge 100$\\,kpc scales not only constrain the efficiency of complex `gastrophysical' processes, but they also empirically calibrate and test halo occupation distribution analyses of galaxy properties \\citep[see for example][]{Yang05}.  Satellite galaxies \\citep{Zaritsky,Prada} and weak gravitational lensing \\citep[see the review by][]{Bible}, are currently the most promising methods available to probe virial masses of dark matter halos. In this letter we focus on weak gravitational lensing, the weak distortion of distant sources by the gravitational field of foreground structure, as it provides a unique way to probe matter on all scales, at all redshifts, irrespective of its dynamical state or nature.  Termed `galaxy-galaxy lensing', the weak tangential shearing of background galaxies relative to foreground galaxies can therefore be used to probe the distribution of dark matter in galaxies, as first demonstrated by \\citet{BBS}.  In this letter we combine information from two surveys of the Chandra Deep Field South; the ground-based COMBO-17 survey \\citep{Wolf03} and the HST GEMS survey \\citep{RixGEMS}, in order to constrain the virial to stellar mass ratio in stellar massive galaxies out to $z \\sim 0.8$. The 17-filter photometry of COMBO-17 provides accurate photometric redshifts with errors $\\sigma_z \\sim 0.02 (1+z)$.  COMBO-17 additionally provides stellar mass estimates $M_{\\rm star}$ by fitting a parameterised star formation history to the low resolution 17-band galaxy spectra \\citep{Borch}, based on a \\citet{Kroupa} initial mass function (Kroupa IMF).  The HST resolution of the GEMS imaging allows for high signal-to-noise measurements of the `galaxy-galaxy' lensing signal which can provide an ensemble estimate of the virial mass for (sub-)sets of galaxies. A detailed description of the galaxy-galaxy lensing analysis of GEMS and the resulting constraints on different dark matter halo profiles for different galaxy types will be presented in a forthcoming paper (Heymans et al. in prep).  For the purposes of this letter however we focus on a complete sample of 626 lens galaxies with high stellar masses. We measure the virial to stellar mass ratio as a function of galaxy redshift to $z \\sim 0.8$, presenting the first redshift dependent galaxy-galaxy lensing analysis to be performed consistently over such a large redshift range.  This letter is organised as follows.  In Section~\\ref{sec:ggtheory} we describe the theory that underpins our measurement of virial masses and review the particulars of our analysis.  We refer the reader to section 3 of \\citet{HymzGEMS} hereafter H05, for a lensing specific description of the GEMS observations and the data reduction.  In Section~\\ref{sec:results} we present constraints on the virial radius and the virial to stellar mass ratio for our full stellar massive galaxy sample and convert this into a stellar fraction of baryons.   We compare our results to the low-redshift results of \\citet{HHbf05}, hereafter HHYLG, and \\citet{RM06}, hereafter M06, and investigate the dependence of the virial to stellar mass ratio on galaxy redshift, setting constraints on the growth of massive galaxies through the conversion of gas into stars from $z \\sim 0.8$ to the present day. We discuss the implications of these results and conclude in section~\\ref{sec:conc}. Throughout this letter we assume a $\\Lambda{\\rm CDM}$ cosmology with $\\Omega_m = 0.3$, $\\Omega_\\Lambda= 0.7$ and $H_0 = 100 \\, h {\\rm km \\, s}^{-1}\\, {\\rm Mpc}^{-1}$.    For our virial to stellar mass ratio measurement we set $h=0.7$. ",
        "conclusions": "\\label{sec:conc}  In this letter, we use a weak gravitational lensing study of the HST GEMS survey, in combination with the 17-colour photometry of the COMBO-17 survey, to determine the ratio between virial and stellar mass in a complete sample of massive galaxies out to $z \\sim 0.8$.  We find a mean virial to stellar mass ratio $M_{\\rm vir}/M_{\\rm star}=53^{+13}_{-16}$ with little or no redshift evolution in this stellar mass fraction.  This result provides a high redshift extension to the analysis of \\citet{HHbf05} (HHYLG) and \\citet{RM06} (M06). The excellent agreement found between the results of this analysis, HHYLG and M06 is important to comment upon as the analyses were performed using very different techniques and came from very different surveys: HHYLG use the Red sequence Cluster Survey, a medium-deep $45.5$ square degree ground-based survey in $BVRz$, M06 use the shallow Sloan Digital all Sky Survey which is ground-based in $ugriz$. This analysis uses GEMS, a deep $0.25$ square degree space-based survey with accompanying 17-band COMBO-17 data.  In this analysis we use stellar masses calculated from the best-fitting models of star formation history to the low resolution 17-band galaxy spectra.  M06 use a similar technique \\citep{Kauffmann03}. HHYLG use the simple relationship of \\citet{BelldeJong} that relates stellar mass to $B-V$ colour.  This relationship has a systematic uncertainty that is often estimated to be $\\sim 30\\%$ \\citep{Bell03}.  Stellar mass uncertainty is indeed hard to avoid with all methods, however, owing to the unknown frequency of bursts of star formation and uncertainty in dust models \\citep{Borch}. The fact all three results are in such close agreement, however, demonstrates that the simple models of HHYLG are consistent with the more complex star formation history models of COMBO-17 and M06, which bodes well for similar studies in future broad-band surveys.  We have applied the \\citet{SchRix} maximum likelihood method to set constraints on the virial radius and the virial to stellar mass ratio using an NFW dark matter profile.  This method takes into account the uncertainty in the redshift of both the source and lens galaxies, using a Monte Carlo technique and includes the effects of the multiple deflection of light by successive foreground lens galaxies. In contrast, HHYLG measure the tangential shear around foreground galaxies for different luminosity bins and fit this result with the shear expected from an NFW profile. The HHYLG galaxies are selected to be isolated such that the lensed background galaxies experience only the gravitational potential of the isolated galaxy. Their larger photometric redshift errors are more likely to scatter faint galaxies into higher luminosity bins than scatter bright galaxies into lower luminosity bins.  This then leads to a bias in the measurement of the mass at fixed luminosity to a lower value.  HHYLG therefore need to apply a correction factor to the virial masses which they determine from an analysis of mock catalogues.  This bias is not a problem for our analysis owing to the higher accuracy of photometric redshifts and the use of the \\citet{SchRix} method.  Our observational systematic errors are therefore minimal, and we are limited by random errors from smaller number statistics. The third analysis method of M06 uses the halo model to extract information from the galaxy-galaxy lensing signal and separate contributions from central and satellite galaxies \\citep{Mandelbaumgg}. The fact that these three very different methods produce consistent results serves as a good indication that the three results are robust.  This analysis is the first time a comparison of the properties of luminous galaxies and their parent dark matter halos has been performed consistently across a substantial span of cosmic time.  The galaxy sample has been selected to be complete out to $z \\sim 0.8$ and we are therefore not subject to redshift-dependent selection biases. We have shown that for this galaxy sample, the majority of which show early-type morphologies, the virial to stellar mass ratio as a function of galaxy redshift remains fairly constant with $M_{vir}/M_{\\rm star} \\sim 50$: specifically, we place a $1 \\sigma$ upper limit on the growth of the stellar to virial mass ratio since $z=0.8$, of a factor $< 2.6$.  Because we probe the ratio of stellar to virial mass, we cannot probe merger-driven growth of galaxies as stellar masses and virial masses grow in lock-step.  Instead, our constraint effectively limits the conversion of gas into stars in the massive galaxy population, i.e. star formation, as a function of redshift. We therefore can conclude that the growth of stellar mass due to star formation is limited, to a factor of $\\lesssim 2.5$ from $z=0.8$ to the present day."
    },
    "0606/astro-ph0606242_arXiv.txt": {
        "abstract": "We present high sensitivity \\twelveCO\\ and \\thirteenCO\\ \\jone\\ molecular line maps covering the full extent of the parsec scale L1551 molecular outflow, including the redshifted east-west (EW) flow. We also present \\twelveCO\\ \\jthree\\ data that extends over a good fraction of the area mapped in the \\jone\\ transition. We compare the molecular data to widefield, narrow-band optical emission in H$\\alpha$. While there are multiple outflows in the L1551 cloud, the main outflow is oriented at 50\\arcdeg\\ position angle and appears to be driven by embedded source(s) in the central IRS~5 region. The blueshifted outflowing molecular gas extends to the edge of the molecular cloud and beyond the last HH object, HH 256. On the contrary, the redshifted molecular gas terminates within the cloud, short of the most distant HH object, HH 286, which lies well beyond the cloud boundary. The \\jthree\\ data indicate that there may be molecular emission associated with the L1551 NE jet, within the redshifted lobe of main outflow. We have also better defined the previously known EW flow and believe we have identified its blueshifted counterpart. We further speculate that the origin of the EW outflow lies near HH 102. We use velocity dependent opacity correction to estimate the mass and the energy of the outflow. The resulting mass spectral indices from our analysis, are systematically lower (less steep) than the power law indices obtained towards other outflows in several recent studies that use a similar opacity correction method. We show that systematic errors and biases in the analysis procedures for deriving mass spectra could result in errors in the determination of the power-law indices.  The mass spectral indices, the morphological appearance of the position-velocity plots and integrated intensity emission maps of the molecular data, compared with the optical, suggest that jet-driven bow-shock entrainment is the best explanation for the driving mechanism of outflows in L1551. The kinetic energy of the outflows is found to be comparable to the binding energy of the cloud and sufficient to maintain the turbulence in the L1551 cloud. ",
        "introduction": "From the emergence of a hydrostatic core, the collapse of a protostellar core is accompanied by winds and mass loss \\citep{lada85} probably driven by magnetospheric accretion \\citep[e.g.][]{koenigl91,edwards94, hartmann94}. Integrated over time, the effect of a wind from a young star is to blow away the material that shrouds it during its earliest evolution. The discovery of bipolar molecular outflows has been a key to the understanding of this process \\citep[e.g.][SLP hereafter]{slp80}. These molecular outflows have dimensions of up to several parsecs, masses one or two orders of magnitude larger than their driving sources, and tremendous kinetic energies \\citep{bally83}. Such massive flows must represent swept-up material as the winds, emerging from the star and/or its circumstellar disk, interact with their ambient medium.  One of the important unanswered questions in outflow studies is the mechanism by which the jets/winds emerging from the embedded star/disk entrain and accelerate the surrounding molecular material. Currently there are three scenarios: the ``wide-angle wind'' model in which a wide-angled magnetized wind expands like a bubble and sweeps up the ambient medium in a momentum conserving manner into the shell at the wind-bubble ambient medium interface \\citep{shu91, shu95, shu00, ls96, mm99, lee00, lee01}. Then there are two types of ``jet'' entrainment models, the ''turbulent'' jet entrainment model, in which ambient material is accelerated through a turbulent viscous mixing layer that surrounds the jet formed by Kelvin-Helmholtz instabilities \\citep{canto91,stahler94, lizano95}; and the ''prompt'' entrainment model in which bow shock formed in the jet head drives material ahead and laterally to the axis of the flow \\citep{deyoung86, raga93, stone93a, stone93b, stone94, masson93, suttner97, zz97, ssy97, dr99, lee01, lee02}.  It is essential to distinguish between these models to not only understand the processes that govern the formation of molecular outflows, but also the nature of the jet/wind launching mechanism and the underlying accretion disk processes. Since it is difficult to observe the jet/wind launching regions directly, the properties of large scale bipolar outflows can shed some light on the nature of their initial conditions. In addition, the energy and momentum deposited by molecular outflows within the molecular cloud can be large enough to power turbulence in the cloud or even disrupt the parts of the cloud. Understanding outflow properties over the entire outflow history is critical in determine their impact on the cloud.  Within the last decade, with the advent of CCD cameras with 20\\arcmin\\ fields-of-view, more than two dozen giant Herbig-Haro (HH) flows with dimensions ranging from 1 to 7 pc have been discovered \\citep{rbd97}. HH flows are chains of optically visible shock-excited nebulae which, demonstrate the existence, at least during certain evolutionary phases, of exceedingly well collimated and highly supersonic wind components \\citep[e.g.][]{reipurth86}. The tremendous size of parsec-scale outflows implies that molecular material is being entrained at parsec-scale distances from the driving source, and thus it may have a significant impact on the kinematics and energy density of a substantial volume of the parent molecular cloud.  Our current estimates of the ''typical'' outflow properties may be far from accurate, since there is no large homogeneous sample of molecular outflows which have been mapped with good resolution and sensitivity \\citep{richer2000}. Consequently, our current understanding of the processes might be heavily subjective, and biased by a few frequently studied examples. We have been involved in a multi-wavelength study of a dozen or so parsec-scale outflows, complementing narrow-band optical observations with wide-field, sensitive, multi-transitional \\twelveCO\\ and \\thirteenCO\\ observations (Stojimirovi\\'c et al. 2006 in prep.). Here, we present the detailed results of our multi-wavelength study on the L1551 molecular outflow.  The bipolar molecular outflow associated with the L1551 small dark cloud in the Taurus molecular cloud complex, was the first recognized bipolar outflow from a young stellar object (SLP) and is considered one of the best examples of its kind. At a distance of 140 pc \\citep{lynds62, kdh94}, the L1551 dark cloud has been the site of many multi-wavelength studies that reveal the dynamical complexity of this low mass star formation region.  Optical images of the L1551 outflow region, show new distant HH object, HH 286 \\citep{drb99}, making the whole extent of the optical emission in the L1551 flow 1.3 pc. We made sensitive, large spatial extent, high spatial and velocity resolution maps at millimeter and submillimeter wavelengths. Large-spatial extent mapping performed on parsec-scale outflows, allows us to study in detail the entrainment mechanisms by which jets/winds drive the observed molecular flows. With the availability of both \\twelveCO\\ and \\thirteenCO\\ data over the full spatial extent of the flow, we are able to take a {\\em full} and proper accounting of the mass and hence energetics of the outflows in the L1551 system.  Comparing optical and millimeter data, we are able to discriminate among entrainment models based on morphological grounds. ",
        "conclusions": "All L1551 molecular outflows have been mapped at high sensitivities in the \\twelveCO\\ and \\thirteenCO\\ \\jone\\ transitions. Follow-up submillimeter observations in \\twelveCO\\ \\jthree\\ emission towards a good fraction of the L1551 outflow system were carried out as well. The millimeter and submillimeter data are combined with large-scale, narrow band, optical H$\\alpha$ images in order to perform a detailed study of the outflows in L1551. The main conclusions of the paper are summarized below.  \\begin{enumerate} \\item The full extent of the main CO outflow is 32\\arcmin\\ (1.3 pc), and is oriented at $\\sim 50$\\arcdeg\\ position angle. The molecular outflow extent is similar to the optical extent of the flow seen by \\citet{drb99}. Outflowing molecular material extends well beyond the last known HH object, HH 256 in the south-west, but stops short of HH 286 in the north-east. Our molecular line maps indicate that HH 286 has completely exited the L1551 cloud. \\item Most of the CO molecular outflow appears to be driven by source(s) in the IRS 5 region. Some features, especially in the redshifted lobe of the outflow may be attributed to outflows originating in L1551 NE and the HL Tau region. \\item Our maps cover a much larger extent of the collimated redshifted EW flow. The EW flow is $\\sim 0.82$ pc in length, and is $\\sim 0.15$ pc in width. In contrast to earlier studies, the sensitivity in our data reveals that the EW flow extends to high velocities (6.5~\\kms\\ to 20~\\kms). It is suggested that a smaller extent blueshifted component seen in the main northwestern redshifted lobe is the blueshifted counterpart of the EW flow, and that the driving source for the EW flow is close to the location of HH 102. \\item We refine the procedure of velocity-dependent opacity correction adopted by recent authors to obtain a more accurate determination of outflow mass with velocity. The resultant mass-spectral power law indices are lower (less steep) than recently obtained indices towards other outflows. We attribute this systematic difference in power-law index to the better quality of our data, and the more careful approach we use in calculating mass. The resulting mass of the L1551 molecular cloud core is $\\sim 110$~M$_\\odot$ and the combined mass of the outflows in L1551 is $\\sim 7.2$~M$_\\odot$. \\item Even for no inclination angle correction, and lower $T_{\\rm ex}$ assumption, the kinetic energy of the outflow is comparable to the binding energy of the cloud and there seems to be enough outflow energy to maintain turbulence in the cloud. Therefore the parsec scale outflow in L1551 has a potential of making a strong impact on the cloud's evolution and fate. \\item Multiple lines of evidence from \\mv\\ spectral indices, position velocity plots, and morphological appearance of our multi-wavelength data indicate that the EW molecular outflow is a good example of jet-driven bow shock entrainment. The main molecular outflow in L1551 shows evidence for entrainment from both jet-driven and wind-driven mechanisms. \\item We suggest that new stars are being triggered to form in the L1551 system as a result of the effects of the multiple outflows in the region. \\end{enumerate}"
    },
    "0606/astro-ph0606352_arXiv.txt": {
        "abstract": "We review the properties and behavior of 20 X-ray binaries that contain a dynamically-confirmed black hole, 17 of which are transient systems.  During the past decade, many of these transient sources were observed daily throughout the course of their typically year-long outburst cycles using the large-area timing detector aboard the {\\it Rossi X-ray Timing Explorer}.  The evolution of these transient sources is complex.  Nevertheless, there are behavior patterns common to all of them as we show in a comprehensive comparison of six selected systems.  Central to this comparison are three X-ray states of accretion, which are reviewed and defined quantitatively.  We discuss phenomena that arise in strong gravitational fields, including relativistically-broadened Fe lines, high-frequency quasi-periodic oscillations (100-450 Hz), and relativistic radio and X-ray jets. Such phenomena show us how a black hole interacts with its environment, thereby complementing the picture of black holes that gravitational wave detectors will provide.  We sketch a scenario for the potential impact of timing/spectral studies of accreting black holes on physics and discuss a current frontier topic, namely, the measurement of black hole spin. ",
        "introduction": "Oppenheimer \\& Snyder (1939) made the first rigorous calculation describing the formation of a black hole (BH).  The first strong evidence for such an object came from X-ray and optical observations of the X-ray binary Cygnus X--1 (Bolton 1972; Webster \\& Murdin 1972). Today, a total of 20 similar X-ray binary systems are known that contain a compact object believed to be too massive to be a neutron star or a degenerate star of any kind (i.e., $M > 3$ \\msun).  These systems, which we refer to as black-hole binaries (BHBs), are the focus of this review.  These 20 dynamical BHs are the most visible representatives of an estimated $10^8$--$10^9$ stellar-mass BHs that are believed to exist in the Galaxy (e.g., Brown \\& Bethe 1994; Timmes, Woosley, \\& Weaver 1996).  Stellar-mass BHs are important to astronomy in numerous ways. For example, they are one endpoint of stellar evolution for massive stars, and the collapse of their progenitor stars enriches the universe with heavy elements (Woosley et al.\\ 2002).  Also, the measured mass distribution for even the small sample featured here is used to constrain models of BH formation and binary evolution (e.g., Fryer \\& Kalogera 2001; Podsiadlowski et al.\\ 2003). Lastly, some BHBs appear to be linked to the hypernovae believed to power gamma-ray bursts (Israelian et al.\\ 1999; Brown et al.\\ 2000; Orosz et al.\\ 2001).  The BHBs featured here are mass-exchange binaries that contain an accreting BH primary and a nondegenerate secondary star. For background on X-ray binaries, see Psaltis (2006).  For comprehensive reviews on BHBs, see McClintock \\& Remillard (2006), Tanaka and Shibazaki (1996) and Tanaka \\& Lewin (1995).  X-ray observations of BHBs allow us to gain a better understanding of BH properties and accretion physics.  In this review, we emphasize those results that challenge us to apply the predictions of general relativity (GR) in strong gravity (\\S8). Throughout, we make extensive use of the extraordinary data base amassed since January 1996 by NASA's {\\it Rossi X-ray Timing Explorer} (RXTE; Swank 1998).  In an astrophysical environment, a BH is completely specified in GR by two numbers, its mass $M$ and its specific angular momentum or spin $a = J / c M$, where $J$ is the BH angular momentum and $c$ is the speed of light. The spin value is conveniently expressed in terms of a dimensionless spin parameter, $a_{*} = a / R_{\\rm g}$, where the gravitational radius is $R_{\\rm g} \\equiv GM/c^{2}$. The mass simply supplies a scale, whereas the spin changes the geometry.  The value of $a_{*}$ lies between 0 for a Schwarzschild hole and 1 for a maximally-rotating Kerr hole.  A defining property of a BH is its event horizon, the immaterial surface that bounds the interior region of space-time that cannot communicate with the external universe.  The event horizon, the existence of an innermost stable circular orbit (ISCO), and other properties of BHs are discussed in many texts (e.g., Shapiro \\& Teukolsky 1983; Kato et al.\\ 1998).  The radius of the event horizon of a Schwarzschild BH ($a_{*} = 0$) is \\RS = 2$R_{\\rm g}$ = 30~km($M/$10\\msun), the ISCO lies at $R_{\\rm ISCO}$~=~6$R_{\\rm g}$, and the corresponding maximum orbital frequency is $\\nu_{\\rm ISCO}=220$~Hz($M/$10\\msun)$^{-1}$.  For an extreme Kerr BH ($a_{*} = 1$), the radii of both the event horizon and the ISCO (prograde orbits) are identical, $R_{\\rm K}~=~R_{\\rm ISCO}$~=~$R_{\\rm g}$, and the maximum orbital frequency is $\\nu_{\\rm ISCO}~=~ 1615$~Hz($M/$10\\msun)$^{-1}$. ",
        "conclusions": ""
    },
    "0606/astro-ph0606487_arXiv.txt": {
        "abstract": "The statistical properties of a complete, flux limited sample of 197 long gamma--ray bursts (GRBs) detected by BATSE are studied. In order to bring forth their main characteristics, care was taken to define a representative set of ten parameters. A multivariate analysis gives that $\\sim70\\%$ of the total variation in parameter values is driven by only three principal components. The variation of the temporal parameters is clearly distinct from that of the spectral ones. A close correlation is found between the half-width of the autocorrelation function ($\\tau$) and the emission time ($\\tem$); most importantly, this correlation is self-similar in the sense that the mean values and dispersions of both $\\tau$ and $\\tem$ scale with the duration of the burst ($\\td$). It is shown that the Amati-relation can be derived from the sample and that the scatter around this relation is correlated with the value of $\\tau$. Hence, $\\tau$ has a role similar to that of the break in the afterglow light curve ($t_{\\rm b}$) in the Ghirlanda-relation. In the standard GRB-scenario, the close relation between a global parameter ($t_{\\rm b}$) and a local one ($\\tau$) indicates that some of the jet-properties do not vary much for different lines of sight.  Finally, it is argued that the basic temporal and spectral properties are associated with individual pulses, while the overall properties of a burst is determined mainly by the number of pulses. ",
        "introduction": "\\label{intro}   The Burst and Transient Source Experiment (BATSE) collected during its operation the most extensive gamma-ray burst (GRB) database to the present day, with more than 2700 triggered bursts. Even with the advent of {\\it Swift} and the future mission {\\it GLAST}, it will remain the most complete burst catalog for many years to come.  With such a wealth of data, there are many statistical techniques that could be used to investigate the sample characteristics. So far this analysis has focused mainly on bivariate methods. Such an approach has some drawbacks for intrinsic relations involving more than two observables, since the analysis is limited to two dimensional projections of the actual relation. Not only do such projections increase the scatter but, most importantly, selection effects and observational biases can, for example, introduce spurious correlations. Hence, multivariate methods are to be preferred when the quality of the data is such that this can be reliably done. In the present work a principal component analysis (PCA) is performed on a complete, flux limited sample of GRBs chosen from the BATSE database. Alternative multivariate methods have been discussed and used by \\citet{Muk98} in connection with attempt to discriminate between distinct classes of GRBs.   The PCA is a well-known statistical tool, widely used in many scientific disciplines. Some areas of astronomy, like the study of active galactic nuclei and quasi-stellar objects, have long benefited from this technique \\citep[e.g., by ][]{Bor92,Bor02}. The method finds, in decreasing order, orthogonal directions of maximal variation.  The most common use of the PCA is to reduce the number of variables needed to reproduce the variability of the data without essential loss of information.  Since in most practical situations there is a significant amount of inter-correlation among the observed quantities, there is then some degree of redundancy in the multi-dimensional data.  This simplifies the data analysis, for example, when looking for subclasses. Previous uses of the PCA in the field of GRBs \\citep[e.g., by ][]{Bag98,Bal03} were done mainly for this purpose.  However, the PCA can also be an important aid to better understand the generally complex correlation structure encountered in multivariate studies and to gain physical insight to the system under study.  It is with this purpose in mind that we have selected a broader range of physical parameters for our sample of GRBs than the ones considered in previous works. We include various basic temporal and spectral parameters in order to characterize temporal variability and spectral evolution.  Throughout this paper only long GRBs will be considered (i.e., those with time duration $>2$ s) since short bursts need to be treated separately. This is not only because they are most likely a different physical class, but also because it would not be possible to derive for them some of the parameters under consideration, which are already hard to measure for the long bursts.  The characterization of the temporal properties is complicated by the remarkable morphological diversity of GRBs light curves and the presence of long emission {\\it gaps} in a significant fraction of them. Hence, a parameter designed to give a measure of some property for a ``typical'' GRB light curve will often fail to give consistent or physically meaningful results in other cases.  The problem extends to any other temporally averaged property which follows the pulse structure of the light curve.  For this reason, seemingly similar measures of a given property but based on different criteria can have a complementary roll, as it will be shown in this work.  The paper is organized as follows. In \\S~\\ref{data} we present the sample and discuss the motivation for the choice of parameters used in the subsequent analysis. A short introduction to PCA is given in \\S~\\ref{pcamethod}. Special attention is given to the estimate of uncertainties in the coefficients of the principal components; in particular, the effects of mixing and reordering of the principal components due to (nearly) degenerate eigenvalues are discussed. The basic statistical results are presented in \\S~\\ref{results}, where also a brief discussion is given of how they compare to previous studies. A more in depth analysis is done in \\S~\\ref{pca}, where it is shown that two redshift independent correlations can be derived from the data. One of the challenges of statistical studies is to distinguish between direct causal connections and those due to association through some common variable not necessarily included in the analysis. In an attempt to make the overall correlation structure of the parameters more distinct, we develop a new method in \\S~\\ref{structure}, in which ranking of correlations is combined with the PCA. It is shown that this reveals considerably more structure than does, for example, the use of only the ranking of correlations.  Finally, in \\S~\\ref{discussion} we discuss our results with a particular emphasis on how they can be used together with results already established to illuminate some of the basic properties of GRBs. ",
        "conclusions": "\\label{discussion}  The statistical properties of a sample of GRBs have been analyzed with the use of ten parameters chosen to bring forth the main characteristics as observed by BATSE. In a multivariate analysis it was found that $\\sim 70\\%$ of the total variation in parameter values is driven by only three principal components. Half of that is due to temporal variations while the other half is roughly divided equally between two principal components describing, respectively, the average spectral properties and spectral evolution. A more refined analysis reveals that the parameters in the latter two principal components are closely related. Together with the similarity of the variance (eigenvalue) of these principal components, this makes it likely that the two principal components needed to describe the spectral variations are not uniquely defined by the present analysis. Hence, the conclusion is that there exist two distinct sub-groups of inter-related parameters, which describe the temporal and spectral properties, respectively, of GRBs. One may note from Figure~\\ref{fig:deca}{\\it b} that the only parameter showing substantial connection to both sub-groups appears to be the fluence, i.e., its value is determined jointly by the two sub-groups.  Complementary information is provided by the principal components with the smallest variance. Since these directions in parameter space are defined by their narrow distributions of parameter values, they directly provide analytical relations between the parameters involved. It was found that the parameters in the temporal sub-group discussed above also obey such a well defined relation. It is noteworthy that this relation is self-similar in the sense that the mean values and dispersions of both the emission time ($\\tem$) and half-width of the autocorrelation function ($\\tau$) scale with the duration of the burst ($\\td$). Since this relation is unaffected by the unknown redshifts, it is a characteristic intrinsic to the light curves of GRBs. It is seen from Figure~\\ref{fig:deca}{\\it b} that the parameter $V$, which measures the small time-scale variability, is the only one (possibly also $\\alpha$) without any obvious connection to either of the two sub-groups. It was shown in \\S~\\ref{pc10} that the observed relation between $\\tem$ and $\\tau$ finds a ready explanation in a toy model where the value of $V$ is determined by the number of pulses in a burst. This implies that the scaling of $\\tem$ with $\\td$ would be the same as that for $\\tau$ and, hence, the fundamental scaling is that between $\\tau$ and $\\td$.  Burst light curves have very diverse morphologies and it has been suggested \\citep[see, e.g.,][]{G83,BR01} that they are a composite of pulses, with each pulse representing a single physical event (e.g., in the internal shock scenario, collisions between shells). In the toy model, the number of pulses in a burst is measured by the value of $V$. Independent support for such an interpretation comes from the prominence of the hardness-intensity correlation. It is well-known that this correlation is most apparent in FRED-like bursts. This association is most clearly seen in our data when these are represented as in Figure~\\ref{fig:deca}{\\it b} (i.e., the close relation between $\\SF$ and $\\REpk$). The effects of $V$ can best be seen in the PCs with the largest variances; in all three, $V$ anti-correlates with both $\\SF$ and $\\REpk$, indicating that large values of $V$ tend to smear out both the FRED-like structure and the hardness-intensity correlation. Hence, a simple interpretation of the data is that not only the self-similarity of the light curves but also the hardness-intensity correlation and the FRED-like structure are properties associated with the individual pulses and that the overall properties of a burst is determined by the number of pulses.  The correlation between fluence ($\\fl$) and peak energy ($\\Epk$) was one of the first to be well established for GRBs. Subsequent analysis indicated that the scatter around this relation could be reduced if a time variable were included, although it remained unclear which time variable gave the best result. Part of the problem, when trying to make more quantitative estimates, was the lack of redshifts. With the use of a small sample of GRBs with known redshifts, \\citet{Ama02} found the rest-frame equivalent of the $\\fl$--$\\Epk$-relation (i.e., the $\\Eiso$--$\\Epk'$-relation), while \\citet{GGL04} showed that the scatterer around this relation could be substantially reduced if account were taken of the rest-frame break-time ($t'_{\\rm b}$) in the light curve of the ensuing afterglow \\citep[see also][]{LZ05}. It is interesting to note that the functional dependence between $\\Epk$ and $\\Eiso$ is the same in the Amati and Ghirlanda relations. This is due to the fact that the values of $t'_{\\rm b}$ and $\\Epk'$ are uncorrelated. The reduction of the scatter in the Amati-relation with the use of a time-scale associated with the afterglow rather than the prompt emission indicates a close connection between these two phases.  Although the lack of redshifts prevents a determination of the relation between the values of $\\Eiso$ and $\\Epk'$ for individual GRBs observed by BATSE, we have developed a method by which this relation can be investigated statistically for the sample as a whole from the observed $\\Epk$--$\\fl$--$\\tau$ correlation. By matching both the mean values and the dispersions of these observables, we show that the Amati-relation can be derived from our sample of GRBs. Furthermore, the scatter around this relation correlates with the value of $\\tau'$, the rest-frame value of the half-width of the autocorrelation function. This is analogous to the findings of \\citet{GGL04} that the scatter in the Amati-relation correlates with $t'_{\\rm b}$. The observed scatter in the $\\Epk$--$\\fl$--$\\tau$ correlation is consistent with that expected from the unknown redshifts, indicating a small intrinsic scatter. During the prompt phase of a GRB, it is likely that the observer sees only a small portion of the jet; hence, $\\tau'$ measures a local property and its value could, in principle, vary over the jet (i.e., for different lines of sight). The value of $t'_{\\rm b}$, on the other hand, is due to the overall properties of the jet. The fact that both $\\tau'$ and $t'_{\\rm b}$ reduce the scatter in the Amati-relation indicates that at least some of the local properties do not vary much over the jet.  \\citet{NP05} and \\citet{Ban05} have argued that the Amati-relation is not consistent with the BATSE data and may be the result of a selection effect. However, as shown here, the Amati-relation can, in fact, be derived from our sample of GRBs selected from the BATSE data. This is a complete, flux limited sample, which suggests that the Amati-relation is valid for the intrinsically most luminous GRBs. These apparently contradicting results may be caused by a double bias against observing afterglows with large values of $t'_{\\rm b}$: (1) Due to the anti-correlation between $t'_{\\rm b}$ and $\\Eiso$ found by \\citet{GGL04}, afterglows with large values of $t'_{\\rm b}$ are intrinsically weak. (2) The decline of the light curve with time makes it increasingly hard to establish a value for $t'_{\\rm b}$ in those afterglows where the break occurs late. Hence, the inconsistency found by \\citet{NP05} and \\citet{Ban05} could be due to a larger number of GRBs in the BATSE sample than assumed by them for which it is observationally hard to measure a break in the light curve of the afterglow. The Amati-relation would then constitute an upper envelope of a continuous distribution, a possibility also discussed by \\citet{NP05}. As a result, associating the break in the light curve with the opening angle of a jet, implies an average opening angle considerably larger than typically deduced from observations.  \\citet{Fir06} have analyzed a sample of 15 GRBs with known redshifts and found a good correlation between peak luminosity ($L_{\\rm iso}$), $\\Epk'$ and $\\tem'$ (actually, they used ${\\cal T}'_{45}$ instead of $\\tem'$). They argue that this is a correlation independent from the one between $\\Eiso$--$\\Epk'$--$t'_{\\rm b}$ on the ground that the scatter in the former correlation increases if $t'_{\\rm b}$ is used instead of $\\tem'$ as the third observable. This raises the question whether $\\tem'$ and $t'_{\\rm b}$ are both related to a third time-scale, not included in the analysis, which could then be the appropriate one to use in both of these correlations. If this were the case, the apparent independency would have no physical ground but due to a too limited analysis.  Since we have used peak flux to select our sample, this observable (and, hence, $L_{\\rm iso}$) was not included in the analysis. However, a few of our results have bearing on the relevance of the various time-scales. Although there is a strong correlation between $\\tau$ and $\\tem$, the full 10-parameter analysis clearly indicates that $\\tau$ is the preferred third observable rather than $\\tem$ in the $\\Epk$--$\\fl$-relation. This is in contrast to a limited 3-parameter analysis which makes no difference between $\\tau$ and $\\tem$. Hence, averaging over (or excluding) parameters which contribute only a small amount to a given correlation, can smear out the distinction between direct and indirect dependencies. Furthermore, before applying our method to account for the effects of the unknown redshifts to the $\\Epk$--$\\fl$--$\\tau$ correlation, we took some care to establish that the observed correlation between $\\Epk$ and $\\tau$ is consistent with being due to the unknown redshifts, i.e., that there is no intrinsic correlation between $\\Epk'$ and $\\tau'$. As noted by \\citet{Fir06}, their low value for the sum of residuals $\\chi_{\\rm r}^{2}$ could be due to a correlation between the errors. In fact, in their sample the value of $\\Epk'$ correlates with that for $\\tem'$, indicating that an observable more fundamental than $\\tem'$ may exist (note that the values of $\\Epk'$ and $t'_{\\rm b}$ are uncorrelated). Given our analysis of the $\\Epk$--$\\fl$--$\\tau$ relation, we suggest that this observable could be $\\tau'$, which was not among the parameters considered by \\citet{Fir06}. It is then possible that $\\tem'$ and $t'_{\\rm b}$ are both related to $\\tau'$ and, hence, that $\\Eiso$--$\\Epk'$--$t'_{\\rm b}$ and $L_{\\rm iso}$--$\\Epk'$--$\\tem'$ are not independent correlations. In this context, one may note that in the toy model used above to account for several of the observed properties of GRBs, $\\tau'$ is the fundamental parameter, while the value of $\\tem'$ is derived from $\\tau'$ and the number of pulses."
    },
    "0606/astro-ph0606214_arXiv.txt": {
        "abstract": "While the high-redshift quasar luminosity function closely parallels the hierarchical growth of dark matter halos, at lower redshifts quasars exhibit an anti-hierarchical turn-off, which moves from the most luminous objects to the faintest.    We explore the idea that this may arise from self-regulating feedback, caused by quasar outflows.  Using a detailed hydrodynamic simulation we calculate the luminosity function of quasars down to a redshift of $z=1$ in a large, cosmologically representative volume.  Outflows are included explicitly by tracking halo mergers and driving shocks into the surrounding intergalactic medium, with an energy output equal to a fixed 5\\% fraction of the bolometric luminosity.   Our results are in excellent agreement with measurements of the spatial distribution of quasars on both small and large scales, and we detect an intriguing excess of galaxy-quasar pairs at very short separations. Our results also reproduce an anti-hierarchical turnoff in the quasar luminosity function, however, this falls short of that observed as well as that predicted by analogous semi-analytic models.  The difference can be traced to the treatment of heating of gas within galaxies and the presence of in-shock cooling. Calculations of the mass fraction  of gas above the critical entropy show a strong redshift dependence with close to 20\\% of the baryons being above this limit at $z=1$. Volume fractions show an even stronger trend with redshift and some sensitivity to resolution due to the tendency of high entropy gas to occupy low density regions. The simulated galaxy cluster $L_X-T$ relationship is close to that observed for $z\\approx 1$ clusters, but the simulated galaxy groups at  $z=1$ are significantly perturbed by quasar outflows, suggesting that measurements of X-ray emission in high-redshift groups could well be a ``smoking  gun'' for the AGN heating hypothesis. ",
        "introduction": "In the low-redshift universe, active galactic nuclei (AGN)  are  not very active.   While at high redshifts,  the quasar luminosity function increases with time, since $z \\approx 2$  the number density of optically-selected AGN has been dropping dramatically (Schmidt \\& Green 1983; Boyle \\etal 1988 Koo \\& Kron 1988; Pei 1995; Boyle et al. 2000; Fan et al. 2001).  Deep X-ray surveys have shown that  this downturn occurs anti-hierarchically, such that the spatial density of AGN with higher X-ray luminosities  peaks earlier than that of lower-luminosity AGN (Steffen \\etal 2003; Ueda \\etal 2003). Complementary emission line studies suggest that this trend is driven by a decrease in the characteristic mass of  actively growing black holes (Heckman \\etal 2004), and is likely to closely parallel the formation history of early-type galaxies (\\eg Granato \\etal 2001, 2004). Furthermore, optical and near-infrared observations indicate that the largest galaxies were already in place by z $\\approx 2$, while smaller ones continued to form stars at much lower redshifts (Pozzetti \\etal 2003; Fontana et al. 2004; Glazebrook et al. 2004; van Dokkum et al. 2004; Treu et al. 2005), and a similar trend is observed in the morphological evolution of galaxies (Bundy, Ellis, \\& Conselice 2005).   Yet, despite these observations, such widespread galaxy ``downsizing'' (Cowie \\etal 1996) was unexpected.  The $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model, while in spectacular agreement with observations (\\eg Spergel \\etal 2003), is a hierarchical theory,  in which gravitationally-bound structures grow by accretion and merging. Superimposed on this distribution is the baryonic component, which falls into the dark-matter potential wells, shock heats, and must radiate this energy away before forming stars (Rees \\& Ostriker 1977; Silk 1977).  The larger the structure, the longer it takes to cool, and thus galaxy evolution should be even more hierarchical than structure formation.  Recently, several theoretical studies have shown that the missing element in this picture could be kinetic feedback from AGN.  As bulge and  black hole masses are closely related (Gebhardt \\etal 2000; Merritt \\& Ferrarese 2001; see also King 2003, 2005), such outflows would have the largest impact on the largest forming elliptical galaxies, suppressing their formation first (Scannapieco \\& Oh 2004, hereafter SO04;  Binney 2004; Di Matteo \\etal 2005; Croton  \\etal 2006).  This would also help to explain the X-ray luminosity-temperature relationship observed in the intracluster medium (ICM) in galaxy clusters.  If nongravitational heating were unimportant, the gas density distribution would be self-similar, resulting in $L_X \\propto T^2$  (Kaiser  1986), but instead the observed slope steepens considerably  for low-temperature clusters (\\eg David \\etal 1993; Arnaud \\& Evrard 1999; Helsdon \\& Ponman 2000). Furthermore the $100$ keV cm$^{2}$ level of preheating necessary to explain this discrepancy (Cavaliere \\etal 1998; Kravtsov \\& Yepes 2000;  Wu \\etal 2000;  Babul \\etal 2002)  is not arbitrary.  Rather it corresponds to the threshold value for cooling within the age  of the universe (Voit \\& Bryan 2001; Oh \\& Benson 2003).  Taken together, these observations are strongly suggestive of a model in which the properties of the ICM, the formation history of elliptical galaxies, and the evolution of the quasar luminosity function are all set by self-regulating AGN feedback.  In fact, S004 have shown that the addition of  AGN outflows into the semi-analytic model developed by  Wyithe \\& Loeb (2002; 2003) can reproduce  the drop in the AGN luminosity at low redshifts, as heating this gas slows the accretion of matter onto the supermassive black hole.   The central AGN engine is then starved of fuel and  a strong suppression at the bright end of the luminosity function occurs.  The main drawback of the model, however, is that precise matching of observational  results requires fine tuning of the model parameters, such as black hole  mass and outflow efficiency.  Several other recent numerical and analytic investigations have sharpened our understanding of various aspects of this process. Binney (2004) conducted an analytic study of the impact of AGN on inhibiting gas cooling in large galaxies.  Di Matteo \\etal (2005) carried out Smoothed Particle Hydrodynamic (SPH) simulations of AGN outflows in individual galaxy mergers and studied the role of feedback in determining the  colors of elliptical galaxies and establishing the relationship between stellar velocity dispersion and black hole mass.  This suite of simulations was later related analytically to the global evolution of  quasars and elliptical galaxies in series of papers by Hopkins \\etal (2005a, 2005b, 2005c, 2006).  Levine \\& Gnedin (2005) combined cosmological simulations with an analytic model  to constrain the filling factor of AGN outflows as a function of  redshift.  Scannapieco, Silk, \\& Bouwens (2005) emphasized the role that quasars may play in the downsizing of the star-forming galaxy population.  Levine \\& Gnedin (2006) studied the impact of AGN outflows on the matter power spectrum. Menci \\etal (2006) used a semianalytical model to study the role of AGN feedback on the color distribution of galaxies from $z=0$ to $z=4.$ The importance of the nature of gas accretion in determining the effectiveness of feedback processes has recently been discussed in Dekel \\& Birnboim (2006) and Cattaneo \\etal (2006). Finally, Croton \\etal (2006) combined semi-analytic models with the dark-matter evolution taken from the Millennium  Simulation (Springel \\etal 2005) to study the impact of a more temporally-extended model of AGN feedback on the bright end of the galaxy luminosity function.  In this paper we undertake the first detailed hydrodynamic simulations of quasar outflows in a general cosmological context.  Adopting a burst model that associates AGN with merger events and a global outflow model that is based on our simulations of  high-redshift starbursts (Scannapieco, Thacker, \\& Davis 2001, hereafter STD01), we are able track in detail the impact that quasars have both on their own formation, and on the properties of galaxy clusters and the intergalactic medium (IGM).   The structure of this work is as follows.  In \\S 2 we describe our overall numerical approach, method for quasar identification, and implementation of outflows.  In \\S 3 we compare our simulation results with measures of quasar clustering and the observed luminosity function in the optical and the X-ray bands. In \\S 4 we study AGN feedback in a more global context, examining its impact on the properties of the intergalactic and intracluster media. In \\S5 we present a discussion and conclusion. ",
        "conclusions": "We have presented results from a suite of cosmological simulations that self-consistently follow the evolution of quasars and the outflows associated with them. By tracking the merger history of halos and applying the quasar model of Wyithe \\& Loeb (2003), we have been able to make direct predictions for the spatial distribution and luminosity function of these objects. Our results are in excellent agreement with the observed correlation function of quasars on both small and large scales, reproducing both the power-law behavior measured by the 2DF redshift survey on $\\geq$ 1 Mpc scales, and the strong break measured from the SDSS on $\\leq$ 1 Mpc scales.  Furthermore, we predict that the quasar-galaxy cross-correlation function should show a small scale up-turn relative to the galaxy-galaxy correlation function.  This occurs on sub Mpc scales, and therefore can only be measured from a very large sample of quasars, making its detection difficult even in the large DEEP2 survey. However, it is  worth noting that the {\\em Dark  Energy Survey} will produce a quasar catalog spanning 5,000 square  degrees which combined with photometric redshifts for 300 million galaxies should be able to suppress statistical uncertainties to a sufficiently low level to measure this upturn. We also note that this clustering excess is in  qualitative agreement with the increase in the integrated galaxy  overdensity at small scales for the SDSS quasar sample (Serber \\etal  2006). Although the origin of the excess is uncertain, there is an interesting possibility that the presence of an ancillary galaxy can accelerate the merger process via a three-body interaction. Given the extended nature of the mass distribution associated with galaxies it is not clear whether this mechanism will work in a similar way to stellar interactions. We plan to investigate this intriguing idea in the near future.   As the correlation function is dominated by more populous low-luminosity quasars, our spatial results should be largely interpreted as lending support to our dark-matter modeling choice of merger model.  Similarly, as it tracks the number of quasars formed as a function of the gas mass accreted  onto galaxies,  the luminosity function is more sensitive to the details of our feedback modeling. In this case, our simulation qualitatively reproduces the observed anti-hierarchical  behavior, but the turn-down is much weaker than observed. Matching the suppression at the brighter end requires that we increase the efficiency  of heating from the AGN outflows to mimic that assumed in semi-analytic models. This is equivalent to suppressing one key physical process, namely in-shock cooling in the presence of substructure, and including the ejection of gas from quasar host galaxies. These results emphasize how sensitive the luminosity function is to issues in the baryon physics and how the treatment of these issues in semi-analytic models is still quite approximate.   Investigation of the mass fractions and filling factors of gas above $S_{\\rm crit}$ showed good convergence in the mass fractions at low redshift, but less so in the filling factors. The differences between runs are most noticeable at high redshift, where successively higher resolution leads to the first generation of AGN outflows occurring at earlier epochs. While mass fractions show fairly strong convergence below $z\\approx 2$, the tendency of gas above $S_{\\rm crit}$ to occupy low overdensity regions makes an accurate calculation of the volume filling  factor difficult due to sampling issues. These results are also further  complicated by the known dependence of shock resolution on particle  number (\\eg Thacker \\etal 2000) where accurate modeling of shock jumps  in  spherical collapse is reached once $N_{collapse}>30,000$.  Nonetheless, the results do elucidate that the impact of  quasar outflows is largely felt at low redshifts, in agreement with the downsizing trend. Perhaps the most interesting issue we have not yet explored with regards to filling factors is the relative impact of including a more bi-polar outflow. However, there is no reason to expect  a difference beyond a factor of two as the gas will flow toward low overdensity regions, as observed in supernova outflow calculations (STD01).  In somewhat denser environments, our results for the cluster $L_X-T$ relation are in broad agreement with observations. This is especially encouraging as our cluster modeling is not as sophisticated as other more targeted studies.  In this case, the predominant view is that AGN heating is best modeled in terms of ``hot bubble'' ejection from the brightest cluster galaxy (BCG) which then mixes with the surrounding material.  Interestingly, the exact nature of this mixing is not well understood, and substantial differences are found between Eulerian and Lagrangian simulations, which by definition differ significantly in their treatment of advection.  While SPH simulations inhibit the development of instabilities that promote mixing due to the necessary use of flow interpenetration suppression, it has the advantage of exhibiting significantly less numerical diffusion than many Eulerian methods. Ultimately, input from laboratory experiments may well be necessary to help determine the correct mixing behavior. We also note that due to our inefficient star formation model, our simulations include a significant amount of cold gas in BCGs that will tend to promote an ``overcooling\" instability, which others have avoided by applying phase decoupling (Pearce \\etal 2000).  We believe that this is a significant contributor to the lack of an entropy core in our radial profiles, and removing this central core produces an $L_X-T$ relationship that is significantly less noisy while not exhibiting a large change in normalization at intermediate cluster mass scales. Additionally, the lack of a strong turn-down in the quasar luminosity function at $z=1.2$ also promotes extremely strong outflows in these clusters that may well be transporting high entropy gas from the core out to the edges of the cluster without significantly raising the entropy of the gas immediately surrounding the BCG.   Perhaps the most intriguing result to come out of our $L_X-T$ study is the prediction that the gas in $z \\approx 1$ galaxy groups should be strongly perturbed by AGN activity, which is in the process of turning-off at that epoch.   This effect is confined to small and high-redshift clusters for two main reasons. At lower redshifts, the comparative paucity of AGN activity will allow group gas to evolve largely adiabatically, decreasing $L_X$ dramatically to establish the steep $L_X-T$ relation ($L_X \\propto T^5$) observed locally. In more massive $z \\approx 1$ clusters, on the other hand, the effect of outflows is largely masked when averaging over the material within $r_{200}.$  Thus galaxy groups at $z \\approx 1$ represent the key mass and redshift range at which the AGN heating hypothesis is most likely to be testable though observations.   Lastly, our simulation results indicate that simplifications in current semi-analytic models may well be downplaying critical physics. While much attention has been paid to drawing broad conclusions from comparisons between observations and these models, the differences we have uncovered in this investigation are troubling.  While, as expected, post processing of our simulation results was able to reproduce the semi-analytic behavior, it is clear that in-shock cooling due to substructure in low-overdensity halos is an issue that must be considered carefully in semi-analytic calculations.   Fortunately,  constraining the effect by simulation would not be difficult,  and ultimately a single parameter could be used to quantify this behavior. On the simulation side, it also is clear that a more detailed understanding of gas ejection from quasar hosts will be necessary for definitive conclusions to be reached. Nonetheless, the insight gained from this initial simulation study clearly serves to highlight the promise of a self-regulating picture of quasar formation.  The simple merger  model of Wyithe \\& Loeb (2003), when supplemented with an outflow model, appears to represent a significant step forward in understanding the evolution of quasar clustering and the cause of their antiheirarchical low-redshift turn-off."
    },
    "0606/astro-ph0606022_arXiv.txt": {
        "abstract": "% We present explicit examples to show that the `compatibility criterion' [recently obtained by us towards providing equilibrium configurations compatible with the structure of general relativity] which states that: for a given value of $\\sigma [\\equiv (P_0/E_0) \\equiv $ the ratio  of  central  pressure  to   central   energy-density], the compactness ratio $u [\\equiv (M/R)$, where $M$ is the total mass and $R$ is the radius of the configuration]   of  any static configuration cannot  exceed the compactness ratio, $u_h$, of the homogeneous density sphere (that is, $u \\leq  u_h$), is capable of providing a {\\em necessary} and {\\em sufficient} condition for any regular configuration to be compatible with the state of hydrostatic equilibrium. This conclusion is drawn on the basis of the finding that the $M-R$ relation gives the necessary and sufficient condition for dynamical stability of equilibrium configurations only when the compatibility criterion for these configurations is appropriately satisfied. In this regard, we construct an appropriate sequence composed of core-envelope models on the basis of compatibility criterion, such that each member of this sequence satisfies the extreme case of causality condition $v = c = 1$ at the centre. The maximum stable value of $u \\simeq 0.3389$ (which occurs for the model corresponding to the maximum value of mass in the mass-radius relation) and the corresponding central value of the local adiabatic index, $(\\Gamma_1)_0 \\simeq 2.5911$, of this model are found fully consistent with those of the corresponding {\\em absolute} values, $u_{\\rm max} \\leq 0.3406$, and $(\\Gamma_1)_0 \\leq 2.5946$, which impose strong constraints on these parameters of such models. In addition to this example, we also study dynamical stability of pure adiabatic polytropic configurations on the basis of variational method for the choice of the `trial function', $\\xi =re^{\\nu/4}$, as well as the mass-central density relation, since the compatibility criterion is appropriately satisfied for these models. The results of this example provide additional proof in favour of the statement regarding compatibility criterion mentioned above. Together with other results, this study also confirms the previous claim that just the choice of the `trial function', $\\xi =re^{\\nu/4}$, is capable of providing the necessary and sufficient condition for dynamical stability of a mass on the basis of variational method. Obviously, the upper bound on compactness ratio of neutron stars, $u \\cong 0.3389 $, which belongs to two-density model studied here, turns out to be much stronger than the corresponding `absolute' upper bound mentioned in the literature.     {\\it PACS Nos.: 04.20.Jd; 04.40.Dg; 97.60.Jd.} ",
        "introduction": "Einstein's   field    equations for static and spherically symmetric mass distribution were first solved  by Schwarzschild (1916). The first solution describes the geometry of the space-time exterior to a prefect fluid sphere in hydrostatic equilibrium. While the other, known as interior Schwarzschild solution, corresponds to the interior geometry of a fluid sphere of constant (homogeneous) energy-density, $E$. The importance of these two solutions in general relativity is well known. The interior Schwarzschild solution provides  two  very  important  features  towards obtaining configurations in  hydrostatic  equilibrium,  compatible with general relativity, namely - (i) It gives an  absolute  upper limit on compactness ratio, $u (\\equiv M/R$, mass  to  size  ratio  of the entire configuration in geometrized units) $\\leq (4/9)$ for any static and spherical configuration (belonging to arbitrary density profiles, provided the density does not increase outwards) in hydrostatic equilibrium (Buchdahl 1959; Weinberg 1972), and (ii) For an assigned value of the  compactness  ratio, $u$, and radius $R$ (or mass $M$),  the  minimum  central  pressure, $ P_0$,   corresponds   to   the homogeneous density solution (see,  e.g., Weinberg 1972).  Despite the non linear and coupled differential equations, various  exact solutions of the field equations for  static  and  spherically  symmetric   metric   are available  in  the  literature (see, e. g., Kramer et al 1980) which may be used to obtain  various  physical  properties  of  spherical  and static compact object (provided they are physically realistic).  Knutsen (1988; 1989) examined  physical properties of the various exact solutions and found that  these  solutions correspond to nice physical  properties  and  also  remain  stable against small radial pulsations upto certain values of $u$.  Another way to explore the  physical  properties  of compact objects like neutron star, one may expect to have some  physically viable equation of state (EOS). However,  for  such  objects  the equations of state (EOSs) are not well known  [empirically]  because  of the lack of knowledge of nuclear interactions beyond the  density $ \\sim 10^{14} {\\rm g\\, cm}^{-3}$ (Dolan 1992), and the only way to obtain EOSs far beyond  this density  range  is  extrapolation.  Various   such extrapolated equations  are available in the literature (Arnett \\& Bowers 1977). As  a  way  out, one  can  impose  some  restrictions  upon  the  known   physical quantities,  such  that,  the   speed   of   sound   inside   the configuration, $v [\\equiv \\sqrt (dP/dE)]$, does  not  exceed  the  speed  of light in vacuum, i.e., $v \\leq c =  1$  (in  geometrized  units),  and obtain various physical properties, like upper bound on stable neutron  star  masses (Rhoades \\& Ruffini 1974; Brecher \\& Caporaso 1976; Hartle 1978; Friedman \\& Ipser 1987). Haensel and Zdunik (1989) have shown that the  only  EOS  which  can describe a submillisecond pulsar and the static mass of $1.442  M_\\odot $ simultaneously, corresponds to the EOS, $(dP/dE) = 1$, however, they emphasized that this EOS represents an `abnormal' state of matter in the sense that pressure vanishes at densities of the order  of nuclear density or even higher.  We have recently proposed a core-envelope model with stiffest EOS, $(dP/dE) = 1$, forming the core and a polytropic equation with constant adiabatic index $\\Gamma_1$ = (d$lnP/$d$ln \\rho)$ [where $P$ is the pressure and $\\rho$ represents the rest-mass density] describing the envelope, such that $P, E$, both of the metric  parameters ($\\nu$, and $\\lambda$), their first derivatives, and the speed of sound are  continuous  at  the core-envelope boundary  and  at  the  exterior  boundary (surface)  of  the structure (Negi \\& Durgapal 2000). The other remarkable feature of this core-envelope model is that not only the `abnormalities' (in the sense discussed in the literature, see, e.g. Lee 1975; Haensel \\& Zdunik 1989) disappear, the maximum value of $u \\cong 0.3574$ for the stable configuration turns out to be as large as that obtained by using the EOS, $(dP/dE) = 1$, alone (see, e. g., Haensel \\& Zdunik 1989). The  maximum  stable mass  of neutron star based upon this model (by using the maximum value of $u \\cong 0.3574$ for stable configuration) turns out to be $7.944 M_\\odot$,  if the (average) density  of  the  configuration  is  constrained  by fastest rotating pulsar, with rotation period, $P_{rot} \\cong 1.558$ ms, known to date.  The  model  gives  pulsationally  stable configurations with compactness ratio $u > (1/3) $, which  are  important  to study Ultra-Compact Objects (UCOs) [see, e. g., Negi \\& Durgapal 1999a, b; 2000; and references therein].  Recently, by using property (ii) of the homogeneous density sphere as mentioned above,  we  have connected  the compactness ratio, $u$, of any static and spherical configuration with the corresponding ratio  of  central pressure  to   central   energy-density $\\sigma [\\equiv (P_0/E_0)$] and worked out  an important criterion which concludes that for a given value of $\\sigma$, the maximum value of compactness ratio, $u(\\equiv u_h)$, should always correspond to  the  homogeneous density sphere (Negi \\& Durgapal 2001).  An examination of this criterion on some well known exact solutions and EOSs indicated  that this criterion, in fact, is fulfilled only by two types of configurations corresponding to a {\\em single} EOS or density variation: (i) the  {\\em regular} (positive finite density  at  the  origin  which  decreases monotonically outwards) configurations which correspond to a vanishing density at  the  surface  together with pressure [so called, the gravitationally-bound structures] (Negi \\& Durgapal 2001; Negi 2004b; Negi 2006), and (ii) the structures which correspond to a non-vanishing surface density but exhibit singularities at the centre, in the sense that both pressure and density become infinity as $r \\rightarrow 0$ [so called, the self-bound singular structures] (Negi 2004b; Negi 2006). On the other hand, it is seen that the EOSs or analytic  solutions,  corresponding  to  a non-zero {\\em finite}, surface density (that is,  the  pressure  vanishes  at finite surface density, and so called the self-bound regular structures), in fact, do  not fulfill this criterion (Negi \\& Durgapal 2001; Negi 2004b; Negi 2006). We have shown this inconsistency particularly for the EOS, $(dP/dE) = 1$ (as it represents the most successful EOS to obtain the various extreme characteristics of neutron stars as discussed above).  In addition to the self-bound regular configurations corresponding to a single density variation, the compatibility criterion may not be satisfied by various two-density, or multiple-density, regular, gravitationally-bound structures. Such structures are widely discussed in the literature, particularly, for determining the upper bound on neutron star (NS) masses. In this connection, we would consider the core-envelope model proposed by Negi \\& Durgapal (2000).  The reason(s) behind non-fulfillment of the criterion obtained in the study of Negi and Durgapal (2001) by various exact self-bound regular solutions and EOSs, as well as the two-density, gravitationally-bound, regular structures, and their further implications are discussed in the following sections. ",
        "conclusions": "We have re-investigated the core-envelope model with  stiffest  equation  of  state [speed of sound equal to  that  of  light]  in  the  core  and  a polytropic  equation  with  constant   adiabatic   index  $\\Gamma_1 = [$d$lnP/$d$ln\\rho]$ in the envelope, based upon the criterion obtained by Negi and Durgapal (2001). We find that the condition of hydrostatic equilibrium is assured only when the minimum ratio of $(P_b/E_b)$ at the core-envelope boundary reaches about $2.9201 \\times 10^{-1}$ [that is, when the value of the adiabatic index, $\\Gamma_1$ at the core-envelope boundary reaches around 4.4246 as compared to the previous case of $\\Gamma_1 \\cong 72.4286$]. Under this condition, the pressure, density, both of the metric parameters  including  their first derivatives, and the speed of sound are  continuous  at  the core-envelope boundary  and  at  the  surface. The mass-radius diagram indicates that the configuration remains dynamically stable upto a $u$ value as large as 0.3389. The corresponding central value of the local $(\\Gamma_1)_0$ is obtained as 2.5911. These values are fully consistent with those of the {\\em absolute} values, $u_{max} \\cong 0.3406$, and $(\\Gamma_1)_{0,{\\rm max}} \\cong 2.5946$, compatible with the structure of general relativity, causality, and dynamical stability, obtained by using a causal configuration of homogeneous energy-density (Negi 2004a). Not just for the particular model considered in the present study, we have also found that the maximum value of $u$ and the corresponding value of $(\\Gamma_1)_0$ upto which the configurations remain pulsationally stable can not exceed the values of $u \\simeq 0.34$ and $(\\Gamma_1)_0 \\simeq 2.5128$ respectively, if the `compatibility criterion' is followed and the envelope of the present model is replaced by the polytropic EOS $ (4/3) \\leq $d$lnP/$d$ln\\rho \\leq 2 $ (Negi 2005).  In addition to the study of two-density models considered in the present study (sec.4), the dynamical stability of the polytropic configurations (sec.5) explicitly shows that the compatibility criterion (Negi \\& Durgapal 2001) alone is capable of providing a necessary and sufficient condition for any regular configuration to be consistent with the structure of general relativity [however, the study of $M-R$ relation of such sequences (corresponding to two-density models with $v = c = 1$ at the centre), consistent with the definition of actual mass (Negi 2004 b; Negi 2006) and the `compatibility criterion' (Negi \\& Durgapal 2001) in the near future will finally settle down this issue].   The two-density  structures   are   dynamically   stable    and gravitationally bound even for the value of compactness ratio, $u \\geq  (1/3)$,  thus  giving  a  suitable  model  for  studying   the Ultra-compact Objects [UCOs] discussed in the literature (see, e. g., Negi \\& Durgapal 1999a, b; 2000; and references therein). The present type of studies may also find application to test various models of NSs based upon EOSs of dense nuclear matter, and the models of relativistic stellar objects like - star clusters."
    },
    "0606/astro-ph0606744_arXiv.txt": {
        "abstract": "The detection of  GRB 060218 at z=0.033 by {\\em Swift} within 1.5 years of operation, together with the detection of GRB 980425 at $z=0.0085$ by BeppoSAX, suggest that these low-luminosity (LL) GRBs have a much higher event rate than the canonical high-luminosity (HL) GRBs, and they form a distinct new component in the GRB luminosity function. We explore the contribution of this previously neglected GRB population to the diffuse neutrino background within the internal shock model and compare it with that of the canonical HL population. By considering a wide range of distributions of various parameters (e.g. luminosity, spectral break energy, duration, variability time, Lorentz factor, redshift) for both populations, we find that although it is difficult to detect neutrinos from the individual LL GRBs, the contribution of the LL population to the diffuse neutrino background is more than the HL population above about $10^{8}$ GeV. ",
        "introduction": "The possibility of high energy neutrino emission from GRBs in the burst and the afterglow phases has been studied earlier by different groups \\cite{wax1,all,razzaque,raz2}.  It has been found that only the nearby high luminosity (HL) GRBs (e.g. GRB 030329) are expected to be detected individually by neutrino telescopes like ICECUBE \\cite{razzaque,ice1}. In general the diffuse neutrino background from the whole GRB population is of great observational interests. Recently a low-luminosity (LL), long-duration event, GRB 060218/SN 2006aj was detected by {\\it Swift} \\cite{campana} at a redshift of 0.0331. Earlier, another low luminosity event, GRB 980425/SN 1998bw was detected by {\\it BeppoSAX} at an even lower redshift of 0.0085 \\cite{galama,kul}. These two LL events share some common characteristics such as low luminosity, long duration, and low isotropic $\\gamma$-ray energy. More importantly, their detections at very low redshifts within a relatively short period of time (1.5 years for {\\em Swift} and 6 years for {\\em BeppoSAX}) imply that they have a much higher event rate than the canonical HL GRBs \\cite{cobb,pian,soderberg,liang}. A straightforward estimate of the local rate of LL-GRBs $\\rho_{0}^{\\rm LL}$ can be derived from \\beq \\rho_{0}^{\\rm LL}V(z<0.033)\\Big(\\frac{\\Omega^{Beppo}}{4\\pi}T^{Beppo} +\\frac{\\Omega^{Swift}}{4\\pi}T^{Swift}\\Big)\\sim 2, \\eeq where $V(z<0.033)\\sim1.2\\times10^{-2}{\\rm Gpc}^3$ is the volume enclosed by $z=0.033$, $\\Omega^{Beppo}=0.123$ and $\\Omega^{Swift}=1.33$ are the solid angles of the GRBM on board {\\it BeppoSAX} and the BAT on board {\\it Swift} respectively. The durations of observation are $T^{Beppo} \\sim 6$ yr and $T^{Swift}\\sim 1.5$ yr (when this paper has been written) for {\\it BeppoSAX} and {\\it Swift}, respectively. The number 2 on the right hand side of Eq.(1) accounts for the detection of GRB 980425 and GRB 060218. This gives a rough estimate of $\\rho_{0}^{\\rm LL}$ as $800 {\\rm Gpc^{-3}yr^{-1}}$. Several groups independently estimated this event rate\\cite{cobb,soderberg}, e.g. $\\sim (300-1000) {\\rm Gpc^{-3} yr^{-1}}$, which are comparable to the rate of LL GRBs $550_{-430}^{+700} {\\rm Gpc^{-3} yr^{-1}}$ obtained by Liang et al.\\cite{liang}. These are all much higher than the event rate of HL GRBs, $\\rho_{0}^{\\rm HL}\\sim 1 {\\rm Gpc^{-3}yr^{-1}}$\\cite{schmidt}, and its simple extrapolation to the LL regime assuming a same population, e.g. $\\rho_{0} \\sim 10 {\\rm Gpc^{-3}yr^{-1}}$\\cite{guetta}. If one assumes that the two LL GRBs within $z \\le 0.033$ belong to the population of HL GRBs, then the expected number of detection of GRBs within $z \\le 0.033$ in the {\\it BeppoSAX} and {\\it Swift} era is nearly 0.0026 for $\\rho_{0}^{\\rm HL}\\sim 1 {\\rm Gpc^{-3}yr^{-1}}$, and is 0.026 for $\\rho_{0} \\sim 10 {\\rm Gpc^{-3}yr^{-1}}$. The Poisson probability of detecting two events during the observational period is $3.4 \\times 10^{-6}$ for the former and $3.3\\times 10^{-4}$ for the latter. Hence, the detection of GRB 060218 strongly suggests the presence of a distinct population (LL), which is not a subset of the canonical HL population.  This implies two distinct components in the GRB luminosity function\\cite{liang}. One may characterize the luminosity function of each population by a broken power law, i.e. \\beq S(L)=S_{0}\\Big[\\Big(\\frac{L}{L_b}\\Big)^{\\alpha_1}+ \\Big(\\frac{L}{L_b}\\Big)^{\\alpha_2}\\Big]^{-1} \\eeq The spectral indices and break luminosities of this function for LL and HL populations have been derived/constrained in \\cite{liang}. The values of $\\alpha_1$, $\\alpha_2$, $L_b$ are $0.6$, $4.5$, $10^{47}$erg/sec and $1.15$, $2.5$, $10^{51}$erg/sec for LL and HL GRBs respectively. The values of $S_{0}$ depend on the upper and lower limits in the values of luminosity for each population.  Although the high energy neutrino flux of these individual LL GRBs is low, the contribution to the diffuse neutrino flux by the whole LL population is non-negligible due to their very high event rate. In this paper we explore the neutrino emission of this hitherto ignored LL-GRB population, and compare the contributions of the LL- and HL- GRB populations to the diffuse neutrino background. ",
        "conclusions": "The neutrino spectra from HL and LL populations are plotted in units of ${\\rm GeV~ cm^{-2}~ s^{-1}~ sr^{-1}}$ in Fig.1. for $\\gamma_1=1$, 1.1, 1.3 and 1.5 and $\\gamma_2=2.25$.  We have assumed $\\epsilon_e=0.3$ and $\\epsilon_B=0.3$ in our calculations. The long dashed (green) lines represent neutrino fluxes from the LL population and the solid (red) lines are for the HL population. The Waxman-Bahcall \\cite{limit1} limit with $z$ evolution (including all three flavors) has been shown by short dashed (blue) lines, and the dotted (pink) lines represent the sensitivity of the IceCube detector \\cite{ice2} after three years of operation. The current limit from AMANDA-II data taken between 2000 and 2003 \\cite{amanda} is also shown by triple dotted line above the WB upper bound. The HL component has an average maximum energy of $6.3\\times10^9$ GeV and for LL population it is $1.5\\times10^9$ GeV. Above about $10^8$ GeV the contribution of LL population to the diffuse neutrino background is more than the HL population. The LL component continues to below $100$ GeV but the HL component has a lower cut-off energy of about $1000$ GeV. Between $1000$ GeV and $4\\times 10^{4}$ GeV their fluxes are equal. We have also plotted the GZK neutrino flux from \\cite{stecker} in thin dashed lines for a comparison.  Comparing with the Icecube sensitivity, we can see that the GZK neutrinos are generally not detectable by Icecube, while the contribution of the LL-component would be detectable if $\\gamma_1\\geq1.1$. In this case we expect 30$km^{-2}sr^{-1}$ neutrino events per year or more in IceCube at an energy of $10^7$GeV. If we compare our derived neutrino spectra from HL GRBs with some of the previous results \\cite{raz2} the width of the intermediate region between the two break energies is narrower. This is because we have used the muon cooling energy instead of the pion cooling energy in determining the second break energy of the spectra to avoid overestimation of neutrino fluxes. However, in this way we slightly underestimate the $\\nu_{\\mu}$ fluxes as they are produced by $\\pi^{+}$ decay unlike $\\bar\\nu_{\\mu}$ and $\\nu_e$.  We also find that our results are very sensitive to variations in the mean value of Lorentz factor. As the break energies in the neutrino spectra depend on $\\Gamma^2$ and $\\Gamma^4$ a slight deviation in the value of $\\Gamma$ changes the neutrino fluxes significantly. In the future it might be possible to set limits on the values of GRB Lorentz factors if the detection or upper limit of high energy neutrino background from GRBs is set up. Overall, our results show that the previously neglected LL population of GRBs significantly contribute to the diffuse neutrino background as their event rate is very high. This conclusion would be strengthened as {\\em Swift} detects more LL GRBs in the years to come.  Finally we notice that after we submitted the first version of this paper to astro-ph, similar results were independently reported by Murase et al.\\cite{murase}.  \\begin{figure*}[t] \\begin{center} \\centerline{ \\epsfxsize=15.cm\\epsfysize=10.cm \\epsfbox{tara1.eps} } \\end{center} \\caption{The diffuse neutrino fluxes expected on earth from low and high luminosity populations of GRBs distributed up to a redshift of 5, being plotted in units of ${\\rm GeV ~cm^{-2}~ s^{-1} sr^{-1}}$ against the observed neutrino energy. The four panels are for photon spectral indices $\\gamma_1=1, 1.1, 1.3, 1.5$ (clockwise from left bottom), and $\\gamma_2=2.25$ are adopted for all of them. The Waxman-Bahcall limit with $z$ evolution \\cite{limit1} has been shown with short dashed (blue) line. The double dotted (pink) line represents IceCube experiment's sensitivity \\cite{ice2} after 3 years of operation and the current limit from AMANDA-II data (2000-2003) \\cite{amanda} triple dotted (orange) line is just above the Waxman-Bahcall limit. These limits include all the three flavors of neutrinos. GZK neutrino flux has been plotted in thin dashed line from \\cite{stecker} for comparison.} \\label{fig:Nnu} \\end{figure*}  We thank the anonymous referee for constructive reports that help to improve the paper.  This work was supported by NASA under grants NNG05GB67G and NNG06GH62G."
    },
    "0606/astro-ph0606058_arXiv.txt": {
        "abstract": "We present new parametrisations of energy spectra of secondary particles, $\\pi$-mesons, gamma-rays, electrons and neutrinos,  produced in inelastic proton-proton collisions. The simple analytical approximations based on simulations of proton-proton interactions  using the public available SIBYLL  code,  provide very good accuracy for energy distributions of  secondary products in the energy range  above 100 GeV. Generally, the recommended analytical formulae deviate from the simulated distributions within a few percent over a large range of $x=E_i/E_p$ - the fraction of energy of the incident proton transferred to the secondaries. Finally, we describe  an approximate procedure of continuation of calculations towards low energies, down to the threshold of $\\pi$-meson production. ",
        "introduction": "Introduction}  Any reliable interpretation of an astronomical observation requires not only high quality experimental information concerning the spectral, temporal and spatial properties of radiation, but also unambiguous identification of  the relevant radiation processes.  In this regard,  the good knowledge of characteristics of radiation mechanisms is a principal issue in astrophysics, in particular  in gamma-ray astronomy where we often face a problem when the same observation can be equally well explained by two or more radiation processes.  Fortunately, all basic radiation processes relevant to high energy gamma-ray astronomy can be comprehensively studied using the methods and tools of experimental and theoretical physics. This concerns, in particular, one of the most principal gamma-ray production mechanisms  in high-energy astrophysics -- inelastic proton-proton interactions with subsequent decay of the secondary $\\pi^0$ and $\\eta$-mesons into gamma-rays. The decays of accompanying $\\pi^\\pm$-mesons and some other (less important) short-lived secondaries result in  production of high energy neutrinos. This  establishes  a deep link between the high energy gamma-ray and neutrino astronomies. Finally, the secondary electrons (positrons) from p-p interactions may compete with directly accelerated electrons and thus significantly contribute to the nonthermal electromagnetic radiation of gamma-ray sources from radio to hard X-rays.  The principal role of this process in high energy astrophysics was recognized long ago by the pioneers of the field (see e.g. \\cite{Morrison,GinzSyr,Hayakawa}), in particular in the context of their applications in gamma-ray (e.g. \\cite{Stecker_book}) and neutrino (e.g. \\cite{BerVol}) astronomies. The first reliable experimental results obtained with the SAS-2 and COS-B gamma-ray satellite missions, initiated new, more detailed studies of $\\pi^0$-production in $p-p$ interactions \\cite{StepBad1981,Dermer86}, in particular for the interpretation of the diffuse galactic gamma-ray background emission. The prospects of neutrino astronomy initiated similar calculations for high energy neutrinos and their links to gamma-ray astronomy  \\cite{Berez_Book,Gaisser,Berez91,Vissani}.  The observations of diffuse gamma-ray emission from different parts of the galactic disk by EGRET revealed a noticeable excess of flux at energies of several GeV. Although this excess can be naturally explained by assuming somewhat harder proton spectrum (compared to the locally measured flux of cosmic rays) \\cite{Mori,AhAt00,StrongMosk00}, this result recently initiated a new study of the process \\cite{Kamae05} based on the approach of separation of the diffractive and non-diffractive channels of interactions as well as incorporating violation of the Feynman scaling.  It is remarkable that while precise calculations of gamma-ray spectra require quite heavy integrations over differential cross-sections measured at laboratory experiments, the emissivity of  gamma-rays for an arbitrary broad and smooth, e.g. power-law, energy distribution of protons can be obtained, with a quite reasonable accuracy, within a simple formalism (see e.g.  Ref.\\cite{AhAt00}) based on the assumption of a constant fraction $\\kappa$ of energy of the incident proton released in the secondary gamma-rays (see \\cite{Torres} for comparison of different approaches). This approach relies on  the energy-dependent total inelastic cross-section of pp interactions, $\\sigma_{\\rm pp}^{\\rm inel}(E)$, and assumes  a fixed value of the parameter $\\kappa \\approx 0.17$, which  provides the best agreement with the accurate numerical calculations over a wide  energy range of gamma-rays.  On the other hand, in the case of sharp spectral features  like pileups or cutoffs in the proton energy distribution, one has to perform accurate numerical calculations based on simulations of inclusive cross-sections of production of secondary particles. Presently three well developed  codes of simulations of \\textit{p-p} interactions are public available  --  Pythia \\cite{Pythia_model}, SIBYLL \\cite{SIBYLL_model}, QGSJET \\cite{QGSJET_model}. The last two as well as some other models are combined in the more general CORSIKA code \\cite{CORSIKA} designed for simulations of interactions of cosmic rays with the Earth atmosphere.  These codes are based on  phenomenological models of  \\textit{p-p} interactions incorporated with comprehensive  experimental data obtained at particle accelerators.  While these codes can be directly used for calculations of gamma-ray spectra for any distribution of primary protons,  it is quite useful to have simple analytical parameterizations which not only significantly reduce  the calculation time, but also allow better understanding of characteristics of secondary electrons, especially for the distributions of parent protons with distinct spectral features. This concerns, for example, such an important question as the extraction of the shape of the proton energy distribution in the cutoff region based on the analysis of the observed gamma-ray spectrum.  Indeed, while the delta-functional approach implies, by definition, similar spectral shapes of gamma-rays and protons (shifted in the energy scale  by a factor of $\\kappa$), in reality the energy spectrum of highest energy gamma-rays appears, as we show below,  smoother than the distribution of protons in the corresponding (cutoff) region. Thus,  the use of simple approximations not always can be justified, and, in fact,  they may  cause misleading astrophysical  conclusions  about the energy spectra of accelerated protons.   Two different parameterizations of inclusive cross-sections for pion-production in proton-proton interactions has been published by Stephens  and Badwar \\cite{StepBad1981} and Blattnig et al \\cite{Blattnig}. However,  recently  it was recognized  \\cite{Torres} that   both parameterizations do not  describe correctly  the gamma-ray spectra in the high energy regime. The parameterization of Stephens and Badwar underpredicts the yield of high energy pions.  The parameterization by Blattnig \\textit{et al.} is valid for  energies below 50 GeV; above this energy  it significantly overpredicts the pions production.  In this paper we present new parameterizations for high energy spectra of gamma-rays, electrons and neutrinos based on the simulations of proton-proton interactions  using  the SIBYLL code \\cite{SIBYLL_model}, and partly (only for distributions of $\\pi$-mesons) the QGSJET \\cite{QGSJET_model} code. We provide simple analytical approximations  for the energy spectra of secondaries with  accuracies  generally better than several percent. ",
        "conclusions": "We present new parameterizations of energy spectra of gamma-rays, electrons and neutrinos produced in proton-proton collisions. The results are based on the spectral fits of secondaries produced in inelastic \\textit{p-p} interactions simulated by the SIBYLL code. The recommended simple analytical presentations provide accuracy better than 10 percent over the energy range of parent protons $0.1-10^5$ TeV and $x=E_i/E_p \\geq 10^{-3}$."
    },
    "0606/astro-ph0606328_arXiv.txt": {
        "abstract": "We measure the mean halo-mass of $z\\!\\simeq\\!0.5$ \\MgII\\ absorbers using the cross-correlation (over co-moving scales 0.05--13$h^{-1}{\\rm \\,Mpc}$) between \\nabstot\\ \\MgII\\ quasar absorption systems and \\ngalsim\\ Luminous Red Galaxies (LRGs), both selected from the Sloan Digital Sky Survey Data Release 3.  The \\MgII\\ systems have \\ma\\ rest-frame equivalent widths $\\EW\\ga 0.3$\\,\\AA.  From the ratio of the \\MgII--LRG cross-correlation to the LRG--LRG auto-correlation, we find that the bias ratio between \\MgII\\ absorbers and LRGs is $\\overline b_{\\rm \\MgII}/\\overline b_{\\rm LRG}=\\Arelcorr$, which implies that the absorber host-galaxies have a mean halo-mass $\\massratio$ times smaller than that of the LRGs; the \\MgII\\ absorbers have haloes of mean mass $\\langle \\log \\Mh (\\msun)\\rangle=\\;$\\LRGmassrangeallsys. We demonstrate that this statistical technique, which does not require any spectroscopic follow-up, does not suffer from contaminants such as stars or foreground and background galaxies.  Finally, we find that the absorber halo-mass is anti-correlated with the equivalent width. If \\MgII\\ absorbers were virialized in galaxy haloes a positive \\Mh--\\EW\\ correlation would have been observed since \\EW\\ is a direct measure of the velocity spread of the \\MgII\\ sub-components. Thus, our results demonstrate that the individual clouds of a \\MgII\\ system are {\\it not} virialized in the gaseous haloes of the host-galaxies. We review past results in the literature on the statistics of \\MgII\\ absorbers and find that they too require an \\Mh--\\EW\\ anti-correlation.  When combined with measurements of the equivalent width distribution (\\dNdzdW), the \\Mh--\\EW\\ anti-correlation naturally explains why absorbers with $\\EW\\ga2$\\,\\AA\\ are not seen at large impact parameters.  We interpret the \\Mh--\\EW\\ anti-correlation within the starburst scenario where strong \\MgII\\ absorbers are produced by supernovae-driven winds. ",
        "introduction": "The connection between quasar (QSO) absorption line (QAL) systems and galaxies is crucial to our understanding of galaxy evolution since QALs provide detailed information about the physical conditions of galaxy haloes out to large impact parameters ($\\rho\\!>\\!100{\\rm \\,kpc}$) with no direct dependence on the host galaxy luminosity.  \\MgII\\ absorbers are ideal for this purpose as the \\MgII\\LL\\ 2796,2803 doublet can be detected from $z\\simeq0.3$ to $z\\simeq 2.2$ at optical wavelengths. Since the ionization potential of \\MgI\\ is less than 13.6\\,eV but the ionization potential of \\MgII\\ is greater than 13.6\\,eV, \\MgII\\ absorbers trace cold gas. In fact \\MgII\\ absorbers with equivalent widths $\\EW\\ga0.03$\\,\\AA\\ have been shown to be associated with \\HI\\ absorbers covering five decades in \\HI\\ column density ($N_{\\HI}$) \\citep[e.g.][]{PetitjeanP_90a}, including sub-Lyman limit systems \\citep{ChurchillC_99a,ChurchillC_00a}, Lyman limit systems \\citep[e.g.][]{BergeronJ_86a,SteidelC_92a} and damped \\lya\\ systems \\citep[DLAs; e.g.,][]{LeBrunV_97a,RaoS_00a,BoisseP_98a,ChurchillC_00b,RaoS_06a}, which means that a large range of galactic environments are likely to be sampled. For example,  strong absorbers with $\\EW\\!\\geq\\!0.3$\\,\\AA\\ are known to be generally associated with galaxies \\citep{LanzettaK_90a,BergeronJ_91a,BergeronJ_92a,SteidelC_92a,DrinkwaterM_93a,SteidelC_94a}. These groups have shown that galaxies responsible for the \\MgII\\ host-galaxies have luminosities consistent with normal field $0.7L^*_B$ galaxies, and \\citeauthor{SteidelC_94a} showed that, from their average colour, they are on average late-type (Sb) galaxies, a result reproduced by \\citet{ZibettiS_05a}.  Medium-resolution spectroscopy of the \\MgII\\ absorbers quickly revealed that such absorbers are composed of several sub-components \\citep{BergeronJ_86a,TytlerD_87a,PetitjeanP_90a} and that the number of sub-components    strongly correlates with equivalent width \\citep{PetitjeanP_90a}. Building on this work, \\citet{ChurchillC_97a} showed that the mean Doppler width of individual components is 5\\,\\kms, with an rms of comparable magnitude, using high resolution spectra (${\\rm FWHM} \\simeq 6$\\,\\kms).  This corresponds to a thermal temperature of $\\sim$30,000~K. \\citet{ChurchillC_97a} also directly constrained the turbulent component to be $\\la 2$\\,\\kms.  \\citet{ChurchillC_03a}, using high-resolution spectra, confirmed that equivalent width does strongly correlate with the number of sub-components as shown by \\citet{PetitjeanP_90a}. This correlation arises because a large equivalent width can only be produced by more components spread over a large velocity range since extremely few components with large Doppler widths are seen. The equivalent width, \\EW, should therefore be correlated with the velocity range, $\\Delta v$, covered by the sub-components and, indeed, this is observed to be the case \\citep[e.g.][see their figure 3]{EllisonS_06a}. The velocity range for strong \\MgII\\ systems are very large, from 50 to 400~\\kms. If the individual clouds were virialized within the haloes of the host galaxies, $\\Delta v$ would represent the velocity dispersion of the many clouds in the host-galaxy halo. In this case, $\\Delta v$ would be directly related to the mass of the host-galaxy. Since a larger \\EW\\ is achieved by having a larger number of clouds spread over a larger $\\Delta v$, $\\EW$ should also be positively correlated with the mass of the host galaxies {\\it if} the individual clouds are virialized within the haloes of the host galaxies.  Among the few correlations observed between the host-galaxy properties and the absorber properties, \\citet{LanzettaK_90a} noted a significant anti-correlation between \\MgII\\ equivalent width and the impact parameter distribution from the sample of \\citet{BergeronJ_91a}. While \\citet{BergeronJ_91a} argued that the anti-correlation was not significant, it was clearly seen in the sample of \\citet{SteidelC_92a} and \\citet{SteidelC_95b}: absorbers with $\\EW \\ga 2$\\,\\AA\\ are observed at small impact parameters but do not exist at large impact parameters ($\\rho\\ga 50$\\,kpc). This is puzzling as large \\EW\\ absorbers at large impact parameters ought to be easily detectable.  Only a slight correlation is observed between absorption cross-section and galaxy luminosity. \\citet{SteidelC_93a} and \\citet{SteidelC_95b} showed that the cross-section of \\MgII\\ absorbers with $\\EW\\!\\geq\\!0.3$\\,\\AA\\ has radius $R_{\\rm phys}\\!\\sim\\!40\\;h^{-1}{\\rm \\,kpc}$ (physical, $70\\;h^{-1}{\\rm \\,kpc}$ co-moving) from the observed absorber--host-galaxy impact parameter distribution.  Furthermore, they concluded that the cross-section radius only slightly increases with luminosity, $R_{\\rm phys}\\propto (L_K/L_K^*)^{0.2}$.   The only other known correlation between galaxy properties and absorber properties is that recently discovered between \\EW\\ and the galaxy morphological asymmetry \\citep{KacprzakG_05a,KacprzakG_06a}. \\citeauthor{KacprzakG_06a} searched for correlations between gas properties, host galaxy impact parameters, inclinations (\\&\\ position angles) and morphological parameters.  Among those parameters, they only found a 3.2-$\\sigma$ correlation between \\EW\\ and the asymmetry in the host galaxy morphology as measured from the residuals of 2D light profile fits.  From studies of individual absorbers, some recent work that compared the absorption kinematics with the host galaxy rotation curve favour the idea that strong \\MgII\\ clouds arise in galactic outflows \\citep{BondN_01a,EllisonS_03a}.  On the other hand, \\citet{SteidelC_02a} showed that, in 4 out of 5 \\MgII\\ absorbers (selected to be edge-on and aligned towards to QSO), the absorption kinematics are consistent with rotation being dominant for the absorbing-gas kinematics. However, a simple extension of the rotation curve fails.  Many authors have tried to put QALs in the context of theoretical models since \\citet{BahcallJ_69a} who proposed that the metal absorption lines are produced by gaseous haloes of intervening galaxies with a large cross-section (up to 100~kpc).  Later, \\citet{YorkD_86a} argued that QALs arise mainly in the haloes of gas rich Magellanic-type dwarfs.  \\citet{MoH_96a} followed the ideas of \\citet{BahcallJ_69a} and produced a detailed model in which \\MgII\\ absorbers are signatures of in-falling (photo-ionized) cold gas embedded in a $T\\simeq10^6$~K gas halo. In this model, the in-falling cold gas should be virialized. \\citet{MallerA_04a} reached similar conclusions.  Despite these numerous past results, a fundamental question remains: What is the physical nature of strong \\MgII\\ absorbers?  Here we constrain one important physical property of \\MgII\\ absorbers -- namely the halo-mass of the host galaxies -- {\\it statistically} by studying the clustering of galaxies around the absorbers.  Clustering studies of metal-line systems, such as the absorber--absorber auto-correlation in velocity space, have been used for some time \\citep[e.g.][]{SargentW_88a,SteidelC_92a,ChurchillC_97a,CharltonJ_98a}. This technique measures the line-of-sight clustering and therefore suffers from strong peculiar velocity effects.  Numerical models are then required to infer the physical properties (halo mass, halo sizes, etc.) of the host galaxies \\citep[see][for a recent example]{ScannapiecoE_06a}.  In order to avoid these limitations, we choose to cross-correlate \\MgII\\ absorbers with `field' galaxies tracing the two-dimensional large scale structure as in \\citet{BoucheN_04c}, \\citet*{BoucheN_04a}, \\cite{BoucheN_05b}, \\citet{CookeJ_06a} and \\citet{Ryan-WeberE_06a}. Specifically, we measure the amplitude {\\it ratio} of the absorber--galaxy cross-correlation to the galaxy--galaxy auto-correlation. In hierarchical galaxy formation scenarios this is a direct measure of the ratio between the bias of the absorber host-galaxies and of the field galaxies used. From this bias ratio, the mass of the absorber host-galaxies can be inferred.  In this paper, we extend our DR1 results \\citep{BoucheN_04a} by cross-correlating $\\nabstot$ \\MgII\\ absorbers with $\\ngalsim$ Luminous Red Galaxies (LRGs), both selected from the Sloan Digital Sky Survey (SDSS) Data Release 3 \\citep[DR3;][]{AbazajianK_05a}, over $\\sim\\!800$\\,square degrees. Thanks to the large SDSS survey, this is a leap forward from other clustering studies such as \\citet{WilligerG_02a} and \\citet{HainesC_04a} which cover a few square degrees.   This paper is organized as follows. In Section~\\ref{section:sample}, we summarise the selection of our \\MgII\\ absorbers and LRGs.  In Section~\\ref{section:method}, we describe our method to measure the halo-mass.  The results are presented in Section~\\ref{section:results}.  We test these results against numerous past results on \\MgII\\ absorbers in Section~\\ref{section:discussion} and discuss a physical interpretation of our results in Section~\\ref{section:interpretation}. Our main results are summarised in Section~\\ref{section:summary}.  For those familiar with absorber--galaxy clustering analyses, a quick read of this paper comprises Fig.~\\ref{fig:xcorr:050505:2halo} and Fig.~\\ref{fig:xcorr:EW}, followed by Section~\\ref{section:interpretation}.  A critical reader should focus on the several consistency tests we performed (Figs.~\\ref{fig:xcorr:test:z_off} \\& \\ref{fig:xcorr:test:Wz}) and on the discussion in Section~\\ref{section:discussion}.   We adopt $\\Omega_{\\rm M}\\!=\\!0.3$, $\\Omega_\\Lambda\\!=\\!0.7$ and $H_0\\!=\\!100 h\\,\\kms{\\rm \\,Mpc}^{-1}$ throughout. Thus, at $z\\!=\\!0.6$, $1\\arcsec$ corresponds to $7.44h^{-1}{\\rm \\,kpc}$ and $\\delta z\\!=\\!0.1$ corresponds to $216 h^{-1}{\\rm \\,Mpc}$, both in co-moving coordinates. ",
        "conclusions": "In this Section, we study the incidence probability of absorbers in more detail.  Specifically, we use two arguments to show that a $M_{\\rm h}$--\\EW\\ anti-correlation is already required by past results, i.e.  regardless of our cross-correlation results.  In Section \\ref{section:discussion:cross}, we show that the cross-section--luminosity relationship combined with the cross-section--equivalent width relationship require an anti-correlation between luminosity and equivalent width.  In Section \\ref{section:discussion:lum}, we show that the cross-section--luminosity relationship, the observed incident probability (\\dNdzdW) and any plausible host-galaxy luminosity function (LF) together require an anti-correlation between luminosity and equivalent width. We also show that the slope of the expected anti-correlation is completely consistent with our new result in Fig.~\\ref{fig:xcorr:EW}.  The starting point for the discussions below is the incidence equation, equation~(\\ref{eq:dNdz}), but in its integral form: \\begin{eqnarray} &&\\int_{W_1}^{W_2}\\frac{\\mathrm d^2 N}{\\mathrm d W_{\\rm r} \\mathrm d z}(W_{\\rm r})\\mathrm d W_{\\rm r} \\nonumber\\\\ &=& \\int_{M_1}^{M_2}\\mathrm d\\log M\\, \\frac{\\mathrm d^2 N }{\\mathrm d \\log M \\mathrm d z}(M) \\nonumber\\\\ &=& \\int_{M_1}^{M_2}\\mathrm d\\log M\\, \\frac{\\mathrm d^2 N }{\\mathrm d \\log M \\mathrm d V}(M) \\,\\sigma_{\\rm co}(M)\\, \\frac{\\mathrm d r}{\\mathrm d z}\\,,\\label{eq:dNdz:integral} \\end{eqnarray} where ${\\mathrm d^2 N }/{\\mathrm d \\log\\!M /\\mathrm d V}$ is the number of haloes (or galaxies) per unit volume per unit mass (or luminosity).  In addition, a covering factor $C_{\\rm f}$ should be included in the right-hand side of this equation since observations imply that it may be less than unity: \\citet{TrippT_05a} studied close galaxy--QSO pairs and found that $\\sim$50\\,per cent of galaxies do not produce any \\MgII\\ absorption down to 0.3\\,\\AA, indicating that the mean covering factor for our sample is $C_{\\rm f}\\sim 0.5$. We note that $C_{\\rm f}$ is likely to depend on the size of the host galaxy and therefore its mass, a complication that we ignore here.  For the past two decades, the main problem in solving equation (\\ref{eq:dNdz:integral}) for the cross-section $\\sigma_{\\rm co}$ was that one had to perform the integrals of the left-hand side over a range of \\EW\\ and the right-hand side over a range of mass $M$ (or luminosity $L$), ignoring any possible -- and probably quite strong -- dependence of \\EW\\ on $L$ or $M$.  We will show in the next two sub-sections that the $M_{\\rm h}$--\\EW\\ anti-correlation is required (independently of our cross-correlation results).  We first show (in Section 5.1) that a $M_{\\rm h}$--\\EW\\ relationship is required from the following simple arguments. Observationally, for a given bin in \\EW, there is a maximum impact parameter, $R_{\\rm phys}$, that defines the cross-section, $\\sigma$. The radius of the cross-section depends on the luminosity of the host galaxy, $R_{\\rm phys}\\propto L^{\\beta}$, i.e.~the cross-section is a function of mass or luminosity.  Since there is a relationship between \\EW\\ and $R_{\\rm phys}$ and between $R_{\\rm phys}$ and $L$, there is a relationship between \\EW\\ and \\Mh\\ (or $L$). Section~5.1 shows that the two physical parameters should be anti-correlated. Section~5.2 arrives at the same conclusion based on different arguments which make one assumption, namely that the relationship between $\\log \\Mh$ and \\EW\\ is linear.  \\subsection{The cross-section distributions require a $\\bmath{\\Mh}$--$\\bmath{\\EW}$ anti-correlation} \\label{section:discussion:cross}  \\begin{figure} \\centerline{ \\includegraphics[width=80mm]{fig9.eps} } \\caption{The \\MgII\\ equivalent width, \\EW, as a function of impact parameter, $\\rho$, for the host-galaxies of absorbers with $\\EW\\ga0.3$\\,\\AA.  The solid squares show the impact parameters for host-galaxies selected from a sample of our SDSS \\MgII\\ absorbers closer than 100\\hkpc\\ (co-moving) to the QSO lines-of-sight and with $|z_{\\rm phot}-z_{\\rm abs}|<0.05$. The open squares, open triangles and solid triangle show the samples of \\citet{SteidelC_95b}, \\citet{KacprzakG_06a} and \\citet{BergeronJ_91a} respectively. All impact parameters are shown in our cosmology.  The solid line shows the fiducial model for the maximum impact parameter $R_{\\rm phys}$ used in Section~\\ref{section:discussion:cross} (see text) to represent the cross-section--equivalent width relationship, $\\sigma_{\\rm co}(\\EW)$.} \\label{fig:impact-EW} \\end{figure}  Here, we investigate the existence of a luminosity (or mass)--\\EW\\ relationship and whether the two parameters should be correlated or anti-correlated given the observed distribution of \\MgII\\ absorbers in the cross-section--equivalent width plane and the dependence of the cross-section on luminosity \\citep{SteidelC_95b}.  The radius of the cross-section, $R_{\\rm phys}$, decreases steeply with increasing \\EW: \\citet{LanzettaK_90a} and \\citet{SteidelC_95b} (their figure~3) showed that the host-galaxies of absorbers with $\\EW>2$\\,\\AA\\ do not exist at large impact parameters. Figure~\\ref{fig:impact-EW} reproduces these results.  The data points show the impact parameters $\\rho$ (in physical units) of the \\MgII\\ host galaxies from the samples of \\citet{BergeronJ_91a} (filled triangle), \\citet{SteidelC_95b} (open squares) and \\citet{KacprzakG_06a} (open triangles) for absorbers with $\\EW>0.3$\\,\\AA.  The solid squares represent the 44 host-galaxies that we have identified in the SDSS/DR3 imaging data with $|z_{\\rm abs}-z_{\\rm phot}|<0.05$ and which were within 100~kpc (co-moving) of the QSO line-of-sight. This sample is truncated at an impact parameter $\\rho\\la 20$\\,\\hkpc\\ (physical; $\\sim$3\\arcsec) due to the large point-spread function of SDSS.  It is clear that the data points do not fill the upper right side of the plot. More precisely, at increasing equivalent width, the maximum impact parameter -- which defines the radius $R_{\\rm phys}$ of the cross-section -- decreases. This dependence can be roughly parametrized as $W_{\\rm r}(R_{\\rm phys}) \\simeq -3 \\log [R_{\\rm phys}]+K$ as represented by the solid line in Fig.~\\ref{fig:impact-EW}.  In addition, a relationship between between $R_{\\rm phys}$ and luminosity $L$ also exists: \\citet{SteidelC_95b} found that $R_{\\rm phys}$ slightly increases with luminosity: $R_{\\rm phys}(L)\\propto (L_K/L_K^*)^{\\beta}$, with $\\beta\\simeq0.2$.  Thus, the simple observational results that there are relationships between $W_{\\rm r}$ \\& $R_{\\rm phys}$ and between $R_{\\rm phys}$ \\& $L$ require a $L$--\\EW\\ relationship.  We emphasize that this $L$--\\EW\\ relationship does not mean that there is a direct correspondence between mass and equivalent width on a galaxy-by-galaxy basis. This relationship simply means that {\\it on average} a sample with a well-defined \\EW\\ (e.g.~a bin of \\EW) will have a well-defined and predictable mean mass (or luminosity). More specifically, for a sample of absorbers selected in a bin of \\EW, the host galaxy will have an impact parameter less than $R_{\\rm phys}$, and a mean mass \\Mh.  One should not interpret the distribution in \\EW\\ as being entirely due to the distribution in halo mass (or in galaxy luminosity).  Finally, if one combines the observed cross-section--equivalent width relation (Fig.~\\ref{fig:impact-EW}) with this cross-section--luminosity relation one finds that $W_{\\rm r}^{\\rm phys}(\\log L_K)\\propto -3\\beta\\log L_K\\simeq -0.6 \\log L_K$, i.e.~$\\log L_K \\propto \\frac{1}{-3 \\beta}W_{\\rm r} \\simeq -1.6 W_{\\rm r}$.  In other words, the cross-section as a function of luminosity and as a function of equivalent width {\\it together require} that $W_{\\rm r}$ (therefore $\\Delta v$) and luminosity (or mass) be anti-correlated.   \\subsection{The incidence probability requires a $\\bmath{\\Mh}$--$\\bmath{\\EW}$ anti-correlation} \\label{section:discussion:lum}  \\begin{figure} \\centerline{ \\includegraphics[width=7cm]{fig10.eps} } \\caption{The symbols represent the absorber $K$-band luminosity function (LF) of \\citet{SteidelC_95b}. The error bars are uniform and represent the average uncertainties.  The solid line represents a Schechter function with $\\alpha=-1$, as observed for the field $K$-band LF \\citep{KochanekC_01a}.  The dashed line represents a Schechter function with $\\alpha=-0.5$.  Both lines are reasonable descriptions of the $K$-band LF.  Both curves were scaled to match the data points.  }\\label{fig:LF} \\end{figure}  Since we have established the existence of a \\Mh--\\EW\\ relationship, the right-hand side of equation (\\ref{eq:dNdz:integral}) can be rewritten as \\begin{eqnarray} && \\int_{W_1}^{W_2}\\mathrm d W_{\\rm r} \\frac{\\mathrm d\\log M}{\\mathrm d W_{\\rm r}}\\, \\frac{\\mathrm d^2 N }{\\mathrm d \\log M \\mathrm d V} \\,\\sigma_{\\rm co}\\,\\frac{\\mathrm d r}{\\mathrm d z}\\,.\\nonumber \\end{eqnarray} By differentiation of both sides of equation (\\ref{eq:dNdz:integral}), it follows that~\\footnote{If $F(x)\\equiv\\int^x f(a) \\mathrm d a$ and $G(x)\\equiv\\int^x g(a) \\mathrm d a$. If $F(x)=G(x)$ for all $x$ then $F=G$ and their derivatives are also equal, $f(a)=g(a)$.} \\begin{eqnarray} \\frac{\\mathrm d^2 N}{\\mathrm d W_{\\rm r} \\mathrm d z} &=&  \\frac{\\mathrm d\\log M}{\\mathrm d W_{\\rm r}}\\, \\frac{\\mathrm d^2 N }{\\mathrm d \\log M \\mathrm d V} \\,\\sigma_{\\rm co}\\,\\frac{\\mathrm d r}{\\mathrm d z}\\,,\\label{eq:dNdz:diff} \\end{eqnarray} where $\\frac{\\mathrm d\\log M}{\\mathrm d W_{\\rm r}}$ is the Jacobian of the transformation between $W_{\\rm r}$ and $M$. We will assume that the relationship between $W_{\\rm r}$ and $M$ is linear, i.e.~$\\frac{\\mathrm d\\log M}{\\mathrm d W_{\\rm r}}$ is a constant.   An independent way to show the existence of a $L$--$W_{\\rm r}$ anti-correlation makes use of equation~(\\ref{eq:dNdz:diff}) in terms of luminosity: \\begin{eqnarray} \\frac{\\mathrm d^2 N}{\\mathrm d z\\mathrm d W_{\\rm r}} &=& \\frac{\\mathrm d\\log L}{\\mathrm d W_{\\rm r}}\\, \\frac{\\mathrm d^2 N}{\\mathrm d V\\mathrm d \\log L}\\;\\sigma_{\\rm co}(L)\\; \\frac{\\rm d r}{\\rm d z} \\label{eq:dNdzdlogL} \\,, \\end{eqnarray} where ${\\mathrm d^2 N}/{\\mathrm d V/ \\mathrm d \\log\\!L}$ is simply the LF in logarithmic form, $\\phi(L)L\\ln(10)$.  The differential incidence probability, \\dNdzdW, is now well constrained observationally thanks to the large SDSS database \\citep[e.g.][]{NestorD_05a,ProchterG_05a}.  For luminosities below $L_*$, $\\phi(L)L\\propto L^{1+\\alpha}$, where $\\alpha$ is the faint-end power-law slope of the LF. Equation~(\\ref{eq:dNdzdlogL}) then implies that \\begin{eqnarray} \\frac{\\mathrm d^2 N}{\\mathrm d z\\mathrm d \\log L} &\\propto&L^{(1+\\alpha+2\\beta)} \\,\\label{eq:dNdzdlogL:approx} \\end{eqnarray} using $\\sigma_{\\rm co}(L)\\propto L^{2\\beta}$ from \\citet{SteidelC_95b}. In addition, \\citet{NestorD_05a} showed that \\dNdzdW\\ is an exponential, i.e. \\begin{eqnarray} \\frac{\\mathrm d^2 N}{\\mathrm d z\\mathrm d W_{\\rm r}}  &\\propto &  \\exp(-W_{\\rm r}/W_*)\\,,\\label{eq:dNdz:approx} \\end{eqnarray} where $W_*$ is the exponential scale length.  Quantitatively, from equations~(\\ref{eq:dNdzdlogL:approx}) and (\\ref{eq:dNdz:approx}), the expected slope of the $L$--$W_{\\rm r}$ anti-correlation is \\begin{eqnarray} \\frac{\\rm d\\log L}{\\rm d W_{\\rm r}}& \\simeq& \\frac{-\\log_{10}(e)/W^*}{1+\\alpha+2\\beta},\\\\ &\\simeq&-1.8\\; \\hbox{for $\\alpha=-1.0$, or}\\nn\\\\ &\\simeq&-0.8\\; \\hbox{for $\\alpha=-0.5$}\\nn, \\end{eqnarray} since \\citet{NestorD_05a} showed that $W^*\\sim0.6$ in the redshift range $\\langle z\\rangle\\sim 0.5$ of our survey, and \\citet{SteidelC_95b} showed that $\\beta=0.2$ .  Therefore, equations~(\\ref{eq:dNdzdlogL:approx}) and (\\ref{eq:dNdz:approx}) imply that $W_{\\rm r}$ and $\\log L$ (or $\\log M$, using a constant $M/L$) must be anti-correlated as long as ${1+\\alpha+2\\beta}>0$, i.e.~$\\alpha>-1.4$ for $\\beta=0.2$. Most plausible LFs easily satisfy this condition.  The $B$-band LF of \\MgII\\ absorbers was constrained by \\citet{SteidelC_95a} and is known to be bell-shaped, with $\\alpha$ shallower than $-1$ \\citep[figure~4 of][]{SteidelC_95b}.  The $K$-band LF of \\MgII\\ is likely to be closer to a mass function. Figure~\\ref{fig:LF} reproduces the absorber $K$-band luminosity function of \\citet{SteidelC_95b}. The error bars are uniform and represent the average uncertainties. The solid line represents a Schechter function with $\\alpha=-1$, as observed for the $z=0$ $K$-band field-galaxy LF \\citep{KochanekC_01a,BellE_03a}. \\citet{SteidelC_95b} argued that the host-galaxy $K$-band luminosity function is similar to the $K$-band field-galaxy LF. The dashed line represents a Schechter function with $\\alpha=-0.5$ and is an equally good match to the data.  One can go a step further and `predict' the \\Mh--\\EW\\ relationship using the mass-to-light ratio.  Since the faint-end slope of the mass function of haloes \\citep[e.g.][]{MoH_02a} is steeper than that of galaxies, the halo mass-to-light ratio, $\\Mh/L$, decreases with increasing $L$ \\citep[see][and references therein]{ShankarF_06a}. Parametrizing $\\Mh/L$ as $L^{q}$ (with $q\\la0$), one finds that the expected slope of the $\\Mh$--\\EW\\ relationship is: \\begin{eqnarray} \\frac{\\rm d\\log M}{\\rm d W_{\\rm r}}&=& (1+q) \\frac{\\rm d\\log L}{\\rm d W_{\\rm r}}\\nn\\\\ &\\simeq&(1+q)\\frac{-\\log_{10}(e)/W^*}{1+\\alpha+2\\beta}\\label{eq:predict:mass-EW} \\end{eqnarray} Fig.~\\ref{fig:test:lum}(left) plots the expected relationship between $\\Mh$ and \\EW\\ for $q=0, -0.5$ \\& $-0.75$ and $\\alpha=-1.0$ \\& $-0.5$, with all curves normalized at $\\log \\Mh(\\msun)=12$.  The data points are as in Fig.~\\ref{fig:xcorr:EW}(right) and agree well with the expected relationships (lines).  \\begin{figure*} \\centerline{ \\includegraphics[width=80mm]{fig11a.eps} \\includegraphics[width=80mm]{fig11b.eps} } \\caption{{\\it Left}: The black squares show the halo-mass ($\\Mh$) as a function of equivalent width (\\EW) as in Fig.~\\ref{fig:xcorr:EW}(right). The grey squares show a different binning.  The various lines represent the expected \\Mh--\\EW\\ relationship, normalized at $\\log \\Mh(\\msun)=12$, for various LFs ($\\alpha=-1.0, -0.5$) and mass-to-light ratios ($q=0, -0.5, -0.75$) [see text and equation~(\\ref{eq:predict:mass-EW})].  The results from our cross-correlation analysis follow the expected anti-correlation.  The error bars represent the 1-$\\sigma$ statistical uncertainties and the size of the equivalent width bin used.  {\\it Right}: \\EW\\ as a function of impact parameter, $\\rho$, for the \\MgII\\ host-galaxies.  The data points are as in Fig.~\\ref{fig:impact-EW}.  The lines are as in the left panel and represent the radii of the cross-section $R_{\\rm phys}$ (i.e. maximum impact parameter allowed).  The numbers along the dot-dash line indicate the logarithm of the halo mass corresponding to the \\EW\\ from the left panel.  } \\label{fig:test:lum} \\end{figure*}  Our conclusion from this exercise is that any plausible LF, combined with \\dNdzdW\\ \\citep[e.g][]{NestorD_05a} and with the observed $R_{\\rm phys}(L_K)$ relation \\citep{SteidelC_95b}, together {\\it require} a $L$--\\EW\\ (or \\Mh--\\EW) anti-correlation.  Since a $\\Mh$--\\EW\\ anti-correlation is required from past results, we can perform a simple consistency check as follows. We use the series of relationships shown in Fig.~\\ref{fig:test:lum}(left) and the halo-mass function, $n(\\Mh)\\equiv{\\mathrm d N}/{\\mathrm d V}/{\\mathrm d \\log M}$, from \\citet{MoH_02a} in equation~(\\ref{eq:dNdz:diff}) to predict the radius $R_{\\rm phys}$ of the cross-section as a function of equivalent width, $R_{\\rm phys}(\\EW)$, that is required by the incidence distribution \\dNdzdW\\ of \\citet{NestorD_05a}.  The result is shown in Fig.~\\ref{fig:test:lum}(right) for each of the relationships on the left plot. The data points are as in Fig.~\\ref{fig:impact-EW}.   Thus, we have shown with two different arguments (in Sections 5.1 \\& 5.2) that a positive $\\Mh$--\\EW\\ correlation is not consistent with the impact parameter distribution ($\\rho$--\\EW) of current \\MgII\\ samples. This conclusion is completely independent of our clustering results presented in Section \\ref{section:results}. One can turn these arguments around and state that the anti-correlation between \\EW\\ and maximum impact parameter (Fig.~\\ref{fig:impact-EW}) is a natural outcome of the \\Mh--\\EW\\ anti-correlation in Fig.~\\ref{fig:xcorr:EW}.   \\label{section:summary}  From the SDSS/DR3 we selected \\nabstot\\ \\MgII\\ absorbers with $\\EW\\!\\geq\\!0.3$\\,\\AA\\ in the redshift range $0.4\\!\\leq\\!\\zabs\\!\\leq\\!0.8$ and $\\ngalsim$ LRGs in the absorber fields selected with photometric redshifts. From the ratio of the \\MgII--LRG cross-correlation to the LRG--LRG auto-correlation, we constrained the halo-mass of \\MgII\\ absorbers in a statistical manner, i.e.~without directly identifying them.  To summarise our main results, we have shown that: \\begin{enumerate} \\item the ratio of the cross- and auto-correlation, $w_{\\rm ag}/w_{\\rm gg}$, is $a=\\Arel$; \\item the corresponding bias-ratio between strong \\MgII\\ absorbers and LRGs is $b_{\\rm \\MgII}/b_{\\rm LRG}=\\Arelcorr$; \\item this bias-ratio implies that the absorber host-galaxies have a halo-mass of $\\langle \\log \\Mh (\\msun)\\rangle =\\;$\\LRGmassrangesys. The 1-$\\sigma$ uncertainty includes the statistical errors in $a$ and possible systematic errors in the LRG halo-mass; \\item the \\MgII\\ equivalent width, \\EW, is significantly anti-correlated with the absorber halo-mass, \\Mh. \\end{enumerate} The main consequence of point (iv) is that the \\MgII\\ clouds are not virialized in the gaseous haloes of the host-galaxy, in the sense that $\\Delta v$, the velocity spread of the individual clouds (ranging from 50 to $\\sim~400$~\\kms), does not positively correlate with halo mass. In other words, had the clouds been in virial equilibrium with the halo dark matter distribution, $\\Delta v$ (as measured by \\EW) would correlate with the mass of the halo. However, we find the opposite: lower-mass haloes harbour the host-galaxies producing the largest number of individual \\MgII\\ clouds spread over the largest velocities.  Since the \\Mh--\\EW\\ anti-correlation may initially seem surprising, we performed several consistency tests on the method itself and tested the \\Mh--\\EW\\ anti-correlation against numerous past results in the literature. For the galaxy clustering method we have shown that: \\begin{itemize} \\item The cross-correlation is solely due to the \\MgII\\ absorbers; \\item The cross-correlation amplitude decreases when the width of the LRG redshift distribution is increased, as expected; \\item The $w_{\\rm ag}/w_{\\rm gg}$ ratio is independent of the LRG redshift distribution, as expected if one uses the same galaxies to calculate both $w_{\\rm ag}$ and $w_{\\rm gg}$. Therefore, the relative amplitude is free of possible systematic errors caused by contaminants (stars or interloping galaxies). \\end{itemize}  In checking past results in the literature for consistency with the observed $\\Mh$--\\EW\\ anti-correlation we found that:  \\begin{itemize} \\item When one combines the observed luminosity- and \\EW-dependence of the cross-section, one finds that luminosity (or mass) and equivalent width must be anti-correlated; \\item When one combines the observed luminosity-dependence of the cross-section with the observed incidence probability of strong \\MgII\\ absorbers, \\dNdzdW, one finds that luminosity (or mass) and equivalent width must be anti-correlated for all plausible luminosity functions with faint-end power-law indexes $\\alpha > -1.4$. \\end{itemize} In other words, a positive correlation between \\EW\\ and halo-mass (or virial velocity) -- expected if the clouds were virialized -- is inconsistent with past results. This is independent of the \\Mh--\\EW\\ anti-correlation from the clustering analysis. When one combines the observed \\Mh--\\EW\\ relationship with the observed \\dNdzdW\\ one predicts a cross-section radius which agrees very well with the distribution of absorber impact parameters observed in previous and current samples.  We interpret our results as a strong indication that a large proportion of strong \\MgII\\ absorbers, particularly those with $\\EW\\ga 2$\\,\\AA, arise in galactic outflows. A conclusion similar was reached by \\citet{SongailaA_06a} for $z\\sim2$ \\CIV\\ absorbers. Supernovae-driven winds drive hot gas out of the galaxy plane, forming cavities if numerous supernova explosions occur in a small region on a similar time-scale. The hot gas breaks out of the disk, forming a hot gaseous corona (traced by high-ions such as \\CIV\\ and \\OVI) and may entrain cold, dusty material.  Part of this gas cools, forming pressure-confined clouds traced by \\MgII\\ which may be the analogs of the HVCs.   This interpretation allows us to make several predictions.  Firstly, it provides a natural explanation as to why absorbers with $\\EW>2$\\,\\AA\\ are not observed at large impact parameters. These systems should be physically closer to the host galaxy and also closer in time to more recent and more active star formation. Thus, they ought to have star formation rates typical of star-bursts. The resonance line O[{\\sc ii}]$\\lambda$3727 may also be detectable in stacked SDSS spectra since the host-galaxies of $\\EW\\ga 2$\\,\\AA\\ absorbers will be very close to the line-of-sight. The detectability of this emission line should be a strong function of \\EW\\ but will also obviously depend on the host-galaxy metallicity. Secondly, the \\NaI~D doublet ought to be detectable in the same stacked spectra. Thirdly, the colours of the stacked light distribution should be a strong function of \\EW. The preliminary results of S.~Zibetti (private communication) confirm this last prediction."
    },
    "0606/astro-ph0606602_arXiv.txt": {
        "abstract": "X-ray cluster measurements interpreted with a universal baryon/gas mass fraction can theoretically serve as a cosmological distance probe.  We examine issues of cosmological sensitivity for current (e.g.\\ Chandra X-ray Observatory, XMM-Newton) and next generation (e.g.\\ Con-X, XEUS) observations, along with systematic uncertainties and biases.  To give competitive next generation constraints on dark energy, we find that systematics will need to be controlled to better than 1\\% and any evolution in $f_{\\rm gas}$ (and other cluster gas properties) must be calibrated so the residual uncertainty is weaker than $(1+z)^{0.03}$. ",
        "introduction": "\\label{sec:intro}  An increasing number of cosmological methods have been suggested to explore the cosmological model and the accelerating universe.  This diversity offers hope in understanding the nature of our universe, if the probes are robust and clear in interpretation.  Astrophysical systematic uncertainties of a technique and its observations limit the cosmological leverage, despite statistical precision. Here we examine the level of control of systematics required for accurate cosmology estimation from the X-ray cluster distance method, sometimes called the baryon or gas mass fraction method, that assumes a universal baryon/gas mass fraction.  X-ray cluster observations of fewer than 50 clusters by the Chandra X-ray Observatory \\cite{chandra} and XMM-Newton \\cite{xmm} have already been claimed to yield cosmological constraints \\cite{ettori,allen04,weller,rapetti}.  For the future, there are prospects of more and deeper clusters, for example from Constellation-X \\cite{conx} and XEUS \\cite{xeus}.  We investigate the possible constraints, their cosmological degeneracies, and complementarity with other probes, giving special attention to the effect of hypothetical systematic floors in the achievable accuracy.  While the X-ray cluster baryon technique requires understanding the influence of cosmology and astrophysics on the ingredients of cosmic geometry, mass distributions, and hydrodynamical gas properties, it also has rich observational data.  A number of recent papers have raised issues concerning the central assumption of universality of the cluster gas mass fraction, from observational, simulation, and theory points of view, e.g.\\ \\cite{vikhlinin,ettorifgas,hallman,loeb,nagai,ferramacho,gastaldello}. Given this uncertainty (which this paper does not try to resolve), it is useful to investigate how an overall uncertainty in the technique, arising from whatever source, affects the cosmological conclusions. We emphasize that this paper presents purely an investigation of cosmological sensitivity to the innate parameter degeneracies and end-of-pipeline level of systematic uncertainties; we do not make claims for the error contribution from specific measurements or what level of end systematics is achievable.  In \\S\\ref{sec:method} we examine the cosmological sensitivity of the constant gas mass fraction technique based on parameter degeneracies and survey depth. Biases in the cosmological results arising from possible systematics are calculated in \\S\\ref{sec:bias}.  A brief review of the gas mass fraction technique, noting assumptions and possible areas for systematic uncertainties, is presented in the Appendix, along with a comment on bias from non-Gaussian errors in the translation to a distance measure. ",
        "conclusions": "\\label{sec:concl}  We have considered the possibility of current and next generation observations of X-ray clusters as cosmological distance probes through the baryon/gas mass fraction method.  In current and future usage, assuming a cosmological constant universe, this technique can provide tight constraints on the matter density.  However, even with the next generation Con-X and XEUS observatories, we find that it will be a considerable challenge for such X-ray cluster measurements to provide a precision dark energy probe.  As with any probe, we cannot ignore possible systematic uncertainties.  While many have been identified (see the Appendix for a brief discussion), the level of systematics control possible is currently not well determined.  Understanding of cosmic geometry, the nonlinear mass distribution, and gas hydrodynamics are simultaneously required.  The amplitude factor $A$ must be fit for, and redshift dependent uncertainties from evolution or population drift are especially degenerate with the dark energy equation of state time variation. Self-calibration only works if the functional form of the evolution is known and if the redshift scaling can be calculated to within $(1+z)^{0.03}$.  X-ray cluster studies from simulations, theory, and observations over the next decade will address what is a realistic level of systematic uncertainties, both as a floor on precision at a given redshift and as a redshift dependent bias.  We emphasize that this paper does not claim what level should be adopted but provides a general investigation of what level will impact cosmological parameter determination.  Even if the X-ray gas mass fraction technique turns out not to be significant as a dark energy probe, the X-ray data is likely to greatly extend and deepen our knowledge of cluster properties and evolution.   \\ack  I thank Jochen Weller and especially the referee for helpful comments. This work has been supported in part by the Director, Office of Science, Department of Energy under grant DE-AC02-05CH11231.  \\appendix"
    },
    "0606/astro-ph0606420_arXiv.txt": {
        "abstract": "We analyze the Millennium run semi-analytic galaxy catalog to explore quantitatively the gravitational pancaking effect on the orientation of galaxy velocity field.  We first calculate the probability density distribution of the cosine of the angle between the velocity of a field galaxy and the direction normal to a local pancake plane which is determined using two nearest neighbor field galaxies. A clear signal of alignment is detected for the case that the pancake scale is in the range of $5-8h^{-1}$ Mpc. The tendency of the velocity-pancake alignment is found to still exist when the pancakes are determined using three neighbor galaxies, indicating that it has a spatial coherence. The degree of the velocity-pancake alignment is shown to increase with the velocity magnitude and the local density, while it decreases with the separation distance from  the galaxy to the pancake and disappears when the pancake has a filamentary shape. A final conclusion is that our work may provide another clue to understanding the large-scale structure in the universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606616_arXiv.txt": {
        "abstract": "We present results of large-scale numerical simulations of the evolution of neutrino and antineutrino flavors in the region above the late-time post-supernova-explosion proto-neutron star. Our calculations are the first to allow explicit flavor evolution histories on different neutrino trajectories and to self-consistently couple flavor development on these trajectories through forward scattering-induced quantum coupling. Employing the atmospheric-scale neutrino mass-squared difference ($|\\delta m^2|\\simeq 3\\times 10^{-3}\\,\\mathrm{eV}^2$) and values of $\\theta_{13}$ allowed by current bounds, we find transformation of neutrino and antineutrino flavors over broad ranges of energy and luminosity in roughly the ``bi-polar'' collective mode. We find that this large-scale flavor conversion, largely driven by the flavor off-diagonal neutrino-neutrino forward scattering potential, sets in much closer to the proto-neutron star than simple estimates based on flavor-diagonal potentials and Mikheyev-Smirnov-Wolfenstein evolution would indicate. In turn, this suggests that models of \\textit{r}-process nucleosynthesis sited in the neutrino-driven wind could be affected substantially by active-active neutrino flavor mixing, even with the small measured neutrino mass-squared differences. ",
        "introduction": "}  In this paper we employ large-scale computational techniques to tackle the vexing problem of neutrino flavor transformation in the core collapse supernova environment. Neutrinos and the weak interaction play pivotal roles in the core collapse/explosion phenomenon. The Chandrasekhar mass core of iron-peak material left at the end of the hydrostatic evolution of a massive star goes dynamically unstable and collapses in $\\sim 1$ s to a proto-neutron star configuration at nuclear density. The amount of gravitational energy promptly converted into trapped seas of neutrinos is $\\sim 1\\%$ ($\\sim 10^{52}$ erg) of the core mass. Within a few seconds after bounce $\\sim 10\\%$ ($\\sim 10^{53}$ erg) of the core mass (the gravitational binding energy) will be emitted as neutrinos.  Nearly all of this gravitational energy is converted into seas of $\\nu_e$, $\\bar\\nu_e$, $\\nu_\\mu$, $\\bar\\nu_\\mu$, $\\nu_\\tau$ and $\\bar\\nu_\\tau$ neutrinos in rough energy equipartition. Though these neutrinos diffuse with short mean free paths in the proto-neutron star, they decouple near the stellar surface where the matter density falls off steeply, the so-called neutrino sphere. Neutrinos propagate nearly coherently above this point, though neutrino-matter interactions, especially the charged current capture reactions $\\nu_e+n\\rightarrow p+e^-$ and $\\bar\\nu_e+p\\rightarrow n+e^+$, can deposit energy and set the local neutron-to-proton ratio, $n/p$.  For this reason, and because the fluxes and energy spectra may be different for $\\nu_e$, $\\bar\\nu_e$ and $\\nu_\\mu \\bar\\nu_\\mu \\nu_\\tau \\bar\\nu_\\tau$, the flavor content of the neutrino field above the proto-neutron star and its evolution in time and space can be important \\cite{Fuller:1992aa,Qian:1993dg,Qian:1994wh}. This can be true both for the supernova shock reheating epoch (where the time post core-bounce is $\\tpb\\lesssim 0.5$ s) and in the later hot bubble, neutrino-driven wind epoch ($\\tpb\\gtrsim 3$ s). In this paper we concentrate on the latter epoch.  Following the development of neutrino and antineutrino flavors in the coherent regime above the proto-neutron star surface is challenging. The potential governing the effective neutrino mass differences in this environment will have contributions from charged current neutrino-electron forward scattering and neutral-current neutrino-neutrino forward scattering. The former contribution \\cite{Wolfenstein:1977ue} is diagonal in the flavor basis, while the latter neutrino-neutrino potential has both flavor-diagonal \\cite{Fuller:1987aa} and flavor off-diagonal \\cite{Pantaleone:1992xh,Sigl:1992fn} components. The neutrino-neutrino forward scattering potential renders the neutrino flavor evolution problem nonlinear in the sense that the potential which governs neutrino flavor transformation is itself dependent on the flavor evolution histories of the neutrinos.  Furthermore, neutrinos propagating on intersecting world lines can have their flavor evolution subsequently quantum mechanically coupled by forward scattering (see Fig.~8 in Ref.~\\cite{Qian:1994wh} and the text beneath it). We sometimes will refer to this coupling as ``entanglement''. By this terminology we do not mean \\textit{quantum entanglement} of momentum states, a phenomenon which has been argued to be unimportant in the supernova environment \\cite{Friedland:2006ke}. In any case, neutrino trajectories coming off the proto-neutron star surface at different angles in general will have different flavor evolution histories which must be self-consistently calculated.  Another issue revolves around the efficacy of a mean field Schr\\\"{o}dinger-like or Boltzmann kinetic equation approach to the evolution of neutrino flavors \\cite{Qian:1994wh,Bell:2003mg,Friedland:2003eh,Friedland:2006ke}. In this paper we will assume that higher order correlations in neutrino-neutrino scattering are unimportant in the coherent regime above the proto-neutron star.  Previous attempts to model neutrino flavor evolution in the coherent regime \\cite{Fuller:1987aa,Fuller:1992aa,Qian:1993dg,Qian:1994wh,% Qian:1995ua,Pastor:2002we,Balantekin:2004ug} have made the approximation that all neutrinos evolve in flavor space the way a radially-propagating neutrino does. This we will term the ``single-angle'' approximation. Since the neutrino-neutrino forward scattering potential is intersection angle dependent, this is not always a good approximation, especially for regions close to the proto-neutron star.  However, these previous studies done with the single-angle approximation have found that it is possible to have large-scale collective behavior in neutrino flavor evolution, where all, or some significant subset of, neutrinos experience similar time/space flavor evolution histories. They also have shown that neutrino flavor transformation can differ significantly from the Mikheyev-Smirnov-Wolfenstein (MSW) \\cite{Wolfenstein:1977ue,Wolfenstein:1979ni,Mikheyev:1985aa} paradigm. Recent work \\cite{Fuller:2005ae} has shown that the expected neutrino fluxes in both the shock reheating and hot bubble epochs could provide the ``necessary'' conditions for large-scale simultaneous collective neutrino and antineutrino flavor transformation over broad ranges of neutrino energy. Whether these expected neutrino fluxes are actually ``sufficient'' to obtain these collective modes has remained an open question, to be answered with an appropriately sophisticated numerical simulation. Likewise, the range of possible collective neutrino behavior \\cite{Duan:2005cp}, be it the ``synchronized'' mode \\cite{Pastor:2001iu} or the ``bi-polar'' mode \\cite{Duan:2005cp}, may depend sensitively on the neutrino flux conditions and on the geometry.  It should be noted that many previous numerical studies employing the single-angle approximation have also used unphysically large values of neutrino mass-squared difference. This is because with straight MSW, and without taking account of the neutrino-neutrino scattering-induced flavor off-diagonal potential, it requires $|\\delta m^2|\\gtrsim 1\\,\\mathrm{eV}^2$ to have significant neutrino flavor transformation deep enough in the supernova envelope to affect shock reheating or the \\textit{r}-process \\cite{Fuller:1992aa,Qian:1993dg,Qian:1994wh,Sigl:1994hc,Mezzacappa:1999co}.  However, recent observations/experiments (see, \\textit{e.g.}, Ref.~\\cite{Fogli:2005cq} for a review) have revealed much about the fundamental flavor mixing parameters of the three known ``active'' neutrinos. (In this paper, we will ignore the effects of speculative additional ``sterile'' neutrino states.) We know the two mass-squared differences, the atmospheric scale, $\\delta m^2_{\\mathrm{atm}}\\simeq 3\\times 10^{-3}\\,\\mathrm{eV}^2$, and the solar scale, $\\delta m^2_\\odot\\simeq 8\\times 10^{-5}\\,\\mathrm{eV}^2$. We as yet do not know the neutrino mass hierarchy related to the atmospheric mixing and we do not know the absolute neutrino mass eigenvalues. Of the four mixing parameters in the unitary transformation between the flavor (weak interaction) eigenstates and the mass (energy) eigenstates, we know two of the three vacuum mixing angles, $\\theta_{12}$ and $\\theta_{23}$, and we have a firm upper limit on $\\theta_{13}$, $\\sin^22\\theta_{13}\\lesssim 0.1$. We do not know the \\textit{CP}-violating phase.  In this paper we study $2\\times2$ neutrino and antineutrino flavor transformation at the $\\delta m^2_{\\mathrm{atm}}$ scale, explicitly following the coupled flavor evolution on neutrino trajectories ranging from radially-directed to those tangential to the neutron star surface. (In other words, we perform ``multiangle'' calculations with many trajectory/angle zones.) Our goal is to study the nonlinear behavior of the neutrino field in the coherent regime and to find out if large-scale (collective mode) neutrino/antineutrino transformation can occur in the late-time supernova environment. We specialize to late time for two reasons: (1) this epoch is when there may be significant differences in flux or energy spectrum between $\\nu_e$, $\\bar\\nu_e$, and the mu and tau flavor neutrinos; and (2) this epoch may have a simpler, more compact, matter density profile near the neutron star surface. We follow Refs.~\\cite{Balantekin:1999dx,Caldwell:1999zk} and argue that $2\\times 2$ mixing is adequate because the $\\nu_\\mu$ and $\\nu_\\tau$ neutrinos are nearly maximally-mixed in vacuum ($\\theta_{23}\\simeq \\pi/4$) and these species experience nearly identical interactions everywhere in the late-time supernova environment.  In Sec.~\\ref{sec:background} we summarize the physical and geometric assumptions in our numerical simulations in what we call the ``neutrino bulb model''. In this section we also present the basic physics of neutrino flavor transformation in the practical formalism used in our numerical simulations. We also review the spin analogy for neutrino flavors, and estimate the (MSW) resonance locations in the hot bubble using both the standard MSW and synchronization mechanisms. In Sec.~\\ref{sec:multiangle} we explain some details of our numerical codes and discuss the particular numerical difficulties and potential pitfalls in multiangle simulations. We also present the main results of our multiangle simulations. The simulations show large-scale flavor transformation different from what would be predicted if the conventional MSW or synchronization mechanisms apply. In Sec.~\\ref{sec:analysis} we identify the flavor transformation in our results as being of  the bi-polar type \\cite{Duan:2005cp}, and we analyze this behavior with the help of single-angle simulations. In Sec.~\\ref{sec:conclusions} we give our conclusions. ",
        "conclusions": "}  We have carried out large-scale multiangle simulations of neutrino flavor transformation in the hot bubble employing the atmospheric neutrino mass-squared difference, $|\\delta m^2|\\simeq 3\\times 10^{-3}\\,\\mathrm{eV}^2$, and effective $2\\times 2$ vacuum mixing angle $\\theta=0.1$. The numerical results we have presented support previous conjecture on the existence of collective neutrino flavor transformation of the bi-polar type in the supernova environment \\cite{Duan:2005cp}. Our simulations also show that both neutrinos and antineutrinos can simultaneously undergo significant flavor conversion, largely driven by flavor off-diagonal potentials, at values of radius much smaller than those expected from ordinary MSW. This is along the lines of what was predicted in Ref.~\\cite{Fuller:2005ae}. We have found that this flavor transformation occurs in both the normal and inverted neutrino mass hierarchy scenarios.  For the normal mass hierarchy case, the full synchronization limit gives a lower bound on the radius where large-scale neutrino flavor transformation begins. (Ref.~\\cite{Pastor:2002we} was the first to point out that, since $\\Esync$ is much smaller than average neutrino energies, synchronized flavor transformation modes operate closer to the neutrino sphere than those driven by the MSW matter potential.) Although an analytical analysis of neutrino flavor transformation in the bi-polar mode has yet to be done, our numerical simulations suggest that, for the normal mass hierarchy, the onset of bi-polar type flavor transformation always occurs at values of radius larger than those required in the full synchronization case. Our simulations also support the prediction of large-scale neutrino flavor transformation in the inverted mass hierarchy scenario \\cite{Duan:2005cp}. (Large-scale neutrino/antineutrino flavor transformation with small mixing angles in the inverted mass hierarchy case previously was seen in the early universe context \\cite{Kostelecky:1993dm} and also in the supernova context  \\cite{Pastor:2002we}.) We have found that this may occur at values of radius even smaller than those seen in the full synchronization mode in the normal mass hierarchy scenario. We have found that single-angle simulations can be used to give a lower bound on the radius where large-scale neutrino flavor transformation occurs in  the inverted mass hierarchy scenario.   Our ``multiangle'' calculations are the first to include self-consistent flavor evolution history entanglement on intersecting neutrino world lines. Although we find that ``single-angle'' simulations in some cases can give the correct \\textit{qualitative} features of large-scale neutrino and/or antineutrino conversion in the late-time, hot bubble region, our simulations clearly show that a \\textit{quantitatively} correct treatment must include coupled flavor development on different neutrino trajectories. Furthermore, since the location where large-scale neutrino and/or antineutrino flavor transformation begins in the supernova envelope can be a crucial issue for supernova shock reheating \\cite{Fuller:1992aa}, \\textit{r}-process nucleosynthesis \\cite{Qian:1993dg,Qian:1994wh,Sigl:1994hc,Qian:1995ua,Pastor:2001iu,Balantekin:2004ug}, and the supernova neutrino signal \\cite{Schirato:2002tg,Tomas:2004gr,Yoshida:2005uy}, it is essential that simulations \\textit{be} quantitatively as accurate as possible.  The simulations we have presented focus on the late-time supernova  environment, \\textit{i.e.}, the regime after the shock has been somehow re-energized. This epoch is a leading candidate for the site of the production of some or all of the \\textit{r}-process elements and will be a major focus of future neutrino detectors/observatories should we be lucky enough to catch a galactic core collapse event. Though our simulations show that large-scale neutrino and antineutrino flavor conversion can take place during this epoch for the expected conditions of neutrino flux and entropy, we must go further than we have in this paper to produce quantitative predictions. There are three principal reasons for this: (1) we do not as yet know the matter density distribution above the proto-neutron star to sufficient accuracy at any epoch; (2) we do not know precisely the neutrino and antineutrino energy distributions and fluxes which are emergent from the proto-neutron star; and (3) the matter composition (\\textit{i.e.}, $Y_e$) can be affected by any changes in the neutrino and antineutrino spectra engendered by flavor transformation and we have not put this feedback in the calculations  presented here.  On point (2), recent work on supernova models at $\\tpb < 1$ s suggests that additional channels for neutrino scattering may weaken or dilute the effects of the charged current opacities \\cite{Keil:2002in,Raffelt:2003en}. This would tend to make the $\\nu_e$, $\\bar\\nu_e$, and $\\nu_\\mu\\bar\\nu_\\mu\\nu_\\tau\\bar\\nu_\\tau$ energy spectra more similar. Of course, if the neutrino energy spectra and fluxes are identical for all flavors, interconversion of these will have no astrophysical effect. We note, however, that reliable neutrino transport calculations at the late time we considered here do not exist, and the core's composition and neutron excess is expected to change considerably between $\\tpb\\simeq 1$ s and $\\tpb\\simeq 10$ s. Clearly, this issue is critical for gauging the astrophysical effect of neutrino flavor mixing.  It is well known that density fluctuations on short length scales and other inhomogeneities can modify coherent neutrino flavor evolution through MSW resonances \\cite{Sawyer:1990tw,Loreti:1995ae}. How these fluctuation-induced modifications could manifest themselves in quantum flavor history entanglement on intersecting neutrino trajectories is not known. This issue may be closely related to the problem of calculating neutrino transport and predicting the emergent neutrino energy spectra in general \\cite{Liebendoerfer:2002xn,Thompson:2002mw,Liebendoerfer:2003es,% Mezzacappa:2004iz,Walder:2004ym} and to the inclusion of neutrino flavor mixing in the core in particular \\cite{Sawyer:2005jk,Strack:2005ux}. Though our simulations are spherically symmetric, they do show that the density and $Y_e$ profiles near the proto-neutron star surface are important for obtaining the correct flavor evolution of the neutrino and antineutrino fields, even well above the proto-neutron star.  These uncertainties aside, our calculations indicate that large modifications of the emergent neutrino and antineutrino energy spectra are likely to occur over most of the range of expected thermodynamic and neutrino emission parameters of relevance in the late-time supernova environment. Furthermore, we have found that these modifications could set in sufficiently deep in the supernova envelope to affect $Y_e$ \\cite{Qian:1993dg} and \\textit{r}-process nucleosynthesis \\cite{Fuller:1995ih,Qian:1996db} through neutrino interactions. However, we have not included charged-current weak interaction [Eq.~\\eqref{eq:nu-interactions}] feedback in the calculations presented here.    We have fixed  $Y_e$ and $g_{\\mathrm{s}}$ in this work, essentially to simplify the computations. In future simulations we will remove these constraints and allow $Y_e$ and $g_{\\mathrm{s}}$ to be calculated consistently with feedback from neutrino capture reactions. However, we expect that the collective neutrino flavor transformation illustrated here will not be changed qualitatively with changing $Y_e$ and $g_{\\mathrm{s}}$. The bi-polar neutrino flavor transformation seen in our simulations is largely independent of the values of $Y_e$ and $g_s$. For example, we have shown that $\\nu_e$ of energy smaller (larger) than a critical energy $\\EC$ could convert to other flavors if $\\delta m^2 > 0$ ($\\delta m^2 < 0$). This critical energy $\\EC$ asymptotically approaches a limit if $L_\\nu$ is large enough, or equivalently, the electron density profile is sufficiently condensed toward the proto-neutron star. The asymptotic limit of $\\EC$ depends only on the neutrino mixing parameters and the initial energy spectra for neutrinos and antineutrinos.  Because the proto-neutron star is neutron-rich, the initial $\\nu_e$ energy spectrum may be softer than those for neutrinos in other flavors. Our simulations suggest that $\\nu_e$ and neutrinos in other flavors may swap the low-energy ($E_\\nu<\\EC$)  or high-energy ($E_\\nu>\\EC$)  parts of their spectra depending on the sign of $\\delta m_{13}^2$. Note that this stepwise swapping is independent of the details of the neutrino energy spectra. With other effects correctly accounted for and a good signal from a galactic supernova, this phenomenon may offer a unique probe of the neutrino mass hierarchy problem.  We have employed $2\\times 2$ neutrino flavor mixing in our simulations. It is possible to extend our codes to implement the neutrino mixing of all three active flavors. However, we expect that neutrino flavor transformation in the hot bubble region will not change much on inclusion of a third neutrino flavor. For one thing, $\\nu_\\mu$ and $\\nu_\\tau$ are almost equally mixed in the hot bubble because they experience the same weak interactions and $\\theta_{23}\\simeq \\pi/4$. For another, the two neutrino mass-squared differences, $\\delta m_{\\mathrm{atm}}^2$ and $\\delta m_\\odot^2$, are separated by over an order of magnitude. Taking $\\delta m^2=\\delta m_\\odot^2\\simeq 8\\times 10^{-5}\\,\\mathrm{eV}^2$ and $\\theta = \\theta_{12} \\simeq 0.6$, we estimate that the onset radius of large-scale neutrino flavor transformation in the full synchronization limit is $\\rXo(\\Esync)\\simeq 227$ km for an entropy per baryon $S=140$. This location is almost outside the range of  the simulation results presented in Sec.~\\ref{sec:results}.  In summary, though many aspects of our calculations are reasonable approximations at best (\\textit{e.g.}, $2\\times 2$ mixing, assumptions of spherical symmetry, an infinitely thin neutrino sphere, neutrino/antineutrino energy spectra of the Fermi-Dirac type, \\textit{etc.}), our computations do mark an important advance in that they self-consistently treat coupled neutrino flavor evolution on different trajectories. We cannot claim generality for our conclusions. However, our assumptions are reasonable, and our results are robust, and so there is nothing to suggest that our results represent an isolated case either. Not only do our results show that a proper treatment of coupled neutrino trajectories is important, but they also indicate that the measured neutrino mass-squared difference values and mixing angles likely imply large-scale flavor conversion of neutrinos and antineutrinos in astrophysically important regions in the post-explosion supernova environment."
    },
    "0606/astro-ph0606085_arXiv.txt": {
        "abstract": "With its ability to look at bright galactic X-ray sources with sub-millisecond time resolution, the Rossi X-ray Timing Explorer (RXTE) discovered that the X-ray emission from accreting compact stars shows quasi-periodic oscillations on the dynamical timescales of the strong field region. RXTE showed also that waveform fitting of the oscillations resulting from hot spots at the surface of rapidly rotating neutron stars constrain their masses and radii. These two breakthroughs suddenly opened up a new window on fundamental physics, by providing new insights on strong gravity and dense matter. Building upon the RXTE legacy, in the Cosmic Vision exercise, testing General Relativity in the strong field limit and constraining the equation of state of dense matter were recognized recently as key goals to be pursued in the ESA science program for the years 2015-2025. This in turn identified the need for a large (10 m$^2$ class) aperture X-ray observatory. In recognition of this need, the XEUS mission concept which has evolved into a single launch L2 formation flying mission will have a fast timing instrument in the focal plane. In this paper I will outline the unique science that will be addressed with fast X-ray timing on XEUS. ",
        "introduction": "The X-ray emission from accreting compact stars (neutron stars and black holes) has been shown to vary on (sub)-millisecond timescales, i.e. on timescales comparable to the dynamical timescales of the innermost regions of the accretion disk. This variability was discovered with RXTE \\citep{bra93} just ten years ago in the form of kilo-Hz quasi-periodic oscillations (QPOs) in Fourier power density spectra (see the review by van der Klis (2005) and Fig. \\ref{fig:sco}). There is a wide consensus that these QPOs probe the motion of matter under strong gravity with the signal originating from within a few Schwarzschild radii of the compact star. This is a region where the spacetime curvature is extreme, orders of magnitude larger than that sampled in weak field tests (e.g. Gravity-Probe B). This is a regime where the most dramatic effects of GR are expected, and where fundamental predictions of GR, such as the existence of an innermost stable circular orbit (ISCO) have yet be tested. With RXTE, X-ray timing has thus become a complementary tool to X-ray spectroscopy, polarimetry and gravitational wave measurements to probe strong field GR. It has even been claimed that studying black hole variability might be used to test alternative theories of gravity \\citep{psa04}. \\begin{figure} \\centerline{\\psfig{figure=col05-22_f1.ps,height=3.0in}} \\caption{Power density spectrum of Sco X-1 in a $\\nu F\\nu$-like representation, showing two kHz quasi-periodic oscillations (data courtesy of Michiel van der Klis). \\label{fig:sco}} \\end{figure} \\subsection{Overall QPO properties} It is is beyond the scope of this paper to review the properties of kHz QPOs, and the reader is referred to van der Klis (2005) for a recent review and {\\it X-ray Timing 2003 - Rossi and Beyond} by Kaaret, Lamb and Swank (2004, AIP Conference Proceedings, Vol. 714, Melville, NY: American Institute of Physics) on which the present paper draws heavily on. Here I will focus on their general properties, relevant to some of the points discussed below. So far, kHz QPOs have been detected from about 20 neutron star low-mass X-ray binaries (about 10\\% of known systems) \\citep{swa04}, and from a handful of binaries containing a black hole candidate.  In general two kHz QPOs are observed. The properties of the QPO (frequency, amplitude, width) vary in a complicated way with parameters such as the source luminosity. For neutron stars, the QPO frequencies can vary by at least two orders of magnitude with the amplitude decreasing with increasing frequency. There is a trend for the QPO fractional amplitude to increase with energy, reaching in some cases 20\\% at $\\sim 20$ keV. In black holes, QPOs are detected closer to the sensitivity limit of RXTE with typical amplitudes of 1\\%: they are transient, show stable frequencies and are detected in the hardest X-ray bands (typically above 10 keV)  \\citep{rem04}.  \\subsection{QPO interpretation} The true nature of the underlying signal is unknown and there is a wide range of possible models to be considered, mostly focussing on predicting QPO frequencies. Because the disk is a natural source of periodicities, from orbital to epicyclic motions and disk oscillations, it is generally agreed that QPOs originate in the cool ($\\sim 1-2$ keV) accretion disk. This implicitly assumes that the oscillator must survive the strong damping forces thought to be present in the disk and that an amplification mechanism operates to account for the fact that QPOs show larger amplitude at higher energies ($\\sim 10-20$ keV). For neutron star QPOs, the main constraint for the models so far comes the observation that the frequency separation between the twin QPOs is close to the spin frequency or half the spin frequency of the neutron star (where that is known). This strongly suggests that the spin is involved in the generation of the QPOs. For black holes, the main constraint for the models comes from the observations that QPOs, when observed in pairs, have commensurate frequencies, with ratios close to small integer ratios (e.g. 3:2) \\citep{rem04}, lending support to the idea that a resonance mechanism is involved in the make up of the QPOs.  Some models (not all of them consistent with the constraints above) relate QPO frequencies directly characteristic frequencies of test particles moving in the strongly curved spacetime (General Relativistic orbital and epicyclic frequencies \\citep{abr04}),  others link the QPO frequencies to the frequency of orbiting clumps interacting with radiation from the neutron star \\citep{mil98}. Some models involve global disk oscillations, relativistic \\citep{kat01}, or not \\citep{tit04}. Yet others associate QPOs with strong field GR effects, e.g. relativistic dragging of inertial frames \\citep{ste99}, see \\citep{kli05} for a more complete list of references.  There is however a trend in the models to associate QPOs with resonance phenomena, involving General Relativistic frequencies. As an example, in the accreting millisecond pulsar, SAXJ1808-3658, it has been proposed that the two QPOs (whose frequency difference is just equal to half the neutron star spin frequency) are generated at a radius in the disk where the difference between the general relativistic orbital frequency and radial epicyclic frequency is equal to half the spin frequency \\citep{wij03}. Similarly, to account for the 3:2 frequency ratios observed in black hole systems, a parametric resonance concept has been put forward, in which the QPOs are produced at a radius in the disk where two of the three general relativistic frequencies (orbital, vertical and radial epicyclic) have commensurate values, matching the observed QPO frequencies \\citep{abr04}. Although this is still very much under discussion, it is interesting to note that the latter two models link QPOs directly to general relativistic frequencies. \\subsection{The innermost stable circular orbit} The ISCO is one of the key predictions of strong gravity GR stating that there exists a region around sufficiently compact stars within which no stable circular orbital motion is possible. In a Schwarzschild geometry, the radius of the ISCO is 6GM/c$^2$, which is larger than the radii of neutron stars, deduced from models constructed with most modern high-density equations of state \\citep{akm98}. Interestingly enough, claims have already been made that the QPO properties have revealed signatures of the ISCO. Prior to the launch of RXTE, it was suggested that the ISCO might induce a frequency cutoff in the power density spectrum \\citep{klu90}.  After the discovery of kHz QPOs, it was proposed that signatures of the ISCO could include a frequency saturation of kHz QPOs with increasing mass accretion rate, or a drop in the quality factor and amplitude of kHz QPOs as they approach a limiting frequency \\citep{mil98}. A saturation of the QPO frequency with count rate \\citep{zha98} or inferred mass accretion rate \\citep{blo00}, as well as a sudden drop of coherence and amplitude at some critical frequency \\citep{bar05b} have now been reported. Although the interpretation of these results is still debated, this demonstrates the potential of fast X-ray timing for testing the existence of the ISCO. Note that identifying securely a frequency with an orbital frequency at the ISCO would yield directly the mass of the neutron star. \\subsection{QPO and the black hole spin and mass} Depending on the interpretation, parameters such as the spin and the mass of black holes can be constrained directly from measuring the QPO frequencies. One example is given in Fig. \\ref{fig:psa} where the most model independent constraint is used. It simply assumes that the QPO frequency is smaller than the orbital frequency at the ISCO (the latter being a function of the black hole mass and spin). More stringent constraints on the black hole mass and spin can be obtained from models involving resonance between relativistic frequencies \\citep{psa04,abr04}. All constraints point to maximally spinning black holes. As an application, because of the 1/M dependency of all general relativistic frequencies, it has been proposed that measuring QPOs from Ultra-Luminous X-ray sources, whose nature is a mystery, might help in constraining the mass of the compact object \\citep{abr04} in these systems. \\begin{figure} \\centerline{\\psfig{figure=col05-22_f2a.ps,height=2.5in}\\psfig{figure=col05-22_f2b.ps,height=2.5in}} \\caption{Left) The azimuthal ($f_\\phi$), radial epicyclic ($\\kappa$), and vertical ($f_\\perp$) frequencies at different radii around a black hole with spin parameter $a/M=0.2$. Right) {\\em Dashed Lines:\\/} the dynamical measurement of the black-hole mass in the source GRO~1655$-$40. {\\em Solid Line:\\/} the minimum spin parameter $a/M$ for each black-hole mass, for which the observed 450~Hz QPO can be produced as a Keplerian frequency of a stable orbit.  According to this argument, the black hole in GRO~1655$-$40 is spinning with $a/M>0.2$ (Both figures from \\citep{psa04}) \\label{fig:psa}} \\end{figure} ",
        "conclusions": "RXTE has opened a brand new window on fundamental physics. Gravity in the strong field limit and matter at supra-nuclear density, which constitute the core of the RXTE legacy, have been identified as key science goals to be pursued in the science program of ESA after 2015. XEUS, which is foreseen as a potential mission for this time period, now has a dedicated timing detector in the focal plane baseline, with capabilities matching the requirements for a follow-up instrumentation to RXTE. XEUS has therefore a great potential for delivering unique science on strong gravity and dense matter, building on what has been learned in the last few years with RXTE."
    },
    "0606/astro-ph0606566_arXiv.txt": {
        "abstract": "Satellite galaxies are tidally disrupted as they orbit the Milky Way.  If dark matter (DM) experiences a stronger self-attraction than baryons, stars will preferentially gain rather than lose energy during tidal disruption leading to an enhancement in the trailing compared to the leading tidal stream. The Sgr dwarf galaxy is seen to have roughly equal streams, challenging models in which DM and baryons accelerate differently by more than 10\\%.  Future observations and a better understanding of DM distribution should allow detection of equivalence violation at the percent level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0606/astro-ph0606750_arXiv.txt": {
        "abstract": "{ By making use of the sub-arcsecond angular resolution of the High Resolution Camera (HRC-I) aboard the Chandra X-ray satellite we have examined the central compact object \\rx\\, in the supernova remnant Puppis-A for a possible proper motion. Using data which span an epoch of 1952 days we found the position of \\rx\\, different by $0.574\\pm0.184$ arcsec, implying a proper motion of $\\mu=107.49\\pm34.46$ mas/yr. For a distance of 2.2 kpc, this proper motion is equivalent to a recoil velocity of $1121.79\\pm 359.60$ km/s. The position angle is found to be $241^\\circ \\pm 24^\\circ$. Since both the magnitude and direction of the proper motion are in agreement with the birth place of \\rx\\, being near to the optical expansion center of the supernova remnant, the result presented in this letter is a promising indication of a fast moving compact object in a supernova remnant. Although the positional shift inferred from the current data is significant at a $\\sim3\\sigma$ level only, one or more future HRC-I observations can obtain a much larger positional separation and further constrain the measurement. ",
        "introduction": "The group of supernova remnant (SNRs) which are known to host a radio-quiet but X-ray bright central compact object (CCO) is a slowly growing one. Thanks to more sensitive X-ray observatories it currently includes Cas$-$A (CXOU J232327.9+584842; Tananbaum 1999), RX J0852.0$-$4622 (CXOU J085201.4$-$461753; Aschenbach 1998), RX J1713.7$-$3946 (RX J1713.4$-$3949; Pfeffermann \\& Aschenbach 1996), RCW\\,103 (1E 161348-5055; Tuohy \\& Garmire 1980), \\Pu\\,(\\rx; Petre et al.~1982; Petre, Becker, \\& Winkler ~1996; Hui \\& Becker 2006), PKS 1209-51/52  (1E 1207.4$-$5209; Helfand \\& Becker 1984) and Kes 79 (CXOU J185238.6$+$004020; Seward et al.~2003; Gotthelf et al.~2005).  \\rx\\, was first noticed in one of the EINSTEIN HRI images of the \\pu\\ supernova remnant (Petre et al.~1982). However, it appeared in this data only as a faint X-ray feature. It is not until the era of ROSAT when it became evident that \\rx\\, is the central compact stellar remnant which was formed in the supernova event (Petre, Becker, \\& Winkler 1996).  Recently, Hui \\& Becker (2006) presented results from a detailed analysis of \\rx\\, which made use of all XMM-Newton and Chandra data available from it by beginning of 2005. The spectral analysis of XMM-Newton data revealed that the X-ray emission from \\rx\\, is in agreement with being of thermal origin. A double blackbody model with the temperatures $T_{1}=(2.35-2.91)\\times10^{6}$ K, $T_{2}=(4.84-5.3)\\times10^{6}$ K and the projected blackbody emitting radii $R_{1}=(2.55-4.41)$ km and $R_{2}=(600-870)$ m gave the best description of the observed point source spectrum among various spectral models tested. The X-ray images taken with the Chandra HRC-I camera allowed for the first time to examine the spatial nature of \\RX\\, with sub-arcsecond resolution. Besides an accurate measurement of the source position, this observation constrained the point source nature of \\rx\\, down to $0.59\\pm0.01$ arcsec (FWHM) for the first time.  Despite the effort in searching for coherent radio pulsations, \\rx\\, has not been detected as a radio pulsar (Kaspi et al. 1996). Similar to many other CCOs, it has no optical counterpart down to a limiting magnitude of $B\\ga 25.0$ and $R\\ga 23.6$ (Petre, Becker, \\& Winkler 1996). Together with the lack of long term flux variation (Hui \\& Becker 2006), all these evidences rule out many types of X-ray sources as a likely counterpart of \\rx, except a neutron star.  \\rx\\, is located about $\\sim 6$ arcmin distant from the optical expansion center of \\pu, which is at RA=$08^{\\rm h}22^{\\rm m}27.45^{\\rm s}$ and Dec=$-42^{\\circ}57'28.6\"$ (J2000) (cf.~Winkler \\& Kirshner 1985; Winkler et al.~1988). The age of the supernova remnant, as estimated from the kinematics of oxygen-rich filaments is $\\sim 3700$ years (Winkler et al.~1988). If these estimates are correct, \\rx\\, should have a proper motion of the order of $\\sim 100$ mas/yr to a direction away from its proposed birth place.  In this letter we test this hypothesis by making use of archival Chandra HRC-I data spanning an epoch of somewhat more than five years. The expected positional displacement for \\RX\\ in this time span is of the order of $\\sim 0.5$ arcsec. This is in the range of the Chandra accuracy given the possibility to correct for pointing uncertainties by using X-ray counterparts of stars with accurate position and which are serendipitously located in the field of view.  \\vspace{-0.5cm} ",
        "conclusions": "The fitted X-ray positions of \\rx\\, are plotted for the observational epochs December 1999 and April 2005 in Figure 2. The size of the error circle for each position is determined by adding the error in correcting the aspect offset and the error of the position of \\rx\\, in quadrature. This gives 0.162 arcsec and 0.088 arcsec for the December 1999 and the April 2005 observations, respectively. The corresponding error circles are indicated in Figure 2.  Comparing the positions of \\rx\\, as deduced for the two observations in 1999 and 2005 we found that they are different by $0.574\\pm 0.184$ arcsec. This is well consistent with the displacement estimated from the kinematic age of the SNR and the positional offset of \\rx\\, from the SNR's optical expansion center (cf.~\\S1). The quoted error is $1-\\sigma$ and combines the positional errors in both epochs in quadrature.  Given the epoch separation of 5.34 years for both Chandra HRC-I observations the observed proper motion of \\rx\\, is $\\mu=107.49\\pm 34.46$ mas/yr. The position angle is PA=$241^{\\circ}\\pm24^{\\circ}$. Taking the position inferred from the December 1999 observation and the position of the optical expansion center of \\pu\\, (Winkler et al.~1988), a position angle of PA$\\simeq 243^{\\circ}$ is implied. Such consistency suggests that \\rx\\, is indeed physically associated with \\pu\\, and its actual birth place is not too far away from the SNR's optical expansion center. Assuming that the distance of 2.2 kpc to \\pu\\, is correct, the observed proper motion implies that \\rx\\, has a projected recoil velocity as high as $1121.79\\pm 359.60$ km/s to the southwest. The inferred transverse velocity seems high compared with the average velocity, $\\sim 250$ km/s, of ordinary radio pulsars (Hobbs et al. 2005) though the large error prevents any constraining conclusions. We would like to mention, though, that there are several fast moving pulsars including PSRs B2011+38 and B2224+64 which have 2-D speeds of $\\sim 1600$ km/s (Hobbs et al. 2005).  The interesting association of \\pu\\ and \\rx\\, provides a unique opportunity to observe the interaction between a compact stellar remnant and a SNR. Also, this system is a test-bed for studying the dynamics of the supernova explosion. Measuring the proper motion of \\rx\\, with higher accuracy than possible with the current data will make it possible to provide constrains for the SN explosion model which formed \\pu.  Although our result is a promising indication of a fast moving CCO in a SNR,  we note that the deduced positional shift is only significant at a $3.1\\sigma$ level. Since the X-ray source A is rather faint, its relatively large positional error predominates in the error budget.  \\begin{figure} \\centerline{\\psfig{figure=5941fig2.eps,width=9cm,clip=}} \\caption{The best-fitted X-ray positions of \\rx\\, in two epochs are marked by crosses. The circles indicate the $1-\\sigma$ error. The arrow shows the direction of proper motion inferred from both positions. } \\end{figure}  In this letter we have presented the first attempt to measure the proper motion of \\rx. In order to further constrain this measurement, one or more Chandra HRC-I observations can attain a much larger positional separation in a few years from today. Furthermore, a more accurate proper motion measurement for the optical counterpart of the reference source is also necessary."
    },
    "0606/astro-ph0606036_arXiv.txt": {
        "abstract": "Mergers of gas-rich galaxies lead to gravitationally driven increases in gas pressure that can trigger intense bursts of star and cluster formation. Although star formation itself is clustered, most newborn stellar aggregates are unbound associations and disperse. Gravitationally bound star clusters that survive for at least 10--20 internal crossing times ($\\sim$20--40 Myr) are relatively rare and seem to contain $<$10\\% of all stars formed in the starbursts. The most massive young globular clusters formed in present-day mergers exceed $\\omega$ Cen by an order of magnitude in mass, yet appear to have normal stellar initial mass functions.  In the local universe, recent remnants of major gas-rich disk mergers appear as protoelliptical galaxies with subpopulations of typically 10$^2$--10$^3$ young metal-rich globular clusters in their halos. The evidence is now strong that these ``second-generation'' globular clusters formed from giant molecular clouds (GMC) in the merging disks, squeezed into collapse by large-scale shocks and high gas pressure rather than by high-velocity cloud--cloud collisions. Similarly, first-generation metal-poor globular clusters may have formed during cosmological reionization from low-metallicity GMCs squeezed by the universal reionization pressure. ",
        "introduction": "\\label{sec1}  When studying the myriads of point-like luminous sources brighter than any individual star on \\emph{HST} images of \\emph{ongoing} mergers (e.g., NGC 4038/39, NGC 3256), one would like to know which ones---or at least what fraction---will survive as globular clusters (GC). Yet, it is very difficult to distinguish gravitationally bound young star clusters from unbound OB associations or even spurious asterisms. As it turns out, the adopted operational definition for ``cluster'' may determine the answers to the scientific questions we ask about these objects.  Modern astronomical dictionaries universally include in their definition of ``star cluster'' (open or globular) the requirement that it be \\emph{gravitationally bound}, thus distinguishing it from any looser, expanding ``stellar association'' (e.g., \\cite{hopk76,ridp97}). As I explain in Sect.~\\ref{sec2} below, I believe that our present inability to make this distinction for many stellar aggregates younger than 10--20~$t_{\\rm cr}$ (internal crossing times) in ongoing mergers leads to a notion of ``infant mortality'' that is seriously exaggerated.  In recent merger \\emph{remnants}, where the merger-induced starburst has subsided (e.g., NGC 3921, NGC 7252), the definition of a \\emph{young globular cluster (YGC)} is more easy and secure. Any young compact stellar aggregate older than 10--20 $t_{\\rm cr}$ ($\\sim$20--40 Myr), more massive than a few $10^4 M_{\\odot}$, and with a half-light radius $R_{\\rm eff}$ comparable to that of a typical Milky-Way globular (say, $R_{\\rm eff} \\lapprox 10$~pc) is most likely gravitationally bound and, hence, a YGC. It is the size requirement that places stringent upper limits on any possible expansion velocity ($\\lapprox$0.2--0.5 \\kms) and thus guarantees that the cluster is gravitationally bound.  An important result to emerge from recent \\emph{HST} and follow-up studies of YGCs concerns their masses. These masses do not only cover the full range observed in old Milky-Way GCs ($\\sim$10$^4$ -- $5\\times 10^6 \\msun$), but also extend to nearly $10^8 \\msun$ or $\\sim$20$\\times$ the mass of $\\omega$~Cen at the high-mass end. The most massive YGCs are invariably found in remnants of gas-rich major mergers such as NGC 7252 \\cite{ss98,mara04}, NGC 1316 \\cite{bast06}, and NGC 5128 \\cite{mart04}. Interestingly, dynamical masses determined from velocity dispersions agree well with photometric masses based on cluster-evolution models with normal (e.g., Salpeter, Kroupa, or Chabrier) initial mass functions (IMFs). Therefore, some earlier worries that YGCs formed in mergers may have highly unusual stellar IMFs (e.g., \\cite{brod98}) seem now unfounded.  Relatively little work has been done so far on the brightness profiles and detailed structural parameters (core and tidal radii) of YGCs in mergers. Yet, the subject looks promising. Radial profiles of selected YGCs in NGC 4038 suggest that the initial power-law envelopes of YGCs may be tidally stripped within the first few 100 Myr, while the core radii may grow \\cite{whit99}. Correlations between core radius and cluster age are known to exist for the young cluster populations of the Magellanic Clouds (e.g., \\cite{mack03}) and deserve further study via the rich cluster populations of ongoing mergers and merger remnants. ",
        "conclusions": "\\label{sec5}  During mergers, increased gas pressure leads to much \\emph{apparent} cluster formation, but most of the stellar aggregates are unbound and disperse. Gravitationally bound globular and open clusters are relatively \\emph{rare} and seem to contain $<$10\\% of all stars formed in the starbursts.  Major gas-rich mergers form not only E and S0 galaxies, but also their metal-rich ``second-generation'' GCs. Specifically, in the local universe young remnants of major such mergers appear as protoelliptical galaxies with subpopulations of young metal-rich halo GCs (NGC 3921, NGC 7252; later NGC 1316, NGC 5128). The evidence is now strong that these second-generation GCs form from giant molecular clouds in the merging disks, squeezed into collapse by large-scale shocks and high gas pressure rather than by high-velocity cloud--cloud collisions.  Similarly, first-generation metal-poor GCs may have formed during cosmological reionization from low-metallicity giant molecular clouds squeezed by the reionization pressure. \\\\ \\\\ \\textbf{Acknowledgement.} I thank Brad Whitmore for his permission to reproduce some figures."
    },
    "0606/astro-ph0606200_arXiv.txt": {
        "abstract": "It is commonly assumed that high mass X-ray binary (HMXB) populations are little-affected by metallicity. However, the massive stars making up their progenitor systems depend on metallicity in a number of ways, not least through their winds. We present simulations, well-matched to the observed sample of Galactic HMXBs, which demonstrate that both the number and the mean period of HMXB progenitors can vary with metallicity, with the number increasing by about a factor of three between solar and SMC metallicity. However, the SMC population itself cannot be explained simply by metallicity effects; it requires both that the HMXBs observed therein primarily sample the older end of the HMXB population, and that the star formation rate at the time of their formation was very large. ",
        "introduction": "The high mass X-ray binaries represent a late stage in the life of a close massive binary, when one star has collapsed to a neutron star (NS) or black hole (BH). Accretion of matter lost in a wind or by Roche lobe overflow from its companion onto the compact object results in X-ray emission. As distinct from the low mass X-ray binaries (LMXBs), which may become luminous considerably later than the formation of the compact object and after considerable orbital evolution (e.g. Verbunt \\& ven den Heuvel 1995), the HMXB phase is though to occur relatively soon after the first SN in a system, and cannot last longer than the nuclear burning time ($ \\ltsimeq 3 \\times 10^{7} {\\rm yr}$) of the massive companion star. The occurence of HMXBs is therefore linked to recent star formation in a similar manner to massive stars.  The population is itself split into two subgroups, depending on the spectral type of the companion star (van Paradijs 1983). Approximately three-quarters of observed Galactic HMXBs have Be star companions. These systems have long periods and eccentric orbits; X-ray emission is often observed at periastron, when the NS interacts with the disk around the Be star. In contrast, the rest of the HMXB population is made up of OB supergiant systems, for which usually persistent emission is powered by the strong wind of the companion star prompting disk formation around the NS, or very occasionally by Roche lobe overflow (RLOF). These are generally shorter-period systems which have low eccentricity, often because the orbits have circularized. OB supergiant systems have significant runaway velocities whereas Be systems do not (Chevalier \\& Ilovaisky 1998), which suggests different evolutionary paths are involved (van den Heuvel et al. 2000). Due to the transient nature of the Be/X-ray binaries, it is likely that the true fraction of the population they represent is in fact significantly larger than that observed.  The question of whether HMXBs display metallicity dependence is an interesting one. It is commonly assumed that binary evolution is little-affected by the metallicity of the stars involved. However, this assumption may break down in the case of massive stars, since the loss of mass and angular momentum through winds during their lifetimes is non-negligible. A massive star which goes through a Wolf-Rayet phase may lose over half its mass in winds even if there is no binary interaction (e.g. Maeder \\& Meynet 2000). The exact metallicity dependence of massive star winds is still a matter of debate (e.g. Vink \\& de Koter 2005) but, since the winds of massive hot stars are line-driven, decreasing the metallicity is likely to lower the mass-loss rate. A ${\\rm Z^{0.5 - 0.7}}$ dependence for mass loss has been suggested (Kudritzki et al. 1989; Vink, de Koter \\& Lamers 2001). This may affect both the evolutionary path the system takes and also the compact object mass, both of which in turn may affect whether the system is split by the first supernova or goes on to undergo a HMXB phase. Theoretically, lower mass loss should mean greater rotation rates for lower-metallicity massive stars, since more angular momentum is retained. If this persists until the HMXB stage, it may increase the incidence of the propeller effect (Shtykovskiy \\& Gilfanov 2005a) which impedes accretion onto the NS. However, it should be noted that observationally there is some evidence that metallicity does not have a strong effect on rotation rates (Penny et al. 2004). A star of lower metallicity will also generally have a smaller radius, which in turn can delay the onset of RLOF with respect to a comparable high-metallicity system.  The population synthesis of binary black holes carried out by Belczynski, Sadowski \\& Rasio (2004, see also Belczynski et al. 2004b) finds an increased binary BH population at low Z, which they attribute both to lower mass-loss rates leading to a greater BH population and also to the greater retention of angular momentum leaving binaries harder to disrupt. This contrasts with the LMXB population, which may show an increase at higher metallicity (Maccarone, Kundu \\& Zepf 2004). Hurley, Tout and Pols (2002) also find an increase in HMXB population between $Z = 0.02$ and $Z=0.0001$. There are a large number of population synthesis calculations for or involving HMXBs carried out at solar metallicity: see e.g. Podsiadlowski et al. (2004) and references therein; Sepinsky, Kalogera \\& Belczynski (2005) and references therein; Dalton \\& Sarazin (1995); Pfahl et al. (2002); Portegies Zwart \\& Verbunt (1996); Brandt \\& Podsiadlowski (1995); Van Bever \\& Vanbeveren (2000); Kalogera (1996).  Observationally, it is hard to pick out trends which cannot be explained by other means. The Small Magellanic Cloud (SMC), with a metallicity of ${\\rm ~ 0.2 Z_{\\odot}}$, appears overabundant in HMXBs by a factor of about 50 per unit stellar mass when compared to the Milky Way (Majid, Lamb \\& Macomb 2004). It is likely, however, that this is primarily due to recent intense star formation. It is notable amongst the large SMC HMXB population that only one source (SMC X-1) has an OB supergiant companion. For the Large Magellanic Cloud (LMC), which is intermediate in metallicity between the SMC and the Milky Way, the HMXB population appears similar to the Galactic one (Neguerela \\& Coe 2002) albeit with a deficit of low-luminosity sources which may be due to the propeller effect (Shtykovskiy \\& Gilfanov 2005a).  To a first approximation, the X-ray luminosity function of HMXBs obeys a universal power law distribution; it has been suggested that X-ray luminosity would therefore be a possible star formation rate (SFR) indicator in distant galaxies (Grimm et al. 2003). This requires the number of HMXBs produced by a given amount of star formation to be roughly constant with metallicity.  We have previously carried out simulations aimed at reproducing the populations of runaway O and WR stars in the Galaxy and LMC (Dray et al. 2005, hereinafter D05). Over half of runaways get their high velocities from supernova explosions in binary systems (Hoogerwerf, de Bruijne \\& de Zeeuw 2001), most of which lead to the splitting of the binary system. It is an inevitable result of such simulations that one also obtains the parameters of a simulated population of binaries which are not separated. These are the HMXB progenitors. In this paper we consider their match to observed HMXB systems and the metallicity trends they reveal. ",
        "conclusions": "From previous sections, it appears that theoretical and observational agreement is reasonable for the Galaxy and LMC, but there remain some questions about the HMXB population of the SMC; it is hard to produce such an overabundence of long-period systems without an extremely intense starburst. This model failure, in the light of the relative success with the other populations, suggests either that there may be other effects coming into play for the SMC population which have not been considered here, perhaps related more specifically to the SMC environment (e.g. star formation in regions of high turbulence) or else that in general there is an extra effect which we have not considered which becomes more important once the metallicity is below some threshold value. However, in that case all starburst environments with populations of comparable age to the young SMC population and similar or lower metallicity should be similarly highly overabundant in HMXBs. One such effect could be, for example, the different evolutionary paths followed by very rapidly-rotating stars, if there are more such stars at low metallicity (Woosley \\& Heger 2006). The easiest way to create much larger HMXB populations is to have smaller SN kicks, because then a larger proportion of systems survive the first supernova. Thus the failure of our model for the SMC could be taken as indirect evidence for smaller kicks at low metallicity, again possibly related to the behaviour of rotation with metallicity.  There remain a few other interesting issues and conclusions which may be drawn from our simulations. It is notable, particularly in the light of D05 in which non-conservative RLOF produced a better match with the observed population, that some systems are difficult to fit with non-conservative mass transfer. This has previously been noted by Wellstein \\& Langer (1999) in the case of GX 301-2, although there is also significant evidence that non-conservative mass transfer must take case in some cases (van den Heuvel et al. 2000). The problem lies in producing a sufficient number of neutron stars close to $1.4 \\msun$ which have $\\sim 20 \\msun$ or greater companions such as Cen X--3 and QV Nor. We would expect therefore if these mass measurements are correct that at least some systems have (quasi-)conservative mass transfer. The better fit of conservative systems to the X-ray binary population than to the runaway population may in fact be entirely consistent, however. In real life, it seems likely that some RLOF is conservative and some is not, depending on, for instance, the rotation rates and mass ratio involved (Langer, Wellstein \\& Petrovic 2003; Petrovic, Langer \\& van der Hucht 2005). It might be considered that those systems which are conservative will end up at the pre-SN stage with more massive secondaries (albeit in generally wider orbits due to the loss of angular momentum in non-conservative mass transfer), and may hence be less likely to be split by the SN. However in D05 rates of binary splitting did not differ significantly between the two scenarios. A conservative system which is not split is also likely to have greater accretion rates than its non-conservative equivalent since the more massive companion star will have greater mass-loss rates. This may make conservative systems brighter in X-rays, suggesting a significant observational bias in favour of their detection.  A further consequence of accretion in compact binaries is that, in some cases, a neutron star may accrete enough mass that it exceeds the maximum NS mass. We have assumed in our simulations that this situation results in the `quiet' formation of a black hole. However, it has been recently suggested that such a scenario might, analagously to type Ia SNe, power a short gamma-ray burst (McFadyen, Ramirez-Ruiz \\& Zhang 2005). If this is the case, we would expect roughly one in $10^{5}$ stars over $10 \\msun$ to undergo this fate if accretion is not Eddington-limited, or as few as one in $10^{7}$ if it is. How many of these would be observable would, of course, depend on the degree of beaming involved and the intrinsic brightness of such an event. However it is notable that, certainly in the Eddington-limited case, the event rate of such transitions is significantly lower than we would expect for long GRBs, assuming  1\\% of SN Ib/c produce an observable long burst (Granot \\& Loeb 2003).  Finally, though we have limited our models to the Eddington luminosity, it is also interesting to note their behaviour when not so limited (Rappaport et al. 2005). In particular, we find that $L_{X}$ values as great as $10^{41} {\\rm erg\\,s^{-1}}$ may be (briefly) attained in some cases. A system observed with this luminosity would be a particularly bright ULX. The association with star-forming regions, particularly in the Cartwheel galaxy, implies that at least some ULXs are relatively young objects and probably an extension of the high-luminosity HMXB distribution (King 2004). Some LMXBs have periods of apparently super-Eddington luminosity (Grimm et al. 2002) and mechanisms have been proposed by which a HMXB containing a stellar mass BH can exceed the Eddington limit, such as beaming (King et al. 2001b) and photon-bubble instabilities in the disk (Ruskowski \\& Begelman 2003). If these mechanisms are responsible for some ULXs then from our models we would expect this ULX population to peak a little before $10^{7}$ years after a burst of star formation, and for each individual ULX to be a relatively short-lived phenomenon. The potential ULX population increases sharply (faster than the HMXB population) with decreasing metallicity, although it is also confined to a shorter time period."
    },
    "0606/astro-ph0606493_arXiv.txt": {
        "abstract": "We report a discovery of a new high redshift quasar at $z=5.96$, observed with the FOCAS long-slit spectrograph on board the Subaru telescope. The spectrum shows strong and broad Ly$\\alpha$+NV emission lines with a sharp discontinuity to the blue side. A Ly$\\beta$+OVI emission line is also detected, providing a consistent redshift measurement with the Ly$\\alpha$+NV emission.  The QSO has an absolute magnitude of $M_{AB,1450}=-26.9$ ($H_{0}=50$km s$^{-1}$ Mpc$^{-1}$, $q_0=0.5$). The spectrum shows significant flux in the region 8000-8300 $\\AA$ and thus does not show a complete  Gunn-Peterson trough in the redshift range 5.58 to 5.82, along the line of sight to this $z=5.96$ QSO. Therefore the Universe was already highly ionized at $z=5.82$. ",
        "introduction": "\\begin{figure*} \\begin{center} \\includegraphics[scale=1]{SDSSJ084119.52+290504.4.ps} \\end{center} \\caption{ The Subaru/FOCAS spectrum of a QSO, SDSSJ084119.52+290504.4, at z=5.96. The spectrum is smoothed using a 10 \\AA\\ box. Expected locations of various emission lines at $z=5.96$ are indicated with the dotted line. }\\label{fig:QSO_spectrum} \\end{figure*}   \\begin{figure*} \\begin{center} \\includegraphics[scale=0.85]{qso_grid_ver3.eps} \\end{center} \\caption{ The sky-subtracted two-dimensional spectrum of the QSO. The width of the vertical axis is 30 pixel or 9 arcsec. The vertical axis is expanded by 50\\% for clarity. }\\label{fig:QSO_2d} \\end{figure*}   High-redshift QSOs are highly important objects in various scientific aspects. They provide direct probes of the epoch when the first generation of galaxies and QSOs formed. The absorption spectra of these QSOs reveal the state of the intergalactic medium (IGM) close to the reionization epoch \\citep{1999ApJ...519..479H,2000ApJ...542L..69M,2000ApJ...542L..75C}. The lack of a Gunn-Peterson trough % \\citep{1965ApJ...142.1633G} % in the spectrum of the luminous QSO at $z=6.43$ \\citep{Fan_z6.43,2003AJ....126....1W,2005AJ....129.2102W} indicates that the universe was already highly ionized at that redshift. Assuming that the QSO is radiating at the Eddington luminosity, this object contains a central black hole of several billion solar masses. The assembly of such massive objects in a timescale shorter than 1 gigayear yields constraints on models of the formation of massive black holes \\citep[e.g.,][]{2001ApJ...552..459H}. The abundance and evolution of such QSOs can provide sensitive tests for models of QSO and galaxy evolution.  The Sloan Digital Sky Survey \\citep[SDSS,][]{2000AJ....120.1579Y} uses a dedicated 2.5m telescope and a large format CCD camera % to obtain images in five broad bands % \\citep[$u$, $g$, $r$, $i$ and $z$ centred at 3551, 4686, 6166, 7480 and 8932 \\AA, respectively;][]{1996AJ....111.1748F} over 10,000 deg$^2$ of high Galactic latitude sky. This unprecedented large sky coverage % provides us a unique opportunity to find a very rare class of objects such as highest redshift QSOs, passive spiral galaxies\\citep{2003PASJ...55..757G}, and E+A galaxies \\citep{2003PASJ...55..771G,2004A&A...427..125G,2005MNRAS.357..937G}. The inclusion of the reddest band, $z$, in principle enables the discovery of QSOs up to $ z \\sim 6.7$ from the SDSS data as a $z$-band only detection \\citep{Fan2000}. In this work, we have used the fourth public data release of the SDSS \\citep{2006ApJS..162...38A}.  Here we report our pilot search for high redshift QSOs with the Subaru telescope and a successful discovery of a new QSO at $z=5.96$. ",
        "conclusions": "We have found a new QSO at $z=5.96$ from the SDSS data. The object was selected for a spectroscopic follow-up with the Subaru telescope as a $z$-band only detection with blue $z-J$ colour. The broad Ly$\\alpha$+NV emission and the Ly$\\beta$+OVI emission provides a consistent redshift measurement of $z=5.96$, making this the 11th highest redshift QSO to date. The spectrum shows significant flux in the region 8000-8300 $\\AA$ and thus does not show a complete  Gunn-Peterson trough in the redshift range 5.58 to 5.82, along the line of sight to this $z=5.96$ QSO. Therefore the Universe was already highly ionized at $z=5.82$. It is our important future work to compute surveyed area and space density of high redshift QSOs once we finish observing all the selected targets."
    },
    "0606/astro-ph0606170_arXiv.txt": {
        "abstract": "Template fits to observed galaxy fluxes allow calculation of $K$-corrections and conversions among observations of galaxies at various wavelengths. We present a method for creating model-based template sets given a set of heterogeneous photometric and spectroscopic galaxy data. Our technique, non-negative matrix factorization, is akin to principle component analysis (PCA), except that it is constrained to produce nonnegative templates, it can use a basis set of models (rather than the delta function basis of PCA), and it naturally handles uncertainties, missing data, and heterogeneous data (including broad-band fluxes at various redshifts).  The particular implementation we present here is suitable for ultraviolet, optical, and near-infrared observations in the redshift range $0 < z < 1.5$. Since we base our templates on stellar population synthesis models, the results are intepretable in terms of approximate stellar masses and star-formation histories. We present templates fit with this method to data from GALEX, Sloan Digital Sky Survey spectroscopy and photometry, the Two-Micron All Sky Survey, the Deep Extragalactic Evolutionary Probe and the Great Observatories Origins Deep Survey. In addition, we present software for using such data to estimate $K$-corrections and stellar masses. ",
        "introduction": "\\label{motivation}  New surveys at low and high redshift have provided us with estimates of galaxy spectral energy distributions (SEDs) for an enormous number of galaxies. When comparing populations of galaxies at different redshifts in these surveys, we need to use comparable measurements of the galaxy SEDs. However, different surveys use different bandpasses and the restframe wavelengths of these bandpasses necessarily vary with redshift. We need to be able to handle this heterogeneity in order to make sensible comparisons among all of these new surveys.  In this paper, we present a method for doing so, by calculating $K$-corrections between observed and desired bandpasses.  The $K$-correction between a bandpass $R$ used to observe a galaxy at redshift $z$ and the desired bandpass $Q$ is defined by the equation (\\citealt{oke68a, hogg02a}): \\begin{equation} m_R = M_Q + \\mathrm{DM}(z) + K_{QR}(z) - 5 \\log_{10} h \\end{equation} where $\\mathrm{DM}(z) = 25 - 5\\log_{10} (d_L / (h^{-1}{\\mathrm{~Mpc}}))$ is the bolometric distance modulus calculated from the luminosity distance $d_L$, and $M_Q$ is the absolute magnitude. The absolute magnitude is defined as the apparent magnitude an object would have if were observed 10 pc away, in bandpass $Q$, at rest.  The traditional definition of the $K$-correction takes $Q=R$. However, we note that in practice many surveys do perform $K$-corrections from one observed bandpass $R$ to another bandpass $Q$ in the rest frame. This practice is particularly common when dealing with high redshift observations. In addition to $K$-corrections, this method also provides an interpretation of the data in terms of a physical model which describes the stellar mass and star-formation history of each galaxy.  The method is designed to work well for a wide range of data sets.  It uses photometry and spectroscopy of the Sloan Digital Sky Survey (SDSS; \\citealt{york00a}) the Galaxy Evolution Explorer (GALEX; \\citealt{martin05a}) in the ultraviolet, and the Two-Micron All Sky Survey (2MASS; \\citealt{skrutskie97a}) in the near infrared (NIR). In addition, at higher redshifts, we use constraints from the Deep Extragalactic Evolutionary Probe 2 (DEEP2; \\citealt{davis03a, faber03a}) and the Great Observatories Origins Deep Survey (GOODS; \\citealt{giavalisco04a}).  These and other data sets provide a huge set of information about galaxy colors and spectra which we can use to help understand their star-formation histories.  We note that some of the algorithmic techniques used here may have applications in other areas of astrophysics. First, nonnegative matrix factorization (NMF; \\citealt{lee00a}), and the extensions to it we describe here, is a particularly useful variant on principal component analysis. Second, the nonnegative least squares algorithm of \\citet{sha02a} has the virtue of being extremely simple to implement. These methods may find other applications in image analysis, spectroscopic analysis, and model fitting in astrophysics.  We also release the templates in electronic form as well as an implementation of the methods used to fit the templates to data. This software {\\tt kcorrect v4\\_1} is distributed on the World Wide Web{\\footnote {\\tt http://cosmo.nyu.edu/blanton/kcorrect/}}. It consists of a core C library which performs most of the complex and computationally intensive tasks, plus an IDL library that provides a high-level interface.  This software is an update of two major earlier releases ({\\tt v1\\_16}; \\citealt{blanton03b}; {\\tt v3\\_2}). The improvement over the previous version is twofold. First, we have improved the templates such that they successfully fit galaxies in the restframe UV, as observed by GALEX, DEEP2, and GOODS. Second, because the templates are now completely model-based, the fits have a physical interpretation in terms of a star-formation history. The IDL library also has a number of new useful functions.  In Section \\ref{algorithm}, we describe how to find suitable templates, given a set of models and a set of data. We also describe the data and models used here. In Section \\ref{results}, we show the results: the best-fit templates and how well they fit the data.  In Section \\ref{kcorrect}, we describe how to convert our results into an estimate of the $K$-correction. In Section \\ref{physical}, we describe the physical interpretation of the templates and of fits to data.  In Section \\ref{sdss2bessell}, we present some simplified, linear transformations between various bandpasses. In Section \\ref{summary}, we summarize our results. Appendix \\ref{nmf} describes the NMF algorithm.  Appendix \\ref{format} describes the electronic format of our results.  Where necessary, we have assumed cosmological parameters $\\Omega_0 = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 100$ $h$ km s$^{-1}$ Mpc$^{-1}$ (with $h=1$), unless otherwise noted. All magnitudes are (unless otherwise noted) AB-relative.  Except where noted, I will always be referring to the version of the {\\tt kcorrect} software labeled {\\tt v4\\_1\\_4}. ",
        "conclusions": "\\label{summary}  We have here described a method for fitting templates to nearly-arbitrary sets of spectra and broad-band fluxes. The basic concept is that we can recover a small set of templates (based on models), nonnegative linear combinations of which can explain a much larger set of inhomogeneous observations.  We have applied this method to SDSS, GALEX, 2MASS, DEEP2, and GOODS data to come up with a set of templates which are effective for describing all of these data sets.  Nonnegative matrix factorization (NMF) is an effective way to reduce the dimensionality of any large data set, and has that in common with PCA.  It is, in some sense, a ``nonnegative PCA.''  However, at least in the current context the method has several advantages over PCA. First, it has a natural physical interpretation associated with that spectral subspace, which is the corresponding subspace of all possible star-formation histories. Second, it naturally handles data uncertainties and missing data, which allows it to ignore that variation which is due purely to statistical errors.  Third, it handles the complications of observing galaxy spectra photometrically using broad-band filters.  We note here that the general NMF method does not {\\it depend} on using a model --- our ``model'' could just have been a set of top hat functions or Gaussians on a wavelength grid. In some situations, such as datasets for which there are not well-developed theoretical models, this approach could be more appropriate.  We have released our results in the form of a templates and a code base called {\\tt kcorrect} which fits those templates to many types of data (SDSS photometric and spectroscopic data, GALEX data, GOODS data, DEEP2 data, and 2MASS data). Furthermore, this code returns $K$-corrections and a physical interpretation of the photometry. All of the plots in this paper were created using code in the repository. The code is available from a web page maintained by one of the authors\\footnote{{\\tt http://cosmo.nyu.edu/blanton/kcorrect}}. This web page is kept updated on improvements in the code and new developments. It consists of a {\\tt C} language library, stand-alone {\\tt C} programs, and IDL language wrappers around the {\\tt C} library. Thus, one can use the basic templates and fitting code using only a machine with a {\\tt C} compiler. However, there is significant functionality which is programmed in the (unfortunately, proprietary) IDL language. These routines depend on the {\\tt idlutils} library\\footnote{{\\tt http://skymaps.info}}."
    },
    "0606/astro-ph0606346_arXiv.txt": {
        "abstract": "This paper studies the effects of dynamical interactions among the planets in observed extrasolar planetary systems, including hypothetical additional bodies, with a focus on secular perturbations. These interactions cause the eccentricities of the planets to explore a distribution of values over time scales that are long compared to observational time baselines, but short compared to the age of the systems.  The same formalism determines the eccentricity forcing of hypothetical test bodies (terrestrial planets) in these systems and we find which systems allow for potentially habitable planets.  Such planets would be driven to nonzero orbital eccentricity and we derive the distribution of stellar flux experienced by the planets over the course of their orbits.  The general relativistic corrections to secular interaction theory are included in the analysis and such effects are important in systems with close planets ($\\sim$4 day orbits). Some extrasolar planetary systems (e.g., Upsilon Andromedae) can be used as a test of general relativity, whereas in other systems, general relativity can be used to constrain the system parameters (e.g., $\\sin i \\gta 0.93$ for HD160691). For the case of hot Jupiters, we discuss how the absence of observed eccentricity implies the absence of companion planets. ",
        "introduction": "The number of discovered planets in other systems is rapidly growing and the observational sample shows an astonishing diversity of architectures (beginning with Mayor \\& Queloz 1995; Marcy \\& Butler 1996). One way to characterize these planetary systems is in terms of their orbital elements, where the semi-major axis $a$ and orbital eccentricity $e$ are most often used. These variables are equivalent to specifying the energy and angular momentum of the orbit. Previous expectations, informed by the structure of our Solar System, predicted planetary orbits with larger values of $a$ and smaller values of $e$ than those found in the current observational sample. A major theoretical effort is now being put forth to provide an explanation of the observed distributions of orbital elements, e.g., the $a-e$ plane. In multiple planet systems, however, it can be important to keep in mind that the orbital eccentricities can vary dramatically through secular interactions on time scales that are long compared to observational baselines but short compared to the system ages. Instead of being described by a single value of eccentricity, the orbits of planets in multiple planet systems should generally be characterized by a complete distribution of eccentricity values (see also Takeda \\& Rasio 2005).  Secular interactions, and the apparent lack of evidence for such interactions, can be used to place constraints on observed planetary systems. For the known multiple planet systems, secular interactions place strong constraints on the possibility of additional as-yet undetected terrestrial planets. For putative single planet systems with zero or low eccentricity, the inferred absence of secular interactions (which tend to prevent small eccentricity values $e \\sim 0$) implies that any additional planets in the system must have small masses, small eccentricities, and/or large semi-major axes; we quantify these constraints for observed systems. For multiple planet systems with hot Jupiters (close planetary companions, e.g., in 4 day orbits), secular perturbations tend to increase the eccentricity of the inner planet, which also experiences tidal stressing forces from the star. This interplay leads to continual energy dissipation in the planet.  This paper follows a large body of previous work. Secular interactions have been studied for centuries, primarily in the context of our Solar System (Brower \\& van Woerkom 1950; Laskar 1988; and many others). With the discovery of extrasolar multiple planet systems (starting with Upsilon Andromedae; Butler et al. 1999), secular perturbation theory can be applied to an ever growing collection of planetary systems (e.g., Wu \\& Goldreich 2002).  General schemes for studying secular interactions have been developed (Mardling \\& Lin 2002) and used to study close (hypothetical) terrestrial planets (Mardling \\& Lin 2004). For hierarchical planetary systems, a higher order (octopole) approximation scheme has been developed and applied to HD168443 and HD12661 (Lee \\& Peale 2003) and to triple star systems (Ford et al. 2000ab); an alternate approximation scheme has been developed and applied to the Upsilon Andromedae planetary system (Michtchenko \\& Malhotra 2004). The effects of disk potentials, and the disappearance of the disk, have been studied (Nagasawa et al. 2003) with the goal of explaining observed orbital eccentricities (see also Lubow \\& Ogilvie 2001). The effects of the Kozai mechanism have been explored for systems containing a binary companion (Takeda \\& Rasio 2005).  This paper builds upon the studies outlined above with the modest goal of using secular interaction theory to extract additional information from extant multiple planet systems.  The basic theory is outlined in \\S 2, which includes the leading order corrections for general relativity (GR).  The results and applications are then presented in \\S 3.  In \\S 3.1, we determine the distributions of eccentricity sampled by extrasolar planets over the course of their secular cycles, determine the secular time scales, and study the signature of secular eccentricity variations in the $a-e$ diagram. We place constraints on the possibility of additional terrestrial planets residing in currently detected extrasolar planetary systems (\\S 3.2) and place constraints on possible additional giant planets in observed hot Jupiter systems (\\S 3.4).  We explore the effects of general relativity in secular interactions and show that some systems can be used to test relativity, while in other systems relativity can be used to place constraints on system parameters (\\S 3.3). We conclude, in \\S 4, with a discussion and summary of results. ",
        "conclusions": "This paper has applied the leading order theory of secular interactions to observed extrasolar planetary systems and used the results to constrain their properties. Our results can be summarized as follows:  [1] Through dynamical interactions, described here using secular theory, the orbital eccentricities in multiple planet systems vary over secular time scales.  The eccentricities measured by ongoing planet searches represent the current eccentricity value, which is drawn from a wider distribution of values sampled by the planet. In other words, the eccentricities in multiple planet systems should not be considered as particular values, but rather as distributions of values.  The widths of these eccentricity distributions can be substantial (Tables 2 and 3) and we have verified that secular theory predicts distribution widths that are in good agreement with direct numerical integration (Figures \\ref{fig:numhd37124} -- \\ref{fig:numhd12661}). This effect implies that interpretations of the $a-e$ diagram must be regarded with caution for multiple planet systems (see Figure \\ref{fig:aeplane}). For the simplest case of two planet systems, the resulting distribution of eccentricity can be found analytically (eqs. [\\ref{eq:dpde} -- \\ref{eq:sigmae}]).  The time scale for secular eccentricity variations is typically thousands of years (Table 1), much longer than observational survey time scales (tens of years) and much shorter than the system lifetimes (few Gyr).  [2] For a sub-sample of observed multiple planet systems, we have calculated the frequencies of free oscillation and amplitudes of forced eccentricity oscillation (Figures \\ref{fig:reshd38529} -- \\ref{fig:reshd12661}) for (hypothetical) test bodies. These results constrain the possibility of multiple planet systems containing additional small (terrestrial) planets. This calculation shows that many of the systems allow the existence of additional planets, but the forced eccentricity amplitudes would be large compared to that of Earth, i.e., $e$(forced) $\\approx$ 0.10 -- 0.15. In addition, numerical integrations (Barnes \\& Raymond 2004) indicate that only a fraction of these systems allow test bodies to remain stable over long time scales, i.e., the condition of small forced eccentricity amplitude is necessary but not sufficient.  [3] This paper generalizes the concept of a habitable zone to include the possibility of larger eccentricity values (\\S 3.2.2). Specifically, we have calculated the distribution of flux experienced by a planet in an eccentric orbit over the course of its year (eq. [\\ref{eq:dpdf}]), as well as the mean flux and the width of the distribution (eqs. [\\ref{eq:fzero} -- \\ref{eq:sigmaf}]). Notice that all of these results are calculated analytically.  [4] For multiple planet systems with favorable architectures, the effects of general relativity can be significant (Fig. \\ref{fig:genrel}). When planetary systems have two relatively massive outer planets and a third inner planet of smaller mass near secular resonance, the effects of general relativity are large enough to move the system in or out of a resonant condition.  Since the resonance condition depends on planetary masses, which in turn depend on $\\sin i$, and since small inner planets cannot survive in exact resonance, this effect can be used to constrain the allowed range of $\\sin i$ in some observed extrasolar planetary systems. For the HD160691 system, e.g., this constraint implies that $\\sin i \\gta 0.93$.  This effect thus provides another means of probing the properties of these planetary systems. For other systems (e.g., Upsilon Andromedae), general relativity can make the inner planet precess forward in its orbit fast enough to compromise eccentricity pumping from a second planet, i.e., the mean eccentricity values are much smaller than they would be in flat space. This latter effect can thus be used as a test of general relativity. For Ups And, the system has a 78\\% chance of being observed with its measured parameters if general relativity is correct, but only a $\\sim 2\\%$ chance if the general relativistic corrections were zero. In the future, additional extrasolar planetary systems and/or higher precision observations can provide more stringent tests of GR (see also Adams \\& Laughlin 2006b).  [5] For single planet systems with small or zero measured eccentricity values (e.g., 51 Peg), the small observed $e$ values imply the absence of large, close planetary companions. For a given system, this line of reasoning implies a limit on the possible orbital elements of any additional planets in the system (shown in Fig. \\ref{fig:hotjup}).  The objective of this paper is to outline the effects of secular interactions in multiple planet systems. We have applied the results to a small collection of multiple planet systems that have been detected, and a number of additional applications are expected in the near future. Although the sample of extrasolar planets ($\\sim 150$) is now large enough to provide some statistical significance, the number of multiple planet systems is still low ($\\sim 20$). Furthermore, the possible parameter space accessible to multi-planet systems is enormous, so it is premature to make statistical statements about multiple planet systems. Nonetheless, future observations will find greater numbers of systems with multiple giant planets (through continued radial velocity searches and other methods) and systems containing smaller, Earth-mass planets (e.g., TPF).  Secular interactions can add to our understanding of these forthcoming multiple planet systems in a variety of ways. In trying to find theoretical explanations for the observed orbital elements, one must take into account the distributions of eccentricities driven by secular interactions (Fig, \\ref{fig:aeplane}).  In systems with known giant planets, the search for Earths can be guided by studying the forced eccentricity variations as depicted in Figures \\ref{fig:reshd38529} -- \\ref{fig:reshd12661}. Some multiple planet systems will provide additional (stronger) tests of general relativity (as in Fig. \\ref{fig:genrel}).  In other systems, we can deduce the presence or absence of additional (undetected) planets --- or at least constrain their properties --- through examination of the properties of the detected planets (Fig. \\ref{fig:hotjup}). Over longer time spans, secular interactions combine with tidal circularization and energy dissipation processes (as outlined in our companion paper Adams \\& Laughlin 2006a).  \\bigskip  This work was supported at U. Michigan (FCA) by the Michigan Center for Theoretical Physics and by NASA through the Terrestrial Planet Finder Mission (NNG04G190G) and the Astrophysics Theory Program (NNG04GK56G0).  This material is based in part upon work supported by the National Science Foundation CAREER program under Grant No. 0449986 (GL), and was also supported at U. C. Santa Cruz (GL) by NASA through the Terrestrial Planet Finder Precursor Science Program (NNG04G191G) and through the Origins of Solar Systems Program (NAG5-13285).  \\newpage"
    },
    "0606/astro-ph0606292_arXiv.txt": {
        "abstract": "We compare the highly clustered populations of very high redshift galaxies with proto-clusters identified numerically in a standard $\\Lambda$CDM universe  ($\\Omega_0=0.3,\\, \\lambda_0=0.7$) simulation. We evolve $256^3$ dark matter particles in a comoving box of side $150h^{-1}\\Mpc$. By the present day there are 63 cluster sized objects of mass in excess of $10^{14}h^{-1}\\msun$ in this box. We trace these clusters back to higher redshift finding that their progenitors at $z=4$--$5$ are extended regions of typically 20--40 Mpc (comoving) in size, with dark halos of mass in excess of $10^{12}h^{-1}\\msun$ and are overdense by typically 1.3--13 times the cosmological mean density. Comparison with the observation of Ly$\\alpha$ emitting (LAEs) galaxies at $z=4.86$ and at $z=4.1$ indicates that the observed excess clustering is consistent with that expected for a proto-cluster region if LAEs typically correspond to massive dark halos of more than $10^{12}h^{-1}\\msun$. We give a brief discussion on the relation between high redshift concentration of massive dark halos and present day rich clusters of galaxies. ",
        "introduction": "The formation and evolution of structures in the universe, such as galaxy clusters and large-scale structures, are part of the most important issue in astrophysics. Since clustering properties of galaxies in the distant universe give us great clues to this issue, many observations have been done in order to detect (proto) clusters of galaxies at high redshifts \\citep{Steidel98, Steidel00,  Campos99, Pentericci00, Rhoads00, Ouchi01, Ouchi03, Miley04}. These deep surveys of galaxies, which have significantly advanced our understanding of the properties and distribution of galaxies at high redshifts, are equally important in constraining the underlying structure formation scenarios.  \\citet{Venemans02} showed a region around a radio galaxy in which number density of Ly$\\alpha$ emitters (LAEs) is much higher than the mean of the universe suggested by \\citet{Rhoads00} at $z=4.1$. The region has size of 2.7 $\\times$ 1.8 Mpc (physical) and mass of $\\sim 10^{15}\\msun$. More recently, \\citet{Shimasaku03} found that LAEs are clustered in an elongated region with a size of 20 $\\times$ 50 Mpc (comoving) at $z=4.86$, which is comparable to the size of present-day large-scale structures. In this elongated region, there is a circular region of high surface density of LAEs with 12 Mpc radius, which may be a progenitor of a galaxy cluster. They also estimated the bias parameter of the LAEs in the range of $b\\sim 3$--$16$  for spatially flat low-density cosmological model ($\\Lambda$CDM model). The elongated distribution of the LAEs in this region is also proposed to be a part of large-scale structures, although observation by \\citet{Shimasaku04} of the same sky area which is closer to us by 39Mpc (comoving) does not show large structure. In addition, \\citet{Hayashino04} reported that LAEs distribution at $z=3.1$ in their observed area shows broad large scale structure. They claim that this structure cannot be explained in the context of the standard CDM model. \\citet{Ouchi05} observed LAEs at $z=5.7$ and discovered filamentary structures and voids.  Although these regions which have much high number densities of LAEs are claimed to be proto-clusters, there is not yet enough detailed study on these possibilities. Therefore, we investigate formation and evolution of galaxy clusters in early universe using cosmological simulation. In this letter we report some properties of simulated proto-clusters at $z=5$ and compare our numerical results with observations by \\citet{Shimasaku03} and \\citet{Venemans02}. ",
        "conclusions": "We have investigated proto-cluster regions and the difference between those regions and background field regions using N-body cosmological code. In order to investigate proto-clusters, we identify 63 galaxy clusters at in our cosmological simulation $z=0$, and trace dark matter particles which belong to those, back to a high redshift epoch. We study properties of the proto-cluster regions by analyzing the mass overdensity, $\\dmass$, and the halo overdensity, $\\dhalo$, of which definitions are given  in \\S\\ref{sec:method}. The results at $z=5$ are shown in Figures \\ref{fig:histogram_dhalo}, \\ref{fig:histogram_dmass}, and \\ref{fig:scat_dm_dh}. As expected from the bias theory \\citep[e.g.][]{TS00}, $\\dhalo$ shows excess over $\\dmass$ and variance of $\\dhalo$ is much larger than that of $\\dmass$.  In order to show difference between the proto-cluster regions and the field, we calculate $\\dmass$ and $\\dhalo$ for the randomly selected regions with the same size of the typical proto-cluster region at $z=5$ (25 Mpc in the comoving scale). We select 630 regions (10 times greater than the proto-cluster regions) for the randomly selected regions. Clustering properties of the dark halos in these regions are significantly different from those in the proto-cluster regions. The halo overdensities for most of the randomly selected regions are much smaller than those for the proto-cluster regions as shown in Figure \\ref{fig:histogram_dhalo}.  It is very interesting to study a critical value of $\\dhalo$ by which we can select a region as the proto-cluster regions. From Figure \\ref{fig:dh_z5_Mmax_z0}, we have found that more than 80\\% of the regions with $\\dhalo \\ge 3$ at $z=5$ contain rich clusters ($M \\ge 10^{14}\\hiMsun$) at the present epoch. Thus, we conclude that $\\dhalo \\ge 3$ is a good criterion to distinguish proto-cluster regions from field regions at $z=5$. This criterion is very useful to judge whether observed galaxies excess regions are actual proto-clusters or not.  It should be pointed out that there are a few regions which do not contain clusters at $z=0$ but have large $\\dhalo$ values (3--7) at $z=5$, although a fraction of those regions are very small. We discuss that the physical origin of the large $\\dhalo$ is non-linearity and stochasticity of bias parameters proposed by \\citet{TS00}. They studied variance of biasing parameters, and conclude that the variance increases strongly with redshift. They also conclude that stochasticity of the biasing is generated by the scatter in the halo mass distribution at higher redshift. \\citet{Yoshikawa01} confirmed these results for $0 < z< 3$ by large $\\mathrm{P^3M}$ simulation. We conclude that our results for $z=3$ are consistent with their results and the trend of evolution of bias variance is still true for $3 < z < 5$.  We compare our results with the observation of \\citet{Shimasaku03}. Properties of the LAEs concentrated region reported by them are as follows: (1) Diameter of the region is 25 Mpc (comoving unit), (2) Projected overdensity of LAEs $\\delta_{\\Sigma} \\sim 2$, (3) Bias parameter is estimated 3--16 for $\\Lambda$CDM model, and the best-fit value is $b\\sim 6$, i.e.~$\\dmass\\sim 0.3$. (4) The number of LAEs in the region is about 20. Properties from (1) to (3) are consistent with our numerical results of proto-cluster regions, if LAEs correspond to dark halos with mass more than $10^{12}h^{-1}\\msun$, which is suggested by \\citet{Hamana04} using correlation function on small-scales. The typical number of dark halos in the simulated proto-clusters is 5--10, which is about half of that of observed LAEs. One explanation of this discrepancy is due to possibility that some dark halos have more than one LAEs. We speculate that some pairs of LAEs in Fig.~3 of \\citet{Shimasaku03} are included in the same dark halo. It is also possible that some of observed sample of LAEs in the region are low redshift interlopers. \\citet{Shimasaku03} estimate the contamination of their sample to be about 20\\%.  \\citet{Venemans02} also reported properties of a LAEs rich region. The size of the region is $\\sim$ 14 $\\times$ 10 Mpc (comoving unit). They estimate the region is overdense in LAEs by a factor of 15 compared with the blank field \\citep{Rhoads00}. The overdensity is much larger than that of proto-clusters at $z=4$ in our simulation and the size of their observed region is smaller than the typical size of the proto-clusters in our simulation. Therefore, we suggest that the region observed by them is not a whole proto-cluster region, but the central region of the proto-cluster because of small size of the observed region and their high value of the overdensity."
    },
    "0606/astro-ph0606547_arXiv.txt": {
        "abstract": "{\\sl Aims.} {The purpose of this study is to investigate the distributions of the isomeric molecules HCN and HNC and estimate their abundance ratio in the protostellar core Cha-MMS1 located in Chamaeleon {\\sc i}.}  {\\sl Methods.} {The core was mapped in the $J=1-0$ rotational lines of \\hcn, \\hnc, and \\hnthic. The column densities of \\hthicn, \\hnthic, \\hfifnc\\ and \\ammo\\ were estimated towards the centre of the core.}  {\\sl Results.} {The core is well delineated in all three maps. The kinetic temperature in the core, derived from the \\ammo\\ (1,1) and (2,2) inversion lines, is $12.1 \\pm 0.1$ K.  The \\hnthic/\\hthicn\\ column density ratio is between 3 and 4, i.e. similar to values found in several other cold cores.  The \\hnthic/\\hfifnc\\ column density ratio is $\\sim 7$.  In case no $^{15}$N fractionation occurs in \\hnc\\ (as suggested by recent modelling results), the \\hnc/\\hnthic\\ abundance ratio is in the range $30-40$, which indicates a high degree of $^{13}$C fractionation in \\hnc.  Assuming no differential $^{13}$C fractionation the \\hcn\\ and \\hnc\\ abundances are estimated to be $\\sim 7 \\, 10^{-10}$ and $\\sim 2\\, 10^{-9}$, respectively, the former being nearly two orders of magnitude smaller than that of \\ammo.  Using also previously determined column densities in Cha-MMS1, we can put the most commonly observed nitrogenous molecules in the following order according to their fractional abundances: $\\chi(\\ammo) > \\chi(\\cyana) > \\chi(\\hnc) > \\chi(\\hcn) > \\chi(\\diaz)$.}  {\\sl Conclusions.} {The relationships between molecular abundances suggest that Cha-MMS1 represents an evolved chemical stage, experiencing at present the 'late-time' cyanopolyyne peak. The possibility that the relatively high HNC/HCN ratio derived here is only valid for the $^{13}$C isotopic substitutes cannot be excluded on the basis of the present and other available data.} ",
        "introduction": "Hydrogen cyanide, \\hcn, and its metastable isomer hydrogen isocyanide, \\hnc, are commonly used tracers of dense gas in molecular clouds. The \\hnc\\ column density has been observed to be similar or larger than that of \\hcn\\ in cold regions, whereas in warm GMC cores the [HNC]/[HCN] ratio is typically much smaller than unity (e.g., \\cite{goldsmith1981}; \\cite{churchwell1984}; \\cite{irvine1984}; \\cite{schilke1992}; \\cite{hirota1998}).  Gas-phase chemistry models, including both ion-molecule and neutral-neutral reactions, predict that the \\hnc/\\hcn\\ abundance ratio should be close to unity in cold gas (see e.g. \\cite{herbst2000}, where also previous work is reviewed).  The two isomers are thought to form primarily via the dissociative recombination of HCNH$^+$, which in turn is produced by a reaction between \\ammo\\ and C$^+$. The recombination reaction regulates the abundance ratio despite the initial production mechanisms of the two isomers, provided that they are efficiently protonated via reactions with ions, such as H$_3^+$ and HCO$^+$ (\\cite{churchwell1984}; \\cite{brown1989}). Recent theoretical studies corroborate the equal branching ratios for \\hcn\\ and \\hnc\\ in dissociative recombination (\\cite{hickman2005}; \\cite{ishii2006}).   \\hnc/\\hcn\\ column density ratios clearly larger than unity have been reported: \\cite{churchwell1984} find that [\\hnc]/[\\hcn] may be as high as $\\sim 10$ in some dark clouds cores, and \\cite{hirota1998} find a ratio of $\\sim 5$ in L1498.  The first result quoted is uncertain because it depends on the assumed degree of $^{13}$C fractionation in HNC, whereas the second value refers to \\hnthic\\ and \\hthicn\\ measurements. Nevertheless, both results suggest that the \\hcn/\\hnc\\ chemistry is not fully understood.  Further determinations of the relative abundances of these molecules towards cores with well-defined chemical and physical characteristics seem therefore warranted.  In this paper we report on observations of \\hcn\\ and \\hnc\\ and some of their isotopologues towards the dense core Cha-MMS1, which is situated near the reflection nebula Ced 110 in the central part of the Chamaeleon I dark cloud.  We also derive the gas kinetic temperature in the core by using the (1,1) and (2,2) inversion lines of \\ammo. Cha-MMS1 was discovered in the 1.3 mm dust continuum by \\cite{reipurth1996}, and was studied in several molecules by \\cite{kontinen2000}. The object is embodied in one of the most massive 'clumps' in Chamaeleon {\\sc i} identified in the large scale \\ceio\\ survey of \\cite{haikala2005} (clump no. 3, $\\sim 12 \\Msun$). \\cite{reipurth1996} suggested that Cha-MMS1 contains a Class 0 protostar. A FIR source (Ced 110 IRS10) was detected near the centre of the core by \\cite{lehtinen2001}. However, no centimetre continuum nor near-infrared sources have been found in its neighbourhood (\\cite{lehtinen2003} and references therein), and the core therefore represents a very early stage of star formation. At the same time, its chemical composition has probably reached an advanced stage (\\cite{kontinen2000}).  In Sect. 2 of this paper we describe the observations. In Sect. 3 we present the direct observational results in the form of spectra and maps, and derive the \\hthicn, \\hnthic, and \\hfifnc\\ column densities. By combining our results with previous observations we derive in Sect.~4 the fractional abundances of several nitrogen containing molecules in Cha-MMS1, and discuss briefly the chemical state of the core. In Sect. 5 we summarize our results. ",
        "conclusions": "\\subsection{Fractional abundances}  The present observations and the LTE method give reasonable estimates for the column densities of \\hthicn, \\hnthic, \\hfifnc\\ and \\ammo\\ near the centre of Cha-MMS1. Unfortunately the column densities of the common isotopologues of \\hcn\\ and \\hnc\\ cannot be directly determined because of the large optical thickness of their lines. \\hfifnc\\ is, however, likely to be useful for this purpose.  According to the modelling results of \\cite{terzieva2000} hardly any $^{15}$N fractionation occurs in dense interstellar clouds.  The isotopic $^{14}$N/$^{15}$N ratio lies probably between 240 and 280, i.e. the values derived for the local ISM (\\cite{lucas1998}), and the terrestrial atmosphere.  Using this range and the \\hfifnc\\ column density listed in Table~\\ref{table:columndensities} we obtain $N(\\hnc) \\sim 6 \\, 10^{13} \\, \\percmsq$. Assuming that no differential $^{13}$C fractionation occurs, we get from the \\hnthic/\\hthicn\\ ratio derived above that $N(\\hcn) \\sim 2 \\, 10^{13}$ \\percmsq.  The dust continuum emission provides the most reliable estimate for the \\molh\\ column density. \\cite{reipurth1996} measured towards Cha-MMS1 a total 1.3 mm flux density of 950 mJy from a nearly circular region with a diameter of about $50^{\\prime\\prime}$ (bordered by the lowest white contour in Fig.~8).  As this size is similar to the SEST beamsize at 3 mm, the resulting average intensity and average $N(\\molh)$ can be compared with the column densities derived here. The dust temperature as derived from ISO observations by \\cite{lehtinen2001} is 20 K (Cha-MMS1 = Ced 110 IRS10). Due to crowding of sources in this region and the limited angular resolution of ISOPHOT, this value may be affected by emission from other sources in its vicinity (see Fig. 2 of \\cite{lehtinen2001}). Therefore, value 20 K is probably an upper limit.  A lower limit is provided by the gas kinetic temperature $\\sim 12$ K derived above. By using the two temperatures and Eq. (1) of \\cite{motte1998} with $\\kappa_{\\rm 1.3 mm} = 0.01$ g\\percmsq\\ appropriate for circumstellar envelopes around Class 0 and Class I protostars, we get $N(\\molh) = 1.9 - 4.0 \\,10^{22}\\, {\\percmsq}$. The corresponding mass range of the dust core obtained from the total flux given by \\cite{reipurth1996} is $0.4-0.8\\, \\Msun$.  The median of the reasonable H$_2$ column density range, $N(\\molh) = 3.0 \\,10^{22}\\, {\\percmsq}$, yields the following values for the fractional \\hcn, \\hnc\\ and \\ammo\\ abundances: $\\chi(\\hcn) \\sim 7 \\, 10^{-10}$, $\\chi(\\hnc) \\sim 2 \\, 10^{-9}$, and $\\chi(\\ammo) \\sim 5 \\, 10^{-8}$. The column densities of two other nitrogen bearing molecules towards Cha-MMS1 were determined by \\cite{kontinen2000}: $N(\\cyana) = (4.5 \\pm 0.8) \\, 10^{14}$ \\percmsq, and $N(\\diaz) = (1.4 \\pm 0.1) \\, 10^{13}$ \\percmsq, which imply the fractional abundances $\\chi(\\cyana) \\sim 2 \\, 10^{-8}$ and $\\chi(\\diaz) \\sim 5 \\, 10^{-10}$.  We find that both \\hcn\\ and \\diaz\\ are about 100 times less abundant than \\ammo, whereas \\cyana\\ is almost equally abundant as \\ammo. The derived \\diaz/\\ammo\\ column density ratio is in accordance with previous observational results (e.g., \\cite{tafalla2002}; \\cite{hotzel2004}). The abundance ratios are discussed in Sect. 4.3. The \\ceio\\ column density in this direction is $(1.5 \\pm 0.1) \\, 10^{15}$ \\percmsq\\ (\\cite{haikala2005}).  The fractional \\ceio\\ abundance is therefore $5 \\, 10^{-8}$ or smaller. This upper limit falls below the 'standard' value $\\sim 2\\, 10 ^{-7}$ of \\cite{frerking1982}, and suggests CO depletion.   \\subsection{Comparison with previous maps}  The superposition of our \\hnthic\\ map, the \\cyana\\ map of \\cite{kontinen2000}, and the 1.3 mm dust continuum map of \\cite{reipurth1996} is shown in Fig.~\\ref{figure:hc3n+dust+hn13c}. The two molecular line maps have similar angular resolution (about $60^{\\prime\\prime}$) whereas the resolution of the continuum map is $22^{\\prime\\prime}$. Both molecules show rather compact, symmetric distributions around the dust continuum source, \\cyana\\ agreeing slightly better with dust than \\hnthic. Nevertheless, the dense core can be clearly localized with the aid of \\hnthic\\ as well as \\hcn\\ and \\hnc. This is in contrast with SO, which shows two peaks on both sides of the core (\\cite{kontinen2000}), and with C$^{18}$O which has a flat distribution in this region (\\cite{haikala2005}), probably due to depletion. This result agrees with the results of \\cite{hirota2003} and \\cite{nikolic2003}, which suggest that \\hcn\\ and \\hnc\\ are less affected by the accretion onto dust grains than some other commonly observed molecules.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{tennekes_fig8.ps}} \\caption{The \\hnthic\\ map (dark contours) superposed on the \\cyana\\ map from \\cite{kontinen2000} (greyscale) and the 1.3 mm dust continuum map of \\cite{reipurth1996} (white contours).} \\label{figure:hc3n+dust+hn13c} \\end{figure}  \\subsection{The chemical state of Cha-MMS1 and the HNC/HCN abundance ratio}  CO depletion and a large \\ammo\\ abundance point towards an advanced stage of chemical evolution. The \\ammo/\\diaz\\ abundance ratio $\\sim 100$ agrees with model predictions for dense cores, but does not put severe constraints on the evolutionary stage (\\cite{aikawa2005}). According to the observational results of \\cite{hotzel2004} this value should be more appropriate for a star-forming than for a prestellar core. However, due to slightly different beamsizes the accuracy of the present determination of the \\ammo/\\diaz\\ ratio is not sufficient for distinguishing between the two cases.  \\cite{kontinen2000} suggested that the large \\cyana\\ abundance observed in Cha-MMS1 indicates either a so called 'late-time' cyanopolyyne peak associated with a large degree of depletion according to the model of \\cite{ruffle1997}, or carbon chemistry revival by the influence of a newly born, so far undetected star.  A large ammonia abundance can be understood in either case.  Ammonia production benefits from the diminishing of CO and can resist depletion longer (e.g. \\cite{nejad1999}). On the other hand, ammonia can possibly form on grain surfaces and be released into the gas phase in shocks (\\cite{nejad1990}; \\cite{aikawa2003}). As there is no clear evidence for star-cloud interaction in Cha-MMS1, the alternative that the core - or the portion traced by the present observations - represents matured pre-stellar chemistry is more likely.  The behaviour of \\hcn\\ and \\hnc\\ under these circumstances is not clear. The main gas-phase production pathway via the HCNH$^+$, which is effective especially at later stages, requires \\ammo\\ and C$^+$ ions. The latter can be produced also in deep interiors of molecular clouds via the cosmic ray ionization of C atoms, or in the reaction between CO and the He$^+$ ion, the latter mechanism becoming less important with CO depletion (\\cite{nejad1999}).  To our knowledge the production of \\hcn\\ and \\hnc\\ on interstellar grains and their desorption has not been studied.  The models of \\cite{aikawa2005} for collapsing prestellar cores include the gas-phase production of \\hcn, \\cyana, \\ammo, and \\diaz. Although the order of fractional abundances, $\\chi(\\ammo) > \\chi(\\hcn) > \\chi(\\diaz)$, and the \\ammo/\\diaz\\ abundance ratio derived from the present observations are well reproduced by this model, especially for the gravity-to-pressure ratio $\\alpha=4$ impying a rapid collapse and less marked depletion, the predicted \\hcn\\ column densities relative to \\ammo\\ and \\diaz\\ are an order of magnitude larger than those observed here.  A likely explanation is that \\hcn\\ is not as centrally peaked than \\ammo\\ and \\diaz\\ (see Fig.~3 of \\cite{aikawa2005}).  Therefore a large fraction of \\hcn\\ resides in on the outskirts of the core, where the density is not sufficient for the collisional excitation of its rotational levels, and the molecule is not seen in emission (see also the next section). Another discrepancy is that the model predicts much lower \\cyana\\ column densities relative to ammonia than observed here. The agreement is slightly better for the model with overwhelmingly large gravitational potential ($\\alpha=4$) than for the slowly evolving model near hydrostatic equilibrium ($\\alpha=1.1$) which results in very high degrees of depletion. Unfortunately the spectral resolution and S/N ratio of the present data do not allow to study the core kinematics.  As compared to the 'principal' N-bearing molecules \\ammo\\ and \\cyana, \\hcn\\ and \\hnc\\ have clearly smaller abundances, and probably only marginal importance to the nitrogen chemistry in the core. Nevertheless, the fact that \\hnc/\\hcn\\ abundance ratio is once again observed to be larger than unity remains a problem. \\cite{talbi1998} have shown that even though the reaction between \\ammo\\ and C$^+$ may produce a significant amount of the metastable ion ${\\rm H_2NC^+}$, which should produce only \\hnc\\ in electron recombination, the energy released in the former reaction leads to the efficient transformation into the linear isomer ${\\rm HCNH}^+$ before the recombination. Likewise, according to \\cite{herbst2000}, the vibrational energies of \\hnc\\ and \\hcn\\ formed in the {\\sl neutral} reactions ${\\rm C} + {\\rm NH_2} \\rightarrow \\hnc + {\\rm H}$, and ${\\rm N} + {\\rm CH_2} \\rightarrow \\hcn + {\\rm H}$ are sufficient to overcome the isomerization barrier, and roughly equal amounts of both isomers are produced.  The theoretical results presented in \\cite{talbi1998} and \\cite{herbst2000}, and more recently by \\cite{hickman2005} and \\cite{ishii2006}, show quite convincingly that the \\hnc/\\hcn\\ abundance ratio resulting from the chemical reactions known to be effective in dark clouds {\\sl cannot} differ significantly from unity. There are two possible explanations for the discrepancy between the observations and theoretical predictions. Either so far unknown processes favour \\hnc\\ in the cost of \\hcn, or then the derived \\hnthic/\\hthicn\\ column density ratio is affected by different degrees of $^{13}$C fractionation in \\hnc\\ and \\hcn. In the latter case $^{13}$C would about three times more enhanced in \\hnc\\ than in \\hcn\\ in Cha-MMS1.  The \\hnthic/\\hfifnc\\ column density ratio 7.4 together with the above quoted range of $^{14}$N/$^{15}$N ratios imply that the \\hnc/\\hnthic\\ ratio in Cha-MMS1 is in the range $30-40$. As the typical $^{12}$C/$^{13}$C isotopic ratio observed in nearby clouds is 60 (\\cite{lucas1998}) the result suggests considerable $^{13}$C enhancement in \\hnc. An \\hcfifn\\ column density determination would readily show, according to the results of \\cite{terzieva2000}, whether also \\hcn\\ has a similar degree of $^{13}$C fractionation.  In the only survey so far including all these isotopologues, \\cite{irvine1984} found very similar \\hnthic/\\hfifnc\\ and \\hthicn/\\hcfifn\\ intensity ratios ($\\sim 1.5$) towards TMC-1. This observation suggests similar $^{13}$C fractionation in \\hcn\\ and \\hnc. On the other hand, this result might not be universally valid. \\cite{ikeda2002} determined substantially different \\hthicn/\\hcfifn\\ column density ratios in the starless dark cloud cores L1521E, TMC-1, and L1498, listed in order of increasing \\hthicn/\\hcfifn\\ ratio (2.6, 6.7, and 11).  The \\hnthic/\\hthicn\\ column density ratios in these sources increase in the same order (0.54, 1.7, and 4.5; \\cite{hirota1998}).  The values obtained towards TMC-1 are consistent with those of \\cite{irvine1984}, but the results towards L1521E and L1498 suggest (if the possibility of $^{15}$N fractionation is neglected) that the $^{13}$C fractionation in \\hcn\\ changes from source to source, and perhaps the variation in the case of \\hnc\\ is even larger.  L1521E is suggested to represent youthful chemistry with little molecular depletion (\\cite{hirota2002}; \\cite{tafalla2004}) whereas the position observed in L1498 is known to be heavily depleted (\\cite{willacy1998}).  Thus, it appears that the freeze-out of molecules favour the $^{13}$C substituted isotopologues of \\hcn\\ and \\hnc.  It seems warranted to further examine the possibility of their differential $^{13}$C fractionation, although no theoretical arguments are in favour of different gas-phase production and destruction rates of \\hnthic\\ and \\hthicn, or their different freezing rates. As no empirical data on this matter is currently available, the best way to proceed is probably to perform similar observations as those done by \\cite{irvine1984} towards a sample of cores with different degree of depletion.  \\subsection{Scattering envelope}  The low \\hcn/\\hthicn\\ and \\hnc/\\hnthic\\ intensity ratios and the line shapes in the \\hcn\\ and \\hnc\\ spectra observed towards the centre of the core require some attention.  The intensity ratios suggest very large optical thicknesses for the \\hcn\\ and \\hnc\\ lines, and consequently, deep self-absorption features and 'anomalous' hyperfine component ratios are to be expected.  A clear absorption feature can be seen in the {\\sl simulated} \\hcn\\ spectrum shown in Fig.~\\ref{figure:montecarlo}, although a very low $^{12}$C/$^{13}$C isotopic ratio is used in order to avoid inverted hyperfine ratios, which would disagree with the observations.  The {\\sl observed} \\hcn\\ and \\hnc\\ spectra are, however, free from absorption dips, and they look more like weakened 'normal' lines.  This weakening can probably be attributed to less dense gas surrounding the core. The \\hcn\\ and \\hnc\\ molecules in the envelope, where the density lies below the critical density of their $J=1-0$ lines, function as scatterers of line emission arising from the core, i.e. they re-emit absorbed photons (into random directions) before colliding with other particles.  The effect is less marked for $^{13}$C isotopologues, and the resulting \\hcn/\\hthicn\\ and \\hnc/\\hnthic\\ intensity ratios decrease, without \\hcn\\ and \\hnc\\ loosing their characteristic line patterns. This model was used by \\cite{cernicharo1984} to explain \\hcn\\ hyperfine anomalies, particularly the intesity of the intrinsically weakest component $F=0-1$ observed towards dark clouds. The \\hcn\\ spectrum shown in Fig.~\\ref{figure:5spectra} conforms with this explanation: the $F=0-1$ component on the left is stronger than $F=1-1$ component on the right, although the optically thin LTE ratio is $1:3$.  The scattering of \\hcn\\ emission from the core has probably lead to underestimation of the excitation temperature and column densities derived using the LTE method in Sect.~3.2. For example, an increase of 1 K in $T_{\\rm ex}$, would mean 30 \\% larger column densities, and 3 K would correspond to a 100\\% increase in the column density. Quantitative estimates of the scattering effects are, however, beyond the scope of this paper, and here we merely note this additional source of error."
    },
    "0606/astro-ph0606221_arXiv.txt": {
        "abstract": "We consider secular perturbations of nearly Keplerian two-body motion under a perturbing potential that can be approximated to sufficient accuracy by expanding it to second order in the coordinates. After averaging over time to obtain the secular Hamiltonian, we use angular momentum and eccentricity vectors as elements. The method of variation of constants then leads to a set of equations of motion  that are simple and regular, thus allowing efficient numerical integration.  Some possible applications are briefly described. ",
        "introduction": "If a binary star, or a planet revolving around a star, is perturbed by one or more distant relatively slowly moving masses the perturbation can be approximated by a tidal field. Fast computation of secular perturbations in such a potential is thus desirable. Potential applications include the Kozai resonance, motion of a well isolated binary in a star cluster and cometary motion under the galactic tidal field.  Several authors have considered the computation of cometary motion perturbations due to the galactic tidal field e.g. \\cite{1986Icar...65...13H}, % \\cite{1999Icar..137...84W}, % \\cite{2001MNRAS.324.1109B}. % More recently \\cite{2004MNRAS.349..347F} % suggested the use of the time averaged secular perturbing function and the Lagrangian equations for computing the evolution of orbital elements of a comet. While \\cite{2005CeMDA..93..229F} % compared various ways to compute the evolution of orbital elements (using Keplerian or Delaunay elements).  \\cite{2005MNRAS.364.1222B} % used the so called vectorial elements, i.e. the areal velocity vector and the eccentricity vector. However, the perturbing potential considered in all these studies was a simplified one, in most cases just the term causing the force perpendicular to the galactic plane was included.  In this paper we consider a general quadratic potential. For more mathematical generality we include the possibility of a constant force (although such a term does not appear in a tidal field) and also a slow rotation of the quadratic perturbing field is allowed. We form the simple regular equations of motion using the vectorial elements.  Finally it is necessary to stress that the main purpose of this paper is to introduce the simple secular equations of motion when a two-body system is perturbed by a (at most) quadratic perturbing potential. ",
        "conclusions": "We have generalized the equations of \\cite{2005MNRAS.364.1222B} who considered the simple case of perturbing function of the form $g_{33} z^2$.  The equations for vectorial elements, even in the case of a general quadratic potential are simple and regular, contrary to what one obtains when using Keplerian or Delaunay elements (\\cite{2005CeMDA..93..229F}).  The integration of the secular theory equations is typically faster than direct coordinate integration  by  two orders of magnitude.  The implicit midpoint method may be even faster than  Burlirsh-Stoer, but there is no simple way of knowing accuracy and optimal stepsize.  The most reliable integration method is Burlirsh-Stoer extrapolation.  For comets the secular theory results are not accurate for $a>>20,000$AU, as comparison with direct integration shows (simply: perturbation can be too large).  For an integration over only one  period at a time the midpoint method and Burlirsh-Stoer are nearly equally fast. Thus in this case the midpoint method may be preferable due to its simplicity."
    },
    "0606/astro-ph0606017_arXiv.txt": {
        "abstract": "A synthesis of the present knowledge on gamma-ray emission from the magnetosphere of a rapidly rotating neutron star is presented, focusing on the electrodynamics of particle accelerators. The combined curvature, synchrotron, and inverse-Compton emission from ultra-relativistic positrons and electrons, which are created by two-photon and/or one-photon pair creation processes, or emitted from the neutron-star surface, provide us with essential information on the properties of the accelerator --- electric potential drop along the magnetic field lines. A new accelerator model, which is a mixture of traditional inner-gap and outer gap models, is also proposed, by solving the Poisson equation for the electrostatic potential together with the Boltzmann equations for particles and gamma-rays in the two-dimensional configuration and two-dimensional momentum spaces. ",
        "introduction": "\\label{sec:nature} When a star has exhausted its nuclear fuel, the interior thermal pressure can no longer support its own weight. If the star is not very massive, it collapses to the point at which electron degeneracy pressure halts gravitational collapse, and a white dwarf is formed. However, beyond the critical mass~\\cite{Chandra31}, even the electron degeneracy pressure becomes insufficient to support the star. In the collapsing stellar core, most of the electrons are absorbed by protons via inverse $\\beta$ decay and eventually the neutron degeneracy pressure halts the gravitational collapse~\\cite{Landau32,Baade34,Oppen39}. For a progenitor star to collapse directly to a neutron star, its mass should be more than $10$ times greater than our Sun. A neutron star, which can rotate with a period as fast as $P \\sim 1$~ms, is considered to be the only candidate for the source of remarkably stable radio pulses, of which fastest period is 1.55~ms (PSR~J1939+2134~\\cite{Lyne87}). When it was determined that the period of the Crab pulsar was slowly increasing~\\cite{Richards69}, the identification of pulsars as rapidly rotating neutron stars was essentially confirmed~\\cite{Gold68,Gold69}.  Having established pulsars as rotating neutron stars, the next thing to discuss is the pulsar characteristics deduced from observations. Since all known isolated pulsars show gradual spin down, there must be some mechanisms causing the neutron star to lose their rotational energy. The rotational energy-loss rate is given by \\begin{equation} \\dot{E}=I\\Omega\\dot\\Omega=-(2\\pi)^2I\\dot{P}/P^3, \\label{eq:Edot} \\end{equation} where the neutron star is rotating with angular frequency $\\Omega=2\\pi/P$ and moment of inertia $I$. For typical high-energy pulsars, $P\\sim 0.1$~s, $\\dot{P}\\sim 10^{-13}\\,\\mbox{s s}^{-1}$, $I\\sim 10^{45} \\mbox{g\\,cm}^2$; thus, we obtain $\\dot{E}=-4\\times 10^{36}\\,\\mbox{ergs s}^{-1}$, which is enough to explain the observed $\\gamma$-ray luminosities, $10^{32.5}\\,\\mbox{ergs s}^{-1} < L_\\gamma < 10^{35}\\,\\mbox{ergs s}^{-1}$. For an object as small as $r_\\ast \\sim 10$~km in radius to experience such a large torque, it must have a strong coupling to their surroundings, most likely through magnetic fields intrinsic to the neutron star. If the neutron star is losing its rotational energy via magnetic dipole radiation, the spin-down luminosity is given by \\begin{equation} L_{\\rm dip}=k\\Omega^4\\mu^2/c^3, \\label{eq:Ldip} \\end{equation} where $k=2\\sin^2\\inc/3$ is commonly used for a vacuum rotating magnetic dipole with inclination $\\inc$ with respect to the rotation axis. Equating $L_{\\rm dip}$ with $-\\dot{E}$, we can infer the magnetic moment, $\\mu$, of the star~\\cite{Ostr69}. For $P=0.1$~s, $\\dot{P}=10^{-13}\\,\\mbox{s s}^{-1}$, $I=10^{45} \\mbox{g\\,cm}^2$, we obtain $\\mu = 2.6\\times 10^{30}\\,\\mbox{G cm}^3$ and the surface magnetic field strength, $B_\\ast = 5.2\\times 10^{12} (r_\\ast/10\\mbox{km})^{-3}$~G, as the Newtonian value at the magnetic pole. In what follows, we consider possible emission models of isolated, rotating neutron stars, after briefly mentioning the high-energy observations. ",
        "conclusions": ""
    },
    "0606/astro-ph0606367_arXiv.txt": {
        "abstract": "{While [\\ion{O}{iii}] narrow-band imaging is commonly used to measure the size of the narrow-line regions (NLRs) in active galactic nuclei (AGNs), it can be contaminated by emission from surrounding starbursts. Recently, we have shown that long-slit spectroscopy provides a valuable alternative approach to probe the size in terms of AGN photoionisation. Moreover, several parameters of the NLR can be directly accessed.} {We here apply the same methods developed and described for the Seyfert-2 galaxy NGC\\,1386 to study the NLR of five other Seyfert-2 galaxies by using high-sensitivity spatially-resolved optical spectroscopy obtained at the VLT and the NTT. } {We probe the AGN-photoionisation of the NLR and thus, its ``real'' size using diagnostic line-ratio diagrams. We derive physical properties of the NLR such as reddening, ionisation parameter, electron density, and velocity as a function of distance from the nucleus. } {For NGC\\,5643, the diagnostic diagrams unveil a similar transition between line ratios falling in the AGN regime and those typical for \\ion{H}{ii} regions as found for NGC\\,1386, thus determining the size of the NLR. For the other four objects, all measured line ratios fall in the AGN regime. In almost all cases, both electron density and ionisation parameter decrease with radius. Deviations from this general behaviour (such as a secondary peak) seen in both the ionisation parameter and electron density can be interpreted as signs of shocks from the interaction of a radio jet and the NLR gas. In several objects, the gaseous velocity distribution is characteristic for rotational motion in an (inclined) emission-line disk in the centre. We compare our results to those of NGC\\,1386 and show that the latter can be considered as prototypical also for this larger sample. We discuss our findings in detail for each object.} {} ",
        "introduction": "The luminous central engine in active galaxies, most likely an accreting supermassive black hole (BH), ionises the surrounding interstellar medium leading to the broad-line region (BLR) in the inner part and the narrow-line region (NLR) further out. Among the central AGN components, the NLR has the advantage to be directly accessible via imaging and spatially resolved spectroscopy, at least for nearby AGNs.  [\\ion{O}{iii}]\\,$\\lambda$5007\\AA~(hereafter [\\ion{O}{iii}]) narrow-band imaging is commonly used to study the NLRs of active galaxies. However, this emission can be contaminated by contributions from star formation, shock-ionised gas or tidal tails, resulting in an apparent increase of the NLR. In addition, the measured size depends on the depth of the images: When comparing ground based [\\ion{O}{iii}] images of Seyfert galaxies from \\citet{mul96} with the HST snapshot survey of \\citet{sch03}, the latter reveal, on average, six times smaller NLR sizes, probably due to the 15 to 20 times lower sensitivity. These considerations question the definition of the ``NLR size'' from [\\ion{O}{iii}] imaging alone.  Spatially resolved long-slit spectroscopy is a valuable alternative approach as it can directly probe the size in terms of AGN photoionisation and discriminate the stellar or shock-ionised contribution. In addition, several physical parameters of the NLR such as reddening, ionisation parameter, electron density, and velocity can be directly accessed and analysed as a function of distance from the nucleus.  In \\citet{ben05} and \\citet{ben06a} (hereafter paper I), we describe methods developed to probe the AGN-photoionisation of the NLR and thus, its ``real'' size as well as to derive physical conditions within the NLR of the nearby Seyfert-2 galaxy NGC\\,1386: We use the galaxy itself for subtracting the stellar template, applying reddening corrections to fit the stellar template to the spectra of the NLR. From spatially resolved spectral diagnostics, we find a transition between central line ratios falling into the AGN regime and outer ones in the \\ion{H}{ii}-region regime. Applying \\texttt{CLOUDY} photoionisation models \\citep{fer98}, we show that the observed distinction between \\ion{H}{ii}-like and AGN-like ratios in NGC\\,1386 represents a true difference in ionisation source and cannot be explained by variations of physical parameters such as ionisation parameter, electron density or metallicity. We interpret it as a real border between the NLR, i.e. the central AGN-photoionised region, and surrounding \\ion{H}{ii} regions. We find that both the electron density and the ionisation parameter decrease with radius. The differences between the reddening distributions determined from the continuum slope and the Balmer decrement argue in favour of dust intrinsic to the NLR clouds with varying column density along the line of sight. The NLR and stellar velocity fields are similar and indicate that the NLR gas is distributed in a disk rather than a sphere.  Here, we apply the same methods to a larger sample of five Seyfert--2 galaxies to probe the size of the NLR. We derive physical properties such as reddening, ionisation parameter, electron density, and velocity, and discuss their variations with distance from the nucleus. In our discussion, we include the results for NGC\\,1386 from paper I. A detailed comparison of our results with literature data is given for each object [see also \\citet{ben05}].  A similar study was carried out for six Seyfert-1 galaxies. The results and the comparison with the Seyfert-2 galaxies presented here will be summarised in \\citet{ben06c}. ",
        "conclusions": "We study high-sensitivity spatially-resolved spectra along the extended [\\ion{O}{iii}] emission of six Seyfert-2 galaxies obtained with the VLT and the NTT. To derive the pure emission-line fluxes, we successfully use the galaxy itself as stellar template to subtract the underlying absorption lines in four out of six objects.  The nuclear spectra reveal the typical strong NLR emission from oxygen at different ionisation states, lines from ionised nitrogen and sulphur, as well as Balmer lines. In most objects, high-excitation iron lines are additionally seen in the central spectra, originating from the powerful and hard ionisation source in the centre.  Plotting line-ratios from our spatially resolved spectra in diagnostic diagrams, we observe a transition of emission-line ratios from the central AGN region to \\ion{H}{ii} region in another Seyfert-2 galaxy (NGC\\,5643), in addition to the Seyfert-2 galaxy NGC\\,1386 already discussed in paper I. The most probable explanation for this transition is that the stellar ionisation field starts to dominate that of the AGN. We are thus able to determine the radius of the NLR independent of sensitivity and excluding [\\ion{O}{iii}] contamination from circumnuclear starbursts. In former spectroscopic studies, the observed [\\ion{O}{iii}] has often been attributed to the extended NLR. We can show that at least part of this ``extended NLR'' emission is actually powered by \\ion{H}{ii} regions and that only the central few arcseconds are indeed gas photoionised by the AGN.  In the other four objects, no such transition is observed but the line ratios fall in the AGN regime in all three diagnostic diagrams. Thus, the determined NLR size (1000-5000\\,pc) is a lower limit, limited by either the S/N of our data or the lack of a strong surrounding stellar ionisation field.  We derive physical parameters of the NLR such as reddening, surface brightness, electron density, and ionisation parameter as a function of projected distance from the nucleus. The differences between the reddening distributions determined from the continuum slope and the Balmer decrement argue in favour of  dust intrinsic to the NLR with a varying column density along the line-of-sight. In most cases, both electron density and ionisation parameter decrease with radius.  In all the objects, the gas rotation curve shows a large-scale velocity gradient suggestive of rotation (though the detailed modelling was not carried out and some outflow might be present too). This is an important hint that the NLR gas is distributed in a disk rather than a sphere, an issue that is still a matter of debate.  We discuss the results for each object (see Appendix). In individual objects, substructures are seen in both the ionisation parameter and electron density and can often be interpreted as signs of shocks from the interaction of a radio jet and the NLR gas.  Our results for the five  Seyfert-2 galaxies show that the NLR properties of the Seyfert-2 galaxy NGC\\,1386 in paper I can be considered as prototypical, putting our conclusions on a larger statistical basis. We performed a similar study of the NLR of six Seyfert-1 galaxies. The comparison with the results presented here will be summarised in \\citet{ben06c}. All results presented here are described and discussed in detail in \\citet{ben05}."
    },
    "0606/astro-ph0606198_arXiv.txt": {
        "abstract": "Nuclear activity in galaxies is closely connected to galactic mergers and supermassive black holes (SBH). Galactic mergers perturb substantially the dynamics of gas and stellar population in the merging galaxies, and they are expected to lead to formation of supermassive binary black holes (BBH) in the center of mass of the galaxies merged. A scheme is proposed here that connects the peak magnitude of the nuclear activity with evolution of a BBH system. The scheme predicts correctly the relative fractions of different types of active galactic nuclei (AGN) and explains the connection between the galactic type and the strength of the nuclear activity.  It shows that most powerful AGN should result from mergers with small mass ratios, while weaker activity is produced in unequal mergers. The scheme explains also the observed lack of galaxies with two active nuclei, which is attributed to effective disruption of accretion disks around the secondary in BBH systems with masses of the primary smaller than $\\sim 10^{10}\\,M_\\odot$. ",
        "introduction": "Important roles played by galactic mergers and binary black holes in galaxy evolution was first recognized several decades ago \\cite{bbr80,roos1981}. A large number of subsequent studies have addressed the problems of evolution of binary black holes in post-merger galaxies (see \\cite{mm05} for a recent review) and connection between mergers and nuclear activity in galaxies \\cite{dimatteo03,haehnelt02,springel05}. The correlations observed between the masses, $M_\\mathrm{BH}$, of nuclear black holes in galaxies, the and masses, $M_\\star$, \\cite{magorrian1998} and velocity dispersions, $\\sigma_\\star$, of the host stellar bulges \\cite{ferrarese2000,gebhardt2000,tremaine2002} suggest a connection between the formation and evolution of the black holes and galaxies. Growth of black holes in galactic centers is self-regulated by outflows generated during periods of supercritical accretion \\cite{fabian1999,haehnelt1998,silk1998}. This mechanism offers a plausible explanation for the observed $M_\\mathrm{BH}$--$\\sigma$ relation \\cite{king2003}.  Supermassive black holes are expected to form in the early Universe, with multiple SBH likely to be common in galaxies \\cite{haehnelt02}. However, the detailed connection between the SBH evolution and the nuclear activity is somewhat elusive. Observational evidence for binary SBH is largely indirect (see \\cite{lobanov2005,mm05} and references therein), with only two double galactic nuclei (NGC\\,6240 \\cite{komossa2003} and 3C\\,75 \\cite{owen1985}) observed directly in early merger systems, at large separations. There are no convincing cases for secondary black holes within active galaxies, although some of them may be hiding among the extranuclear X-ray point sources detected by ROSAT and Chandra \\cite{colbert2002,miller2004}. This implies that the activity of secondary companions is quenched at early stages of the merger, possibly due to disruption of the accretion disk.   The nuclear activity depends strongly on the availability of accreting material in the immediate vicinity of a black hole, and AGN episodes are believed to last for $\\sim 10^7$--$10^8$\\,yr. This is likely to be smaller than typical lifetimes of nuclear binary black holes in galaxies \\cite{mm05}. This suggests that nuclear binary black holes systems may provide a mechanism necessary for instilling and supporting high accretion rates over timescales implied by large-scale relativistic outflows produced in AGN \\cite{dokuchaev91,polnarev94}. Evolutionary stages of BBH systems can also be connected phenomenologically to different types of AGN \\cite{lobanov2005b}. Here, an analytical model is proposed that connects the evolution of central SBH in to the nuclear activity in galaxies.   \\subsection{BBH evolution}  The main constituents of the model are: 1)~binary system of supermassive black holes, 2)~accretion disk, 3)~central stellar bulge. The BBH is described by the masses $M_1$, $M_2$ ($M_1\\ge M_2$) of the two black hole and their separation $r$.  The accretion disk is assumed to be a viscous Shakura-Sunyaev disk, with a mass $M_\\mathrm{d}$.  The disk extends from $\\rho_\\mathrm{in}\\,R_\\mathrm{g}$ to $\\rho_\\mathrm{out}\\,R_\\mathrm{g}$, where $R_\\mathrm{g}$ is the gravitational radius and $\\rho_\\mathrm{in}\\approx 6$ and $\\rho_\\mathrm{out} \\approx 10^4$ \\cite{ivanov+1998}.  The central bulge extends over a region of radius $r_\\star$ and has a mass $M_\\star$ ($M_\\star > M_{12} = (M_1 + M_2)$) and a velocity dispersion $\\sigma_\\star$.  The evolution of the BBH is described in terms of \\emph{reduced mass}, $\\tilde{M}$, and \\emph{reduced separation}, $\\tilde{r}$ of the binary. The reduced mass is defined as $\\tilde{M} = 2\\, M_2\\, /\\,M_{12}$.  This definition implies $\\tilde{M} = 0$ for $M_2 = 0$ and $\\tilde{M} = 1$ for $M_2 = M_1$. If $q = M_2\\,/\\,M_1$ is the mass ratio in the system, then $\\tilde{M} = 2\\,q \\,/(1+q)$.  The reduced separation is given by $\\tilde{r} = r\\,/\\,(r+r_\\mathrm{c})$, where $r_\\mathrm{c}$ is the separation at which the two black holes become gravitationally bound (this happens at $\\tilde{r} = 1/2$). Binary systems have $\\tilde{r}\\le 1/2$, while unbound pairs of SBH have $\\tilde{r}>1/2$.  \\subsection{Accretion disk disruption in binary black holes}  Two SBH in a merger galaxy are expected to form a binary system at a separation $r_\\mathrm{c} = r_\\star\\, (M_{12}\\, /\\,M_\\star)^{1/3}$, with an initial orbital speed $v_\\mathrm{init} = \\sigma_\\star (M_{12}\\,/\\,M_\\star)^{1/3}$ \\cite{bbr80}.  Assuming that the relative speed of the two black holes reaches asymptotically its Keplerian value, the \\emph{approach speed} can be defined as $v_\\mathrm{appr}(r) = \\sigma_\\star \\,(r_\\star\\,/\\, r)^{1/2}\\,(M_{12}\\,/\\, M_\\star)$.  Both black holes are assumed to have active accretion disks at early stages of the merger. The separations $r_\\mathrm{d}$ at which the accretion disks are disrupted and eventually destroyed can be estimated for each of the two black holes by equating the approach speed to the Keplerian velocity at the outer edge of the disk: $v_\\mathrm{k,out} = c/\\sqrt{\\rho_\\mathrm{out}}$. This yields $r_\\mathrm{d} = \\rho_\\mathrm{out}\\, r_\\mathrm{c}\\, (\\sigma_\\star\\,/\\,c)^{2} \\, (M_{12}\\,/\\,M_\\star)^{2} \\,$.  The bulge mass and velocity dispersion must satisfy the $M_\\mathrm{BH}$--$\\sigma_\\star$ and $M_\\mathrm{BH}$--$M_\\star$ relations \\cite{king2003}. The resulting reduced separation is $\\tilde{r}_\\mathrm{d} = 1/(1+\\xi)$, where $\\xi = M_1 / [1.86\\times 10^7\\,M_\\odot\\, \\rho_\\mathrm{out}\\, \\phi^2 \\, (2-\\tilde{M})^3]$ and $\\phi$ is the collimation angle of the outflow carrying the excess energy and angular momentum from the immediate vicinity of the black hole. This corresponds to a critical mass $M_\\mathrm{eq} = 1.86\\times 10^7\\, M_\\odot\\, \\rho_\\mathrm{out}\\, \\phi^2$ for which an equal mass binary system will undergo disk destruction at the time of gravitational binding (at $\\tilde{r}_\\mathrm{d} = \\tilde{r}_\\mathrm{c} = 1/2$. In systems with $M_1 < M_\\mathrm{eq}$ the destruction of accretion disk around the secondary will occur before the formation of a gravitationally bound system. For typical values of $\\rho_\\mathrm{out} \\approx 10^4$ and $\\phi = 0.1$--$0.3$, $M_\\mathrm{eq}$ reaches $10^9$--$10^{10}\\,M_\\odot$. It implies that most of active galaxies formed by galactic mergers should undergo destruction of the disk around the secondary BH before or during the formation of a gravitationally bound systems. Since masses of the nuclear black holes in galaxies rarely exceed $10^{10}\\, M_\\odot$, this offers a natural explanation for the observed lack of active galaxies with double nuclei, since it predicts that \\emph{in most galaxies with binary black hole systems the secondary companion will be inactive}.  Denoting $\\epsilon_1 = M_1/M_\\mathrm{eq}$, the disruption distances are \\[ \\tilde{r}_{d1} = \\left(1 + \\frac{\\epsilon_1}{\\tilde{M}^2 (2-\\tilde{M})} \\right)^{-1} \\quad \\mathrm{and} \\quad \\tilde{r}_{d2} = \\left(1 + \\frac{\\epsilon_1}{(2-\\tilde{M})^3} \\right)^{-1}\\,, \\] for the primary and secondary black hole, respectively. A circumbinary disk can exist at orbital separations smaller than $\\sim G\\, M_1 \\rho_\\mathrm{out}^{1/2} c^{-2}$. These three characteristic distances are shown in left panel of Fig.~\\ref{lobanov1-fig1}, for $M_1 = M_\\mathrm{eq}$.  \\begin{figure}[t] \\centering \\includegraphics[width=0.49\\textwidth]{lobanov1F1a.eps} \\includegraphics[width=0.49\\textwidth]{lobanov1F1b.eps} \\caption{Properties of binary black holes in the $\\tilde{M}$--$\\tilde{r}$ plane ($\\tilde{r}=1/2$ signifies the capture distance at which a pair of black holes becomes gravitationally bound). \\textbf{Left:}~Reduced separations for the disruption distance of the accretion disks around the secondary (solid line) and the primary (dashed line) black holes. The separations are calculated for $M_1= M_\\mathrm{eq}$. Above the $\\tilde{r}_\\mathrm{d2}$ line, both black holes retain accretion disks, while only one accretion disk (around the primary) exists. The dotted line shows the limiting distance below which a circumbinary disk may exist. This becomes feasible at $\\tilde{M}\\lesssim 0.2$. \\textbf{Right:}~Peak luminosities of AGN calculated for a range of values of the reduced mass $\\tilde{M}$ and reduced distance $\\tilde{r}$ in binary systems of SBH in the centers of galaxies. Equal luminosity contours are drawn at a logarithmic step of 0.1, starting from a unit luminosity $L_0$ marked by the vertical line at $\\tilde{M}=0$. } \\label{lobanov1-fig1} \\end{figure}  \\subsection{Peak luminosity of AGN}  The peak magnitude of the nuclear activity in a galaxy hosting a binary black hole system can also be connected with the reduced mass and orbital separation of the two black holes. Assuming that the accretion rate increases proportionally to the tidal forces acting on stars and gas on scales comparable to the accretion radius, $2\\, G\\, M_\\mathrm{bh}/\\sigma_\\star^2$, the peak luminosity from an AGN can be crudely estimated from \\[ L_\\mathrm{peak} = L_0 \\left( 1 + \\frac{\\tilde{M}}{2-\\tilde{M}}\\frac{\\tilde{M}} {\\tilde{r}^2}\\right)\\,, \\] where $L_0$ is the ``unit'' luminosity of a typical single, inactive galactic nuclei.  The peak luminosities calculated in this fashion are plotted in the right panel of Figure~\\ref{lobanov1-fig1} for the entire range of $\\tilde{M}$ and $\\tilde{r}$.  The peak luminosity increases rapidly with increasing $\\tilde{M}$ and decreasing $\\tilde{r}$, and it reaches $L_\\mathrm{peak}=1000\\,L_0$ for an equal mass binary SBH at $r\\approx0.03\\,r_\\mathrm{c}$. This corresponds most likely to powerful quasars residing in elliptical galaxies. At the same $\\tilde{r}_\\mathrm{c}$, an unequal mass binary, with $\\tilde{M}=0.15$, will only produce $L_\\mathrm{peak} \\approx 10\\,L_0$, which would correspond to a weak, Seyfert-type of active nucleus. Assuming that galaxies are distributed homogeneously in the $\\tilde{M}$--$\\tilde{r}$ diagram, this scheme implies that about 70\\% of all galaxies should be classified as inactive, while the Seyfert-type of galaxies, with $L_\\mathrm{peak}=10$--$100\\,L_0$, should constitute 25\\% of the galaxy population, and the most powerful AGN, with $L_\\mathrm{peak}>100\\,L_0$, should take the remaining 5\\%. It shows that the most powerful AGN with $L_\\mathrm{peak}>1000\\,L_0$ should be found in binary SBH with nearly equal masses of the primary and secondary black holes. Binary SBH with smaller secondary companions should produce (at the peak of their nuclear activity) weaker, Seyfert-type AGN. Evidence exists in the recent works \\cite{laine2003,letawe2006} that the nuclear luminosity does indeed increase with the progression of the merger, but more systematic and detailed studies are required. ",
        "conclusions": "The model described above can be applied effectively to high-resolution optical studies and data from large surveys that can be used to obtain estimates of the nuclear luminosities and black hole masses in active galaxies. The most challenging task is to assess the state of the putative binary, since the secondary black holes are very difficult to detect. For wide binaries, Direct evidence may be sought in galaxies with double nuclei and extranuclear compact sources. Close binaries can probably be identified only indirectly, through periodic perturbations caused by the secondary companion. Other indicators, such as flattening of the galactic nuclear density profile due to BBH \\cite{mm05}, can also be considered. Once the binary separations have been estimated, it would be possible to populate the $\\tilde{M}$--$\\tilde{r}$ diagram and study whether different galactic and AGN types occupy distinctively different areas in the diagram."
    },
    "0606/astro-ph0606151_arXiv.txt": {
        "abstract": "We consider a mechanism for the formation of superhumps in the TV~Col system, based on the possible existence of a precessional spiral wave in the accretion disk of the system. This mechanism can act in binaries with arbitrary component-mass ratios, and our precessional spiral wave model can be applied to explain observed superhumps of all types. ",
        "introduction": "Superhumps are modulations of the light curves of binary systems with periods that differ from the orbital periods by several percent, and are observed mainly during superoutbursts in SU~UMa systems. The main observational features of superhumps are described in~\\cite{Warner:95}. Various authors have put forward models to explain the superhump phenomenon (a brief description of various models and their problems are given, for instance, in~\\cite{Warner:95}). Currently, the most popular model explains these light variations as an effect of precession of the outer regions of the accretion disk. The presence of the Lindblad 3:1 resonance in the disk results in an instability that leads to precession of the outer part of the disk, with a period that is appreciably longer than the orbital period. The beating of the orbital and precessional periods gives rise to the periodic variations that are identified with superhumps. This model has several shortcomings, the most important being that it implies a limit on the maximum component mass ratio. In order for the Lindblad 3:1 resonance to be located inside the accretion disk, the ratio of the donor and accretor masses $q$ must be lower than $\\sim 0.33$~\\cite{Paczynski:77}, while there are several observed systems with superhumps that have appreciably higher $q$ values, such as TV~Col.  TV~Col is a ``permanent superhumper'', or one of a class of systems in which superhumps are always observed. It was suggested in~\\cite{Osaki:96} that cataclysmic variables with high mass-transfer rates (TV~Col is such a system) may persist in a superhump regime that is not interrupted by quiescent periods. Among stars with superhumps, TV~Col has an unusually long orbital period, $P_{orb}\\sim 5.5$~h~\\cite{Retter-et-al:2003}. Since there is a direct correlation between the orbital period of a system and the donor mass (see, for instance,~\\cite{Osaki:96}), this may be as high as $\\sim 0.56 M_{\\odot}$, while the component-mass ratio is probably in the range $q\\sim0.6\\div 0.9$~\\cite{Hellier:93}.It is obvious that the Lindblad 3:1 resonance cannot be located inside the accretion disk with such a $q$ value. This casts doubt on the popular superhump model indicated above.  In 2004, a new mechanism for the formation of superhumps in SU~UMa systems was suggested in~\\cite{AR:2004:PrecWave,AR:2004:SUUMa}. The basis of this mechanism is the idea that a precessional-type density wave can form in the accretion disk.  The precessional wave in the accretion disks forms as a result of interactions between elliptical streamlines. The asphericity of the gravitational field of the binary system forces the semimajor axes of the streamlines to precess opposite to the direction in which the matter flows. The rate of this precession is proportional to the semimajor axis of the corresponding streamline. Since the streamlines in the disk cannot intersect, their interaction results in the establishment of a certain equilibrium precession rate of all the streamlines, with their apastrons lining up so that they form a spiral pattern. Because the matter velocity is minimum at the apastrons, the density of the matter at these points grows, and a precessional spiral density wave is formed in the disk.  After the formation of the precessional density wave in the disk, the rate of accretion grows sharply (by up to an order of magnitude). Matter approaches the surface of the accretor along the precessional wave, and the region of accretion is localized in azimuth, and, hence, forms a radiating spot at the surface of accretor. The increase in the accretion rate due to the density wave explains both the development of a superoutburst and the amplitude of the superhump. The wave slowly precesses in the stationary (observer's) coordinate frame, leading to a shift of the region of enhanced accretion (i.e., of the superhump) with each rotation of the system. The beating of the orbital period and the precessional period of the wave result in the superhump period, which is slightly larger than the orbital period.  This superoutburst model based on the presence of a precessional spiral wave in the accretion disk~\\cite{AR:2004:SUUMa} made it possible for the first time to explain all important observational manifestations of superoutbursts and superhumps in SU~UMa systems, including their period, duration, and energy, the anticorrelation of the brightness and color temperature at the maximum of an ordinary superhump, the late-superhump phenomenon, etc. Moreover, this model does not place strict constraints on the component mass ratio, and so can be applied to the superhumps in the TV~Col system.  The main aim of the present paper is to investigate the possible formation of precessional spiral waves in accretion disks in systems with $q > 0.33$. We will also consider the possibility of explaining all types of superhumps with this model. ",
        "conclusions": "Our numerical modeling has shown that precessional density waves can form in the accretion disks of systems with large component mass ratios, up to $q=0.93$.  \\begin{figure}[t] \\begin{center} \\epsfig{file=fig6.eps, width=12cm} \\end{center} \\caption{ Dependence of the superhump period excess on the component mass ratio. The bold dashed line corresponds to the theoretical dependence given in~\\cite{Hirose-Osaki:90}. The thin dashed line corresponds to the same dependence for a disk whose radius is 10\\% larger. The open circles show the period excess of the precessional density wave we have found for various possible values of $q$ for TV~Col. The triangle corresponds to the period excess found for OY~Car in~\\cite{AR:2004:SUUMa}, and the square to the period excess found for IP~Peg in~\\cite{AR:2004:PrecWave}. }\\label{prec} \\end{figure}   Our derived dependence of the precession period on~$q$ is complex. Figure~\\ref{prec} shows the theoretical dependence~$\\epsilon^+(q)$ obtained in~\\cite{Hirose-Osaki:90} (bold dashed line). The open circles in this figure show the period excesses found for three different sets of parameters of TV~Col. We also show the model~$\\epsilon^+(q)$ values for OY~Car and IP~Peg found in~\\cite{AR:2004:SUUMa,AR:2004:PrecWave}. Despite the significant scatter, all the~$\\epsilon^+(q)$ values are fairly close to the theoretical line, although the values for $q=0.49$~(IP~Peg) and $q=0.93$~(TV~Col) are outliers. Note that the rotational rate of the precessional density wave depends not only on the component mass ratio, but also on the size of the region occupied by the wave. The wave can exist only in a gas-dynamically unperturbed region of the disk, whose size is limited by the depth of penetration of stationary shocks -- the two arms of the tidal shocks and the hot line (extended region of interaction of the stream from $L_1$ and the circumdisk halo) -- into the disk. The high masstransfer rate in TV~Col may result in an increase in the extent of the hot line and, consequently, to precession rates that are lower than expected theoretically. It is interesting that, in the case of high component mass ratios, the size of the unperturbed region may be significantly limited by tidal waves, which would result in an even larger decrease of the precession rate with increasing $q$, compared to the theoretical value. Nevertheless, the dependence is close to the theoretical dependence, and has a peak in the range $q\\sim 0.4\\div 0.6$.  In TV~Col, superhumps with the period $P_{sh}\\sim 6.3$~h~\\cite{Retter-et-al:2003} are observed, so that the superhump period excess is $\\sim 15\\%$. This value is out of the range of the theoretical estimates of the precession rate (but does not exceed the value obtained in~\\cite{AR:2004:PrecWave} for IP~Peg). The theoretical dependence for~$\\epsilon^+(q)$ was obtained in the model that did not take into account the gas pressure; i.e., it effectively estimates the precession rate of the semimajor axis of the orbit of a free particle. The excesses of the observed and simulated precession rates over the theoretical values may indicate that the estimated disk radius used in the simulations, $r_{disc}=\\dfrac{0.6}{1+q}$~\\cite{Paczynski:77}, is too low (in fact, the derivations in~\\cite{Paczynski:77} did not take into account the gas pressure). If the disk radius were 10\\% larger ($r_{disc}=\\dfrac{0.66}{1+q}$), all the derived precession rates would be below the theoretical line (shown by the thin dashed line in Fig.~\\ref{prec}).  In summary, we conclude that the presence of a precessional density wave in the disk can lead to superhumps both in SU~UMa stars and in binaries whose component mass ratios prohibit the location of the 3:1 Lindblad resonance inside the disk. The precessional wave model is able to explain all types of observed superhumps."
    },
    "0606/hep-th0606091_arXiv.txt": {
        "abstract": "We extend previous work showing that violation of the null energy condition implies instability in a broad class of models, including gauge theories with scalar and fermionic matter as well as any perfect fluid. Simple examples are given to illustrate these results. The role of causality in our results is discussed.  Finally, we extend the fluid results to more general systems in thermal equilibrium. When applied to the dark energy, our results imply that $w = p / \\rho$ is unlikely to be less than $-1$. ",
        "introduction": "\\label{introduction}   Every spacetime is a possible solution to Einstein's equations for some particular choice of energy-momentum tensor $T_{\\mu \\nu}$. It is therefore important to understand what general restrictions can be placed on the energy-momentum of a physical system. Such restrictions are referred to as energy conditions, and play an important role in general relativity. In an earlier paper \\cite{buniy}, two of the present authors demonstrated a direct connection between stability and the null energy condition (NEC) \\cite{null}, $T_{\\mu \\nu} n^\\mu n^\\nu \\ge 0$ for any null vector $n^\\mu$ (satisfying $g_{\\mu\\nu}n^\\mu n^\\nu = 0$). The main results of the study in Ref.~\\cite{buniy} are: (1) classical solutions of both minimally and non-minimally coupled scalar-gauge models which violate the NEC are unstable, (2) a quantum state in which the expectation of the energy-momentum tensor violates the NEC cannot be the ground state (includes models with fermions), (3) perfect fluids which violate the NEC are unstable. These results suggest that violations of the NEC in physically interesting cases are likely to be ephemeral.  Using the Raychaudhuri equation for the expansion of a hypersurface null congruence, one can show that null geodesics have non-increasing expansion as long as the NEC is obeyed \\cite{Hawking-Ellis}. This result plays an important role in the classical singularity theorems of general relativity \\cite{Hawking-Ellis} and in proposed covariant entropy bounds \\cite{bousso}. It also implies that NEC violating matter is necessary for the construction of Lorentzian wormholes~\\cite{wormholes}.  This work is an extension of Ref.~\\cite{buniy}. We give some explicit examples of the classical stability analysis (Sec. \\ref{classical}), additional details regarding the local instability in quantum states which violate the NEC (Sec. \\ref{quantum}), a simplified and more general proof of the result that the fermionic effective action does not lead to violation of the NEC (Sec. \\ref{fermions}), and a generalization of our analysis of fluids to any system in thermal equilibrium (Sec. \\ref{fluids}). In Ref.~ \\cite{dubovsky} it was shown that models with superluminal modes can evade our results in some cases. However, in Ref.~\\cite{nima} it was noted that such models are acausal. In the appendix, we show that in causal models, the NEC is sufficient to determine stability. In Sec. \\ref{classical} it is shown that considerations involving causality or superluminality are necessary only in cases where the classical background violates all rotational symmetries. Cosmological models of the dark energy must exhibit isotropy and hence are covered by the earlier results of Ref.~\\cite{buniy}.   The NEC has direct implications for the dark energy equation of state, often given in terms of $w=p/\\rho$. Dark energy has positive energy density $\\rho$ and the energy-momentum tensor is $T_{\\mu \\nu} = \\text {diag\\,}(\\rho, p, p, p)$ in the comoving cosmological frame. Therefore, $w < -1$ implies violation of the NEC. Instability as a consequence of $w < -1$ was studied previously in scalar models~\\cite{instability} and in general in Ref.~\\cite{buniy}. The analysis applies to dark energy models such as k-essence~\\cite{k-essence}, rolling tachyon condensate~\\cite{tachyon}, and phantoms~\\cite{phantom}. ",
        "conclusions": "\\label{conclusion}   We have exhibited a direct connection between violation of the NEC and instability in a wide variety of models. One important application is to the dark energy, whose equation of state $w$ should be no more negative than $-1$ in order not to violate the NEC.  In classical field theory---including all Lorentz-invariant models involving both minimally and non-minimally coupled scalar and gauge fields with second-order equations of motion---violation of the NEC generally implies the existence of an instability toward the formation of gradients. In the case of configurations that possess at least some rotational symmetry (e.g., those of interest in cosmology), violation of the NEC always implies instability (Sec. IIC). In cases with less symmetry, it is possible for NEC-violating models to be stable if they exhibit superluminal excitations \\cite{dubovsky}; however, such models are not causal \\cite{nima}. In the appendix we show that in any causal scalar model, NEC violation implies instability.  These results can be generalized to quantum theories, including fermions. Analogous results for perfect fluids relate violation of the NEC to either a complex speed of sound or a clumping instability."
    },
    "0606/astro-ph0606145_arXiv.txt": {
        "abstract": "We describe our work towards the manufacture of micro-optical arrays using freeform diamond machining techniques. Simulations have been done to show the feasibility of manufacturing micro-lens arrays using the slow-tool servo method. Using this technique, master shapes can be produced for replication of micro-lens arrays of either epoxy-on-glass or monolthic glass types. A machine tool path programme has been developed on the machine software platform DIFFSYS, allowing the production of spherical, aspherical and toric arrays. In addition, in theory spatially varying lenslets, sparse arrays and dithered lenslet arrays (for high contrast applications) are possible to produce. In practice, due to the diamond tool limitations not all formats are feasible. Investigations into solving this problem have been carried out and a solution is presented here. \\footnote{Copyright 2006 Society of Photo-Optical Engineers. This paper will be published in SPIE Conf. Series 6273 and is made available as an electronic preprint with permission of SPIE. One print or electronic copy may be made for personal use only. Systematic or multiple reproduction, distribution to multiple locations via electronic or other means, duplication of any material in this paper for a fee or commercial purposes, or modification of the content of the paper are prohibited.} ",
        "introduction": "The demands of modern astronomical micro-optical components in terms of surface shape, number of elements and packing density are increasingly more difficult to meet with classical optics manufacture. Apart from the difficulty of producing a large number of identical spherical or possibly aspheric surfaces, the alignment of the elements onto a common carrier is a difficult task. Furthermore the stability of this alignment, e.g. in cryogenic situations, is critical. Alternatively, machining of optical components out of a single block, creating the array or a negative array replication master, becomes increasingly popular especially in context of free-form machines. Once the numerical model is derived, the part can be CNC machined without interaction e.g. for alignment needs between the different surfaces. The CfAI has been involved in design and manufacture of different mirror arrays used for image-slicing integral-field-units (IFUs), based on the advanced image slicer design (Content 1998, \\cite{Content1998}). Those arrays incorporate apheric and toric surfaces, that are achievable with 3D-machining techniques. However, experience with image-slicing arrays has shown that monolithic manufacture is not always possible due to tooling constraints. A similar constraint is described in context of a dithered lenslet array master. \\label{sect:intro}  % ",
        "conclusions": "The theoretical work has shown that free-form machined arrays may be used to tackle the problem of lenslet crosstalk due to spikes. However, the slow-tool servo is not a feasible method for creation of a monolithic dithered lenslet array. The design had to be adapted to the manufacturing process, on the expense of two more elements and one additional substrate. While the micro-prism array may have some potential for general situations, where beams have to be translated depending on the element they are entering, this is a tradeoff between contrast gain on one, but complexity and scatter on the other side. It is planned to test the manufacture of a dithered array using the slow-tool servo method and to compare it with a monolithic array that can be produced by use of a diamond-machined composite of linear arrangements. The comparision of both techniques will enable a final judgement if the slow-tool servo method is the preferable way of manufacture for dithered lenslet array masters."
    },
    "0606/astro-ph0606690_arXiv.txt": {
        "abstract": "Short-period binaries represent extreme cases in the generation of stellar coronae via a rotational dynamo.  Such stars are important for probing the origin and nature of coronae in the regimes of rapid rotation and activity saturation. VW~Cep ($P=0.28$~d) is a relatively bright, partially eclipsing, and very active object.  Light curves made from \\chan/\\hetgs\\ data show flaring and rotational modulation, but no eclipses.  Velocity modulation of emission lines indicates that one component dominates the X-ray emission.  The emission measure is highly structured, having three peaks. Helium-like triplet lines give electron densities of about $3-18\\times10^{10}\\, \\mathrm{cm\\mthree}$. We conclude that the corona is predominantly on the polar regions of the primary star and compact. ",
        "introduction": "\\object{VW~Cep} (\\object{HD~197433}), one of the X-ray-brightest of contact binaries, is a W-type W~UMa system --- one in which the more massive and larger star has {\\em lower} mean surface brightness such that the deeper photometric eclipse occurs during the occultation of the smaller star.  VW~Cep has an 0.28 d (24 ks) orbital period, is partially eclipsing ($i=63^\\circ$), and has component spectral types of K0~V and G5~V \\citep{Hill:1989,Hendry:Mochnacki:2000}, or, according to \\citet{Khajavi:al:2002}, F5 and G0.  Among the coronally active binaries, activity is strongly correlated with the rotation rate, and in particular, the Rossby number, which is a measure of the relative importance of Coriolis forces in the convective layer, which is in turn related to the magnetic dynamo strength \\citep{Pallavicini:1989}. At periods below one day, activity saturates \\citep{Vilhu:Rucinski:1983,Cruddace:Dupree:1984}.  Since the activity level, as defined by $L_x/L_\\mathrm{bol}$, actually decreases with increasing rotation rate, this trend has been referred to as ``super-saturation'' \\citep{Prosser:Randich:al:1996, Randich:1998,Jardine:Unruh:1999, James:Jardine:al:2000,Stepien:Schmitt:Voges:2001}. \\REV{The saturation level of $\\log(L_x/L_\\mathrm{bol})\\sim -3$ extends down to periods of 0.4 days.  For contact binaries with periods between $0.2$ to $0.4$ days, the median $L_x/L_\\mathrm{bol}$ is lower by a factor of four.} \\citet{Buzasi:1997} showed that in rapidly rotating low-mass stars, magnetic loops would be swept to the poles.  Super-saturation may occur because loops are large and are unstable to the Coriolis forces as they exceed the co-rotation radius: extended loops get swept to the poles \\citep{Jardine:Unruh:1999}.  Or loops could be compact (relative to the stellar radius), and the dynamics of surface flows clear equatorial regions \\citep{Stepien:Schmitt:Voges:2001}.  The two scenarios are similar in that activity is predominantly polar, but they differ in an important respect: X-ray sources are either large volume and rarefied or low volume and dense.  Assuming correlation between photospheric spots and coronal emission, optical light curve modeling is consistent with either scenario: Doppler image maps of VW~Cep \\citep{Hendry:Mochnacki:2000} showed large polar spots.  X-ray light curve modeling has been difficult because of the high probability of confusion by flaring, lack of phase redundancy, and the unavailability of high-resolution spectral diagnostics.  Previous high-energy observations have characterized the coronae of VW Cep.  \\cite{Choi:Dotani:1998} analyzed \\asca\\ spectra, and found a flux of about $1\\times10^{-11}$ $\\mathrm{ergs\\,cm^{-2}\\,s^{-1}}$, and two component model temperatures of $7$ and $22\\times10^6$K ($\\log T\\sim 6.8$ and $7.3$) with about equal emission measures.  They obtained significantly reduced abundances of Fe, Si, Mg, and O, but a Solar value for Ne.  A flare also occurred during the observation, with a factor of three increase in the count rate.  They used the flare emission measure and loop-scaling models to estimate a density of about $5\\times10^{10}\\mathrm{cm^{-3}}$.  The Ginga observations \\citep{Tsuru:al:1992} showed a thermal plasma temperature in excess of $10^8\\,$K, a flux similar to that determined by \\cite{Choi:Dotani:1998}, no rotational modulation, and Fe~K flux lower than expected.  The value of $L_x/L_\\mathrm{bol}\\sim -3.6$ places VW~Cep in the super-saturated regime (using the bolometric value implied by \\citet{Stepien:Schmitt:Voges:2001}). ",
        "conclusions": "Super-saturation, at least in the cases of VW~Cep and 44~Boo, appears to be manifested in small area of coverage by compact, near-polar structures.  Given roughly a factor of two for coronal structures occurring in only one hemisphere, and a factor of two in expected emission per area (the secondary is weak relative to the primary), the X-ray flux can be significantly depressed below the saturated limit of $L_\\mathrm{x}/L_\\mathrm{bol}\\sim10^{\\mthree}$.  {\\em Why} this occurs is still unexplained, but we favor the general scenario of \\citep{Stepien:Schmitt:Voges:2001} which forms polar, but compact, structures.  Given the small area of corona on the secondary, we predict that the secondary's coronal signature be highly changeable with phase over different epochs.  This would show as changes in the depth of primary eclipse and phase of perturbation in the velocity profile.  The \\hetgs\\ is the only X-ray spectrograph currently capable of making these measurements.  Obtaining high signal-to-noise per phase bin is of utmost importance to surface geometry reconstruction techniques. Further multi-orbit spectroscopy of these or other short period systems is certainly important to determination of transient or common features.  A factor of 10 increase in exposure would permit line profile diagnostics in a single line (\\eli{Ne}{10}) and greatly facilitate modeling."
    },
    "0606/astro-ph0606003_arXiv.txt": {
        "abstract": "We explore the dynamics of dark energy models based on a Dirac--Born--Infeld (DBI) tachyonic action, studying a range of potentials. We numerically investigate the existence of tracking behaviour and determine the present-day value of the equation of state parameter and its running, which are compared with observational bounds. We find that tachyon models have quite similar phenomenology to canonical quintessence models. While some potentials can be selected amongst many possibilities and fine-tuned to give viable scenarios, there is no apparent advantage in choosing a DBI scalar field instead of a Klein--Gordon one. ",
        "introduction": "Most dark energy modelling using scalar fields has followed the quintessence paradigm of a slowly rolling canonical scalar field. However, there has been increasing interest in loosening the assumption of a canonical kinetic term. In its most general form, this idea is known as k-essence \\cite{AMS}. A more specific choice is the `tachyon' \\cite{gib02}, which can be viewed as a special case of k-essence models with Dirac--Born--Infeld (DBI) action \\cite{chi03}. This kind of scalar field is motivated by string theory as the negative-mass mode of the open string perturbative spectrum, though its use in the dark energy sector is primarily phenomenological. One goal of such studies is to investigate whether there are any distinctive signatures of non-canonical actions available to be probed by observations. For a recent comprehensive review of dark energy dynamics, see Ref.~\\cite{CST}.  Tachyon dark energy has been explored by many authors, for example Refs.~\\cite{CGJP,HL,BJP,GST,CGST,car06}.\\footnote{Here we do not consider the tachyon either as a dark matter candidate \\cite{SW,PC} or as the inflaton field.}  Two papers are particularly closely related to the present work. Bagla \\emph{et al.}~\\cite{BJP} focussed on two specific choices of tachyon potential, and carried out numerical analysis of the cosmological evolution in order to constrain them against supernova data and the growth rate of large-scale structure. Copeland \\emph{et al.}~\\cite{CGST} studied a wider range of potentials, concentrating mainly on analytical inspection of attractor behaviour and the critical point structure without making comparison to specific observations. In this paper, we aim to merge some of the positive features of each analysis, by studying a wide range of potentials and testing them directly against current observational constraints as given in Ref.~\\cite{WM}.  The mechanism of slow rolling is the key ingredient in order to get an accelerating evolution driven by a scalar field. Since the DBI action can be expanded to match the Klein--Gordon one in this regime, one does not expect to find radically different features from the traditional quintessence models. However, we point out that, as compared to canonical quintessence, tachyon models require more fine-tuning to agree with observations. This is consistent with the properties of the slow-roll correspondence between the tachyon and an ordinary scalar \\cite{ben02,GKMP}. ",
        "conclusions": "None of the models we have discussed are very satisfactory. Those which carry reasonable theoretical motivation all end up with a high degree of fine-tuning. Either the potential normalization $U_c$ has to be set to match the observed dark energy density, or the initial conditions tuned to the creeping regime, meaning that the present density was already set in place during the early Universe. In either case, this tuning amounts to a restatement of the cosmological constant problem, rather than a resolution. The models which avoid fine-tuning of initial conditions (while still subject to tuning of the normalization), such as the inverse power-law, are constrained by observations into parameter regimes with no theoretical motivation.  The string tachyon with DBI action seems only weakly competitive as a dark energy candidate, and another formulation of the effective theory might have more successful applications. The tachyonic effective action as the lowest-order level truncation of cubic string field theory has been studied only very recently and its cosmological impact has yet to be fully assessed \\cite{are04,AJ,cut,AK}. An alternative would be to abandon the pure four-dimensional picture (ideally corresponding to a low-energy single or coincident brane configuration in a higher-dimensional spacetime) and consider more general brane--anti-brane setups where open string modes naturally live. However, any modification to Einstein gravity would be severely constrained after nucleosynthesis.  Finally, it was shown that the addition of electric and magnetic fields in the non-BPS brane world volume can slow down the evolution of the inflationary tachyon and relax the fine tuning of the parameters \\cite{CQS,cre05}. It might be worth checking whether the insertion of non-trivial fluxes plays some role in the late universe.  Still, the era of high-precision cosmology opened up by microwave background and large-scale structure observations is allowing us to constrain inflationary and dark energy models in a more and more stringent way, selecting some of them from a plethora of possibilities. Here we have given an example of the first stage of such a procedure. The second one will be to develop and refine those models that seem particularly promising, by embedding them in a comprehensive and consistent picture of the cosmological history and its particle theory content."
    },
    "0606/astro-ph0606235_arXiv.txt": {
        "abstract": "The AMANDA neutrino detector has been in operation at the South Pole for several years. A number of searches for extraterrestrial sources of high energy neutrinos have been performed. A selection of results is presented in this paper. The much larger IceCube detector will extend the instrumented ice volume to a cubic kilometer and 9 out of 80 planned IceCube strings have been deployed to date. We present the  status for both detectors. ",
        "introduction": "\\subsection{Motivation}  The detection of cosmic neutrinos poses great challenges, but has attractive features as well. As opposed to other particles and to electromagnetic radiation,  neutrinos can travel through huge amounts of matter without being absorbed. They can thus lead to the discovery of distant production sites hidden from observation with conventional means. Furthermore, they do not get deflected by magnetic fields during their journey and so their path points straight back to their point of origin. Some known extra-terrestrial sources that  are believed to accelerate particles to very high energies are prime candidates for neutrino production. In that scenario, neutrinos are produced as a result of the interaction of accelerated charged particles with surrounding matter or photons, producing pions, which themselves decay into muons and neutrinos. The actual observation of neutrinos would thus shed light on the hadronic nature of the accelerating process. The energy feeding that process is of gravitational origin and possible candidates are Active Galactic Nuclei (AGNs),SuperNova Remnants (SNRs) or Gamma Ray Bursters (GRBs). Models predict a hard energy spectrum $ E^{-2}$, which at high enough energies has a flux detectable above the background of the much steeper atmospheric neutrino flux. Another possible scheme for neutrino production is one where neutrinos come from the decay of massive relic particles \\cite{learned-mannheim}.    \\subsection{Detection principle}  High energy neutrinos interact via charged and neutral currents with ordinary matter and produce charged particles in the process. These particles in turn send out Cherenkov light when they travel through a medium with a refractive index greater than one. In order to detect the weakly interacting neutrino, large volumes of matter must  be instrumented with light sensors. In the cases of AMANDA and IceCube, the ice sheet at the South Pole serves as the Cherenkov medium. Notice, that the clarity of this ice allows for a sparse instrumentation, reducing the cost of the detector for a given geometry. Muon neutrinos produce long range secondary muons (1 km at 200 GeV) sending out light along their path through the detector. The direction of the parent neutrino can be directly inferred from that of the muon since the mean angle between the two particles is proportional to $E^{-0.5}$, with a value of of $\\approx 0.7^\\circ$ at 1 TeV. Electron and tau neutrinos will produce cascades, reconstructed as point sources using the Cherenkov light, but from which some directional information can be extracted as well. The main background to  these events comes from atmospheric muons produced in the atmosphere above the detector, so only events reconstructed as up-going are considered as signal, since atmospheric muons from this direction get absorbed by the Earth. At PeV energies, the Earth becomes opaque to neutrinos, including the background made by those produced in the atmosphere and one has to look for events coming from close to the horizon and above. At those energies, the $\\nu_\\tau$ has a characteristic ``Double bang'' signature: after the first cascade, produced when the $\\nu_\\tau$ interacts with the ice, the resulting tau particle decays after hundreds of meters, producing a second cascade. Thus the topologies of the three neutrino flavors become quite distinguishable \\cite{topologies}.  \\begin{figure}[htb] \\centering\\includegraphics[height=10cm]{amanda.ps} \\caption{\\it Schematic of the AMANDA detector, with a zoomed in view of one optical module. \\label{amandadetector} } \\end{figure} ",
        "conclusions": "We have presented a selection of the AMANDA results from data taken over the past ten years. The transition from AMANDA to the much larger IceCube detector is under way . The performance of the first IceCube string is as expected. The experience gained in deploying eight additionalstrings with new equipment bodes well for the higher rate of deployment in coming years. \\begin{figure}[H] \\centering \\includegraphics[width=3cm]{event.ps} \\caption{\\it An event acquired with the current nine-string IceCube detector. A shower passes through IceTop and the muon component leaves a track in IceCube. \\label{event}} \\end{figure}"
    },
    "0606/hep-ph0606120_arXiv.txt": {
        "abstract": "We investigate the impact of quantum fluctuations on a light rolling quintessence field from three different sources, namely, from a coupling to the standard model and dark matter, from its self-couplings and from its coupling to gravity. We derive bounds for time-varying masses from the change of vacuum energy, finding $\\Delta m_e/m_e\\ll 10^{-11}$ for the electron and $\\Delta m_p/m_p\\ll 10^{-15}$ for the proton since redshift $z\\sim 2$, whereas the neutrino masses could change of order one. Mass-varying dark matter is also constrained. Next, the self-interactions are investigated. For inverse power law potentials, the effective potential does {\\it not} become infinitely large at small field values, but saturates at a finite maximal value. We discuss implications for cosmology. Finally, we show that one-loop corrections  induce non-minimal gravitational couplings involving arbitrarily high powers of the curvature scalar $R$, indicating that quintessence entails modified gravity effects. ",
        "introduction": "A possible explanation for the observed \\cite{Riess:1998cb,Perlmutter:1997zf} acceleration of the universe is given by a light rolling scalar field \\cite{Wetterich:1987fm,Ratra:1987rm}, usually called quintessence field. The dynamics of this field can lead to a decaying dark energy, and thus address the question why the cosmological ``constant\" is small but non-zero today. Presently, there is a huge number of dark energy models based on  scalar fields, see e.g. ref. \\cite{Copeland:2006wr} for a review. Models admitting  e.g. tracking solutions \\cite{Steinhardt:1999nw} or general scaling solutions \\cite{Copeland:2006wr} additionally possess some appealing properties like attractors which wipe out the dependence on the initial conditions of the field in the early universe, or a dynamical mechanism naturally yielding an extremely small classical mass of the quintessence field of the order of the Hubble parameter. The latter can be necessary e.g. to inhibit the growth of inhomogeneities of the scalar field \\cite{Ratra:1987rm}.  However, the rolling quintessence field usually cannot be regarded as completely independent of other degrees of freedom. If the quintessence dynamics are for example governed by a low-energy effective theory which is determined by integrating out some unknown high energy degrees of freedom, involving e.g. quantum gravity, string theory or supergravity \\cite{Copeland:2006wr,Chung:2002xj}, the low-energy theory should generically contain couplings and self-couplings of the quintessence field suppressed by some large scale, e.g. the Planck scale. In some cases such couplings can be directly constrained observationally, like for a coupling to standard model gauge fields \\cite{Carroll:1998zi}, whereas a significant interaction with dark matter seems to be possible \\cite{Amendola:2002bs} and is used in many models, e.g. \\cite{Amendola:2000uh,Farrar:2003uw,Zhang:2005rg,Huey:2004qv,Srivastava:2004xt,Zimdahl:2001ar}, often accompanied by a varying dark matter mass (VAMP) \\cite{Rosenfeld:2005pw,Franca:2003zg,Comelli:2003cv,Hoffman:2003ru}. Models leading to time-varying standard model masses and couplings, including mass-varying neutrinos (MaVaNs), from a corresponding coupling to a rolling dark energy field are also frequently considered, see e.g. \\cite{Wetterich:2002wm,Doran:2002ec,Doran:2002bc,Fardon:2003eh,Chiba:2001er,Wetterich:2003jt,Anchordoqui:2003ij,Lee:2004vm,Copeland:2003cv,Brax:2006dc}, which can lead to potentially observable effects like a variation of the electron to proton mass ratio \\cite{Reinhold:2006zn,Ivanchik:2005ws} and the fine-structure constant \\cite{Webb:2000mn,Chand:2004ct} or violations of the equivalence principle \\cite{Wetterich:2003jt,Uzan:2002vq}, and could have an effect on BBN \\cite{Campbell:1994bf,Uzan:2002vq}. Furthermore, higher order self-interactions of the field seem to be a typical feature necessary for a successful dark energy model, involving e.g. exponentials or inverse powers of the field \\cite{Ratra:1987rm,Wetterich:1987fm,Albrecht:1999rm,Amendola:1999er,Steinhardt:1999nw,Binetruy:1998rz}. Non-minimal gravitational couplings of the scalar field have also been studied in various settings \\cite{Perrotta:1999am,Chiba:1999wt,Carvalho:2004ty,deRitis:1999zn,Faraoni:2000wk,Catena:2004ba,Boisseau:2000pr,Gannouji:2006jm}, constrained e.g. by solar system tests of gravity and BBN.  Because of the presence of quantum fluctuations of the standard model degrees of freedom as well as of dark matter and of dark energy itself, the dynamics of the quintessence field receive radiative corrections. Therefore, it is important to study the robustness against these corrections, see e.g. refs. \\cite{Kolda:1998wq,Brax:1999yv,Doran:2002bc,Uzan:1999ch,Bartolo:1999sq,Riazuelo:2001mg,Onemli:2002hr} for previous work. Apart from the long-standing problem of the overall normalization of the effective quintessence potential (i.~e. the ``cosmological constant problem\"), which is not addressed here, quantum corrections can influence the dynamics e.g. by distorting the {\\it shape} or the flatness of the scalar potential. In the case of a coupling to standard model or heavy cold dark matter particles, this leads to tight upper bounds for the corresponding couplings, due to a physically relevant field-dependent  shift in the corresponding contribution to the vacuum energy \\cite{Donoghue:2001cs,Banks:2001qc}. The quantitative bounds obtained in this way can be translated into bounds for time-varying masses and for the coupling strength to a long-range fifth force mediated by the quintessence field, as will be shown in section \\ref{sec:2}. The self-coupling and gravitational coupling of the dark energy field are necessary ingredients of basically any given model. The corresponding quantum corrections can lead to significant modifications of the shape of the scalar potential, discussed in section \\ref{sec:3} using an infinite resummation of bubble diagrams. This accounts for the fact that the effective theory for a scalar field does not decouple from the high energy regime due to the presence of quadratic divergences. We also discuss implications for cosmology. In section \\ref{sec:4} we investigate which kind of non-minimal gravitational couplings are induced by quantum fluctuations of the dark energy scalar field. ",
        "conclusions": "Summary and Conclusions}   In this work we have investigated the effect of quantum fluctuations in the context of typical dynamical dark energy scenarios like quintessence models, where an extremely light rolling scalar field supplies the present cosmic acceleration, in some sense similar to the inflaton in the early universe.  First couplings between the quintessence field and heavier degrees of freedom, like the standard model fermions or dark matter, have been discussed. We constrained the discussion to couplings that can effectively be written as a field-dependent mass term. These couplings have to be extremely small  even though we fine-tune the energy density, slope and mass of the quintessence field at its todays value by appropriate renormalization conditions for the quartic, quadratic and logarithmic divergences in the induced effective potential. This leads to a bound on time-varying masses between $z\\sim 2$ and now of the order $10^{-11}$ for the electron and scaling proportional $m^{-4/3}$ with mass, assuming the mass variations are not themselves finely tuned in such a way that the total shift in vacuum energy is negligible. Moreover, we found that the coupling strength to a fifth force mediated by the quintessence field has to be suppressed by a number of the same order relative to its gravitational coupling strength. Only neutrinos could have a large mass variation and interact with the quintessence field as strong as with gravity.  Second we introduced a suitable approximation scheme to investigate the impact of quintessence self-couplings on the {\\it shape} of the effective potential, while an undetermined additive constant has been fine-tuned to be zero, thus bypassing the unresolved ``cosmological constant problem\". We showed that the quantum fluctuations to the scalar potential can be consistently renormalized in leading order in $V''/\\mu^2$, where $\\mu$ is a high energy scale characteristic for an underlying theory and $V''$ the square of the quintessence mass assumed to be of the order of the Hubble parameter. While potentials involving exponentials just get rescaled, inverse power law potentials are flattened at small field values. The effective potential approaches a finite maximal value, thus truncating the singular behaviour of the inverse power law in the field range of interest. This distortion of the potential can directly play a role cosmologically if $\\mu$ is large, roughly $\\mu\\gtrsim M_{pl}/10$, and moreover affect general properties like tracking behaviour.   Third we investigated non-minimal gravitational couplings induced by quantum fluctuations. Since quintessence potentials usually cannot be taken to be a quartic polynomial in the field, non-minimal couplings apart from a term of the form $\\phi^2R$ can be induced. We showed that at one-loop all couplings of the form $\\phi^nR^m$ with integers $n$ and $m$ have to be included, and will be induced by quantum corrections unless the field is exactly conformally coupled. Moreover, we showed that this type of non-minimal coupling which is nonlinear in $R$ cannot be simply removed by a Weyl rescaling, but corresponds to a theory with two interacting scalars in the Einstein frame. Altogether, this may indicate that the origin of cosmic acceleration could not purely be the effect of one rolling scalar field, but also involves modified gravity effects similar to $f(R)$-theories when quantum fluctuations of the scalar are taken into account.  We conclude that quantum fluctuations do play an important role (i) in coupled dynamical dark energy scenarios even if we allow fine-tuning in the form of renormalization conditions, (ii) for the shape of the quintessence potential and (iii) for its interplay with gravity.  \\vspace*{1cm}  \\begin{appendix}"
    },
    "0606/astro-ph0606079_arXiv.txt": {
        "abstract": "We present a database of $BVRI$ time-series photometry of the Orion Nebula Cluster obtained with two ground-based telescopes at different longitudes to provide simultaneous coverage with the 13-d {\\it Chandra} observation of the cluster. The resulting database of simultaneous optical and X-ray light curves for some 800 pre--main-sequence (PMS) stars represents, by a factor of hundreds, the largest synoptic, multi--wavelength-regime, time-series study of young stars to date. This database will permit detailed analyses of the relationship between optical and X-ray variability among a statistically significant ensemble of PMS stars, with the goal of elucidating the origins of PMS X-ray production. In this first paper, we present the optical observations, describe the combined X-ray/optical database, and perform an analysis of time-correlated variability in the optical and X-ray light curves. We identify 40 stars (representing 5\\% of our study sample) with possible time-correlated optical and X-ray variability. Examples of both positive and negative time-correlations are found, possibly representing X-ray flares and persistent coronal features associated with both cool and hot surface spots (i.e.\\ magnetically active regions and accretion shocks). We also find two possible examples of ``white-light\" flares coincident with X-ray flares; these may correspond to the impulsive heating phase in solar-analog flares.  However, though interesting, these represent unusual cases.  More generally, we find very little evidence to suggest a direct causal link between the sources of optical and X-ray variability in PMS stars. The conclusion that accretion is a primary driver of X-ray production in PMS stars is not supported by our findings. ",
        "introduction": "} The X-ray luminosities of solar-type stars, when they are very young, are prodigious. Indeed, it is now well established that low-mass pre--main-sequence (PMS) stars can emit X-rays at levels up to $\\sim 10^4$ times that of the present-day Sun \\citep[e.g.][]{feig-mont}. But after two decades of research since T~Tauri stars (TTSs) were first identified as strong X-ray sources \\citep[e.g.][]{feig-decamp}, we still do not know the origins of X-ray production in young stars.  While flares observed in the X-ray light curves suggest that the X-rays are produced mainly in solar-analog magnetic reconnection flares, the absence of an obvious rotation-activity relation in these stars prevents a clear association of the X-rays to dynamo-generated fields \\citep{flac-basic,feig03,stass04}. The flares can be so powerful that non-solar geometries of the magnetic fields must be considered, including field lines linking the star with the protoplanetary disk \\citep{favata05}. In addition, the role of accretion is unclear, and may affect the production of X-rays either directly or indirectly. For example, energetic shocks associated with accretion flows at or near the stellar surface may enhance X-ray production, as has been suggested in the detailed spectroscopic analyses of TW~Hya \\citep{kastner02,stelzer04} and BP~Tau \\citep{schmitt05}. At the same time, the mass-loading of magnetic field lines may actually inhibit X-ray production in accreting systems, as suggested by \\citet{preibisch05}. These issues are reviewed more fully in \\citet{feig-ppv}.   We have undertaken a synoptic study of the Orion Nebula Cluster (ONC) combining a nearly continuous, 13-d {\\it Chandra} observation with simultaneous, multi-wavelength, time-series photometry in the optical. A specific motivation is to establish the incidence of time-correlated X-ray and optical variability in order to begin disentangling the roles of magnetic and accretion activity in the production of X-rays. For example, accretion shocks on the surfaces of classical TTSs (CTTSs; TTSs that still possess accretion disks) produce strong variability at optical wavelengths, caused by time-variable accretion, rotational modulation of accretion ``hot spots,\" or both \\citep{herbst94,bouvier97,stasswood99}. If X-rays are produced near the sites of accretion shocks, then one might expect that changes in the strength of these shocks will induce changes in the strength of the X-ray emission, and that X-ray and optical variability might therefore be correlated in time. Alternatively, X-rays may originate in coronal structures associated with magnetically active regions on these stars, in which case we might expect that X-ray emission will be more pronounced when the magnetically active regions (which appear as dark photospheric spots in the optical) face toward us, and least pronounced when these regions face away from us; that is, optical and X-ray variability will be anti-correlated in time. More complex behavior may result from the superposition of these effects, emphasizing the need for time-resolved measurements. More generally, by establishing the frequency with which X-ray and optical variations are co-temporal, we seek to constrain the extent to which the sources of the X-ray and optical variability may be co-spatial or otherwise causally linked.  Detailed analyses of simultaneous, time-resolved X-ray and optical variability have previously been reported for three PMS stars: the CTTSs V773~Tau \\citep{feig94} and BP~Tau \\citep{gullbring97}, and the weak-lined TTS (WTTS) V410~Tau \\citep{stelzer03}. In none of these were the optical and X-ray variability found to be time-correlated. Unfortunately, in all three cases the statistical significance of the results were limited by the short duration of the X-ray observations: V773~Tau was observed by ROSAT for 6.7~hr, BP~Tau was observed by ROSAT for 8~hr, and V410~Tau was observed by {\\it Chandra} for 7.2~hr.  The combined X-ray/optical dataset reported here includes some 800 members of the ONC and thus represents, by a factor of hundreds, the largest attempt to date to study the relationship between X-ray and optical variability in PMS stars. We will report results from this database in stages, focusing first on statistical properties of the study sample as a whole. In this first paper, we establish the incidence of {\\it time-correlated} optical and X-ray variability. A companion paper \\citep[][hereafter Paper~II]{stass06} focuses on {\\it time-averaged} measures of optical and X-ray variability, and on the relationship between these variability indicators and basic stellar properties. This database will then form the basis for follow-up studies focusing on fine-grained analyses of individual objects.  In \\S \\ref{data}, we present the X-ray and optical time-series database that forms the basis for these studies and describe the methods that we employ in this paper to search for time-correlated X-ray/optical variability. In \\S \\ref{corr-time}, we identify 40 candidate stars with possible time-correlated optical and X-ray variability, and classify them according to common light-curve characteristics to guide follow-up study of these interesting objects. Importantly, these candidates represent at most $\\sim 5$\\% of our study sample; $\\sim 95$\\% of the stars in our sample exhibit no evidence for time-correlated optical and X-ray variability. As we discuss in \\S \\ref{discussion}, these results imply that the sites of optical and X-ray variability in PMS stars are not---in the vast majority of cases---instantaneously one and the same. A direct causal link between the sources of optical and X-ray variability in PMS stars is not supported by our study. We summarize our findings in \\S \\ref{conclusions}. ",
        "conclusions": "} We have presented a combined analysis of simultaneous optical and X-ray light curves for a sample of more than 800 pre--main-sequence (PMS) stars in the Orion Nebula Cluster (ONC). This dataset represents an increase of several hundred-fold, both in the number of objects studied and in the total duration of the X-ray and optical time-series data, over all previous investigations of simultaneous optical and X-ray variability in PMS stars.  We identify 40 stars that exhibit the first examples to our knowledge of time-correlated optical and X-ray variability in PMS stars. Among these 40 stars, we find examples of both positive and anti-correlations. Positive correlations may represent X-ray--dark coronal structures associated with cool, magnetically active surface spots, or may represent X-ray--bright structures associated with hot accretion spots. Anti-correlations may represent, e.g., magnetic flaring events associated with cool surface spots. In any event, the implication is that the X-ray emitting regions in these stars are in close proximity to---and thus likely to be physically linked with---cool and/or hot spots on their surfaces. Follow-up modeling and analysis of the color information contained in our multi-band optical light curves should help to clarify whether we are observing cool or hot spots in each object. In addition, we identify two possible examples of ``white-light\" flares, perhaps corresponding to the impulsive heating phase in solar-analog flares.  However, these cases represent at most $\\sim 5$\\% of our study sample. The remaining $\\sim 95$\\% of our study sample exhibits no such time-correlated variability. Evidently, time-correlated optical and X-ray variability is rare in PMS stars. We conclude that the sites of optical and X-ray variability in PMS stars are not---for the vast majority of stars in our study sample---instantaneously one and the same. Surface accretion shocks and cool magnetic spots are evidently not the exclusive or dominant sites of X-ray production in most PMS stars."
    },
    "0606/astro-ph0606286_arXiv.txt": {
        "abstract": "The self-accelerating braneworld model (DGP) can be tested from measurements of the expansion history of the universe and the formation of structure. Current constraints on the expansion history from supernova luminosity distances, the CMB, and the Hubble constant exclude the simplest flat DGP model at about 3$\\sigma$. Even including spatial curvature, best-fit open DGP model is a marginally poorer fit to the data than flat $\\Lambda$CDM.   Moreover, its substantially different expansion history raises serious challenges for the model from structure formation. A dark-energy model with the same expansion history would predict a highly significant discrepancy with the  baryon oscillation measurement due the high Hubble constant required. For the DGP model to satisfy this constraint new non-linear phenomena must correct this discrepancy. Likewise the large enhancement of CMB anisotropies at the lowest multipoles due to the ISW effect would require either a cut off in the initial power or new phenomena at the cross over scale.    A prediction that is robust to both possibilities is that high-redshift galaxies should be substantially correlated with the CMB through the ISW effect.  This correlation should provide a sharp test of the DGP model in the future. ",
        "introduction": "Cosmic acceleration may arise either from a new form of energy density, dubbed dark energy, or from a modification of gravity on cosmological scales. Dvali, Gabadadze and Porrati (DGP) \\cite{dvali00} have proposed a braneworld modification of gravity to explain acceleration (see \\cite{lue05} for a recent review). In this model our universe is a (3+1)-dimensional brane embedded in an infinite Minkowski bulk.  The cosmological solution of this theory exhibits self-acceleration on the brane classically \\cite{deffayet00,deffayet01}.  Various cosmological probes which can distinguish DGP from dark energy models (DE) have been investigated in the literature \\cite{lue04,song04,schimd04,knox05,ishak05,amarzguioui05,linder05,sawicki05,koyama05,zhang05,stabenau06,maartens06}.  These come in two classes: those that test DGP modifications of the expansion history and those that test changes in how cosmological structures are formed. The advantage of the former class is that the predictions of the DGP model are well understood theoretically.  The main disadvantage is that they currently lack the power to distinguish the DGP model from its dark energy counterpart the $\\Lambda$CDM model at high significance.  Moreover, expansion rate tests cannot distinguish between the DGP model and an arbitrary form of smooth dark energy that is constructed to mimic it in the background.  Structure formation tests on the other hand can in principle solve both these problems.  However due to complexities in the theory they currently can only be rigorously implemented on a rather small range of scales.   In this paper, we examine the combination of these two types of tests. We employ three tests of the expansion history: the angular diameter distance to recombination from the three-year WMAP data~\\cite{spergel06}, luminosity distance to high-redshift supernovae (SNGold\\cite{riess04} and SNLS\\cite{astier05}) and the local expansion rate from Hubble constant measurements.  These tests do in fact strongly disfavor the simplest version of the DGP model: a modification of gravity in a spatially flat cosmology with no bulk cosmological constant.   The best fit DGP model has a significant negative curvature and a high Hubble constant.  Even with the addition of spatial curvature, it is a slightly poorer fit to the data than flat $\\Lambda$CDM.  The best fit DGP model yields an expansion history that is strikingly different from flat $\\Lambda$CDM.   The dark energy model that mimics this expansion history has negative curvature and a dark energy equation of state that increases sharply with redshift.   In the dark energy context such an expansion history would be strongly excluded by structure formation tests in weakly to strongly non-linear regime.  These include the baryon acoustic oscillations (BAO) in galaxy clustering \\cite{eisenstein05,maartens06}, the abundance of galaxy clusters,  and fluctuations in the Lyman-$\\alpha$ forest.   While these probes likely also test the DGP model, none of these phenomena are sufficiently well understood in the model to provide a definitive test. Indeed the presence of ghost degrees of freedom in the DGP model around a de-Sitter background \\cite{Luty:2003vm,Nicolis:2004qq,Gorbunov:2005zk,Charmousis:2006pn} require that the nature of gravitational interactions change dramatically in the strongly-coupled non-linear regime for the model to remain a viable explanation of cosmic acceleration \\cite{Deffayet:2006wp}.  Under this assumption, structure formation in DGP is currently well understood on scales between the Hubble scale and the scale radius of a typical dark matter halo.  Large scale CMB anisotropy is substantially enhanced by the Integrated Sachs-Wolfe (ISW) effect.   At the lowest multipoles, this excess poses a significant challenge for the model.  Possible resolutions of this problem still leave a robust prediction and sharp test of the DGP model: high redshift galaxies should be substantially correlated with the CMB. ",
        "conclusions": "The self-accelerating DGP braneworld model is challenged by expansion history constraints from the CMB, supernovae and the Hubble constant.   Unlike a model with a cosmological constant as the dark energy, the DGP model requires spatial curvature to satisfy the CMB and supernovae constraint and even so mildly exceeds bounds on the Hubble constant.  In a dark energy model, the DGP expansion history would be ruled out by structure formation tests.   These are serious but perhaps not insurmountable challenges for the model. Among the most powerful are the baryon oscillations and the large angle CMB anisotropy from the ISW effect. In the DGP model, these two phenomena lie close to the two critical scales: $r_*$ where gravity transitions from obeying the usual Einstein equations to a scalar-tensor or weak brane phase and $r_c$ when gravity exits the weak brane phase and becomes fully 5 dimensional. At the $r_{*}$ transition, the theoretical framework to calculate structure formation is currently incomplete.  At the $r_{c}$ scale, changes in the initial power spectrum or new gravitational physics may ameliorate the problem.   The enhancement of the ISW effect from the decay of the gravitational potential extends to intermediate scales where the theoretical calculations are robust. The main difference between the DGP  and the flat cosmological constant expansion history is the prediction that the gravitational begins its decay at high redshift. A sharp test of the DGP scenario would be to isolate this decay at high redshift by cross correlating the CMB with a high-redshift galaxy population.  The cross correlation coefficient is also relatively robust to changes in the initial amplitude of fluctuations. We show that a galaxy survey that covers a substantial fraction of the sky at $z>1$ can detect this correlation at high significance.   Such surveys are currently being planned for the measurement of high-redshift baryon oscillations.  Alternately, radio sources, $X$-ray sources and quasars provide other high-redshift populations for cross correlation.  A possible source of confusion for this test is the apparent number density fluctuation induced by magnification from gravitational lensing by the same structures associated with the ISW effect \\cite{Hui06}.  Magnification both lowers the number density for a fixed population behind a lens and also raises it for a fixed flux limit by bringing faint galaxies into the sample.  This effect may be separated from a true spatial correlation by examining different populations of objects where the competition between these two effects differ and the cross-correlation between high and low redshift galaxy populations.  A full examination of this issue is beyond the scope of this paper but it is likely that the galaxy-ISW cross correlation will become a sharp test of the DGP braneworld acceleration model in the future.   \\vspace{1.cm}  \\noindent {\\it Acknowledgments}: We thank Sean Carroll, Yu Gao, Nemanja Kaloper, Arthur Lue, and Xiaomin Wang for useful conversations. This work was supported by the U.S.~Dept. of Energy contract DE-FG02-90ER-40560. IS and WH are additionally supported by the David and Lucile Packard Foundation. This work was carried out at the KICP under NSF PHY-0114422."
    },
    "0606/astro-ph0606415_arXiv.txt": {
        "abstract": "We investigate the orientation of the axes and angular momentum of dark matter halos with respect to their neighboring voids using high resolution N-body cosmological simulations. We find that the minor axis of halos tends to be aligned along the line joining the halo with the center of the void, and that the major axis is perpendicular to this line. However, we find that the angular momentum of halos does not have any particular orientation. These results may provide information about the mechanisms whereby the large-scale structure of the Universe affects galaxy formation and cast light upon the issue of the orientation of galaxy disks with respect to their host halos. ",
        "introduction": "The knowledge of the intrinsic alignments of the angular momentum of dark matter halos with respect to their axes and the alignments of these axes with respect to the galaxy disks is of great importance in order to understand the processes involved in galaxy formation. Moreover, the alignment of halo axes, angular momentum and galaxy spins with respect to the large-scale structure may provide further clues to this end; for example, they could be determinant to understand the origin of the angular momentum of galaxies (see e.g. Navarro et al. 2004). They also may help to understand several observational results like the distribution of galactic satellites (e.g. Zentner et al. 2005, Azzaro et al 2005), the formation of galactic warps (e.g. Dekel \\& Shlosman 1983; Toomre 1983; Bailin 2004) and the tidal streams observed for example in the Milky Way (e.g. Ibata et al. 2002). Furthermore, they could be a source of contamination in weak lensing studies (e.g. Heymans et al. 2006).  Dark matter halos are in general triaxial being defined by their three semi-axes, $a \\geq b \\geq c$ (Dubinski \\& Carlberg 1991; Warren et al. 1992; Cole \\& Lacey 1996; Jing \\& Suto 2002; Allgood et al. 2006). Recently, some effort have been devoted to measure their intrinsic alignments with respect to the angular momentum. It has been found that the angular momentum of dark matter halos tends to be parallel to their minor axis (Bailin \\& Steinmetz 2005; Allgood et al 2006). This situation is typical for rotationally supported objects. However, as dark matter halos are found to be not rotationally supported, this alignment could have a non trivial origin (Allgood et al 2006).  On the other hand, the measurement of the alignments of galaxy disks with their host dark matter halos is both theoretical and observationally difficult. Using cosmological hydrodynamic simulations of disk galaxy formation, Bailin et al. (2005) found that the halo minor axis within 0.1$R_{vir}$ matches the disk axis. Beyond 0.1$R_{vir}$, the halo orientation is uncorrelated with the disk axis. Moreover, Hoekstra, Yee \\& Gladders 2004, using galaxy-galaxy lensing measurements find that galaxies (disks plus ellipticals combined) are aligned with the major axis of their dark matter halo. Other works show that the angular momentum of disks and dark matter are not perfectly aligned (van den Bosch et al. 2002; Sharma \\& Steinmetz 2005). Hence, the final orientation of the galactic disk in a cosmological dark matter halo is still an unresolved problem.  One way of gaining insight about these alignments is to study the alignments of both, halos and disks, with respect to the large-scale structure. The basic idea is that the angular momentum of halos and galaxies is generated by tidal torques operating on the primordial material destined to form them (Peebles 1969; White 1984). In this scenario, the primordial angular momentum of barions and dark matter are aligned and the final galactic angular momentum is a `relic' of the primordial one. According to this, the current misalignment between galactic and dark matter halo angular momentum is due to subsequent tidal torques affecting dark matter halos. However, this scenario could be more complex because additionally to the primordial tidal torques operating on the material which eventually becomes the disk, the mass accretion by the disk is a clumpy and stochastic process (see Vitvitska et al. 2002). Along this line, Navarro et al. (2004) found, using high resolution numerical simulations, that galaxy disks are rather well aligned with the angular momentum of the surrounding matter at the epoch of turnaround. Bailin \\& Steinmetz (2005) found that the minor axes of halos of groups and clusters show a strong tendency to lie perpendicular to the direction of filaments. However, it has been suggested that although filaments cause alignments of cluster-size halos, they have not discernible influence on smaller halos (Altay et al. 2006). On the other hand, it has recently been found using observations from the 2dFGRS and the SDSS, that galactic disk planes tend to be perpendicular to the walls of cosmic voids (Trujillo, Carretero \\& Patiri 2006).  In this letter we study, using high resolution N-body simulations, the orientation of axes and angular momentum of galactic size dark matter halos embedded in the walls of large voids, those with radius larger than $10 \\mpch$. We will discuss the observational implications of our results. ",
        "conclusions": "In Figure 1 we show the orientations of the three halo axes and angular momentum with the cosmic voids. In the top-left panel we show the probability distribution of the $\\alpha$ angle which is defined between the major axis of the dark matter halos and the vector pointing to the center of the void (black dots). Similarly, in the top-right and bottom-left panels we present the angles $\\beta$ and $\\gamma$ which are the ones between the middle and minor axis with the direction to the center of the void respectively. Finally, in the bottom-right panel we show the angle $\\theta$ between the halo angular momentum and the direction to the center of the void. The solid line corresponds to the probability distribution of randomly distributed angles, which is a sine function for geometrical reasons.  In these plots we clearly see that halos are not randomly oriented within the shells of voids. The major axis of halos tends to lie along the plane perpendicular to the radial direction, i.e. there is an excess of halos with $\\alpha$ angles near 90 degrees. On the contrary, the minor axis tend to point to the direction of the center of the void. There is a marginal orientation of the middle axes toward the plane perpendicular to the radial direction. However, we do not see any particular orientation of the halo angular momentum with respect to the large-scale structure, confirming recent results obtained also from N-Body simulations using similar methodology (Heymans et al. 2006).  To test the robustness of our results we have conducted several tests. First of all, we have analyzed the orientations of dark matter matter halos selected as described above, but in this case we used ``fake'' voids, i.e. we randomly placed spheres of fixed radius over the sample and calculated the angles between the axis (and angular momentum) of the dark matter halos located in the shells of these ``fake'' voids and the direction to their center. As expected, we found that the angles of the halos chosen to be in the shells of random spheres are uniformly distributed, i.e. they follow a sine distribution. Moreover, we applied two different statistical tests: the average of the cosine test and a standard Kolmogorov-Smirnov test. If the angles are randomly distributed, the average of the cosine of these angles is $0.5$. So that, the bigger is the deviation from this value the larger will be the effect. In Table 1 we show the results of the statistical tests. In the upper part of the table we show the results for the mean cosine (column 1) and the probabilities that the results be compatible with isotropy as given by the mean cosine and the KS tests (columns 2 and 3) for halos located in the shells of voids and, in the bottom part the same results but for halos inside the shells of the random spheres (i.e. the `fake' voids). From these results, we conclude that we can reject with high confidence (more than 7 sigmas) the null hypothesis (i.e. that the axis of dark matter halos are randomly distributed) for the halos in the shells of voids. As expected, the results of the statistical test for the halos located in the shells of the random spheres show that the orientation of these halos is compatible with the random distribution.  On the other hand, we used our data to gain some insight into the nature of the internal alignment of the minor axis with angular momentum found in previous works (e.g. Bailin \\& Steinmetz 2004). We selected those halos where both directions form an angle smaller that $20\\deg$. For these halos, we found that the alignment between the minor axis and the radial direction of the void is weaker than for randomly chosen halos, being even compatible with a random distribution.  Comparing these results with the observed tendency for the spin of spiral galaxies to point along the plane perpendicular to the direction of the center of the void found by Trujillo, Carretero \\& Patiri (2006), three scenarios are possible: 1) The non-isotropy of the orientation of the spin of the galaxies is entirely induced by the halo alignments, 2) both alignments are two effects with common origin (namely, the tidal field in the shells of voids), or 3) a combination of the previous ones. In the first scenario the probability distribution for the angle between the major axis of the host halo and the spin of their galaxy is independent of the angle between the spin and the radial direction of the void. In this case, this probability distribution may be uniquely determined by comparing the distribution of the angle between the major axis and the radial direction with that of the angle between the spin of the galaxies and the radial direction. This would imply that the spin is very closely aligned with the major axis of the host halo. However, recent hydrodynamical simulations (Bailin et al. 2005) does not show this effect (they found that the spin of galaxies is uncorrelated with the halo). Therefore, the other two scenarios seem more plausible."
    },
    "0606/astro-ph0606623_arXiv.txt": {
        "abstract": "The development program of the flight model imaging camera for the PACS instrument on-board the Herschel spacecraft is nearing completion. This camera has two channels covering the 60 to 210 microns wavelength range. The focal plane of the short wavelength channel is made of a mosaic of 2x4 3-sides buttable bolometer arrays (16x16 pixels each) for a total of 2048 pixels, while the long wavelength channel has a mosaic of 2 of the same bolometer arrays for a total of 512 pixels. The 10 arrays have been fabricated, individually tested and integrated in the photometer. They represent the first filled arrays of fully collectively built bolometers with a cold multiplexed readout, allowing for a properly sampled coverage of the full instrument field of view.  The camera has been fully characterized and the ground calibration campaign will take place after its delivery to the PACS consortium in mid 2006. The bolometers, working at a temperature of 300~mK, have a NEP close to the BLIP limit and an optical bandwidth of 4 to 5 Hz that will permit the mapping of large sky areas. This paper briefly presents the concept and technology of the detectors as well as the cryocooler and the warm electronics. Then we focus on the performances of the integrated focal planes (responsivity, NEP, low frequency noise, bandwidth). ",
        "introduction": "\\label{sect:intro}  The Herschel Space Observatory is the third ``corner stone'' mission of the European Space Agency. It will be launched by an Ariane 5 rocket in the course of 2008. Herschel will be equipped with the largest telescope ever sent in space (\\O\\,3.5m) and will carry out spectroscopic and imaging observations in the 60 $\\mu$m to 670 $\\mu$m wavelength range. Herschel's payload consists of three instruments. (1) HIFI is a very high resolution heterodyne spectrometer ($R\\sim10^7$), (2) SPIRE is an imager and an imaging spectrometer, operating in the 210-670 $\\mu$m band, using spider-web bolometers coupled to Winston cones (see ref.~\\citenum{turner} for details), and (3) PACS covers the 60-210 $\\mu$m range and is both an imaging spectrometer using photo-conducting detectors, and an imager using novel technology bolometers described in this paper.  The main science objectives of Herschel are twofold. First, Herschel will perform large scale surveys of nearby dark clouds, regions where stars form, in order to identify the mechanisms responsible for the distribution of stellar masses. Indeed we now realize that a star's mass, when it enters the main sequence, is in fact determined when the gas and dust cloud in which it will later form separates itself from its parent cloud and starts to collapse. At this early stage, it is mostly heated by the contraction and its temperature is such that it radiates most of its energy in the Herschel band. By observing very large numbers of these prestellar cores, we will shed light upon the processes that lead to their formation and to their mass distribution.  Herschel will also peer into the distant Universe. Half of the extragalactic background light reaches earth in the infrared, with a peak in Herschel's bandpass. Herschel will perform deep surveys in dark regions of the sky to identify and locate the galaxies responsible for this background. This will allow the reconstruction of the star formation history of the Universe during the last $\\sim$10 Gyr. This star formation history is in fact the result of the galaxy formation process, thus Herschel will participate in the construction of a plausible scenario that leads from the very homogeneous Universe of the Big Bang epoch to the highly structured Universe of galaxies that we see now. For more detailed descriptions, see ref.~\\citenum{pilbratt} for the Herschel mission, and ref.~\\citenum{poglitsch} for the PACS instrument. ",
        "conclusions": "\\label{sect:futur}  The development of the bolometers for the Herschel/PACS instrument demonstrates that it is now possible to build fairly large ($>$ 1000 pixels) focal planes based on filled arrays. The grid + resonant cavity concept works, the science requirements on the Blue focal plane have been reached and we will improve the performances of the Red focal plane before delivery to the PACS consortium.  The future of this type of bolometers looks very bright. We have already started a number of developments that will make use of these detectors in the sub-mm and millimeter wavelength ranges, opening the possibility of wide field imaging in these spectral domains. The two main axes of development are (1) to redesign the packaging of the arrays to make them 4-side buttable, opening the way for very large focal planes of bolometers, and (2) to adapt the grid + cavity concept for absorption at longer wavelength.  Table~\\ref{tab:projets} lists the different projects in terms of spectral range, size of focal planes, and telescopes. \\begin{table}[htbp] \\caption{\\label{tab:projets}On-going developments of filled arrays of bolometers.} \\begin{center} \\begin{tabular}{|c|c|c|c|c|} \\hline Name & Spectral range & Focal plane & Telescope & Operational in \\\\ \\hline P-ARTEMIS & 2 channels 200 + 450 $\\mu$m & 1 array per channel & KOSMA + Chile & 2006 \\\\ PILOT & 2 channels 240 + 550 $\\mu$m & 2x2 arrays per channel & balloon & 2009 \\\\ ARTEMIS-1 & 3 bands 200 -- 350 -- 450 $\\mu$m & 4x4 arrays & APEX & 2009 \\\\ ARTEMIS-2 & 3 bands 0.85 -- 1.2 -- 2 mm & 4x4 arrays & open (IRAM...) & 2010 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The adaptation of the bolometers to longer wavelengths is simply a tuning of the depth of the resonant cavity, in order to keep it at a quarter of the wavelength of interest. This tuning can be done in several ways: keeping the same cavity and putting a dielectric layer with the proper thickness on top of the pixel, inserting a spacer in between the absorbing grid and the reflector to adjust the depth of the cavity, or a mix of these two methods. This adaptation to longer wavelengths is presented in ref.~\\citenum{reveret}, and the ARTEMIS project is presented in greater details in ref.~\\citenum{talvard}.  A first observing run of the P-ARTEMIS camera on the KOSMA telescope occurred in March 2006, which demonstrated that these detectors actually work at longer wavelengths. Figure \\ref{fig:kosma} shows a scanned map of Jupiter taken at 450 $\\mu$m.  \\begin{figure}[htbp] \\begin{center} \\begin{tabular}{c} \\includegraphics[height=8cm]{jupiter_db_scan.eps} \\end{tabular} \\end{center} \\caption{\\label{fig:kosma}Scanned map of Jupiter at 450 $\\mu$m.} \\end{figure}"
    },
    "0606/astro-ph0606309_arXiv.txt": {
        "abstract": "{We derive black hole masses for a sample of about 300 AGNs in the redshift range $0 < z < 2.5$. We use the same virial velocity measure (FWHM \\hbbc) for all sources which represents a significant improvement over previous studies. We review methods and caveats for determining AGN black hole masses via the virial assumption for motions in the gas producing low ionization broad emission lines. We derive a corrected FWHM(\\hbbc) measure for the broad component of \\hb\\ that better estimates the virialized line emitting component by comparing our FWHM measures with a sample of reverberated sources with \\hb\\ radial velocity dispersion measures. We also consider the FWHM of the \\feiiq\\ blend as a potential alternative velocity estimator. We find  a range of black hole mass between $\\log $\\mbh\\ $\\sim$ 6.0--10.0, where \\mbh\\ is in solar masses. Estimates using corrected FWHM(\\hb), as well as FWHM(\\feii) measures, reduce the number of sources with $\\log$ \\mbh\\ $>$ 9.5 and suggest that extremely large \\mbh\\ values ($\\log$ \\mbh $\\ga$ 10) may not be realistic. Derived \\lledd\\ values show no evidence for a significant population of super-Eddington radiators especially after correction is made for sources with extreme orientation to our line of sight. Sources with  FWHM(\\hbbc) $\\la$ 4000 \\kms\\ show systematically higher \\lledd\\ and lower \\mbh\\ values than broader lined  AGNs (including almost all radio-loud sources). ",
        "introduction": "}  Gravitational accretion onto supermassive black holes is generally accepted as the ultimate energy source of Active Galactic Nuclei (AGNs). The last decade has seen a major effort to derive reasonable estimates of black hole masses (\\mbh) by assuming virialized motions in the broad line emitting gas:  \\begin{equation} M_{\\rm BH} =   \\frac{v^2 r_{\\rm BLR}}{G} \\label{eqn:virial} \\end{equation}  where $ v $\\ is the velocity dispersion of the emitting gas at distance \\rb. The velocity dispersion can be written as $ v = f {\\rm FWHM} $\\ where FWHM is the full width half maximum measured for a suitable emission line. The factor $ f $\\ depends on the geometry and details of the kinematics \\citep{krolik01,mcluredunlop01,onkenetal04}. The usefulness of the virial assumption for \\mbh\\ determination is best exploited using reverberation-mapping studies, especially for the Balmer lines \\citep{koratkargaskell91,petersonwandel99,wandeletal99, kaspietal00,petersonetal04}. \\citet{kaspietal00,kaspietal05} derive a relation between Broad Line Region (BLR) distance $r_{\\rm BLR}$\\ and continuum luminosity  \\begin{equation} r_{\\rm BLR} \\propto (\\lambda L_\\lambda )^{\\alpha} \\label{eqn:kaspi} \\end{equation} where the exponent $\\alpha $\\ is constrained between 0.5 and 1.0, and is most likely  $\\approx$ 0.6--0.7 for broad \\hb\\ (\\hbbc) and optical continuum luminosity.  There are caveats associated with this method. The virial assumption is not likely to be generally valid for the emission line gas in AGNs. It has been known for several decades that different emission lines in a source can show different width and profile shape \\citep[e.g.,][]{derobertis85,sulentic89}. The virial assumption implies that the velocity dispersion will steadily decrease with distance from the central black hole $\\propto r^{-\\frac{1}{2}}$. The time lag between continuum fluctuations and corresponding emission line responses will therefore anti-correlate with line width. This trend  has been confirmed in a few objects \\citep{petersonwandel00}.  Observations however show that profile width and shape depend on the ionization potential. The strongest high-ionization lines (HILs; e.g. \\civ) often display blueward asymmetric profiles, or even centroid blueshifts (up to several 1000 \\kms) with respect to the best estimates of the rest frame of the source \\citep[e.g., ][]{gaskell82, marzianietal96,richardsetal02,bachevetal04}. Blueshifts like those observed for \\civ\\ are indicative of obscuration and radial motions which invalidate the virial assumption.  Low ionization lines (LILs) like \\hb\\ are the best candidates for emission arising from a virialized medium. All or much of the Balmer (and \\feii) line emission is thought to arise from an accretion disk or a flattened cloud distribution near the disk \\citep[e.g.,][]{collinsouffrinetal88,marzianietal96}. A major caveat, especially for the Balmer lines, involves the possibility that: (1) there are two or more emission components in a line, and (2) only one of them may arise in a region where the virial assumption holds. The Balmer lines do not usually show very large line shifts (i.e., shift $\\leq$ FWHM) although profiles can be very asymmetric. Both red and blue asymmetries are observed for \\hbbc\\ which is the most studied line because it is relatively unblended and is observable with optical spectrometers up to $z \\approx 1$\\ \\citep{osterbrockshuder82,sulentic89,stirpe90, sulenticetal90,corbin91, corbinfrancis94}.  The most ambiguous sources from the point of view of \\mbh\\ determination show FWHM(\\hbbc) $\\ga$ 4000 \\kms\\ \\citep[Population B, following][]{sulenticetal00b} and redshifted profiles and/or red asymmetries. Not all parts of the H$\\beta$\\ profile respond to continuum changes in the same way implying that some of the line emitting gas may be optically thin or, less likely, is not exposed to the variable \\hi\\ ionizing radiation (we will return to this issue in \\S \\ref{improvemass}). The broad line profile in Pop. B sources may be due to two distinct emitting regions: (1) an optically thick classical BLR and (2) a broader and redshifted very broad component that may be optically thin or marginally optically thick to the Lyman continuum, originating in a distinct Very Broad Line Region (VBLR) \\citep{marzianisulentic93,shieldsetal95,sulenticetal00c}. The redshift of the VBLR component raises doubts that it arises from virialized gas. A strong BLR response to continuum fluctuations, coupled with a weak or absent response of the VBLR component, can lead to an overestimate of FWHM for the virialized BLR component resulting in an overestimate of \\mbh\\ \\citep{wandeletal99,kaspietal00,vestergaard02}.  Sources with FWHM(\\hbbc) $\\la$ 4000 \\kms\\ \\citep[Population A,][]{sulenticetal00b} should provide more reliable \\mbh\\ estimates. The \\hbbc\\ profile is usually well fit with  a symmetric  function \\citep{veroncettyetal01,sulenticetal02} and the BLR emission is thought to arise from a Keplerian disk. The most unreliable Pop. A sources, in a disk emission scenario, should be those observed near face-on where the rotational (i.e. virial) contribution to FWHM(\\hbbc)\\ is minimal. At least some of the face-on sources may be identified as the so-called ``blue outliers\" which show a weak and significantly blueshifted \\oiiiopt\\ lines \\citep{zamanovetal02,marzianietal03b,aokietal05,boroson05}. The inferences in this, and the preceding paragraph have emerged from the Eigenvector 1 (E1) scenario that we have pursued for the past 5+ years \\citep[][following \\citet{borosongreen92}]{sulenticetal00a,sulenticetal00b}. The results reported in this paper appear to confirm them.  There are additional caveats connected  with using FWHM(\\hbbc) for \\mbh\\ derivations. Line asymmetries can affect \\mbh\\ estimates for both Pop. A and B sources. The \\vr\\ displacement of the line centroid at fractional intensities (at half maximum, $c(\\frac{1}{2}$)) can give useful information about the uncertainty of FWHM(\\hbbc) measurements. The average displacement value is a few hundred \\kms\\ with median $c(\\frac{1}{2}$) $\\approx 400$ \\kms\\ \\citep{marzianietal03a}. This implies that deviations from a symmetric profile can affect \\mbh\\ estimates by $\\approx$ 40 \\%. Uncertainties are also introduced by: (1) contamination from overlapping/nearby lines such as \\feii, \\heii, \\ and \\oiiiopt\\ \\citep{osterbrockshuder82,derobertis85,joly88,jacksonbrowne91}, (2) FWHM measures based on single-epoch observations, (3) low $S/N$\\ spectra  and (4) spectra without \\hb\\ narrow component (\\hbnc) subtraction. The \\rb\\ - $L_\\lambda$\\ relation is also not free from uncertainties. Reverberation mapping-based \\mbh\\ determinations are certainly affected by the non-negligible radial extent of the optically thick BLR. The derived \\rb\\ is  not a very well defined quantity and $\\alpha$\\ is also somewhat uncertain because of the intrinsic scatter in the correlation. Finally, reverberation data does not exist for high luminosity/redshift quasars requiring an extrapolation of the \\rb\\ - $L_\\lambda$\\ relation  in order to estimate \\mbh\\ for these sources.  Uncertainties for \\mbh\\ derivations using single profile observation of \\hbbc\\ are estimated to be a factor of 2--3 (at a 1 $\\sigma$\\ confidence level), but may be as low as $30$\\%\\ if the velocity dispersion of the variable part of \\hbbc\\ profile is employed as a virial estimator \\citep{petersonetal04}. \\mbh\\ estimates based on the virial relation retain a statistical validity considering that AGNs span a 5dex range in \\mbh\\ (10$^5$ \\msol $\\la$ \\mbh\\ $\\la 10^{10}$ \\msol) and that the relation has now been applied to large samples of objects  \\citep[$ \\sim$ 10$^3$;][]{mcluredunlop04}.  With these considerations in mind, and supported by previous results, we use the virial relation to compute \\mbh\\ for $\\approx$ 280 AGNs in our E1 sample of low redshift (0 $\\la z \\la 1$) sources \\citep{marzianietal03b} supplemented with 25 intermediate redshift/high luminosity (1 $\\la z \\la 2.5$) quasars. We derive \\mbh\\ and \\lledd\\ values in several ways. New VLT-ISAAC data are presented for 9 sources (\\S \\ref{obs}) which supplement data already published for sixteen quasars \\citep{sulenticetal04}. Line measures are presented  in \\S \\ref{results}. We use \\hbbc\\ and \\feiiq\\ line widths in a consistent way over the redshift range $0 \\la z \\la 2.5$  (a range of 10$^5$\\ in luminosity). We compute black hole mass \\mbh, Eddington ratio \\lledd\\ (\\S \\ref{mass}) and we discuss how mass determinations might be improved (\\S \\ref{improvemass}) so that the evolution of \\mbh\\ and \\lledd\\ with redshift can be considered (\\S \\ref{evolution} and \\S \\ref{discussion}). ",
        "conclusions": "Nine intermediate $z$\\ VLT/ISAAC spectra with high resolution and $S/N$\\ supplement an earlier sample of 17 sources. Emission line measurements on \\hbbc\\ and \\feii\\ presented in this paper strengthen the conclusion of \\citet{sulenticetal04} that luminosity effects are weak or absent in the low-ionization lines of AGNs. Results on the \\hbbc\\ profile are consistent with the population A-B hypothesis and Eigenvector 1 parameter space concept developed for low $z$\\ AGNs. We computed virial masses and Eddington ratios for the 25 intermediate-$z$ objects plus about 280 lower $z$\\ sources, using the same emission line over the entire redshift range for the first time.  We also have  how the distributions of \\mbh\\ and \\lledd\\ vs. $z$\\ are shaped by selection effects intrinsic to any flux-limited survey at the low \\mbh\\ end. At the high \\mbh\\ end, masses exceeding a few 10$^9$ \\msol\\ may be rare if corrections for non-virial broadening and statistical errors are taken into account. This suggestion is based on just 25 objects distributed over the entire redshift range $0.9 \\la z \\la 2.5$. Confirmation from a larger sample of intermediate-$z$\\ observations is needed.   \\newpage"
    }
}